source,target
 Using flux weighted means over those channels of each group. the positions and velocities of spectrum features were obtained.," Using flux weighted means over those channels of each group, the positions and velocities of spectrum features were obtained."
 The uncertainties in relative positions are typically 10 mas., The uncertainties in relative positions are typically 10 mas.
" The mmaser line at 22GHz was observed using 4C39.25 as bandpass and flux calibrator,", The maser line at 22GHz was observed using 4C39.25 as bandpass and flux calibrator.
 The flux density of 4C39.25 at this frequency was taken to be 7.8 Jy (Terasranta priv., The flux density of 4C39.25 at this frequency was taken to be 7.8 Jy (Terasranta priv.
 comm.)., comm.).
 The phase calibrator was mapped with a total of two rounds of phase self-calibration and the resulting corrections applied to the IRAS 20126+4104 data., The phase calibrator was mapped with a total of two rounds of phase self-calibration and the resulting corrections applied to the IRAS $20126+4104$ data.
 The spectral bandwidth was 4 MHz corresponding to 54 km s! velocity range with channel separation of 0.25 km s!., The spectral bandwidth was 4 MHz corresponding to 54 km $^{-1}$ velocity range with channel separation of 0.25 km $^{-1}$.
 Maps of all the spectral channels were generated and de-convolved using the AIPS task IMAGR., Maps of all the spectral channels were generated and de-convolved using the AIPS task IMAGR.
 The restoring data beam had a FWHM of 40 x 8 mas at a position angle of —377., The restoring data beam had a FWHM of 40 $\times$ 8 mas at a position angle of $-37^\circ$.
" The rms noise was typically 11 mJy/beam but up to 40 mJy/beam in the spectral channels with the brightest emission,", The rms noise was typically 11 mJy/beam but up to 40 mJy/beam in the spectral channels with the brightest emission.
 The 6.7-GHz methanol line was observed with just the two antennas in the MERLIN array equipped with the appropriate receivers at the time of the observations., The 6.7-GHz methanol line was observed with just the two antennas in the MERLIN array equipped with the appropriate receivers at the time of the observations.
 The correlator was configured to give a velocity resolution of 0.21 km s! for a total of 2 MHz spectrum bandwidth corresponding to 90 km s! velocity range., The correlator was configured to give a velocity resolution of 0.21 km $^{-1}$ for a total of 2 MHz spectrum bandwidth corresponding to 90 km $^{-1}$ velocity range.
 3C84 was used as the bandpass and amplitude. calibrator., 3C84 was used as the bandpass and amplitude calibrator.
 Its amplitude at the time of the, Its amplitude at the time of the
a 1nore modest conduction efficiency. (O¢ond=1.0 and 0.1 respectively).,a more modest conduction efficiency $\alpha_{\rm cond}=1.0$ and 0.1 respectively).
 In these models. conduction is not sufficient to suppress cooling in the larger halos adequately.," In these models, conduction is not sufficient to suppress cooling in the larger halos adequately."
 If we adopt a lower value for ox. however. a lower conduction efiicicucy gives a reasonable match to the observed luminosity. function.," If we adopt a lower value for $\sigma_8$, however, a lower conduction efficiency gives a reasonable match to the observed luminosity function."
 Model 7.1. shows the DIuunositv function for the case σς=0.7 and oc; T., Model 7.4 shows the luminosity function for the case $\sigma_8=0.7$ and $\alpha_{\rm cond}=7$ .
 This conduction cficiency could be achieved if the temperature eradieut was Toxrt} aud the conduction was ouly slightly suppressed below the Spitzer value., This conduction efficiency could be achieved if the temperature gradient was $T\propto r^{1.3}$ and the conduction was only slightly suppressed below the Spitzer value.
 Although this is still a high rate of conduction. it offers a promising route for explaining the bright end of the galaxy Iunimositv function.," Although this is still a high rate of conduction, it offers a promising route for explaining the bright end of the galaxy luminosity function."
 The expulsion of gas frou halos at hieh cnerey can. In principle. strongly suppress the formation of later eoncrations of galaxies. hence affecting the shape of the huninosity function.," The expulsion of gas from halos at high energy can, in principle, strongly suppress the formation of later generations of galaxies, hence affecting the shape of the luminosity function."
 Starting frou Model 5.2. we add further feedback cucrey that expels cold gas completely. not oulv from the disk but also from the halo.," Starting from Model 5.2, we add further feedback energy that expels cold gas completely, not only from the disk but also from the halo."
 The superwind aust have high cenereyv in order that the expelled inaterial not be recaptured by nore massive halos., The superwind must have high energy in order that the expelled material not be recaptured by more massive halos.
 The effect of à low power superwind is illustrated by Model 8.1 (dashed Lue in Fie. 5))., The effect of a low power superwind is illustrated by Model 8.1 (dashed line in Fig. \ref{fig:models_7}) ).
" This model. with €,=0.27 and Jw=3. has a relatively weak superwind."," This model, with $\esw=0.27$ and $\beta_{\rm SW}=3$, has a relatively weak superwind."
" This eas expulsion is in addition to the reheating of cold disk gas (€:chear= 0.13): we have assuined that there is no heating of the diffuse halo (e,a,,= 0.0).", This gas expulsion is in addition to the reheating of cold disk gas $\erh=0.13$ ); we have assumed that there is no heating of the diffuse halo $\eho=0.0$ ).
 Although the winds eject a large amount of eas. most of the material is recaptured as larger halos collapse aud the Inunuinositv function differs little from Model 5.2.," Although the winds eject a large amount of gas, most of the material is recaptured as larger halos collapse and the luminosity function differs little from Model 5.2."
 Iu Model 8.2 (dotted line). we have set e4;=5.0 aud jay=L. corresponding to a mean energy per superwiud article of Ly.=15keV. Such an euergetie wind is required to ensure that very little material is recaptured by eroup halos.," In Model 8.2 (dotted line), we have set $\esw=5.0$ and $\beta_{\rm
SW}=1$, corresponding to a mean energy per superwind particle of $E_{\rm av}=15$ keV. Such an energetic wind is required to ensure that very little material is recaptured by group halos."
 In this model. the superwind dominates the ecdhack energy budget: indeed. the total euergy required (5.13.«ο1) significantly exceeds that available roni supernovae alone.," In this model, the superwind dominates the feedback energy budget; indeed, the total energy required $5.13 \times10^{49}\ergMsol$ ) significantly exceeds that available from supernovae alone."
 The model comes much closer to natching the luminosity function. but still overproduces xieht ealaxies.," The model comes much closer to matching the luminosity function, but still overproduces bright galaxies."
 We can increase the superwiud cuerey urther. but a factor of 2 increase oulv results iu a small Huprovement iu the huninosity function.," We can increase the superwind energy further, but a factor of 2 increase only results in a small improvement in the luminosity function."
 If we inclhide conduction as well as superwiuds. it is. of course. possible o nmuprove the match but a hieh coucuction efficicucy (Ocona291) as still needed.," If we include conduction as well as superwinds, it is, of course, possible to improve the match but a high conduction efficiency $\alpha_{\rm cond}\gg 1$ ) is still needed."
 Increasing the mass loading of he wind substantially (bv increasing Jui) results iu too ew ealaxies around the knee of the πιοτν fiction., Increasing the mass loading of the wind substantially (by increasing $\beta_{\rm SW}$ ) results in too few galaxies around the knee of the luminosity function.
 As we found in the case of conduction. the luuinosity unction can be matched more casily if we adopt a lower value for ay.," As we found in the case of conduction, the luminosity function can be matched more easily if we adopt a lower value for $\sigma_8$."
 The case ey=5.0. Ow=Ll. og=Tod illustrated by Model 8.3 (solid line).," The case $\esw=5.0$, $\beta_{\rm SW}=1$, $\sigma_8=0.7$ is illustrated by Model 8.3 (solid line)."
 Caven the uucertainties of our recapture prescription. this model eives a reasonable match to the huninosity function: it has a strong break at the correct Inninosityv. aud overproduces bright objects onlv marginallv.," Given the uncertainties of our recapture prescription, this model gives a reasonable match to the luminosity function; it has a strong break at the correct luminosity, and overproduces bright objects only marginally."
 There are a variety of wavs to further improve the match to observations: we could Increase the energy injection (but eg.210.0 is required) or merease the mass loadiug of the wind so that the curve is shifted faintwards., There are a variety of ways to further improve the match to observations: we could increase the energy injection (but $\esw > 10.0$ is required) or increase the mass loading of the wind so that the curve is shifted faintwards.
 An alternative strateey is to allow for a low level ofconduction: μα~Lis sufficient to produce a significant iiproveinent iu the match to the Iuninositv function when σς=0.7 aud superwiuds are prescut., An alternative strategy is to allow for a low level of conduction: $\alpha_{\rm cond}\sim1$ is sufficient to produce a significant improvement in the match to the luminosity function when $\sigma_8=0.7$ and superwinds are present.
 The parameters Oeonq aud esy are highly degencra5 in their effects ou the luminosity function G.c. lucreasing either suppresses the bright end}., The parameters $\alpha_{\rm cond}$ and $\epsilon_{\rm SW}$ are highly degenerate in their effects on the luminosity function (i.e. increasing either suppresses the bright end).
 While curent conrputational Linitations make it prolibitively expoeusive o perform an accurate 4? fit of the model parameters to he data. a crude estimate of the 4? surface for these two xuaneters shows that the data prefer mocels with stroug superwinds (egy2 6) and hieh conductivity (A¢onq2:30) or a model with e=0.93.," While current computational limitations make it prohibitively expensive to perform an accurate $\chi^2$ fit of the model parameters to the data, a crude estimate of the $\chi^2$ surface for these two parameters shows that the data prefer models with strong superwinds $\epsilon_{\rm
SW}\approx 6$ ) and high conductivity $\alpha_{\rm cond}\approx 30$ ) for a model with $\sigma_8=0.93$."
 Lowering ex to 0.7 reduces he requirements to ἐν23 and eon325.," Lowering $\sigma_8$ to $0.7$ reduces the requirements to $\epsilon_{\rm SW}\approx 3$ and $\alpha_{\rm
cond}\approx 25$."
 Further investigation of the 47 surface would require a Bayesian xior to specify formally a physically allowed range for his parameter., Further investigation of the $\chi^2$ surface would require a Bayesian prior to specify formally a physically allowed range for this parameter.
 While our priuarv goal iu this paper is to examine the Iuuiuositv fiction of galaxies. it is prudent to check whether our models are iu reasonable agreement with other basic properties of the galaxy population such as the Tull-Fisher relation aud the ealaxy autocorrelation function.," While our primary goal in this paper is to examine the luminosity function of galaxies, it is prudent to check whether our models are in reasonable agreement with other basic properties of the galaxy population such as the Tully-Fisher relation and the galaxy autocorrelation function."
 We compare the models that best fit the ealaxy DIuuinositv function to these observables., We compare the models that best fit the galaxy luminosity function to these observables.
 We retain the model parameters found earlier and we do not attempt to adjust these or anv other parameters in this comparison., We retain the model parameters found earlier and we do not attempt to adjust these or any other parameters in this comparison.
 We defer a more detailed comparison of our models with a wider range of observational constraints to a future paper., We defer a more detailed comparison of our models with a wider range of observational constraints to a future paper.
 Suinultaucously matching the galaxy Iuuinositv fiction and the Tully-Fisher relation has been a long-standing, Simultaneously matching the galaxy luminosity function and the Tully-Fisher relation has been a long-standing
For comparison. the Alfvénn velocity frou the simulation (Figure 5)) is 200 l,"For comparison, the Alfvénn velocity from the simulation (Figure \ref{fig:f5}) ) is $\approx$ 200 $^{-1}$."
 The simulations show a predominantly two-level bright frout., The simulations show a predominantly two-level bright front.
 Figure 3 shows that the highest intensity. ds concentrated im patches located within a weaker. broadening font.," Figure \ref{fig:f3} shows that the highest intensity is concentrated in patches located within a weaker, broadening front."
 Although these two conrponeuts expand iu a coupled manner during the carly phase of the CATE. during the later frames. the siauulated wave frout is increasinelv dominated by the weaker intensity component. which coutiuues to expaud(densityfrontdif.mov: densitysidedif.mov).," Although these two components expand in a coupled manner during the early phase of the CME, during the later frames, the simulated wave front is increasingly dominated by the weaker intensity component, which continues to expand; )."
 The observations become increasingly noisy as the event progresses aud the coronal wave becomes more and more difficult to detect., The observations become increasingly noisy as the event progresses and the coronal wave becomes more and more difficult to detect.
 In EUVLA base differcuce data. 06:25 UT is the last frame where the bright fout to the west of the active region is discernable (195a_diff.mov).," In EUVI-A base difference data, 06:25 UT is the last frame where the bright front to the west of the active region is discernable )."
 Iu the EUVI-À base differcuce simulation. the higher iuteusitv patch to the west of the active region is visible until 06:35 UT (densitysidedif.mov).," In the EUVI-A base difference simulation, the higher intensity patch to the west of the active region is visible until 06:35 UT )."
 Tn the EUVLED base difference data. the furthest reaches of the bright front Guost obviously ucar to the south polar coronal hole) cau be identified: until 06:55 UT.," In the EUVI-B base difference data, the furthest reaches of the bright front (most obviously near to the south polar coronal hole) can be identified until 06:55 UT."
 The bright frout is approximately stationary af this time diff.mov)., The bright front is approximately stationary at this time ).
 The EUVI-A and D base difference simmlations show that isolated higher iuteusity patches still exist at 07:15 UT., The EUVI-A and B base difference simulations show that isolated higher intensity patches still exist at 07:15 UT.
 Near to the south polar coronal hole. the higher iutensitv patch is located at the same place from 06:15 UT.," Near to the south polar coronal hole, the higher intensity patch is located at the same place from 06:45 UT."
 Near to the north polar coronal hole. a new higher intensity patch develops from 07:05 UT. also remaining at the same location.," Near to the north polar coronal hole, a new higher intensity patch develops from 07:05 UT, also remaining at the same location."
 Figure 6 compares snapshots at 06:05 UT from CORI-A. the simulated CME. and EUVI-A. all scaled to the same size.," Figure \ref{fig:f6} compares snapshots at 06:05 UT from COR1-A, the simulated CME, and EUVI-A, all scaled to the same size."
" The CORI-A aud EUVI images are runing difference mages. the simulation is a base difference Πμαρο,"," The COR1-A and EUVI images are running difference images, the simulation is a base difference image."
 Patsourakos&Vourlidas(2009) show a fit of the 3D CALE inedel of Theruisieuetal.(2006.2000) to the CORI-A data for this event at 06:05 UT.," \cite{Patsourakos09b} show a fit of the 3D CME model of \cite{Thernisien06,Thernisien09} to the COR1-A data for this event at 06:05 UT."
 This leads them to deteriuue an extent for the CME in the low corona that is much too small to match the coronal wave in the corresponding EUVI data. since the 3D model is essentially mace up of a spherical bubble attached to a conical leg.," This leads them to determine an extent for the CME in the low corona that is much too small to match the coronal wave in the corresponding EUVI data, since the 3D model is essentially made up of a spherical bubble attached to a conical leg."
 We note that a simular approach was also used in Patsourakosetal.(2009).., We note that a similar approach was also used in \cite{Patsourakos09a}.
 Tn both papers. the authors interpret the appareut musft between the CME model extension iu the low corona aud the EUVI coronal wave as evidence that the coronal wave aud CME are different structures aud couclude that the coronal wave is à fast-iiode MITD wave.," In both papers, the authors interpret the apparent misfit between the CME model extension in the low corona and the EUVI coronal wave as evidence that the coronal wave and CME are different structures and conclude that the coronal wave is a fast-mode MHD wave."
 Our simulation results at 06:05 UT are shown iu the nuüddle pauel of Figure 6.., Our simulation results at 06:05 UT are shown in the middle panel of Figure \ref{fig:f6}. .
 The simulation gives information about the low corona below 1.3 R.. (200 Alu). the region obscured by the occulting disk iu the CORI data.," The simulation gives information about the low corona below 1.3 $_{\odot}$ (200 Mm), the region obscured by the occulting disk in the COR1 data."
 Comparison of the middle aud right paucls of Figure 6 show that the extension of the CALE in the low corona maps vorv well to the coronal wave in the EUVI base difference data., Comparison of the middle and right panels of Figure \ref{fig:f6} show that the extension of the CME in the low corona maps very well to the coronal wave in the EUVI base difference data.
 Iu particulary. the simulation results show a secondary cavity located to the north of the main CAIE cavity Quarked by the white arrow in the middle panel of Figure 6)).," In particular, the simulation results show a secondary cavity located to the north of the main CME cavity (marked by the white arrow in the middle panel of Figure \ref{fig:f6}) )."
 Comparison with the corresponding EUVI base difference data shows secondary. coronal dinuiuimngs developing at the same location (Grisht panels. EUVI (A). Figure. 1)).," Comparison with the corresponding EUVI base difference data shows secondary coronal dimmings developing at the same location (right panels, EUVI (A), Figure \ref{fig:f4}) )."
 Despite the lack of spectral diagnostics for secondary dinunines. this correlation between the secondary CAE cavity and the secondary coronal cinuuiues is consistent with plasma evacuation.," Despite the lack of spectral diagnostics for secondary dimmings, this correlation between the secondary CME cavity and the secondary coronal dimmings is consistent with plasma evacuation."
 A time-series iiovie of the simulated CORI-À schite-light data (COR1jointi.mov) shows that the secondary cavity expands and merges with the main CALE cavity. so that the low corona is really “opened” to a laree lateral exteut.," A time-series movie of the simulated COR1-A white-light data ) shows that the secondary cavity expands and merges with the main CME cavity, so that the low corona is really “opened” to a large lateral extent."
 This analysis demonstrates the important role of the simulation in developing au understanding of the true lateral extent of the CALE in the low corona., This analysis demonstrates the important role of the simulation in developing an understanding of the true lateral extent of the CME in the low corona.
 Tn 1.2.0 we noted that the higher intensity patches of the coronal wave front no longer expaucd as of ~ 06:55 UT., In \ref{subsec:two_component_bf} we noted that the higher intensity patches of the coronal wave front no longer expand as of $\sim$ 06:55 UT.
 Correspondingly. the simulation results show that the CALE has stopped expanding significantly iu a lateral direction by this time. (," Correspondingly, the simulation results show that the CME has stopped expanding significantly in a lateral direction by this time. ("
The reader is referred to movie CO,The reader is referred to movie ).
R1joint1.mov). Figures 6 and 1 show time-serics plots of the simulation results matched to the STEREO-B (ou-cisl) and STEREO-À (Πο). viewing angles. respectively. (," Figures \ref{fig:f7} and \ref{fig:f8} show time-series plots of the simulation results matched to the STEREO-B (on-disk) and STEREO-A (limb) viewing angles, respectively. ("
The reader is encouraged to view the novies that correspond to these figures. aud SB.mov).,"The reader is encouraged to view the movies that correspond to these figures, and )."
 The inner sphere shows the photosphere with the radial magnetic field streneth from the maguetogram data., The inner sphere shows the photosphere with the radial magnetic field strength from the magnetogram data.
 The outer sphere (ight exev) is at height of 1.18.. (70 Nina) and represents the altitude at which coronal waves are observed., The outer sphere (light grey) is at height of $1.1R_\odot$ (70 Mm) and represents the altitude at which coronal waves are observed.
 The white contours represent the densitv-enhanced frout (same as displaved in Figure 3))., The white contours represent the density-enhanced front (same as displayed in Figure \ref{fig:f3}) ).
 The ereen shade represents an iso-surface of mass densitv with a base ratio (ratio between the current image aud pre-event image) of 1.1., The green shade represents an iso-surface of mass density with a base ratio (ratio between the current image and pre-event image) of 1.1.
 Selected core field lines of the magnetic flux rope are drawn in Red and some surrounding field lines of a range of sizes are drawn in Blue., Selected core field lines of the magnetic flux rope are drawn in Red and some surrounding field lines of a range of sizes are drawn in Blue.
 Where the core fiux rope field (Red) recounects with a surrounding field line (Blue). the Blue field line changes to Red. indicating the new extended counectivity of the core flux rope field.," Where the core flux rope field (Red) reconnects with a surrounding field line (Blue), the Blue field line changes to Red, indicating the new extended connectivity of the core flux rope field."
 Recounectious between surmrouuding magnetic field lines (i.e. Blue and Blue). are shown by the newly recounected field line changing to Yellow.," Reconnections between surrounding magnetic field lines (i.e. Blue and Blue), are shown by the newly reconnected field line changing to Yellow."
 The same magnetic field lues have been plotted in both Figures 6 aud 12.. so that we can study the same development frou the two differcut perspectives.," The same magnetic field lines have been plotted in both Figures \ref{fig:f7} and \ref{fig:f8}, so that we can study the same development from the two different perspectives."
 Referring to Figure 6.. we see that the exeeu iso-surtace of increased mass density approximately maps to the white contour at cach time frame.," Referring to Figure \ref{fig:f7}, we see that the green iso-surface of increased mass density approximately maps to the white contour at each time frame."
" Figure 135. shows a line profile of the density base aud ruuniue differences, as well as the temperature along the path ofthe coronal wave at r=L1R. (shown by the black arrow in the top panel)."," Figure \ref{fig:f9} shows a line profile of the density base and running differences, as well as the temperature along the path of the coronal wave at $r=1.1R_\odot$ (shown by the black arrow in the top panel)."
 It can be seen that the temperature jump Gudicatiug the shock frout) is ahead of the density increase associated with the coronal wave., It can be seen that the temperature jump (indicating the shock front) is ahead of the density increase associated with the coronal wave.
 This meaus that the ereeu shade represcuts the CME frout and uot the shock., This means that the green shade represents the CME front and not the shock.
 Dudeed. conrparison of the ercen iso-surtace with the white-leht siuulation and observational data (Figure 6)) further sugeests that it representsthe outer-most shell of the expanding CALE.," Indeed, comparison of the green iso-surface with the white-light simulation and observational data (Figure \ref{fig:f6}) ) further suggests that it representsthe outer-most shell of the expanding CME."
 The existence of a major reconnectiou (discussed below). further suggests that the ereen iso- represents the actual CALE rather than a shock.," The existence of a major reconnection (discussed below), further suggests that the green iso-surface represents the actual CME rather than a shock,"
system via a parallax measurement would similarly yield the absolute luminosities of the stellar components. thus provide independent access to both Rys and Rwp and the associated parameters.,"system via a parallax measurement would similarly yield the absolute luminosities of the stellar components, thus provide independent access to both $R_\mathrm{MS}$ and $R_\mathrm{WD}$ and the associated parameters."
 Nevertheless. a P.y below the period gap Is consistent with the M5-6V spectral type of the secondary star1998).," Nevertheless, a $P_\mathrm{sd}$ below the period gap is consistent with the M5–6V spectral type of the secondary star."
. Since it can be assumed that the secondary 1s fully convective. we use Eq.8 from for angular momentum loss dominated by gravitational radiation to calculate the time it will take LTT 560 to start mass transfer via Roche-lobe overflow to ~3.5 Gyrs.," Since it can be assumed that the secondary is fully convective, we use Eq.8 from for angular momentum loss dominated by gravitational radiation to calculate the time it will take LTT 560 to start mass transfer via Roche-lobe overflow to $\sim$ 3.5 Gyrs."
 This is much less than the Hubble time. thus LTT 560 can be regarded as representative of the progenitors of todays CVs.," This is much less than the Hubble time, thus LTT 560 can be regarded as representative of the progenitors of todays CVs."
 Since the system contains a non-magnetic white dwarf and the mass ratio g<0.33. itis likely that the future CV LTT 560 will belong to the SU UMa subelass of dwarf novae.," Since the system contains a non-magnetic white dwarf and the mass ratio $q < 0.33$, it is likely that the future CV LTT 560 will belong to the SU UMa subclass of dwarf novae."
spherical.,spherical.
 In addition. low mass AGB stars with low metallicity might never become totally obscured.," In addition, low mass AGB stars with low metallicity might never become totally obscured."
 Although mass loss rate is tightly coupled to pulsation. it is not clear what pulsation period should be used.," Although mass loss rate is tightly coupled to pulsation, it is not clear what pulsation period should be used."
 Also. AGB stars with different total masses and Iuminosity can have the same pulsation period at very different. effective temperatures.," Also, AGB stars with different total masses and luminosity can have the same pulsation period at very different effective temperatures."
 The usage of pulsation (e.g.. Dlóccker 1995) includes in it the dynamical time of the star. but no ‘natural’ transition value exists.," The usage of pulsation (e.g., Blöccker 1995) includes in it the dynamical time of the star, but no `natural' transition value exists."
 In Cie presently proposed criterion the dvnamical (ime is included wilh a quantilalive measure., In the presently proposed criterion the dynamical time is included with a quantitative measure.
 On the AGB the mass loss process is (hat of pulsations coupled with radiation pressure on dust. while for the central stars of PNs il is mainly radiation pressure on ious.," On the AGB the mass loss process is that of pulsations coupled with radiation pressure on dust, while for the central stars of PNs it is mainly radiation pressure on ions."
 The idea is (hat the (transition is defined when (he dominate mass loss process switelies from pulsation and radiation pressure on dust (o raciation pressure on ions., The idea is that the transition is defined when the dominate mass loss process switches from pulsation and radiation pressure on dust to radiation pressure on ions.
 There are (wo main problems with this., There are two main problems with this.
 Firstly. the physics is nol well understood to connect this transition to stellar evolutionary codes.," Firstly, the physics is not well understood to connect this transition to stellar evolutionary codes."
 Secondly. interaction with a companion ean be the dominate mass loss mechanism in many post-AGD stars. either via tidal interaction or a common envelope.," Secondly, interaction with a companion can be the dominate mass loss mechanism in many post-AGB stars, either via tidal interaction or a common envelope."
 It is not clear what temperature to use., It is not clear what temperature to use.
 There is no natural temperature for any physical effect. although the transition occurs around an effective temperature ol (e.g. 5chónnberner 1981).," There is no `natural' temperature for any physical effect, although the transition occurs around an effective temperature of $T \simeq 5000 \K$ (e.g. Schönnberner 1981)."
 Even dust formation can cease at dillerent temperatures. depending on the metallicity of the envelope.," Even dust formation can cease at different temperatures, depending on the metallicity of the envelope."
 Vassiliadis Wood (1994) took (he transition to occur when the elfective temperature is (vice the minimum temperature the star can reach on the AGB. but no physical reason is given for that.," Vassiliadis Wood (1994) took the transition to occur when the effective temperature is twice the minimum temperature the star can reach on the AGB, but no physical reason is given for that."
 Contraction cannot be used because (he star starts to contract before it leaves the ACL., Contraction cannot be used because the star starts to contract before it leaves the AGB.
 The criterion of a rapid contraction. with time or with decreasing envelope mass. captures (he essence of (he (ransilion. but a quantitative value is not easy to define.," The criterion of a rapid contraction, with time or with decreasing envelope mass, captures the essence of the transition, but a quantitative value is not easy to define."
 One can use the logarithmic derivative of the stellar radius with envelope mass Bul 9 changes monotonically in the relevant temperature (radius) range. and it is not clear what value should be used. although 9=I might be a natural choice (Frankowski. A. 2007. private communication).," One can use the logarithmic derivative of the stellar radius with envelope mass But $\delta$ changes monotonically in the relevant temperature (radius) range, and it is not clear what value should be used, although $\delta=1$ might be a natural choice (Frankowski, A. 2007, private communication)."
 Alternatively. one can define the transition to occur when the magnitude of the second logarithmic derivative of the stellar radius with envelope mass," Alternatively, one can define the transition to occur when the magnitude of the second logarithmic derivative of the stellar radius with envelope mass"
The observed dependence of fm on µ is well parametrized as where fm is the major merger fraction (u>Umm= 1/4).,The observed dependence of $f_{\rm m}$ on $\mu$ is well parametrized as where $f_{\rm MM}$ is the major merger fraction $\mu \geq \mu_{\rm MM} = 1/4$ ).
 This dependence was predicted by the cosmological simulations of and used by in mass-selected spectro-photometric close pairs., This dependence was predicted by the cosmological simulations of and used by in mass-selected spectro-photometric close pairs.
" We set the value of fm to the observed one and used Generalized Least Squares (GLS) to estimate the power-law index s (see AppendixAppendixΑ:., for details)."," We set the value of $f_{\rm MM}$ to the observed one and used Generalized Least Squares (GLS) to estimate the power-law index $s$ (see Appendix\ref{mcfit}, for details)."
 The GLS fit to the Table 1 data yields s=—0.60+0.08 at z=0.8 and s=—1.02+0.13 at z=0.5., The GLS fit to the Table \ref{ffmutab} data yields $s = -0.60\pm 0.08$ at $z = 0.8$ and $s = -1.02 \pm 0.13$ at $z = 0.5$.
" To obtain a robust value of s at each redshift range under study, we determine s for different r5*."," To obtain a robust value of $s$ at each redshift range under study, we determine $s$ for different $r_{\rm p}^{\rm max}$."
 We summarize our results in Table 2 and show them in Fig. 3.., We summarize our results in Table \ref{srptab} and show them in Fig. \ref{svszfig}. .
" The values of s measured at nme=100/:! are representative of the median of all the values at different ερ, that are s=—0.59 at z=0.8 and s=—0.96 at z=0.5."," The values of $s$ measured at $r_{\rm p}^{\rm max} = 100h^{-1}$ are representative of the median of all the values at different $r_{\rm p}^{\rm max}$, that are $s = -0.59$ at $z = 0.8$ and $s = -0.96$ at $z = 0.5$."
" We find that the value of s decreases with cosmic time, reflecting a differential evolution in the merger fraction of major and minor companions."," We find that the value of $s$ decreases with cosmic time, reflecting a differential evolution in the merger fraction of major and minor companions."
" We checked that our incompleteness in the range z,2 (Sect. ??))"," We checked that our incompleteness in the range $z_{\rm r,2}$ (Sect. \ref{ncs}) )"
 does not bias our results with the following test., does not bias our results with the following test.
 We define a companion sample with Mg€—17.17—2.8z., We define a companion sample with $M_B \leq -17.17 - 2.8z$.
" This sample becomes artificially incomplete for companions with u>1/10 and u>1/5 at z20.2 and z20.65, respectively; that is, in our first redshift bin, and mimic the completeness behaviour of our companion sample at ζ:2."," This sample becomes artificially incomplete for companions with $\mu \geq 1/10$ and $\mu \geq 1/5$ at $z \geq 0.2$ and $z \geq 0.65$, respectively; that is, in our first redshift bin, and mimic the completeness behaviour of our companion sample at $z_{\rm r,2}$."
" Then, we repeat the previous analysis with the artificially incomplete sample, obtaining s=—0.99+0.08, which is similar to the original value measured in the complete sample."," Then, we repeat the previous analysis with the artificially incomplete sample, obtaining $s = -0.99\pm0.08$, which is similar to the original value measured in the complete sample."
 This implies that the weights Whag properly account for the missing faint companions and that the observed evolution of the index s with redshift in VVDS-Deep is a robust result., This implies that the weights $w_{\rm mag}^{k}$ properly account for the missing faint companions and that the observed evolution of the index $s$ with redshift in VVDS-Deep is a robust result.
 We also study how the luminosity function assumed in Whag determination affects the measured merger fractions., We also study how the luminosity function assumed in $w_{\rm mag}^{k}$ determination affects the measured merger fractions.
" We used the B—band luminosity functions from??;; and?,, finding a variation lower than 3% in the values of the mergerfraction for every r5 compared to our results."," We used the $B-$ band luminosity functions from; and, finding a variation lower than $3$ in the values of the mergerfraction for every $r_{\rm p}^{\rm max}$ compared to our results."
" Hence, assuming a different luminosity function would have only a limited impact on our results."," Hence, assuming a different luminosity function would have only a limited impact on our results."
" We then studied the dependency of the major merger fraction, fim, on the search radius."," We then studied the dependency of the major merger fraction, $f_{\rm MM}$, on the search radius."
 We summarize the fuw values for all 7* under study in Table 3 and show them in Fig. 4.., We summarize the $f_{\rm MM}$ values for all $r_{\rm p}^{\rm max}$ under study in Table \ref{f1tab} and show them in Fig. \ref{brpfig}.
 The value of fyw increases with the search radius and is well described in both redshift ranges by a power-law with index ᾳ=0.95+0.20., The value of $f_{\rm MM}$ increases with the search radius and is well described in both redshift ranges by a power-law with index $q = 0.95\pm0.20$.
" Regarding redshift evolution, the major merger fraction increases with redshift, in agreement with previous results in the literature222222)."," Regarding redshift evolution, the major merger fraction increases with redshift, in agreement with previous results in the literature."
. Westudy this evolution in more details in Sect. ??.., Westudy this evolution in more details in Sect. \ref{ffmmevol}. .
" We can estimate the minor-to-major merger fraction ratio, denoted fiji, as where Jw and mm are the luminosity ratios for major and minor mergers, respectively."," We can estimate the minor-to-major merger fraction ratio, denoted $f_{m/M}$, as where $\mu_{\rm MM}$ and $\mu_{\rm mm}$ are the luminosity ratios for major and minor mergers, respectively."
" This definition does not dependon the normalization of the merger fraction, that varies with n (Fig. 4))."," This definition does not dependon the normalization of the merger fraction, that varies with $r_{\rm p}^{\rm max}$ (Fig. \ref{brpfig}) )."
 We assume wm=1/4 and Umm= 1/10., We assume $\mu_{\rm MM} = 1/4$ and $\mu_{\rm mm} = 1/10$ .
" We find that η=0.73+0.13 at z= 0.8, and fuu= atz= 0.5."," We find that $f_{m/M} = 0.73 \pm 0.13$ at $z = 0.8$ , and $f_{m/M} = 1.55\pm0.30$ at$z = 0.5$ ."
" Therefore, minor companions become more numerous than major ones as one is going to lower"," Therefore, minor companions become more numerous than major ones as one is going to lower"
Grave&Ixumar(2009) discuss the implementation of this tool on a large sample of massive protostellar objects.,\citet{grave09} discuss the implementation of this tool on a large sample of massive protostellar objects.
 For fitting purposes a error was assumed on all fluxes., For fitting purposes a error was assumed on all fluxes.
" The weishted mean (weights being the inverse of 47) of each physical parameter was computed for all models that satisfied the criteria V7—V2,<3. where \7 is the statistical goodness of fit parameter measured per data point."," The weighted mean (weights being the inverse of $\chi^2$ ) of each physical parameter was computed for all models that satisfied the criteria ${\chi}^2 -{{\chi}^2}_{best} < 3$, where $\chi^2$ is the statistical goodness of fit parameter measured per data point."
 A rich cluster of stars can be seen in our A-band image (see Fig., A rich cluster of stars can be seen in our $K$ -band image (see Fig.
 D). the brightest stars in the center of the field coinciding with the peak of the MIPS 24 jmi contours.," 1), the brightest stars in the center of the field coinciding with the peak of the MIPS 24 $\mu$ m contours."
 The two straight lines in Fie., The two straight lines in Fig.
 1 mark (he slit positions used to obtain the spectra of the stars numbered 1 to 7., 1 mark the slit positions used to obtain the spectra of the stars numbered 1 to 7.
 Stars 1. 3 and 8 were modelled using 1 - 24 ji SEDs.," Stars 1, 3 and 8 were modelled using 1 - 24 $\mu$ m SEDs."
 Star 9 was not detected in the MIPS 24 yam band., Star 9 was not detected in the MIPS 24 $\mu$ m band.
 The brightest stars in the field are expected to be massive stars. while the fainter population represents the low mass members.," The brightest stars in the field are expected to be massive stars, while the fainter population represents the low mass members."
 In Fie., In Fig.
 3 (see the online electronic version of (his article for a colour plot) a colour composite of theSpi/zer infrared images is shown., 3 (see the online electronic version of this article for a colour plot) a colour composite of the infrared images is shown.
 The three-colour composite image was made using theSpitzer IRAC 3.6 pan. IRAC 8.0 yam and MIIPS 24 sau images coded as blue. green. and red. respectively.," The three-colour composite image was made using the IRAC 3.6 $\mu$ m, IRAC 8.0 $\mu$ m and MIPS 24 $\mu$ m images coded as blue, green and red, respectively."
 Notice the bipolar shape of the nebula and the FIR excess emission appearing as red., Notice the bipolar shape of the nebula and the FIR excess emission appearing as red.
 The stars 8 9 are associated with excess emission at 24 jim. as can be seen from the red colour surrounding these (wo stars.," The stars 8 9 are associated with excess emission at 24 $\mu$ m, as can be seen from the red colour surrounding these two stars."
 Also. the rich cluster of stars seen in Fig.," Also, the rich cluster of stars seen in Fig."
 1 is not visible in this composite image., 1 is not visible in this composite image.
 We will comment further on this issue in Sect., We will comment further on this issue in Sect.
 4., 4.
 In the following. we first discuss the photometric analvsis of the whole sample. then the IME of the cluster region. followed bv spectra of the seven representative sources (stus 1-7). and finally the SED fitting analvsis ol the brighter massive star candidates.," In the following, we first discuss the photometric analysis of the whole sample, then the IMF of the cluster region, followed by spectra of the seven representative sources (stars 1-7), and finally the SED fitting analysis of the brighter massive star candidates."
 The JHEN photometry of the point sources in the IRAS 19343+2026 region was used to construct colour-colour (CC) and CM cliagrams., The $JHK$ photometry of the point sources in the IRAS 19343+2026 region was used to construct colour-colour (CC) and CM diagrams.
 A combination of CC and CM diagrams made using various optical and infrared bands were then used to evaluate the membership of the cluster against a control field. ancl to evaluate the cluster properties.," A combination of CC and CM diagrams made using various optical and infrared bands were then used to evaluate the membership of the cluster against a control field, and to evaluate the cluster properties."
 The total number of point sources (wilh magnitude errors < 0.1 mag). detected in (he region shown in Fig.," The total number of point sources (with magnitude errors $<$ 0.1 mag), detected in the region shown in Fig."
 1 was 333. 688 and 375 in the J. Hand A bands. respectively.," 1 was 333, 688 and 875 in the $J$, $H$ and $K$ bands, respectively."
 307 stars are common to all three bands: GOT stars appear in the // and A bands., 307 stars are common to all three bands; 607 stars appear in the $H$ and $K$ bands.
 The ΕΕΤΤ 77A. photometric data.," The UFTI $JHK$ photometric data,"
les eewherelhecoe [ficient! is defined in Novikov Thorne (1973. herealter NT73). and the dimensionless «quantities are defined by—— and108m where nay=0.1 is adopted.,", where the coefficient$c_{2}$ is defined in Novikov Thorne (1973, hereafter NT73), and the dimensionless quantities are defined by, and m, where $\eta_{\rm eff}=0.1$ is adopted."
 The dimensionless scale-height of a disk ///HR is in principle a function of 2. ancl it reaches a maximal value in the inner region of the disk (Laor Netzer 1989).," The dimensionless scale-height of a disk $H/R$ is in principle a function of $R$, and it reaches a maximal value in the inner region of the disk (Laor Netzer 1989)."
 We adopt the maximal value of /7/ Rin the estimate of large-scale field strength Ώραat the disk surlace as done by Cao(2002b)., We adopt the maximal value of $H/R$ in the estimate of large-scale field strength $B_{\rm pd}$at the disk surface as done by \citet{cao02b}.
. As L99. the strength of the magnetic field produced by dynamo processes in the disk is given byODE where WV is the integrated shear stress of the disk. and / is the scale-height of the disk.," As L99, the strength of the magnetic field produced by dynamo processes in the disk is given by, where $W$ is the integrated shear stress of the disk, and $H$ is the scale-height of the disk."
 For a relativistic accretion disk. the integrated shear stress is given by Eq. (," For a relativistic accretion disk, the integrated shear stress is given by Eq. ("
5.6.14a) in NT13.,5.6.14a) in NT73.
 Equation (7)) can be re-written as Lr PPA g," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
 Equation (7)) can be re-written as Lr PPA ga," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
 Equation (7)) can be re-written as Lr PPA gau," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
 Equation (7)) can be re-written as Lr PPA gaus," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
 Equation (7)) can be re-written as Lr PPA gauss," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
 Equation (7)) can be re-written as Lr PPA gauss.," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"
We calculate the second order dervatives of JJ in order to examine the third. order behavior of the lens equation (4)).,We calculate the second order dervatives of $J$ in order to examine the third order behavior of the lens equation \ref{eqDLens}) ).
 For visual simplicity. we use the following denotations.," For visual simplicity, we use the following denotations."
 For the quasi-analvtic lenses we are interestedin here. J=1—> |&|7. and," For the quasi-analytic lenses we are interestedin here, $J = 1 - |\kappa|^2$ , and"
"the uncertainties, we have to further consider the nominal pointing uncertainty of the spacecraft.","the uncertainties, we have to further consider the nominal pointing uncertainty of the spacecraft."
 The uncertainty can be estimated from the distribution of aspect offset for a sample of point sources with accurately known celestialpositions?., The uncertainty can be estimated from the distribution of aspect offset for a sample of point sources with accurately known celestial.
. There is 68% of 70 sources imaged on ACIS-I have offsets smaller than ~0.4”.," There is $68\%$ of 70 sources imaged on ACIS-I have offsets smaller than $\sim0.4""$."
 We adopted this value as the astrometric uncertainty and added to the aforementioned quoted errors in quadrature for each coordinate., We adopted this value as the astrometric uncertainty and added to the aforementioned quoted errors in quadrature for each coordinate.
" This gives the resultant lo positional errors as dRA=1.18” and óDec-0.84""."," This gives the resultant $1\sigma$ positional errors as $\delta$ RA=1.18"" and $\delta$ Dec=0.84""."
 Although the search for the optical counterpart of $21 by Weisskopf et al. (, Although the search for the optical counterpart of S21 by Weisskopf et al. (
"2006) yield null-detection, we independently looked for any optical identification of this ray source in the United States Naval Observatory (USNC)-B1.0 catalogue (Monet et al.","2006) yield null-detection, we independently looked for any optical identification of this X-ray source in the United States Naval Observatory (USNC)-B1.0 catalogue (Monet et al."
 2003) and the Digitized Sky Survey with the improved X-ray position., 2003) and the Digitized Sky Survey with the improved X-ray position.
" Within our estimated 3e X-ray positional uncertainty, we cannot identify any optical counterpart of S21 down to the limiting magnitude of USNO-B1.0 catalogue (i.e. 21; cf."," Within our estimated $3\sigma$ X-ray positional uncertainty, we cannot identify any optical counterpart of S21 down to the limiting magnitude of USNO-B1.0 catalogue (i.e. 21; cf."
 Monet et al., Monet et al.
 2003)., 2003).
 This confirms the result reported by Weisskopf et al. (, This confirms the result reported by Weisskopf et al. (
2006).,2006).
" To examine the X-ray emission nature of ((S21), we make use of both and observations that cover it."," To examine the X-ray emission nature of (S21), we make use of both and observations that cover it."
 We extract the source counts from a circle with a radius of 15 arcsec and 20 arcsec in and datasets respectively., We extract the source counts from a circle with a radius of 15 arcsec and 20 arcsec in and datasets respectively.
" The extraction regions are chosen to optimize the signal-to-noise ratio which correspond to an encircled energy fraction of 5,90% and Z7596 at its location in the corresponding detectors in and respectively.", The extraction regions are chosen to optimize the signal-to-noise ratio which correspond to an encircled energy fraction of $\goa90\%$ and $\goa75\%$ at its location in the corresponding detectors in and respectively.
 The background spectrum is extracted from a nearby source-free circular region with a radius of 40 arcsec in the corresponding detectors., The background spectrum is extracted from a nearby source-free circular region with a radius of 40 arcsec in the corresponding detectors.
" After the background subtraction, there are 34+6 and 24+5 net source counts available from the and respectively."," After the background subtraction, there are $34\pm6$ and $24\pm5$ net source counts available from the and respectively."
 We compute the response files with the XMMSAS tasks RMFGEN and ARFGEN for and with the CIAO tools MKACISRMF and MKARF forChandra., We compute the response files with the XMMSAS tasks RMFGEN and ARFGEN for and with the CIAO tools MKACISRMF and MKARF for.
". With the aids of PIMMS, we can compare the count rates from different detectors."," With the aids of PIMMS, we can compare the count rates from different detectors."
 Adopting the best-fit spectral parameters (cf., Adopting the best-fit spectral parameters (cf.
" Tab. 1)),"," Tab. \ref{spec_par}) ),"
 we found the count rates obtained from different detectors are consistent., we found the count rates obtained from different detectors are consistent.
" However, as the parameters are poorly constrained, it is difficult to properly constrain the source variability."," However, as the parameters are poorly constrained, it is difficult to properly constrain the source variability."
" While the nominal observed flux is about ~2x10! erg cm? s~*, the lo upper limit is at the level of ~3x10? erg cm? s."," While the nominal observed flux is about $\sim2\times10^{-14}$ erg $^{-2}$ $^{-1}$, the $1\sigma$ upper limit is at the level of $\sim3\times10^{-12}$ erg $^{-2}$ $^{-1}$."
" Therefore, ascribing to the limited photon statistics of the existing data, we are not able to unambiguously conclude whether there is any flux variability fromJ202131."," Therefore, ascribing to the limited photon statistics of the existing data, we are not able to unambiguously conclude whether there is any flux variability from."
"0+402645.. To further investigate whether this source is a promising pulsar candidate, we examined its hardness ratio and constrain its properties by means of a color-color diagram with the combined net counts obtained from both satellites."," To further investigate whether this source is a promising pulsar candidate, we examined its hardness ratio and constrain its properties by means of a color-color diagram with the combined net counts obtained from both satellites."
 Following Elsner et al. (, Following Elsner et al. (
"2008), we used three energy bands in this analysis S (0.5—1 keV), M (1—2 keV) and H (2—8 keV).","2008), we used three energy bands in this analysis $S$ $0.5-1$ keV), $M$ $1-2$ keV) and $H$ $2-8$ keV)."
" Figure 5 shows the plot of (H—S)/T versus M/T, where T is the energy band 0.5—8 keV. The filled circle with the 1c error bars attached represents the location of iin this MMplot."," Figure \ref{ccd} shows the plot of $(H-S)/T$ versus $M/T$, where $T$ is the energy band $0.5-8$ keV. The filled circle with the $1\sigma$ error bars attached represents the location of in this plot."
 We have also computed the predicted values for a power-law spectrum with photon index varying from Τ-1 to [=6 for different values of hydrogen column absorption., We have also computed the predicted values for a power-law spectrum with photon index varying from $\Gamma=1$ to $\Gamma=6$ for different values of hydrogen column absorption.
 The results are plotted as the curves in Figure 5.., The results are plotted as the curves in Figure \ref{ccd}.
" By definition, all classes of X-ray sources should lie in the triangular boundary formed by S=MH0."," By definition, all classes of X-ray sources should lie in the triangular boundary formed by $S=M=H=0$."
" The soft sources which lie close to the line H=0 are most likely field stars in the Milky Way, and the hard sources lie close to the line S—0 are likely the background active galactic nuclei (AGNs) or pulsars with non-thermal dominant X-ray emission."," The soft sources which lie close to the line $H=0$ are most likely field stars in the Milky Way, and the hard sources lie close to the line $S=0$ are likely the background active galactic nuclei (AGNs) or pulsars with non-thermal dominant X-ray emission."
 iis marginally located at theright side of the color-color diagram yet close to the center., is marginally located at theright side of the color-color diagram yet close to the center.
 Its hardness is found to be too hard for a field star., Its hardness is found to be too hard for a field star.
" On the other hand, its location in the color-color diagram shows that it is modeled by a power-law with photon index generally larger than 3."," On the other hand, its location in the color-color diagram shows that it is modeled by a power-law with photon index generally larger than 3."
 This suggests that the X-ray emission of iis unlikely to be non-thermal dominant., This suggests that the X-ray emission of is unlikely to be non-thermal dominant.
 The same inference is obtained from the spectral analysis (see below)., The same inference is obtained from the spectral analysis (see below).
" Owing to the small numbers of the collected photons, we adopt the C-statistic (Cash 1979) for the spectral analysis which is insensitive to the binning."," Owing to the small numbers of the collected photons, we adopt the $C-$ statistic (Cash 1979) for the spectral analysis which is insensitive to the binning."
 The spectral analysis is performed with XSPEC 12.5 in the energy band of 0.3—10 keV and 0.5—8 keV for the data obtained from and respectively., The spectral analysis is performed with XSPEC 12.5 in the energy band of $0.3-10$ keV and $0.5-8$ keV for the data obtained from and respectively.
" In view of the small photon statistics, we limited our spectral analysis with simple single component model."," In view of the small photon statistics, we limited our spectral analysis with simple single component model."
 The best-fit parameters of all the tested models are tabulated in Table 1.., The best-fit parameters of all the tested models are tabulated in Table \ref{spec_par}. .
 All, All
accretion flow.,accretion flow.
" In particular. we adopt c, and ep computed using the Bernoulli function and Inass accretion rate for spherically svnunetric Bondi accretion with the PW potential."," In particular, we adopt $v_r$ and $\rho$ computed using the Bernoulli function and mass accretion rate for spherically symmetric Bondi accretion with the P–W potential."
" We set px = à22x10237 gem?.3 and specify.κ c4. through. RYm=Ry/Ry (note that Ry=$2265,fc.23 and Ry=€ i the Bondi radius)."," We set $\rho_\infty$ = $2.2 \times 10^{-23}$ $\rm{g/cm^3}$ and specify $c_{s,\infty}$ through $R'_S\equiv R_S/R_B$ (note that $R'_S=2c^2_{s,\infty}/c^2$, and $R_B=\frac{GM}{c_{s,\infty}^2}$ is the Bondi radius)."
 Thus Py characterizes the gas temperature in our simulations., Thus $R'_S$ characterizes the gas temperature in our simulations.
 We specify the initial conditions by adopting a non-zero / for the outer boundary ol the flow., We specify the initial conditions by adopting a non-zero $l$ for the outer boundary of the flow.
" We consider à case where (he angular momentum at the outer radius r, depends on the polar angle via We express (he angular momentum on (he equator as where . is the circularization radius on the equator in units of Ry for the Newtonian ⋅⋅ ↽↝⋅↽≽ ↴∏∐↲∣↽≻∪∏∐≺⇂≀↧↴↕⋅∡∖↽≺∢∪∐≼∐∐∪∐⋝∖⊽≀↧↴↕⋅≼↲⊳∖⊽↕↽≻≼↲"," We consider a case where the angular momentum at the outer radius $r_o$ depends on the polar angle via We express the angular momentum on the equator as where $R'_C$ is the circularization radius on the equator in units of $R_B$ for the Newtonian potential (i.e., $GM/r^2= v^2_\phi/r$ at $r= R'_C R_B$ )."
≺∢∐∎∐↲≼⇂≀↧↴⊳∖⊽↓⋟∪∐∪∖∖⇁⋟∖⊽⋅∶∖↥⊔∐↲↕↽≻∪↥≼↲⋟∖⇁⋅⊔⋅≼↲⋅⋅∣, The boundary conditions are specified as follows.
⇥∶∩≓≀↧↴∐≺⇂⊥⋖↽∖⋚∪≓ ∖∖⇁≼↲≀↧↴↕↽≻↕↽≻↥⋡∖↽≀↧↴∐≀↧↴⇀↸↕⋟∖⊽−∪↓⋟−⋟∖↕∖↽↕⊔∐∐↲↥↕⋅⋡∖↽∣↽≻∪∏∐≺⇂≀↧↴↕⋅∡∖↽≺∢∪∐≼∐∐∪∐⋅↼≚↥∣↽≻∪⊔↥⊔∐↲↕," At the poles, (i.e., $\theta=0^\circ$ and $180^\circ$ ), we apply an axis-of-symmetry boundary condition."
∐∐≼↲↕⋅≀↕↴↕∐⇂∪⋯≼↲↕⋅↕⋅≀↧↴≺∐≀↧↴↥ ∣↽≻∪∏∐≼⇂≀↧↴↕⋅↥≼↲⋝∖⊽⋅∖∖↽≼↲≀↧↴↕↽≻↕↽≻↥⋡∖↽≀↧↴∐∪∏⋯∪∖∖↽∣↽≻∪∏∐≼⇂≀↧↴↕⋅∡∖⇁≺∢∪∐≼∐∐∪∐↓⋟∪↕⋅≀↧↴∐≼⇂∡," At both the inner and outer radial boundaries, we apply an outflow boundary condition for all dynamical variables."
"∖↽∐≀↧↴∐↓↕≺∢≀↧↴↥∖↽≀↧↴↕⋅↕≀↧↴∣↽≻↥≼↲⊳∖⇁⋅↼≚⊳∖⊽↕↥ ↕↴∐∪⋮⋅↽⊰≀↧↪↥∪↕⋅≼↲↕↽≻↕⋅≼↲⋟∖⊽≼↲∐↥⋟∖⊽∩↲≀↧↴≼⇂⋡∖↽≺∢∪∐≼∐∐∪∐⋟∖⊽≀↧↴↥⊔∐↲∪∏∩↲↕⋅↕⋅⋯∐≀↧↴↥∣↽≻∪∏∐≺⇂≀↧↴↕⋅⋡∖⇁⋅≺⇂⋯⋅↕∐≸↽↔↴⊔∐↲≼↲∖↽∪↥∏∐∪∐∪↓⋟ each model we continue to apply the constraints that in the last zone in the radial direction. (eg—0. o,—fGr.ü)frsin9. and the density is fixed at the outer boundary at all times."," As in PB03a, to represent steady conditions at the outer radial boundary, during the evolution of each model we continue to apply the constraints that in the last zone in the radial direction, $v_\theta=0$, $v_\phi=l(r,\theta)/ r \sin{\theta}$, and the density is fixed at the outer boundary at all times."
" Note that we allow v, to float.", Note that we allow $v_r$ to float.
 To solve eqs.( 1))-( 3)), To solve eqs.( \ref{eq:con}) )-( \ref{eq:en}) )
 we use the ZEUS-2D code described by Stone Norman (1992). modified to implement the P.W potential.," we use the ZEUS-2D code described by Stone Norman (1992), modified to implement the P–W potential."
 llere we present results of ten simulations specified by four different. values of 5 (Le. y= 5/3. 4/3. 1.2. LOI). and three values of RG (ie. RY=10. 10!. and 10 31.," Here we present results of ten simulations specified by four different values of $\gamma$ (i.e., $\gamma$ = 5/3, 4/3, 1.2, 1.01), and three values of $R'_S$ (i.e., $R'_S = 10^{-5}$, $10^{-4}$, and $10^{-3}$ )."
 The simulations for + —1.2 were performed only for RG=107., The simulations for $\gamma=$ 1.2 were performed only for $R'_S = 10^{-3}$.
 Summary. of all runs is presented in Tab 1.., Summary of all runs is presented in Tab \ref{tab:1}.
" The table columns (1)-(8) show respectively the name of the run. the numerical radial resolution used in the simulation. the value of RY parameter. the value of cireularization radius in units of Bondi radius At. 5 adiabatic index. the end time at which we stopped each simulation /,; (the Gime is given in units of the dynamical time at the inner boundary /;,,2 595 s al r—1.5 RY for a mass of a black hole to be My,—3.6x10""M.. ). the"," The table columns (1)-(8) show respectively the name of the run, the numerical radial resolution used in the simulation, the value of $R'_S$ parameter, the value of circularization radius in units of Bondi radius $R'_C$, $\gamma$ adiabatic index, the end time at which we stopped each simulation $t_f$ (the time is given in units of the dynamical time at the inner boundary $t_{dyn}$ = 595 s at r=1.5 $R'_S$ for a mass of a black hole to be $M_{bh}=3.6\times10^6 \MSUN$ ), the"
The discovery of a hot Jupiter orbiting PPeg by ? by radial velocity (RV) measurements triggered the quest for extrasolar planets.,The discovery of a hot Jupiter orbiting Peg by \citet{1995Natur.378..355M} by radial velocity (RV) measurements triggered the quest for extrasolar planets.
 This breakthrough was only made possible through the usage of very precise wavelength calibration systems., This breakthrough was only made possible through the usage of very precise wavelength calibration systems.
 The Th-Ar emission lamp. used with the cross-correlation function (CCF) method (?).. and the I» cell. explored with the deconvolution procedure (?) were extensively used to find planets by RV.," The Th-Ar emission lamp, used with the cross-correlation function (CCF) method \citep{1996A&AS..119..373B}, and the $_2$ cell, explored with the deconvolution procedure \citep{1996PASP..108..500B} were extensively used to find planets by RV."
 Recent technological developments allowed for more precise spectrographs to be built. such as HARPS (2).. and reduction and analysis methods have been perfected through the years (?)..," Recent technological developments allowed for more precise spectrographs to be built, such as HARPS \citep{2003Msngr.114...20M}, and reduction and analysis methods have been perfected through the years \citep{2007A&A...468.1115L}."
 As of today. HARPS yields the most precise RV measurements. with sub-m/s precision. allowing for a succession of ground-breaking detections of the lightest planets known (??).. ," As of today, HARPS yields the most precise RV measurements, with sub-m/s precision, allowing for a succession of ground-breaking detections of the lightest planets known \citep{2006Natur.441..305L, 2009A&A...493..639M}."
Given the proven stability of these two well-established wavelength references. little investigation was made on other viable alternatives.," Given the proven stability of these two well-established wavelength references, little investigation was made on other viable alternatives."
 With time. RV expanded into the infra-red (IR) domain. where wavelength calibration is still in its infancy and no method has established itself as the paradigm.," With time, RV expanded into the infra-red (IR) domain, where wavelength calibration is still in its infancy and no method has established itself as the paradigm."
 In our attempt to measure RV with CRIRES at a very early stage of the instruments’ life. we used atmospheric lines as wavelength reference (?)..," In our attempt to measure RV with CRIRES at a very early stage of the instruments' life, we used atmospheric lines as wavelength reference \citep{2008A&A...489L...9H}."
 In a recent paper deseribing an improved data reduction (?) we reached a precision of mm/s over a time scale of one week., In a recent paper describing an improved data reduction \citep{2009arXiv0912.2643F} we reached a precision of m/s over a time scale of one week.
 This result was obtained in à RV standard star using CO» lines as wavelength reference., This result was obtained in a RV standard star using $_2$ lines as wavelength reference.
 The usage of telluric lines as a precise wavelength reference goes back to the first attempts of precise RV measurements. by ?.. on Arcturus and Procyon.," The usage of telluric lines as a precise wavelength reference goes back to the first attempts of precise RV measurements, by \cite{1973MNRAS.162..255G}, , on Arcturus and Procyon."
 Ten years later. the studies of ?.. 2.. and ?.. also explored the usage of atmospheric lines as a viable alternative for wavelength calibration.," Ten years later, the studies of \cite{1982A&A...114..357B}, \cite{1982ApJ...253..727S}, and \cite{1985A&A...149..357C}, also explored the usage of atmospheric lines as a viable alternative for wavelength calibration."
 Using Os lines. these authors showed back in the 80s that a precision of 5mmy/s was within reach.," Using $_2$ lines, these authors showed back in the 80's that a precision of m/s was within reach."
 The value is made even more relevant by the fact that they used different RV determination methods (different instrumentation. different line fitting approaches. etc.).," The value is made even more relevant by the fact that they used different RV determination methods (different instrumentation, different line fitting approaches, etc.)."
 Recently ? used the same principle on H>O lines and reached a precision of mm/s on UVES data., Recently \cite{2004MNRAS.353L...1S} used the same principle on $_2$ O lines and reached a precision of m/s on UVES data.
 In light of these results. two questions follow: In order to answer these two questions. we turned to HARPS archive data. now spanning 6 years.," In light of these results, two questions follow: In order to answer these two questions, we turned to HARPS archive data, now spanning 6 years."
 The high internal stability of HARPS leads to an unequalled precision in RV measurements., The high internal stability of HARPS leads to an unequalled precision in RV measurements.
 The RV variations of atmospheric lines can then be assessed against the very precise wavelength calibration provided by Th-Ar., The RV variations of atmospheric lines can then be assessed against the very precise wavelength calibration provided by Th-Ar.
 In this paper we answer the two previous questions and conclude on the suitability of atmospheric lines as a wavelength anchor., In this paper we answer the two previous questions and conclude on the suitability of atmospheric lines as a wavelength anchor.
 The paper is structured as follows., The paper is structured as follows.
 In Sect., In Sect.
 2 we describe HARPS anc the data-sets used in our investigation., $2$ we describe HARPS and the data-sets used in our investigation.
 Section 3 describes the principles of our method and data reduction., Section $3$ describes the principles of our method and data reduction.
 The results are presented 1n Sect., The results are presented in Sect.
 4 and discussed in Sect. 5., $4$ and discussed in Sect. $5$.
 We conclude in Sect., We conclude in Sect.
 6 with the lessons to learn from this campaign., $6$ with the lessons to learn from this campaign.
 HARPS (?) is a fiber-fed cross-dispersed echelle spectrograph installed at the 3.6m telescope in La Silla., HARPS \citep{2003Msngr.114...20M} is a fiber-fed cross-dispersed echelle spectrograph installed at the 3.6m telescope in La Silla.
 The main dispersion is provided by an R4 echelle grating in. Littrow configuration., The main dispersion is provided by an R4 echelle grating in Littrow configuration.
 The orders are then dispersed in a direction perpendicular to the dispersion direction. by a grism and imaged on a 2x2k4k CCD mosaic., The orders are then dispersed in a direction perpendicular to the dispersion direction by a grism and imaged on a $\times$ 2k4k CCD mosaic.
 This optical design creates 72 orders thàt span the whole optical range form 3785 to A., This optical design creates 72 orders that span the whole optical range form 3785 to $\AA$.
Thespectralresolutionwasmeasuredasbeingo f R-110000andthiei (820 m/s)., The spectral resolution was measured as being of $=$ 110 000 and the mean dispersion as of $\AA$ /pxl (820 m/s).
 The sampling ts of 3.3 pixel per resolution element., The sampling is of 3.3 pixel per resolution element.
 The instrument is located in à vacuum vessel to avoid spectral drift due to temperature and air pressure effects. which are kept below KK and mmbar. respectively.," The instrument is located in a vacuum vessel to avoid spectral drift due to temperature and air pressure effects, which are kept below K and mbar, respectively."
 A Th-Ar emission lamp is used for wavelength calibration., A Th-Ar emission lamp is used for wavelength calibration.
 The very high stability of HARPS allows for a precision of mm/s to be reached routinely., The very high stability of HARPS allows for a precision of m/s to be reached routinely.
 When a precision better than mm/s is required. a second channel can be used to image the Th-Ar simultaneously with the science target.," When a precision better than m/s is required, a second channel can be used to image the Th-Ar simultaneously with the science target."
 HARPS proven intrinsic IP. stability permits to study spectral lines. profile variations as well. which can bedone using the well-know," HARPS proven intrinsic IP stability permits to study spectral lines profile variations as well, which can bedone using the well-know"
"and in turn. that 54=#03dxra) where for small enough5 ZZ. and ܣ,.","and in turn, that $\gamma_{\perp} = \pm (\beta \pm
 i\alpha )$ where for small enough $\Enu$ and $\Eeta$."
 The solutions. which must be bounded. can then be written out as where €2 is either of the quantities S. 6. L or c. and ( and (Quy are integration constants.," The solutions, which must be bounded, can then be written out as where $Q$ is either of the quantities $S$ , $b$, $L$ or $\psi$, and $Q_c$ and $Q_0$ are integration constants."
 Note that the boundary laver thickness ο=1/2 is. as expected. the same as the one obtained by MacCregor Charbonneau (1999) in their open-field calculations.," Note that the boundary layer thickness $\delta_{\perp} = 1/\beta$ is, as expected, the same as the one obtained by MacGregor Charbonneau (1999) in their open-field calculations."
 In order to compare rigorously the simulations to the analvtical solutions derived. above. the matching of the solutions obtained in the boundary [aver to those in the bulk of the Γιά should. be performed.," In order to compare rigorously the simulations to the analytical solutions derived above, the matching of the solutions obtained in the boundary layer to those in the bulk of the fluid should be performed."
 However. the solution in the bulk of the Εις. in particular in the polar regions. is dominated by geometric ellects (as the latitucinal derivatives are not necessarily negligible) anc dilfusive ellects (as the Ekman numbers used in the simulations are not small enough to justify neglecting the cdilfusive terms): as à result. it is beyond to scope of this analysisto derive," However, the solution in the bulk of the fluid, in particular in the polar regions, is dominated by geometric effects (as the latitudinal derivatives are not necessarily negligible) and diffusive effects (as the Ekman numbers used in the simulations are not small enough to justify neglecting the diffusive terms); as a result, it is beyond to scope of this analysisto derive"
outflow. should be observed with millimeter/submillimeter interferometers. at scales of a lew arcseconds. to be able to discriminate the possible individual outllows driven by these sources.,"outflow should be observed with millimeter/submillimeter interferometers, at scales of a few arcseconds, to be able to discriminate the possible individual outflows driven by these sources."
 Being able to trace the molecular outflow close to its driving source(s) would provide direct evidence of whether there is a superposition of outIlows or if one source is responsible for it., Being able to trace the molecular outflow close to its driving source(s) would provide direct evidence of whether there is a superposition of outflows or if one source is responsible for it.
 We presented new interferometric NIL;CI.1). ancl (2.2) observations of the region of high-mass star formation ARGL 487. with the aim to understand the nature of the low-collimation bipolar molecular outllow: previously detected in CO (CGoémmez οἱ al.," We presented new interferometric $_3$ (1,1) and (2,2) observations of the region of high-mass star formation AFGL 437, with the aim to understand the nature of the low-collimation bipolar molecular outflow previously detected in CO (Gómmez et al."
 1992)., 1992).
 These observations were complemented with archive data of radio continuum emission at 2 em. 3.6 em. and 450 p/m. Our main conclusions are as follow: GAL JPG and ld: acknowledge partial support. from Alinisterio de Ciencia e| Innovaciónn (Spain). erant AYA2008-06189-CO3-01.," These observations were complemented with archive data of radio continuum emission at 2 cm, 3.6 cm, and 450 $\mu$ m. Our main conclusions are as follow: GM, JFG and IdG acknowledge partial support from Ministerio de Ciencia e Innovaciónn (Spain), grant AYA2008-06189-C03-01."
 JEG is also supported by Junta de Andaluctaa CE1C-126)., JFG is also supported by Junta de a (TIC-126).
 This research used the facilities of the Canadian Astronomy Data Centre operated by the the National Research Council of Canada with the support of the Canaclian Space Agency., This research used the facilities of the Canadian Astronomy Data Centre operated by the the National Research Council of Canada with the support of the Canadian Space Agency.
with svuthetic colors of hydrogeu-deficicut stars. Evres et al. (,"with synthetic colors of hydrogen-deficient stars, Eyres et al. ("
1998) derive 1.15 from the observed versus expected II? flux.,1998) derive 1.15 from the observed versus expected $\beta$ flux.
 Pollacco (1999). using high. S/N spectra of the planetary nebula. derived 0.71+0.09 from the Balmer decrement.," Pollacco (1999), using high S/N spectra of the planetary nebula, derived $0.71 \pm 0.09$ from the Balmer decrement."
 Iunesweueer lIkerber (1998) determined the interstellar extinction as a function of distance. and arive at a value zx0.8 for distances above 1 kpc.," Kimeswenger Kerber (1998) determined the interstellar extinction as a function of distance, and arrive at a value $\approx 0.8$ for distances above 1 kpc."
 We assune in this paper extinction values = 0.53 and 0.7. aud check which one vields better agreement between computed SEDs aud observations. dereddened for the above values.," We assume in this paper extinction values = 0.53 and 0.7, and check which one yields better agreement between computed SEDs and observations, dereddened for the above values."
 In 1998 and 1999. increasing amounts of circtunstellar dust also modified the observed SED of the ceutral object (Iipper 1999).," In 1998 and 1999, increasing amounts of circumstellar dust also modified the observed SED of the central object (Kipper 1999)."
 At the time of our observation. the influence of eirciuustellar extinction appears to be marginal (sce also section 1).," At the time of our observation, the influence of circumstellar extinction appears to be marginal (see also section 4)."
 Opacity sapling model atiiospheres of differentζω. logg aud abundances are computed with the wroeral SAMOII (Pavlenko 1999).," Opacity sampling model atmospheres of different, $\log\, g$ and abundances are computed with the program SAM941 (Pavlenko 1999)."
 Chemical abundances derived bv Asplund et al. (, Chemical abundances derived by Asplund et al. (
1997) are used as “normal input” for VI331 Ser.,1997) are used as “normal input” for V4334 Sgr.
 Asplund et. al. (, Asplund et al. (
1999) derived abundauces for Octyer 1996 which are probably best to be adopted in the absence of estimates for subsequent dates.,"1999) derived abundances for October 1996, which are probably best to be adopted in the absence of estimates for subsequent dates."
 Nevertheless. we do uot expect our results to be crucially affected by abundances which ciffer from those of Aspluud et al. (," Nevertheless, we do not expect our results to be crucially affected by abundances which differ from those of Asplund et al. ("
1997).,1997).
 We varied a few abundances to study. the nupact of abundance changes on the emitted spectrin., We varied a few abundances to study the impact of abundance changes on the emitted spectrum.
 Results are given in section 3.1., Results are given in section 3.4.
 The jonization-dissociation equilibriun (IDE) was calculated for a μάς of 70 atoms. dons and cdiatonuc molecules.," The ionization-dissociation equilibrium (IDE) was calculated for a mix of 70 atoms, ions and diatomic molecules."
 Constauts for IDE computations were taken mainly from Tsuji (1973)., Constants for IDE computations were taken mainly from Tsuji (1973).
 Absorption by atoms aud ions as well as absorption in frequencies of 20 band systems of diatonuc molecules were taken into account (Table ], Absorption by atoms and ions as well as absorption in frequencies of 20 band systems of diatomic molecules were taken into account (Table 1).
 The atomic line list was taken from the VALD database (Piskunov etal., The atomic line list was taken from the VALD database (Piskunov etal.
 1995). includiug lines of s-process elements which are prescut in the spectrum of V1331 Ser," 1995), including lines of s-process elements which are present in the spectrum of V4334 Sgr."
 Molecular opacities were comped in the approach of he just overlapping approximation (JOLÀ)., Molecular opacities were computed in the approach of the just overlapping approximation (JOLA).
 The JOLA approach ijs based on the asstuption hat the neal Ine separation within a molecular baud is comparable or sinaller than the line widths., The JOLA approach is based on the assumption that the mean line separation within a molecular band is comparable or smaller than the line widths.
 By definition. JOLA overestimates molecular absorption produced by wea- nolecular bauds.," By definition, JOLA overestimates molecular absorption produced by weak molecular bands."
 For stroug (saturated) molecular bands. he results appear to be satisfactorv (Pavleunko 1997).," For strong (saturated) molecular bands, the results appear to be satisfactory (Pavlenko 1997)."
 We used the DIGEL program of Yaremchuk (Nersisvan et al., We used the BIGF1 program of Yaremchuk (Nersisyan et al.
 1987) which realizes the method of Nameushikov et al. (, 1987) which realizes the method of Kamenshikov et al. (
1971).,1971).
 Yaremchuk’s approach takes iuto account the splitting of molecular bands on the P-O-R or P-R 1xanches for svstems with A=1.0. respectively.," Yaremchuk's approach takes into account the splitting of molecular bands on the P-Q-R or P-R branches for systems with $\Lambda =1, 0$, respectively."
 The computations were carried out fora set of vibrational quanti ος 0xv.9 (see Abia et al.," The computations were carried out for a set of vibrational quantum numbers $0 \leq v^{\prime}, 
v^{\prime\prime} \leq 9 $ (see Abia et al."
 1999 aud Pavleuko Yakovina 1999 for more details)., 1999 and Pavlenko Yakovina 1999 for more details).
 Couvection was computed using the ATLASS (Ixurucz 1993) scheme with 7/77 , Convection was computed using the ATLAS9 (Kurucz 1993) scheme with $l/H = 1.6$.
Convective overshooting was not considered., Convective overshooting was not considered.
 Our uwmuerical experiments show that the impact of convective overshooting ou the temperature structure of the model atinosphieres is rather weak., Our numerical experiments show that the impact of convective overshooting on the temperature structure of the model atmospheres is rather weak.
 The main opacity sources in the atinosphiere of V1331 Ser differ ποια those iun atnmiospheres with solar-like almucdauces: To take iuto account both nuon-solu abuudauces aud ypacity sources. we computed tables of Rosseliux opacities 7... for a exid of temperatures T and pressures P.," The main opacity sources in the atmosphere of V4334 Sgr differ from those in atmospheres with solar-like abundances: To take into account both non-solar abundances and opacity sources, we computed tables of Rosseland opacities $\tau_{\rm ross}$ for a grid of temperatures $T$ and pressures $P$."
 The table Tos=FOLP) was used for the temperature COYICCion iu SAMO.," The table $\tau_{\rm ross} = 
f (T,P)$ was used for the temperature correction in SAM941."
 We used the same erid of opacity sources as in model atuosphere computations. 1.6. continuous. molecular baud and atomic line absorptions.," We used the same grid of opacity sources as in model atmosphere computations, i.e. continuous, molecular band and atomic line absorptions."
 This procedure is important since the photosplere of VI331I Ser les at a different pressure height that in the case of solar abundances (cf., This procedure is important since the photosphere of V4334 Sgr lies at a different pressure height that in the case of solar abundances (cf.
 the computatious for R. CrB. DPoavleunko 1999).," the computations for R CrB, Pavlenko 1999)."
iis likely to be significantly larger than the typical neutron star mass (1.35M. 0 obtained from measurements in double neutron star binaries. given that this system has evolved through a long (~ Gyr) LMXB phase with sub-Eddington mass transfer (see. ?)..,"is likely to be significantly larger than the typical neutron star mass $1.35{\rm \,M}_{\odot}$ ) obtained from measurements in double neutron star binaries, given that this system has evolved through a long $\sim$ Gyr) LMXB phase with sub-Eddington mass transfer \citep[see,
 e.g.][]{prp02}."
 Recent work on the mass determination of the original black widow pulsar (PSR BI957-20) by ὁ confirms that the pulsars in these systems can accrete significant amounts of matter., Recent work on the mass determination of the original black widow pulsar (PSR B1957+20) by \citet{vbk11} confirms that the pulsars in these systems can accrete significant amounts of matter.
 For a discussion of the effect of irradiation on the accretion efficiency in LMXBs. see ?..," For a discussion of the effect of irradiation on the accretion efficiency in LMXBs, see \citet{rit08}."
 The Roche-lobe radius of the companion star in echanges slightly with the estimated stellar masses and is found to be RL=0.150.011...," The Roche-lobe radius of the companion star in changes slightly with the estimated stellar masses and is found to be $R_{\rm L} = 0.15 \pm 0.01{\rm
 \,R}_{\odot}$."
 However. the size (radius) of the irradiated companion star is difficult to determine accurately and in the discussion further on we shall assume two different values for the tilling factor (the ratio of the volume-equivalent radius of the companion star to its Roche lobe) of 0.43 and 0.95 (?).. corresponding to stellar radii of about 0.064 R.. and 0.14 . respectively.," However, the size (radius) of the irradiated companion star is difficult to determine accurately and in the discussion further on we shall assume two different values for the filling factor (the ratio of the volume-equivalent radius of the companion star to its Roche lobe) of 0.43 and 0.95 \citep{svbk01b}, , corresponding to stellar radii of about $0.064\,$ $_{\odot}$ and $0.14\,$ $_{\odot}$ respectively."
" In order to monitor their variations over time. values for £7, and wr were derived for each year of data. where three months of overlap were kept between adjacent years."," In order to monitor their variations over time, values for $P_{\rm b}$ and $x$ were derived for each year of data, where three months of overlap were kept between adjacent years."
" In doing so. all model parameters besides .r. £4, and 7c were held tixed and the timing reference epoch was detined to be the centre of each year-long interval."," In doing so, all model parameters besides $x$, $P_{\rm b}$ and $T_{\rm ASC}$ were held fixed and the timing reference epoch was defined to be the centre of each year-long interval."
 The fractional changes of these measurements are shown in Figure 5.., The fractional changes of these measurements are shown in Figure \ref{fig:period}.
" Clearly. tive different epochs can be identified. in which variations of both £4, and .r can be described using only a linear trend. as shown in the figure."," Clearly, five different epochs can be identified, in which variations of both $P_{\rm b}$ and $x$ can be described using only a linear trend, as shown in the figure."
" a, is the observable rate of change of the orbital period and is caused by a variety of effects. both intrinsic to the system and caused by kinematic effects relative to the observer."," $\dot{P_{\rm b}}$ is the observable rate of change of the orbital period and is caused by a variety of effects, both intrinsic to the system and caused by kinematic effects relative to the observer."
" The most important contributions are: As a point of reference. the values for £4, in the two most extreme epochs are 47,—1.51(3)190 Hoand 10 ."," The most important contributions are: As a point of reference, the values for $\dot{P}_{\rm b}$ in the two most extreme epochs are $\dot{P_{\rm b}} = -1.81(3) \times 10^{-11}$ and $\dot{P_{\rm b}} = 1.8(3) \times 10^{-11}$ ."
" The first. term. /4,5QW.—. is the contribution. due to gravitational. wave emission."," The first term, $\dot{P_{\rm b}}^{\rm GW}$, is the contribution due to gravitational wave emission."
 In general relativity. for circular orbits it is given The mass ratio q=36 has been calculated assuming a pulsar mass ny= 1.5MM. and a companion mass nm.=0.05 MM. for an inclination angle of¢=40°.," In general relativity, for circular orbits it is given by \citep{pet64}: The mass ratio $q = 36$ has been calculated assuming a pulsar mass $m_{\rm
 p}=1.8$ $_{\odot}$ and a companion mass $m_{\rm c}=0.05$ $_{\odot}$ for an inclination angle of $i=40^{\circ}$."
" For wwe tind £4,GN~7.5.101."," For we find $\dot{P_{\rm
 b}}^{\rm GW} \simeq -7.5 \times 10^{-14}$."
 Thisvalue is about three orders of magnitude less than the observed value of 73., Thisvalue is about three orders of magnitude less than the observed value of $\dot{P_{\rm b}}$.
 The second term. Bo. is the Doppler correction. which is the combined effect of the proper motion of the system (2) and a correction term for the Galactic acceleration.," The second term, $\dot{P_{\rm b}}^{\rm D}$ , is the Doppler correction, which is the combined effect of the proper motion of the system \citep{shk70} and a correction term for the Galactic acceleration."
 The contribution for the Galactic acceleration at the location of05927. BSan is of order 1.110.4% (2).," The contribution for the Galactic acceleration at the location of, $\dot{P_{\rm b}}^{\rm Gal}$, is of order $1.1 \times 10^{-15}$ \citep{lwj+09}."
 Using numbers from Table 2.. we also calculate the «contribution due to the Shklovskii effect according to the following: By summing we yield the Doppler correction: four orders of magnitude smaller than the measured value.," Using numbers from Table \ref{tab:par}, we also calculate the contribution due to the Shklovskii effect according to the following: By summing we yield the Doppler correction: four orders of magnitude smaller than the measured value."
 An acceleration of the binary system with respect to the Solar System Barycentre (SSB) could also be caused by a third massive body orbiting the binary system., An acceleration of the binary system with respect to the Solar System Barycentre (SSB) could also be caused by a third massive body orbiting the binary system.
 However. it would affect the orbital period derivative andthe spin period derivative in the same way.," However, it would affect the orbital period derivative andthe spin period derivative in the same way."
" Assuming that the spin period derivative is caused entirely by the acceleration. wecan estimate the maximal effect this would have on D: (EN""LEVEDc3.10Pss +."," Assuming that the spin period derivative is caused entirely by the acceleration, wecan estimate the maximal effect this would have on $\dot{P}_{\rm b}$ : $(\dot{P}_{\rm b}/P_{\rm b})^{\rm
 acc}=(\dot{P}/P)\simeq -3 \times 10^{-18}$ $^{-1}$ ."
 This is, This is
about50%.. which leaves some residual instrumental response of the order of[,"about, which leaves some residual instrumental response of the order of."
σοι Our sensitivity limit on average (3 x rms) is 1.2 mK Ty (or 7.3 mJy/beam)., Our sensitivity limit on average (3 $\times$ rms) is 1.2 mK $\rm T_{B}$ (or 7.3 mJy/beam).
 Only strong sources exceeding about 0.7 Jy will show weak polarization response of instrumental origin. which will. however. in most cases confuse with diffuse Galactic emission.," Only strong sources exceeding about 0.7 Jy will show weak polarization response of instrumental origin, which will, however, in most cases confuse with diffuse Galactic emission."
 Compact sources of the NVSS catalogue (?) were used to check the position accuracy., Compact sources of the NVSS catalogue \citep{Condon98} were used to check the position accuracy.
" Finally. all individually edited maps were combined by applying the ""PLAIT-algorithm (?).. where the Fourier transforms of the maps were added and the final map is obtained by an inverse Fourier. transform. “"," Finally, all individually edited maps were combined by applying the “PLAIT”-algorithm \citep{Emerson88}, where the Fourier transforms of the maps were added and the final map is obtained by an inverse Fourier transform. “"
PLAIT™. in addition. is able to suppress remaining low-level scanning effects still visible in a few individual maps.,"PLAIT”, in addition, is able to suppress remaining low-level scanning effects still visible in a few individual maps."
 We have observed scans of up to 10° in length aiming to recover extended structures as large as possible., We have observed scans of up to $\degr$ in length aiming to recover extended structures as large as possible.
 All maps have a relative zero-level by arbitrarily setting the edge values of each scan to zero., All maps have a relative zero-level by arbitrarily setting the edge values of each scan to zero.
 Total intensity maps (Stokes /) always miss a positive temperature offset. while the offsets for Stokes U anc Q maps may be positive or negative.," Total intensity maps (Stokes $I$ ) always miss a positive temperature offset, while the offsets for Stokes $U$ and $Q$ maps may be positive or negative."
 Polarized intensity. P7. of unknown intensity originating from Faraday rotated diffuse emission im the interstellar medium may exist everywhere and is not related to total intensity.," Polarized intensity, $PI$, of unknown intensity originating from Faraday rotated diffuse emission in the interstellar medium may exist everywhere and is not related to total intensity."
 Thus the true zero-level of the observed U and Q maps remains unknown., Thus the true zero-level of the observed $U$ and $Q$ maps remains unknown.
 Therefore P/ anc the polarization angle. PA. as calculated from U and Q. neec to be corrected as well.," Therefore $PI$ and the polarization angle, $PA$, as calculated from $U$ and $Q$, need to be corrected as well."
 We note that relative polarization zero-level setting may be done in different ways. e.g. setting the mean value of U and Q of each scan to zero (?)..," We note that relative polarization zero-level setting may be done in different ways, e.g. setting the mean value of $U$ and $Q$ of each scan to zero \citep{Junkes87}."
 After we combined maps observed along £ direction withmaps along 5 direction. the edge areas of the final combined maps differ from zero.," After we combined maps observed along $\ell$ direction withmaps along $b$ direction, the edge areas of the final combined maps differ from zero."
 Since polarized components are vectors. a missing large-scale component may lead to a misinterpretation of the observed data (?)..," Since polarized components are vectors, a missing large-scale component may lead to a misinterpretation of the observed data \citep{Reich06}. ."
 This is in particular important for polarized emission resulting from Faraday rotation. which clearly dominates the Galactic polarization maps at.t6 cem.," This is in particular important for polarized emission resulting from Faraday rotation, which clearly dominates the Galactic polarization maps at $\lambda$ cm."
 In Paper I. ? already presented a solution for this problem by adopting the three-year K-band (22.8 GHz) polarization data from WMAP (?).. which have a correct zero-level.," In Paper I, \citet{Sun07} already presented a solution for this problem by adopting the three-year K-band (22.8 GHz) polarization data from WMAP \citep{Page07}, which have a correct zero-level."
 Missing large-scale U and Q emission at 26 cem is restored by scaling the K-band data by a factor of (4.8/22.8)7. according to a temperature spectral index of 8=—2.8.," Missing large-scale $U$ and $Q$ emission at $\lambda$ cm is restored by scaling the K-band data by a factor of $(4.8/22.8)^{-2.8}$, according to a temperature spectral index of $\beta = -2.8$."
 This procedure also assumes that the RM of the diffuse emission is not significant., This procedure also assumes that the $RM$ of the diffuse emission is not significant.
 For this second much larger section of the 26 cem polarization survey we slightly modified the method applied in Paper I by taking meanwhile available additional information into account., For this second much larger section of the $\lambda$ cm polarization survey we slightly modified the method applied in Paper I by taking meanwhile available additional information into account.
 We now use the five-year release of the WMAP observations (?).., We now use the five-year release of the WMAP observations \citep{Hinshaw09}.
" We calculated the spectral index distribution between the polarized emission at 1.4 GHz (?) and the K-band data for the entire survey section, smoothed to à common angular resolution of 27."," We calculated the spectral index distribution between the polarized emission at 1.4 GHz \citep{Wolleben06} and the K-band data for the entire survey section, smoothed to a common angular resolution of $\degr$."
 We obtained a mean spectral index of p=-2.92+0.25., We obtained a mean spectral index of $\beta = -2.92\pm0.25$.
 We note that this spectral index ts largely biased by the dominating polarized emission from the bright Fan-region. which is Faraday thin at 1.4 GHz.," We note that this spectral index is largely biased by the dominating polarized emission from the bright Fan-region, which is Faraday thin at 1.4 GHz."
 This. however. is likely not the case for the Galactic plane emission at 1.4 GHz from large distances.," This, however, is likely not the case for the Galactic plane emission at 1.4 GHz from large distances."
 Current estimates of the synchrotron total intensity spectrum quote very similar spectral values between 1.4 GHz and 23.8 GHz (see ? for a recent discussion). which we expect to be valid for the extrapolation of Faraday thin diffuse large-scale polarized emission from 22.8 GHz to 4.8 GHz as well.," Current estimates of the synchrotron total intensity spectrum quote very similar spectral values between 1.4 GHz and 23.8 GHz (see \citet{Dickinson09} for a recent discussion), which we expect to be valid for the extrapolation of Faraday thin diffuse large-scale polarized emission from 22.8 GHz to 4.8 GHz as well."
 We compared the WMAP K-band (22.8 GHz) and Ka-band (33 GHz) polarization data (?) at 2° angular resolution for common extended polarization structuresin the present survey area., We compared the WMAP K-band (22.8 GHz) and Ka-band (33 GHz) polarization data \citep{Hinshaw09} at $\degr$ angular resolution for common extended polarization structuresin the present survey area.
 Clearly. the vast majority of patchy. weak polarization features in the two WMAP maps were not correlated and thus do not show patches of polarized emission.," Clearly, the vast majority of patchy, weak polarization features in the two WMAP maps were not correlated and thus do not show patches of polarized emission."
 This i turn means that an extrapolation of the polarized K-band emission towards 26 cem becomes questionable as it might introduce spurious features specific to the K-band map., This in turn means that an extrapolation of the polarized K-band emission towards $\lambda$ cm becomes questionable as it might introduce spurious features specific to the K-band map.
 Large-scale polarization gradients. however. are common in the K-band and Ka-band maps.," Large-scale polarization gradients, however, are common in the K-band and Ka-band maps."
 We therefore decided to convolve the 6 cem U and ο survey maps and the corresponding K-band maps to 2° angular resolution after having removed a few strong anc compact polarized sources. The convolved maps were split into sections. scaled by a factor of (4.8/22.8)7? and the difference values in their corner areas were determined.," We therefore decided to convolve the $\lambda$ cm $U$ and $Q$ survey maps and the corresponding K-band maps to $\degr$ angular resolution after having removed a few strong and compact polarized sources, The convolved maps were split into sections, scaled by a factor of $(4.8/22.8)^{-2.9}$ and the difference values in their corner areas were determined."
 These difference values were used to define correction hyper planes in UÜ and Q for each 16 cem survey section and were applied to the data at their original resolution., These difference values were used to define correction hyper planes in $U$ and $Q$ for each $\lambda$ cm survey section and were applied to the data at their original resolution.
 In Table 2 we list the LU and Q intensitiesof the Urumqi observations and the corresponding scaled K-map values together with the resulting correction values., In Table 2 we list the $U$ and $Q$ intensitiesof the Urumqi observations and the corresponding scaled K-map values together with the resulting correction values.
 The maximum error introduced by assuming a constant spectral index ofB=—2.9 will occur at £=129° and is estimated to be about +1.5 mK Ty in case the assumed spectral index varies by AB= x0.1., The maximum error introduced by assuming a constant spectral index of $\beta = -2.9$ will occur at $\ell = 129\degr$ and is estimated to be about $\pm$ 1.5 mK $\rm T_{B}$ in case the assumed spectral index varies by $\Delta\beta = \pm$ 0.1.
 A significant RM will change the extrapolated corrections for U and Q. while PI remains unchanged.," A significant $RM$ will change the extrapolated corrections for $U$ and $Q$ , while $PI$ remains unchanged."
 Numerous RMs from extragalactic sources in the Galactic plane are available (?).., Numerous $RM$ s from extragalactic sources in the Galactic plane are available \citep{Brown07}. .
 On average high RM-values are observed in thesurveyed area with a clear gradient along £. but also a significant scatter," On average high $RM$ -values are observed in thesurveyed area with a clear gradient along $\ell$ , but also a significant scatter"
correct picture due to the lack of such hydrogen masses in DB stars.,correct picture due to the lack of such hydrogen masses in DB stars.
 Gravitational accretion of ISAT particles by a star of mass AZ and radius A. in the supersonic regime follows the mathematical form (Aleock&Larionov1980) where Ay is the accretion radius. s=Ver|es? for relative stellar velocity.  and ambient sounc velocity. ὃς. while py is the unperturbecl density of material being accreted.," Gravitational accretion of ISM particles by a star of mass $M$ and radius $R$, in the supersonic regime follows the mathematical form \citep{alc80}
 where $R_A$ is the accretion radius, $s=\sqrt{v^2 + {c_s}^2}$ for relative stellar velocity $v$ and ambient sound velocity $c_s$, while $\rho_{\infty}$ is the unperturbed density of material being accreted."
 The accretion radius is defined as and is often referred to the as the Bondi or Boncli-Llovle racius., The accretion radius is defined as and is often referred to the as the Bondi or Bondi-Hoyle radius.
 For all possible speeds considered here. £245A. and the accretion rate is given by Or This is the Edcington rate (LEddington19960)... the accretion induced. on non-interacting particles by the ecometrical-eravitational cross section of the star às it travels through the ISM.," For all possible speeds considered here, $R_A\gg R$, and the accretion rate is given by or This is the Eddington rate \citep{edd26}, the accretion induced on non-interacting particles by the geometrical-gravitational cross section of the star as it travels through the ISM."
" 3ondi-lHlovle theory νο, including gas pressure: IExlgar demonstrates that the elective. cross section in Equation 4.0. zH4? becomes zd in the [uid dynamical limit (Bondi1952).. vielding an acerction rate for interacting particles This mass infall rate represents the maximal. idealized case where the mean free path of the particles is. such that collisions are important. ancl transverse momentun is ellectively destrovecd: clownstream from the star."," Bondi-Hoyle theory (i.e. including gas pressure; \citealt{edg04}) ) demonstrates that the effective cross section in Equation \ref{eqn4}, $\pi R_AR$ becomes $\pi {R_A}^2$ in the fluid dynamical limit \citep{bon52}, yielding an accretion rate for interacting particles This mass infall rate represents the maximal, idealized case where the mean free path of the particles is such that collisions are important and transverse momentum is effectively destroyed downstream from the star."
 Lt is also physically unrealistic in perhaps all situations excepting an ionizecl plasma. as it assumes no net angular momentum between the accreting star andits surrounding mecdiuni (IxoesterLOTG).. and it certainly does not apply to neutral atoms or large particles (Alcock&Larionov1980).," It is also physically unrealistic in perhaps all situations excepting an ionized plasma, as it assumes no net angular momentum between the accreting star andits surrounding medium \citep{koe76}, and it certainly does not apply to neutral atoms or large particles \citep{alc80}."
. “Phe ratio of the Boneli-Llovle to Eddington accretion rates is ο eo>6s or around. LO! for tvpical. white cbwarf sizes ancl speeds (Ixoester1976).," The ratio of the Bondi-Hoyle to Eddington accretion rates is $v^2R/2GM$ for $v\gg c$, or around $^4$ for typical white dwarf sizes and speeds \citep{koe76}."
. Table 3. lists typical densities and other relevant parameters for four fundamental types of ISM: molecular elouds. diffuse clouds. warm ionizect. and hot ionized.," Table \ref{tbl3} lists typical densities and other relevant parameters for four fundamental types of ISM: molecular clouds, diffuse clouds, warm ionized, and hot ionized."
 Listed. also are the expected: high-end mass infall rates for white cwarls moving through these regions following either. Boncli-Llovle (uid) or Extdington (eceometric) ivpe accretion., Listed also are the expected high-end mass infall rates for white dwarfs moving through these regions following either Bondi-Hoyle (fluid) or Eddington (geometric) type accretion.
" The two-phase aceretion-dilfusion scenario as laid. out w (Dupuisctal.1993a).. invokes 10"" vvr within a cloud (accretion). and 510' vyvr between clouds (dilfusion)."," The two-phase accretion-diffusion scenario as laid out by \citep{dup93a}, invokes $^6$ yr within a cloud (accretion), and $5\times10^7$ yr between clouds (diffusion)."
 For disk stars with Galactic orbits similar to the Sun. these imescales imply about five cloud encounters per MMyr orbit.," For disk stars with Galactic orbits similar to the Sun, these timescales imply about five cloud encounters per Myr orbit."
 The motions of interest are the relative motions »etween stars. and the ISM. the latter which should. be moving within the relatively low velocity spiral arms of the Galactic disk CJabreiB&Wielen1997).," The motions of interest are the relative motions between stars and the ISM, the latter which should be moving within the relatively low velocity spiral arms of the Galactic disk \citep{jah97}."
. 1E correct. this implies a typical helium-rich white dwarl moving at + relative to the LSB and (presumably) the ESAL travels. on average. ppc. within clouds during a single Galactic rotation.," If correct, this implies a typical helium-rich white dwarf moving at $^{-1}$ relative to the LSR and (presumably) the ISM travels, on average, pc within clouds during a single Galactic rotation."
 This timescale corresponds. to the cooling age for a logg=8.0. WKN helium-rich white chart. and hence such a star should. according to this scenario. obtain up to 5107 gg of hivdrogen via geometric accretion. or Πο —4.4.," This timescale corresponds to the cooling age for a $\log\,g=8.0$, K helium-rich white dwarf, and hence such a star should, according to this scenario, obtain up to $5\times10^{22}$ g of hydrogen via geometric accretion, or [H/He] $=-4.4$."
 This value is commensurate with the highest hydrogen abundances observed in DB stars. and. well above the lower limit of detectability (Vossctal. 2007).," This value is commensurate with the highest hydrogen abundances observed in DB stars, and well above the lower limit of detectability \citep{vos07}."
. This) comparison implies 1) the clensities and corresponding Lclelington accretion rates in. ‘Table 3 are too high. or 2) cloud encounters last less than LO? vvr or are less frequent than once per 5«107 vvr. but is otherwise consistent with the geometric capture of hydrogen: within the ISM. and. inconsistent with Uuid accretion.," This comparison implies 1) the densities and corresponding Eddington accretion rates in Table \ref{tbl3} are too high, or 2) cloud encounters last less than $^6$ yr or are less frequent than once per $5\times10^7$ yr, but is otherwise consistent with the geometric capture of hydrogen within the ISM, and inconsistent with fluid accretion."
 LIgnoring the fact that all hydrogen will be ionized within the 3oneli-Llovle radius of the white dwark and. hence should acerete as a plasma at the Uuid rate (Alcock&BHlarionov 1980).. it is clear that [uid aceretion of hydrogen does," Ignoring the fact that all hydrogen will be ionized within the Bondi-Hoyle radius of the white dwarf, and hence should accrete as a plasma at the fluid rate \citep{alc80}, , it is clear that fluid accretion of hydrogen does"
Tot dust will also emit au isotropic infrared echo due o thermal cussion from dust at the rapid sublimation eniperature ~2300 Is. peaking at an observed waveleugth A~2(1|2)yon. Wasximan,"Hot dust will also emit an isotropic infrared echo due to thermal emission from dust at the rapid sublimation temperature $\sim2300$ K, peaking at an observed wavelength $\lambda \sim 2\,(1+z)\,{\rm \mu m}$."
&Draine(1990) argue that ouly the photous iu the 1.7.5eV. range will contribute o dust heating.," \fcitet{wad99} argue that only the photons in the $1-7.5\,{\rm eV}$ range will contribute to dust heating."
 For 7.3~7 the absorption efficiency. for photous iu this euergv range is 20.8: and moreover. such Notons are Likely to carry a considerable fraction of the otal OT emission.," For $\tau_{0.3}\sim 7$ the absorption efficiency for photons in this energy range is $>0.8$; and moreover, such photons are likely to carry a considerable fraction of the total OT emission."
 Therefore the integrated infrared flux Is Adopting our sinple model of dust scattering. Eqs. (2-1.. 2-2))," Therefore the integrated infrared flux is Adopting our simple model of dust scattering, Eqs. \ref{rsubl}, \ref{time}) )"
" allow us to relate the sublimation radius aud OT power. LOMELys ere s.l. to the observed echo delay. [oLolquemob6 os""* where €q=(1μα)0.00 allows for beaming or characteristic scattering aneles different from 207. aud Cy=RfB. should be uxed if the dust is located bevoud Tas"," allow us to relate the sublimation radius and OT power, $10^{47}L_{47}$ erg $^{-1}$, to the observed echo delay, $t^{\rm
E}_{\rm ob} \equiv 10^6 t^{\rm E}_{\rm ob\,6}$ s. where $C_1=(1-\mu)/0.06$ allows for beaming or characteristic scattering angles different from $20^\circ$, and $C_2=R/R_{\rm sub}$ should be used if the dust is located beyond $R_{\rm sub}$."
 For simplicity. we now suppose that the spectral iudex of the OT i a—1.," For simplicity, we now suppose that the spectral index of the OT is $\alpha\sim1$."
 This is quite close to the spectral index of the observed afterelows., This is quite close to the spectral index of the observed afterglows.
" We cau then use Eqxt2-3)) to relate the R-band (0.654541) echo fux density to the escape probability where the observed duration of the optical trausicut is 10AOT,ol} «, Note the strong dependence on redshitt which iuplies that accurate measurements of both the optical transient aud the echo flux could lead to a fairly precise redshift prediction."," We can then use \ref{Fsc}) ) to relate the R-band $0.65{\rm \mu m}$ ) echo flux density to the escape probability where the observed duration of the optical transient is $10^3\Delta t^{\rm
OT}_{\rm ob,3}$ s. Note the strong dependence on redshift which implies that accurate measurements of both the optical transient and the echo flux could lead to a fairly precise redshift prediction."
" The ratios of the optical trausient fux density. FS=Lyfutli)‘pu|2) >. aud infrared echo fiux density to the optical echo flux deusity are likewise eiven by where fy, is the fraction of incident OT photous. clucreine uuscattered from the dust cloud."," The ratios of the optical transient flux density, $F^{\rm OT}_{\nu_{\rm ob}} = 
L_{\nu} f_{\rm ns} (4 \pi)^{-1} D_{\rm A}^{-2} (1+z)^{-3}$ , and infrared echo flux density to the optical echo flux density are likewise given by where $f_{\rm ns}$ is the fraction of incident OT photons, emerging unscattered from the dust cloud."
" For CRD 950326. an excess Ro fux FEO.LpJy was nmieasnred a tine fh,—20 d 19993)."," For GRB 980326, an excess R flux $F^{\rm E}_{\nu_{\rm ob}}[0.65{\rm \mu
m}]\sim0.4\,{\rm \mu Jy}$ was measured a time $t^{\rm E}_{\rm
ob}\sim20$ d )."
" If we make the simplest assumptious. &4—C€»~l. then Ra,~0.301|2)ipe and L9«1070|2)?eres1."," If we make the simplest assumptions, $a_{-1}\sim Q_{\rm abs}\sim C_1\sim C_2\sim1$ , then $R_{\rm
sub}\sim0.3(1+z)^{-1}{\rm pc}$ and $L\sim9\times10^{45}(1+z)^{-2}\;
{\rm erg\, s}^{-1}$."
 Comparing the reported spectral slope. a~2.8 of the putative echo to that of the afterglow (a~0.8). we estimate that του~7(ef Fig.," Comparing the reported spectral slope, $\alpha\sim2.8$ of the putative echo to that of the afterglow $\alpha\sim0.8$ ), we estimate that $\tau_{0.3}\sim7$ Fig."
 2)., 2).
" This. in turn. implies that ο/dQ in the observed R baud ~(02(1|2)t and fü,0405(112)+ (sce Fig. 2))."," This, in turn, implies that $dP^{sc}/d\Omega$ in the observed R band $\sim0.2(1+z)^{-1}$ and $f_{\rm ns}\sim0.05(1+z)^{-4}$ (see Fig. \ref{fig2}) )."
" We can then use Eq.(3-3)) to deduce that AOL,~BCL|uM aud FE~200011:19pd."," We can then use \ref{recho}) ) to deduce that $\Delta t^{\rm OT}_{\rm ob,3}\sim3(1+z)^7$ and $F^{\rm OT}_{\nu_{\rm
ob}}[0.65{\rm\mu m}] \sim200(1+z)^{-10}\;{\rm\mu Jy}$."
" Tf: ~0.1, then the enerev associated0.65; with the first optical weasurement of the afterglow ~LOqe. after 0.5 d). suffices to account for(FOT|0.65;nu the observed excess after 20 d as a dust echo."," If $z\sim0.4$, then the energy associated with the first optical measurement of the afterglow $F^{\rm OT}_{\nu_{\rm ob}}[0.65{\rm\mu m}] \sim10\;{\rm\mu
Jy}$ after 0.5 d), suffices to account for the observed excess after 20 d as a dust echo."
 It: >0.1. then the optical trausieut would have had to be preseut aud create a larger flucuce at earlier times.," If $z>0.4$, then the optical transient would have had to be present and create a larger fluence at earlier times."
 This is uot unreasonable as the OT flux was measured to satisfy POTXf3.," This is not unreasonable as the OT flux was measured to satisfy $F^{\rm
OT}\propto t^{-2}$."
 Tn view of the large muuber of simplifvine assuniptious that we have made. this estimate can only be regarded as illustrative.," In view of the large number of simplifying assumptions that we have made, this estimate can only be regarded as illustrative."
 Dowever it suffices to demonstrate that dust scatteriue is consistent with all of the available data., However it suffices to demonstrate that dust scattering is consistent with all of the available data.
 A somewhat similar story can be told for CRB 970228. where the redshift. 2=0.695. is known (Djorsovskictal.19993).," A somewhat similar story can be told for GRB 970228, where the redshift, $z=0.695$, is known \fcitep{dea99}) )."
 The carliest R-baud measureuicut is ~23047] 0.7 d after the CRB: aud after ~30 d there red excess fiux 0.34Tv. was observed. with the spectral slope (a~3.0) very simular to that seen in CRB 9850326 (Calamactal.1999)).," The earliest R-band measurement is $\sim30{\rm\mu Jy}$ 0.7 d after the GRB; and after $\sim30$ d there red excess flux $\sim0.3{\rm\mu Jy}$ was observed, with the spectral slope $\alpha \sim 3.0$ ) very similar to that seen in GRB 980326 \fcitep{gea99}) )."
 For this object. again. within the uncertainties. the fluence iieasured in the first stages of the optical transicut is sufficicut to account for the energy in the optical excess.," For this object, again, within the uncertainties, the fluence measured in the first stages of the optical transient is sufficient to account for the energy in the optical excess."
 Iu both examples above. the mass of dust required to produce an optical depth το~7 with our siuplest assumptions and assunuünue that it is spherically sviuiumuetricallv distributed with respect to the CRB is 0.1 ML.," In both examples above, the mass of dust required to produce an optical depth $\tau_{0.3}\sim7$ with our simplest assumptions and assuming that it is spherically symmetrically distributed with respect to the GRB is $\sim0.1$ $_\odot$."
 This amount of dust could fori in iui expanding lüeh-anetallicityv wind associated with an earlier stage in the evolution of the CRB progenitor as we lave assumed in our simple model., This amount of dust could form in an expanding high-metallicity wind associated with an earlier stage in the evolution of the GRB progenitor as we have assumed in our simple model.
 Alternatively the dust might be associated with a molecular cloud if CRBs are associated with massive star formation or a molecular torus should they be located im obscured galactic nuclei., Alternatively the dust might be associated with a molecular cloud if GRBs are associated with massive star formation or a molecular torus should they be located in obscured galactic nuclei.
 Iu this letter we present au alternative explanation for the reddened excess cussion observed in CRB 970228 and CRB 980326. which we attribute to dust scattering of the carly-time. afterglow cussion.," In this letter we present an alternative explanation for the reddened excess emission observed in GRB 970228 and GRB 980326, which we attribute to dust scattering of the early-time, afterglow emission."
 This scenario is predictive enough to be confirmed or ruled out with observations of future GRBs., This scenario is predictive enough to be confirmed or ruled out with observations of future GRBs.
 Tn particular. in contrast to the supernova explanation (Bloomctal. 1999: Calamaictal. 1999: Reichart 19993). if the excess Cluission ds due to dust scattering. then its properties will depeud ou the huuinositv of the optical transient.," In particular, in contrast to the supernova explanation \fcitep{bea99}; \fcitep{gea99}; \fcitep{rei99}) ), if the excess emission is due to dust scattering, then its properties will depend on the luminosity of the optical transient."
 IIETE II (http://space.nüt.edu/IIETE/) scheduled to be launched im early 2000. aud Swift (http://switteste.uasa.ecov/homepage.tial). scheduled for 2003. should provide real-time localization of CRD N-rav afterglows with sufficient precision to perit faster follow-up and better measurements of its total fiueuce.," HETE II (http://space.mit.edu/HETE/) scheduled to be launched in early 2000 and Swift (http://swift.gsfc.nasa.gov/homepage.html), scheduled for 2003, should provide real-time localization of GRB X-ray afterglows with sufficient precision to permit faster follow-up and better measurements of its total fluence."
derived peak flux and fluence) should be treated with care. however. since theMVdof of the spectral fits cau reach values as high as lor5.,"derived peak flux and fluence) should be treated with care, however, since the$\chi^2/dof$ of the spectral fits can reach values as high as 4 or 5."
 sinibus:The problem stems frou au excess at high cucreics., The problem stems from an excess at high energies.
 Issues were reported bv Callowayetal.(20081) for other bursts fro this source. and the physical cause is not vet understood.," Similar issues were reported by \citet{gal08b} for other bursts from this source, and the physical cause is not yet understood."
" Adding a power law to our model. for example. reduces (7/dof to below 2: the evolution of Tj, and y, i this case still ugeests a PRE burst."," Adding a power law to our model, for example, reduces $\chi^2/dof$ to below 2: the evolution of $T_{bb}$ and $R_{bb}$ in this case still suggests a PRE burst."
 Photospheric touchdown occurs 3.5 5s after the start of the burst. consistent with the first detection ofthe burst oscillations.," Photospheric touchdown occurs $3.5-5$ s after the start of the burst, consistent with the first detection of the burst oscillations."
 The paramecters of the April 2 burst do not differ dramatically from those of other bursts from IIETE J1900. which also show evidence for PRE subject to uucertainties about the spectral fits (Callowayctal. 20082.b).," The parameters of the April 2 burst do not differ dramatically from those of other bursts from HETE J1900, which also show evidence for PRE subject to uncertainties about the spectral fits \citep{gal08a,gal08b}."
. The burst with oscillations docs not however have the extended or double peak structure exhibited by most of the other bursts (Figure 21)., The burst with oscillations does not however have the extended or double peak structure exhibited by most of the other bursts (Figure \ref{blc}) ).
 Theburst liehteurve showsa i.fast rise. zmO.L 8 defined as in Callowayctal.," The burst lightcurve shows a fast rise, $\approx 0.4$ s defined as in \citet{gal08a}."
(20 The decay can be modelled with a double exponential with decay timescales as of 7.3 s and 8.15 respectively., The decay can be modelled with a double exponential with decay timescales as of 7.3 s and 8.4 s respectively.
 Total burst duration is z60 s (burst eud is defined as the time when flux falls to of the peak flux. corrected for persistent enission).," Total burst duration is $\approx 60$ s (burst end is defined as the time when flux falls to of the peak flux, corrected for persistent emission)."
 Figure 3 shows the intensity. defined as the couutrate in the 2.016.0 keV band. for all publicly available observations from 2005-2009.," Figure \ref{lc} shows the intensity, defined as the countrate in the 2.0–16.0 keV band, for all publicly available observations from 2005-2009."
 Fieure Lo shows the correspouding color-color diagram., Figure \ref{ccd} shows the corresponding color-color diagram.
 Data were deadtimoe corrected. background subtracted. and N-ray bursts removed.," Data were deadtime corrected, background subtracted, and X-ray bursts removed."
 For each observation we calculated N-rav colors from the Standard? data., For each observation we calculated X-ray colors from the Standard2 data.
 We defined soft color as the ratio between count rates iu the 3.56.0 and 2.03.5 keV bands and hard color as the ratio between count rates in the 9.716.0 aud 6.09.7 keV. bands., We defined soft color as the ratio between count rates in the 3.5–6.0 and 2.0–3.5 keV bands and hard color as the ratio between count rates in the 9.7–16.0 and 6.0–9.7 keV bands.
 We normalized colors aud intensity to the Crab values nearest im finie (IXuulkersetal.L991) audin the same PCA eain epoch (sce for example vanStraatenetal. 2003))., We normalized colors and intensity to the Crab values nearest in time \citep{kuu94} and in the same PCA gain epoch (see for example \citealt{vans03}) ).
 Tt is clear from Figures 3-- that the observation where the burst oscillations were detectedl was rather uuimsual., It is clear from Figures \ref{lc}- \ref{ccd} that the observation where the burst oscillations were detected was rather unusual.
 The burst happened when source intensity was at its highest recorded level (2263 imiCrab) aud when it was iu the soft (hanana) state., The burst happened when source intensity was at its highest recorded level $\approx 63$ mCrab) and when it was in the soft (banana) state.
 All previous bursts have been detected in harder states., All previous bursts have been detected in harder states.
 Fitting the PCA spectiun with an absorbed disk-blackbody plus power-law model. aud assuming a standard bolometric correction factor of 2 (iitZandctal.2007) we fud an unabsorbed bolometric fux of 3.5«10P eres st cian? Guterstellar absorption was fixed to 1.6«1022 cu 7).," Fitting the PCA spectrum with an absorbed disk-blackbody plus power-law model, and assuming a standard bolometric correction factor of 2 \citep{int07} we find an unabsorbed bolometric flux of $3.5 \times 10^{-9}$ ergs $^{-1}$ $^{-2}$ (interstellar absorption was fixed to $1.6 \times 10^{21}$ cm $^{-2}$ )."
 At a distance of 5 kpe. this corresponds to of the Eddineton huuinosity if we asstune {πα=2-5«107 eres |.," At a distance of 5 kpc, this corresponds to of the Eddington luminosity if we assume $L_\mathrm{Edd}=2.5 \times 10^{38}$ erg $^{-1}$ ."
 What causes burst oscillations is still not understood., What causes burst oscillations is still not understood.
 Flame spread from a point should lead to asviuncetries in the carly phase of the burst., Flame spread from a point should lead to asymmetries in the early phase of the burst.
 However. while this," However, while this"
GRBO030329/SN2003clh (Staneketal.2003:Ljorth2003) and GRD930425/5N1993bw (Galamaοἱal.1998) show that Type Ib/c supernovae are the parent population of long GRBs.,"GRB030329/SN2003dh \citep{sta03,hjo03} and GRB980425/SN1998bw \citep{gal98} show that Type Ib/c supernovae are the parent population of long GRBs."
 Type Ib/c SNe are believed {ο represent core-collapse events of massive stus in compact binaries (Wooslev1993:Paczvuski1998;Brownetal.2000: 2003).," Type Ib/c SNe are believed to represent core-collapse events of massive stars in compact binaries \citep{woo93,pac98,bro00,bet03}."
". Thev are probably part of a continuous sequence adjacent to Tvpe II SNe. ordered by increasing compactness of the binary in which the hydrogen (Ib/c) and the helium (1c) envelope are removed in a common envelope phase (Nomoto,Iwamotoratto 2003).."," They are probably part of a continuous sequence adjacent to Type II SNe, ordered by increasing compactness of the binary in which the hydrogen (Ib/c) and the helium (Ic) envelope are removed in a common envelope phase \citep{nom95,tur03a}."
 The remaining naked star rotates rapidly at the orbital period by (dal spin-up., The remaining naked star rotates rapidly at the orbital period by tidal spin-up.
 As the inactive iron-core succumbs to its own weight and that of the surrounding IHe-envelope. a rotating black hole nucleates during core-collapse (Dethe.Drown&Lee2003).," As the inactive iron-core succumbs to its own weight and that of the surrounding He-envelope, a rotating black hole nucleates during core-collapse \citep{bet03}."
. Some of the binding enerev liberated during gravitational collapse will be channeled to eject matter. producing an accompanvinge lverogen (and helium) deficient Type Ib (Ivpe Ic) supernova (MacFadyen2003).," Some of the binding energy liberated during gravitational collapse will be channeled to eject matter, producing an accompanying hydrogen (and helium) deficient Type Ib (Type Ic) supernova \citep{mac03}."
. The branching ratio of Type Ib/c SNe to GRB-SNe can be caleulated from the ratio (1-22)x10© of observed GRDs-to-Tvpe II supernovae (Porciani&Madan 2001)... a beaming," The branching ratio of Type Ib/c SNe to GRB-SNe can be calculated from the ratio $(1-2)\times 10^{-6}$ of observed GRBs-to-Type II supernovae \citep{por01}, , a beaming"
have incorporated this component.,have incorporated this component.
 The hard X-ray background is largely produced by QSOs and lower-luminosity AGN., The hard X-ray background is largely produced by QSOs and lower-luminosity AGN.
 Combining the the total background at z=0 with the observed X-ray spectra of individual sources and the redshift dependence of AGN output permitted to estimate the 10?—10° keV background as a function(2004) of redshift.," Combining the the total background at $z=0$ with the observed X-ray spectra of individual sources and the redshift dependence of AGN output permitted to estimate the $10^0
- 10^5$ keV background as a function of redshift."
" This background has a Compton temperature of 10? K due to the peak in vJ, around 30—50 keV, which penetrates regions that are optically thick to UV many simulations, including ours, treat the UV as (althoughoptically thin also) and can provide a considerable source of both ionization and heating."," This background has a Compton temperature of $10^{7.3}$ K due to the peak in $\nu J_{\nu}$ around $30 - 50$ keV, which penetrates regions that are optically thick to UV (although many simulations, including ours, treat the UV as optically thin also) and can provide a considerable source of both ionization and heating."
" Indeed, found that including an X-ray background increases(1999) the equilibrium temperature of the IGM by ~20%."," Indeed, found that including an X-ray background increases the equilibrium temperature of the IGM by $\sim 20$."
". Finally, since the formation of massive ellipticals where star formation has been effectively quenched since z~1—2 is still not well understood and is the subject of ongoing studies a study to examine the extent to which results are 2008),,sensitive to the assumed ionizing radiation background is warranted. ("," Finally, since the formation of massive ellipticals where star formation has been effectively quenched since $z\sim1-2$ is still not well understood and is the subject of ongoing studies, a study to examine the extent to which results are sensitive to the assumed ionizing radiation background is warranted. ("
"Naturally, further work could also be done with disk galaxies, Lyman-break populations at higher redshift and ","Naturally, further work could also be done with disk galaxies, Lyman-break populations at higher redshift and others.)"
The paper is organized as follows., The paper is organized as follows.
" In others.) refsect:bkg,, we describe the numerical methods and parameters of our simulations, in particular the ionizing radiation backgrounds which we apply, comprising a recent version of(1996),, a new rescaled version of the UV background which falls off much more slowly at high redshift, this new UV background with an additional X-ray component, and the more realistic, recently calculated UV background of(2009)."," In \\ref{sect:bkg}, we describe the numerical methods and parameters of our simulations, in particular the ionizing radiation backgrounds which we apply, comprising a recent version of, a new rescaled version of the UV background which falls off much more slowly at high redshift, this new UV background with an additional X-ray component, and the more realistic, recently calculated UV background of."
". In refsect:results we describe the results obtained from those simulations, in particular the effects of the backgrounds on the gas properties, star formation and stellar dynamics."," In \\ref{sect:results} we describe the results obtained from those simulations, in particular the effects of the backgrounds on the gas properties, star formation and stellar dynamics."
" refsect:disc is a discussion of the implications of these results when taken collectively, and 855 is conclusion."," \\ref{sect:disc}
 is a discussion of the implications of these results when taken collectively, and 5 is conclusion."
" As a baseline, we use the UV background as used incitetnaab07;; we call this the *Old UV? model."," As a baseline, we use the UV background as used in; we call this the “Old UV” model."
" To create an upper bound on high-z UV, we keep the same assumed spectral shape but set the intensity to decline as (1+2) in physical units above the peak, which1 occurs at z2.4; we call this the “New UV” model."," To create an upper bound on high-z UV, we keep the same assumed spectral shape but set the intensity to decline as $(1+z)^{-1}$ in physical units above the peak, which occurs at $z\approx 2.4$; we call this the “New UV” model."
" For a more realistic case, we use the background calculated by(2009),, which has a similar z-dependence of intensity but whose spectrum softens markedly for z>3 as the quasar contribution dies out; we call this “FG UV”."," For a more realistic case, we use the background calculated by, which has a similar z-dependence of intensity but whose spectrum softens markedly for $z\gtrsim 3$ as the quasar contribution dies out; we call this “FG UV”."
" The ionization rates of these three models are compared in Fig.H 1,,"," The H ionization rates of these three models are compared in Fig. \ref{fig:spectra},"
 along with the data of mentioned above., along with the data of mentioned above.
" It is apparent that New UV and (2008)FG UV provide a better fit to the observations than does Old UV, in particular at zd."," It is apparent that New UV and FG UV provide a better fit to the observations than does Old UV, in particular at $z>4$."
 We also implement an X-ray background., We also implement an X-ray background.
" This component uses the spectral shape given in(2004), which represents an average quasar background considering both obscured and unobscured sources, and is strongly peaked around 30 keV in EFzg."," This component uses the spectral shape given in, which represents an average quasar background considering both obscured and unobscured sources, and is strongly peaked around 30 keV in $EF_E$."
" The intensity normalization comes from(2007),, who modeled the AGN/QSO X- background using both deep pencil-beam pointings and shallow surveys."," The intensity normalization comes from, who modeled the AGN/QSO X-ray background using both deep pencil-beam pointings and shallow surveys."
" This is converted into heating and ionization rates using Cloudy1998),, which includes(v07.02, photo and Compton heating as well as secondary ionizations; the heating rates are then increased by a factor of 1.5 to better agree with the more recent model of(2005)."," This is converted into heating and ionization rates using Cloudy, which includes photo and Compton heating as well as secondary ionizations; the heating rates are then increased by a factor of 1.5 to better agree with the more recent model of."
. Heating rates due to the UV and X-ray backgrounds are roughly equal for virialized gas at z=0., Heating rates due to the UV and X-ray backgrounds are roughly equal for virialized gas at $z=0$.
" The redshift dependence of this background is taken to be similar to New UV: intensity in physical units increases as (1+z)? to z=2, and declines as (14-z)~! thereafter (which is admittedly unrealistic; see below)."," The redshift dependence of this background is taken to be similar to New UV: intensity in physical units increases as $(1+z)^3$ to $z=2$, and declines as $(1+z)^{-1}$ thereafter (which is admittedly unrealistic; see below)."
 We call this the “New UV+X” model., We call this the “New UV+X” model.
 Notice in Fig., Notice in Fig.
 1 that New UV+X has only a negligibly higher ionization rate than NewUV: photons at keV energies and above deposit >99% of their energy as heat in a highly ionized medium through electron-electron collisions1985).," \ref{fig:spectra}
 that New UV+X has only a negligibly higher ionization rate than NewUV: photons at keV energies and above deposit $\gtrsim 99\%$ of their energy as heat in a highly ionized medium through electron-electron collisions."
. The X-ray background also (Shullcontributes at the ~30% level to the higher-energy Hell ionizations., The X-ray background also contributes at the $\sim 30\%$ level to the higher-energy HeII ionizations.
" As an additional motivation for the revised ionizing background, we calculate a simple (homogeneous) semianalytic model of reionization using Cloudy."," As an additional motivation for the revised ionizing background, we calculate a simple (homogeneous) semianalytic model of reionization using Cloudy."
" We create bins at successive epochs and apply the corresponding background on our z-dependence formulas) to gas of the corresponding(based (physical) density (i.e., the mean present density scaled as z)?); Cloudy outputs the electron-scattering and (14-Gunn-Peterson optical depths (re, and tap, respectively)."," We create bins at successive epochs and apply the corresponding background (based on our z-dependence formulas) to gas of the corresponding (physical) density (i.e., the mean present density scaled as $(1+z)^3$ ); Cloudy outputs the electron-scattering and Gunn-Peterson optical depths $\tau_{\text{es}}$ and $\tau_{\text{GP}}$ , respectively)."
" For Tes this method is replaced above redshift 10 by an analytic integral of the electron density (i.e., Tes=στfnedl over"," For $\tau_{\text{es}}$ this method is replaced above redshift 10 by an analytic integral of the electron density (i.e., $\tau_\text{es}=\sigma_T \int n_e dl$ over"
imply 4|.=1.5. compatible with what we obtain. allowing for large error bars.,"imply $\lambda \, = \, 1.5$, compatible with what we obtain, allowing for large error bars."
 It seems possible that the merger cross section rises more steeply with mass than what we assumed here. and an exponent close to unity would allow a better match with data.," It seems possible that the merger cross section rises more steeply with mass than what we assumed here, and an exponent close to unity would allow a better match with data."
 We conclude that mergers between black holes might be able to explain the entire mass distribution., We conclude that mergers between black holes might be able to explain the entire mass distribution.
 However. we have to ask whether such a growth process could operate in a very similar way for nuclear star clusters. as for black holes.," However, we have to ask whether such a growth process could operate in a very similar way for nuclear star clusters, as for black holes."
 All such merger arguments may work as well for nuclear star clusters as for black holes surrounded by stars., All such merger arguments may work as well for nuclear star clusters as for black holes surrounded by stars.
 In such a picture the upper end of the distribution ts just the maximum that can be reached given the density of galaxies. and the mass of the central black holes.," In such a picture the upper end of the distribution is just the maximum that can be reached given the density of galaxies, and the mass of the central black holes."
 Why is there a minimum mass in super-massive black holes?, Why is there a minimum mass in super-massive black holes?
 As we argued earlier. the data and much work by ?. and others show. that there are very few black holes near to and below 10°M...," As we argued earlier, the data and much work by \citet{2008ApJ...688..159G} and others show, that there are very few black holes near to and below $10^{6} \, M_{\odot}$."
 There is a variety of exploratory ideas why this is so (??2)..," There is a variety of exploratory ideas why this is so \citep{2004Natur.428..724P,2006A&A...458L...9M,2005A&A...436..805M}."
 [t seems plausible to assume. that the transition from massive black holes to nuclear star clusters holds a clue to solving this question.," It seems plausible to assume, that the transition from massive black holes to nuclear star clusters holds a clue to solving this question."
 A possible transition from a nuclear star cluster to a super-massive star has been discussed by ?.. and again by ?.. with the latter team arguing for a transition to a super-massive black hole (222)..," A possible transition from a nuclear star cluster to a super-massive star has been discussed by \citet{1970ApJ...162..791S}, and again by \citet{2004Natur.428..724P}, with the latter team arguing for a transition to a super-massive black hole \citep{1972AAa,1972AAb,2003ApJ...591..288H}."
 However. the results of (2?) preclude any contribution from super-massive stars near or above 10°M... since such stars explode completely due to an instability given by General Relativity. leaving no black hole behind. we do need a mechanism that manages to give black holes right below this cutoff.," However, the results of \citep{1972AAa,1972AAb} preclude any contribution from super-massive stars near or above $10^{6} \, M_{\odot}$, since such stars explode completely due to an instability given by General Relativity, leaving no black hole behind, we do need a mechanism that manages to give black holes right below this cutoff."
 ? have shown that wind mass loss effectively competes with agglomeration. and so would limit massive stars of a few hundred M. to below 100 M: this implies that it would be difficult to get past this barrier 1 mass.," \citet{2008A&A...477..223Y} have shown that wind mass loss effectively competes with agglomeration, and so would limit massive stars of a few hundred $M_{\odot}$ to below 100 $M_{\odot}$; this implies that it would be difficult to get past this barrier in mass."
 On the other hand. agglomeration is à runaway process. while stellar winds are a quasi-steady process. basically limited to the Eddington luminosity: so perhaps more extreme conditions are required to get a serious run-away in agglomeration.," On the other hand, agglomeration is a runaway process, while stellar winds are a quasi-steady process, basically limited to the Eddington luminosity; so perhaps more extreme conditions are required to get a serious run-away in agglomeration."
" There are a number of processes. that contribute: One is the simple momentum exchange between stars. with a time scale of (2222): where ον is the velocity dispersion of the stars in. the system. assumed to be in virial equilibrium. v7, is the mass of the stars. s, 1s the density of the stars. and A is the Coulomb logarithm. typically of value 20."," There are a number of processes, that contribute: One is the simple momentum exchange between stars, with a time scale of \citep{1942psd..book.....C,1962pfig.book.....S,1987degc.book.....S,1987gady.book.....B}: where $\sigma_{\star}$ is the velocity dispersion of the stars in the system, assumed to be in virial equilibrium, $m_{\star}$ is the mass of the stars, $n_{\star}$ is the density of the stars, and $\Lambda$ is the Coulomb logarithm, typically of value 20."
 Considering nuclear star clusters it is hard to see why this process by itself would lead to a sudden transition at a specific mass. although a gravo-thermal catastrophe could in principle do this (?):: however. this process by itself would suggest. that the lower masses become a black hole. and the higher masses remain à star cluster. contrary to observation.," Considering nuclear star clusters it is hard to see why this process by itself would lead to a sudden transition at a specific mass, although a gravo-thermal catastrophe could in principle do this \citep{1987degc.book.....S}; however, this process by itself would suggest, that the lower masses become a black hole, and the higher masses remain a star cluster, contrary to observation."
" On the other hand. the agglomeration of stars is governed by their collision time scale. which is where N, is the total number of stars. 72, 1s the density of stars. and X, is the cross section of typical stars."," On the other hand, the agglomeration of stars is governed by their collision time scale, which is where $N_{\star}$ is the total number of stars, $n_{\star}$ is the density of stars, and $\Sigma_{\star}$ is the cross section of typical stars."
" We need only one star to start a runaway coalescence. and that is why also have the factor N,."," We need only one star to start a runaway coalescence, and that is why also have the factor $N_{\star}$."
 The question is whether either of these two processes or a combination of the two would allow for a transition at a specific mass of a nuclear star cluster such. that a short range of masses Is pinpointed.," The question is whether either of these two processes or a combination of the two would allow for a transition at a specific mass of a nuclear star cluster such, that a short range of masses is pinpointed."
 ? suggest. that there is à very long mass range. in which the process of agglomeration can give a large variety of masses. intermediate mass black holes.," \citet{2007ASPC..367..697M} suggest, that there is a very long mass range, in which the process of agglomeration can give a large variety of masses, intermediate mass black holes."
 Therefore their conclusion is that this would not lead to a relatively sharp transition., Therefore their conclusion is that this would not lead to a relatively sharp transition.
 Some have suggested (?).. that in the case of a binary black hole merger of equal masses a gravitational rocket effect could eject black holes from galactic centers. and one could so imagine. that all lower mass black holes might form. but no longer be in galactic centers.," Some have suggested \citep{2008ApJ...677..146P}, that in the case of a binary black hole merger of equal masses a gravitational rocket effect could eject black holes from galactic centers, and one could so imagine, that all lower mass black holes might form, but no longer be in galactic centers."
 In such a speculation these black holes below the transition point would exist. but be invisible.," In such a speculation these black holes below the transition point would exist, but be invisible."
 However. in ?. we show that this process is unlikely to be statistically relevant.," However, in \citet{2009ApJ...697.1621G} we show that this process is unlikely to be statistically relevant."
 In summary. accepting the process of agglomeration. what could modify the conclusion of previous authors. that a variety of masses is formed. and 1n contrast narrow down the mass range for the transition mass?," In summary, accepting the process of agglomeration, what could modify the conclusion of previous authors, that a variety of masses is formed, and in contrast narrow down the mass range for the transition mass?"
 There are several avenues to consider: First. galaxies grow by merging. starting from some minimum size: Could this minimum size of a central black hole correspond to the minimum size of a galaxy?," There are several avenues to consider: First, galaxies grow by merging, starting from some minimum size: Could this minimum size of a central black hole correspond to the minimum size of a galaxy?"
 That is hard to maintain. even considering. that ? have identified a minimum mass of order 5-107M. most of it in dark matter.," That is hard to maintain, even considering, that \citet{2007ApJ...663..948G} have identified a minimum mass of order $5 \cdot 10^{7} \, M_{\odot}$, most of it in dark matter."
 To grow a galaxy like ours. with a central black hole close to the low mass cut-off. would require so many merger events. that it is hard to see that much of any connection to the minimum galaxy could survive with à signature. except for properties that survive in all galaxies. independent of whether they contain a black hole at their center.," To grow a galaxy like ours, with a central black hole close to the low mass cut-off, would require so many merger events, that it is hard to see that much of any connection to the minimum galaxy could survive with a signature, except for properties that survive in all galaxies, independent of whether they contain a black hole at their center."
 Second. in a merger process we do obtain a central spike in dark matter from the merger density profile (222):: pumQU8xEy! where x=r/r.. a scaled radial coordinate. and γω is of order unity: various variants of this formula have been discussed (?)..," Second, in a merger process we do obtain a central spike in dark matter from the merger density profile \citep{1997ApJ...490..493N,1998ApJ...502...48K,1998ApJ...499L...5M}: $
\rho_{dm} \, \sim \, (x^{\gamma_{sp}} ( 1+ x)^{2})^{-1}
$ where $x = r/r_c$, a scaled radial coordinate, and $\gamma_{sp}$ is of order unity; various variants of this formula have been discussed \citep{1998ApJ...502...48K}."
" This implies a central dark matter component with a mass enclosed with R of My,Rr."," This implies a central dark matter component with a mass enclosed with $R$ of $M_{R, dm} \, \sim \, R^{3 - \gamma_{sp}}$."
" Applying this first just to stars as well implies à radial dependence of n,-.... of velocity dispersion of σι~ye-y. ""7 "," Applying this first just to stars as well implies a radial dependence of $n_{\star} \, \sim \, 1/x^{\gamma_{sp}}$, of velocity dispersion of $\sigma_{\star} \, \sim \, x^{(2 - \gamma_{sp})/2}$ ."
The combination yields to τω~πο.... so to an arbitrarily short time scale of stellar agglomeration Tye. at the center of merged galaxies.," The combination yields to $\tau_{aggl} \, \sim \, x^{ (3 \gamma_{sp} -1)/2}$, so to an arbitrarily short time scale of stellar agglomeration $\tau_{aggl}$ at the center of merged galaxies."
 This would then suggest. that all galaxies should have a central super-massive black hole. and not just those above a specific mass: this 1s again in contradiction to data.," This would then suggest, that all galaxies should have a central super-massive black hole, and not just those above a specific mass; this is again in contradiction to data."
 However. including the process of massive star formation. e.g..(?) near the center of a galaxy might require a certain minimum amount of gaseous turn- by star (star formation. mass ejection by winds. and," However, including the process of massive star formation, \citep{2009ApJ...697.1741B} near the center of a galaxy might require a certain minimum amount of gaseous turn-over by star (star formation, mass ejection by winds, and"
target galaxies in our program will be measured.,target galaxies in our program will be measured.
 Finally. we will complete our clistance Ladder by addiugD> LL superuova-host fiekl [n]galaxies.," Finally, we will complete our distance ladder by adding 44 supernova–host field galaxies."
 These steps will allow us to measure accurate ancl consistent clistauces [rom the closest Cepheids to the most distaut galaxies., These steps will allow us to measure accurate and consistent distances from the closest Cepheids to the most distant galaxies.
 This work deals with Cepheids in the Large Magellanic Cloud (LMC)., This work deals with Cepheids in the Large Magellanic Cloud (LMC).
 The LMC Cepheic population has been studied exteusively for over half a ceutury., The LMC Cepheid population has been studied extensively for over half a century.
 For example. presented well-samplecd light[n] curves [or LO Cepheids in the LMC. aud study of the population continues to this day with large-scale projects such as OGLE III (Soszvuskietal.2008)...," For example, \citet{1940AnHar..90..253S} presented well–sampled light curves for 40 Cepheids in the LMC, and study of the population continues to this day with large-scale projects such as OGLE III \citep{2008AcA....58..163S}."
 Its proximity means we can readily observe Cepheids over a rauge of periods. (rom those that are sinl: “to many of the Galactic parallax sample (ie. 2€10 days). to the loug period Cepheids with P>10 days. which more generally overlap Cepheids most easily observed in our most distant targets.," Its proximity means we can readily observe Cepheids over a range of periods, from those that are similar to many of the Galactic parallax sample (i.e. $P \leq 10$ days), to the long period Cepheids with $P \geq 10$ days, which more generally overlap Cepheids most easily observed in our most distant targets."
 Until relatively. recently. the majority of distance measurements have been undertaken at optical waveleneths (e.g. FOL).," Until relatively recently, the majority of distance measurements have been undertaken at optical wavelengths (e.g. F01)."
 Cepheid studies at optical waveleugths have their shortcomines. the main one being extinction.," Cepheid studies at optical wavelengths have their shortcomings, the main one being extinction."
 Although the effect. can. to first order. be removed by use of the reddeuiug-free Weseuheit iudex (MadoreL976) this techuique still requires prior knowledge of the extinction law. as well as au assumption that it is universal.," Although the effect can, to first order, be removed by use of the reddening–free Wesenheit index \citep{1976RGOB..182..153M} this technique still requires prior knowledge of the extinction law, as well as an assumption that it is universal."
 By moving to the mid-iufrarec. reddening aud extinetion are dimiuished by arouud a [actor of twenty (Rieke&Lebolsky1985).. making their absolute contribution and their uncertainties negligible.," By moving to the mid–infrared, reddening and extinction are diminished by around a factor of twenty \citep{1985ApJ...288..618R}, making their absolute contribution and their uncertainties negligible."
 In. adcition to the drop in extinction ellects. McGoueealetal.(1982) describe two other advantages of the infrared over the optical.," In addition to the drop in extinction effects, \citet{1982ApJ...257L..33M} describe two other advantages of the infrared over the optical."
 In the infrared the amplitudes of the Cepheids’ lighteurves decrease. as does the intrinsic width of the instability strip. because these wavelenetls are less sensitive to temperature clhauges.," In the infrared the amplitudes of the Cepheids' lightcurves decrease, as does the intrinsic width of the instability strip, because these wavelengths are less sensitive to temperature changes."
 IudeedMceConegaletal.(1982) demonstrated that the width of the near-infrared A baud huninosity relation from observations is less than the width [rom time-averaged B-baud observatious., Indeed\citet{1982ApJ...257L..33M} demonstrated that the width of the near–infrared $H$ band period--luminosity relation from observations is less than the width from time–averaged $B$ –band observations.
 Near-infrared observatious of LMC Cepheids were more recently obtained by (2001).., Near–infrared observations of LMC Cepheids were more recently obtained by \citet{2004AJ....128.2239P}.
 By moving to longer wavelengths iu the mid-iulrared. in combination with well phased observations. we can decrease the measured width even further.," By moving to longer wavelengths in the mid–infrared, in combination with well phased observations, we can decrease the measured width even further."
 The IRAC imager on is a superb instrument for undertaking a recalibration of the Cepheid distance scale., The IRAC imager on is a superb instrument for undertaking a recalibration of the Cepheid distance scale.
 LikeHubble.. has the advantage of operating in a stable euviromment without weather or seeing variations. aud with great flexibility in scheduling.," Like, has the advantage of operating in a stable environment without weather or seeing variations, and with great flexibility in scheduling."
 As such. we have been able to obtain precise aud deterministically well-saimplecd light. curves for 85 Cepheids.," As such, we have been able to obtain precise and deterministically well–sampled light curves for 85 Cepheids."
 The observatious of the LMIC Cepheids were one of the first programs unclertaken by post-crvogeni¢ “Wari Spitzer’.," The observations of the LMC Cepheids were one of the first programs undertaken by post-cryogenic “Warm Spitzer""."
 The target selection aud observations are described in Section 2. and the photometry aud calibration are discussed iu Section 3.., The target selection and observations are described in Section \ref{sec:observations} and the photometry and calibration are discussed in Section \ref{sec:data_reduction}.
 Light aud color curves are presented for each Cepheid in Section {.., Light and color curves are presented for each Cepheid in Section \ref{sec:results}.
 PL relations. inclucing a discussion ou their use in determining the tilt of the LMC. are eiven in Section 5..," PL relations, including a discussion on their use in determining the tilt of the LMC, are given in Section \ref{sec:pl_relations}."
 The period-color relation at mean light has been measured Lor the first time at these waveleugths: it is discussed iu Section 6.., The period–color relation at mean light has been measured for the first time at these wavelengths; it is discussed in Section \ref{sec:co_absorption}. .
 Section 7 provides a summary., Section \ref{sec:summary} provides a summary.
"yy aati the probability that a galaxy with a mass between Mag and μμ!dAdo at a redshift 2 was the result. of the merging of two galaxies with a mass ratio between Auer, and [ers|€Ruyere in the preceding time interval df.",") = , the probability that a galaxy with a mass between $\mgal$ and $\mgal + d\mgal$ at a redshift $z$ was the result of the merging of two galaxies with a mass ratio between $\rmerg$ and $\rmerg + d\rmerg$ in the preceding time interval $dt$."
 We define Dues<Lo Le. it is the ratio of the barvonic mass of the smaller parent to that of the larger parent.," We define $\rmerg < 1$, i.e., it is the ratio of the baryonic mass of the smaller parent to that of the larger parent."
 Unfortunately. we do not resolve a large enough. dynamic range to Lully characterize the function. V. so instead we study several interesting “projections” of it to learn about its overall parameter dependence.," Unfortunately, we do not resolve a large enough dynamic range to fully characterize the function $\Psi$, so instead we study several interesting “projections” of it to learn about its overall parameter dependence."
 With the morphological assignment recipe of refsec:btd.. à galaxy’s Hubble type is determined. by the merger history of its “main” progenitor. the most. massive progenitor at each output.," With the morphological assignment recipe of \\ref{sec:btd}, a galaxy's Hubble type is determined by the merger history of its “main” progenitor, the most massive progenitor at each output."
 LP the. Hubble. types of the merging galaxies change the result of the merger. bbecause bulges behave dillerentlv. from. disks. then the merger history of the minor progenitors also matters.," If the Hubble types of the merging galaxies change the result of the merger, because bulges behave differently from disks, then the merger history of the minor progenitors also matters."
 However. this effect is likely to be of secondary importance. and in this paper we will not distinguish between mergers of disk and bulge galaxies.," However, this effect is likely to be of secondary importance, and in this paper we will not distinguish between mergers of disk and bulge galaxies."
 Table 1. lists the number of resolved main branch mergers (mergers involving a galaxy's main progenitor) and the total number of resolved: mergers for all of our merger trees., Table \ref{tab:prop} lists the number of resolved main branch mergers (mergers involving a galaxy's main progenitor) and the total number of resolved mergers for all of our merger trees.
 Only for the high mass sample. ⋅Lu ⋠⋠⋅ ∆∪⊔∣∆↥↙⊽↻⋅≟↓≱⇀∪⋅⊳∠⇂∪∖∖⋎∢⊾↓⋅⋖⋅≱∖∪⇂∖⇁∢⋅⋜↧≱∖↓⋏∙≟⊔↓∐≼⇍⋜⋃∐⊔⊔⊔↓∣⋡∢⋅↓⋅∪⇂ oll-main branch mergers.," Only for the high mass sample, $M_{\rm gal} > 6.4\times10^{10} \Msun$, do we resolve a significant number of off-main branch mergers."
 In most cases. therefore. we present results only for main branch mergers. but we occasionally compare the statistics of main branch mergers to those of all mergers for the high mass saniple.," In most cases, therefore, we present results only for main branch mergers, but we occasionally compare the statistics of main branch mergers to those of all mergers for the high mass sample."
 Figures 3. and 4 present our main characterizations of WM.οναμα the average number of mergers per galaxy per Gyr above fh thresholds of 0.125. 0.25. and 0.5 for ealaxies in the high mass sample (Fig. 3))," Figures \ref{fig:resh} and \ref{fig:resl} present our main characterizations of $\Psi(\mgal,z,\rmerg)$: the average number of mergers per galaxy per Gyr above $\rmerg$ thresholds of 0.125, 0.25, and 0.5 for galaxies in the high mass sample (Fig. \ref{fig:resh}) )"
 and the medium and low mass samples (Fig. 4))., and the medium and low mass samples (Fig. \ref{fig:resl}) ).
" The relation of «Nuusfdi to VMziers) is where fü is the mass ratio threshold. Αα is the number of simulated. galaxies in the mass range Mii, to Αν. παλ is the galaxy barvonic mass function. and nds the mean space density of galaxies in the mass range."," The relation of $dN_{\rm merg}/dt$ to $\Psi(\mgal,z,\rmerg)$ is where $R_{\rm min}$ is the mass ratio threshold, $N_{\rm gal}$ is the number of simulated galaxies in the mass range $M_{\rm min}$ to $M_{\rm max}$, $dn/dM$ is the galaxy baryonic mass function, and $\bar{n}$ is the mean space density of galaxies in the mass range."
 We compute μονοdE by counting mergers in a redshift interval As and dividing by the corresponding time interval Af. with typical values AfzmdOvr.," We compute $dN_{\rm merg}/dt$ by counting mergers in a redshift interval $\Delta z$ and dividing by the corresponding time interval $\Delta t$, with typical values $\Delta t \approx 1 {\rm Gyr}$."
 The lower mass ratio mergers (Lo... Risinl/8 or L/4) can only be resolved for higher mass galaxies.," The lower mass ratio mergers (i.e., $R_{\rm min}=1/8$ or $1/4$ ) can only be resolved for higher mass galaxies."
 Merger rates are substantially higher for massive galaxies e.g.. the average rate of At770.5 mergers ab 2=0.3 is 0054Gyr.+ for the high mass sample and 0.018Gyr.1 for both the medium and low niass saiples.," Merger rates are substantially higher for massive galaxies — e.g., the average rate of $\rmerg>0.5$ mergers at $z=0.3$ is $0.054\,{\rm Gyr}^{-1}$ for the high mass sample and $0.018\,{\rm Gyr}^{-1}$ for both the medium and low mass samples."
 Our main limitation in computing these statistics is that we do not reliably. resolve SINID. groups with fewer than G4 particles., Our main limitation in computing these statistics is that we do not reliably resolve SKID groups with fewer than 64 particles.
 Thus. when a galaxys mass increases by 63 particles we cannot tell without detailed examination whether this growth was the result of a merger or of smooth accretion.," Thus, when a galaxy's mass increases by 63 particles we cannot tell without detailed examination whether this growth was the result of a merger or of smooth accretion."
 The directly calculated rates shown by solid [ines in Figures 3. and 4 are therefore lower limits to the true rates., The directly calculated rates shown by solid lines in Figures \ref{fig:resh} and \ref{fig:resl} are therefore lower limits to the true rates.
 Dashecl curves accompanying the solid. curves show rates that include the maximum contribution of unresolved mergers. assigning all growth of up to 63 SPILL particle masses that is not in resolved mergers to unresolved mergers.," Dashed curves accompanying the solid curves show rates that include the maximum contribution of unresolved mergers, assigning all growth of up to 63 SPH particle masses that is not in resolved mergers to unresolved mergers."
 At low redshift there are no unresolved mergers. but as we eo back in redshift and the ealaxy masses decrease. the number of possible unresolved mergers increases.," At low redshift there are no unresolved mergers, but as we go back in redshift and the galaxy masses decrease, the number of possible unresolved mergers increases."
 We stop each line when the number of possible unresolved: mergers is greater than the Poisson uncertainty in the number of, We stop each line when the number of possible unresolved mergers is greater than the Poisson uncertainty in the number of
he warmn/hot eas in the simmlation (Figure 6)) to he measurements of soft N-rav. background. (?)..,the warm/hot gas in the simulation (Figure \ref{fig:xspec}) ) to the measurements of soft X-ray background \citep{kuntz_etal01}.
 The predicted dux a energies X0.5 keV is 10 o LOO times lower than the observed flux.," The predicted flux at energies $\lesssim
0.5$ keV is $10$ to $100$ times lower than the observed flux."
 The LSC eas therefore cannot explain this soft N-rav enudssion Which is likely to be associated with the rot halo of the Milkv Way., The LSC gas therefore cannot explain this soft X-ray emission which is likely to be associated with the hot halo of the Milky Way.
 At cucreics around LkeV. the predicted flux coustitutes z:5LOM of he total NRB fiux aud is therefore au important component of the NRB at these energies.," At energies around $1$ keV, the predicted flux constitutes $\approx 5-10\%$ of the total XRB flux and is therefore an important component of the XRB at these energies."
 At higher enereies. the contribution of he LSC eas to the ARD is insignificaut.," At higher energies, the contribution of the LSC gas to the XRB is insignificant."
 I ls no clear whether the N-ray cussion of the intergalactic gas can be reliably detected iu the near future., It is not clear whether the X-ray emission of the intergalactic gas can be reliably detected in the near future.
 Such detection is difficult because oue needs a survey that covers a laree skv area and is suffiieutlv sensitive to detect the sieual ouly ~1 of the XRD., Such detection is difficult because one needs a survey that covers a large sky area and is sufficiently sensitive to detect the signal only $\sim 1\%$ of the XRB.
 The diffuse LSC emission should be correlated with the superealactic plane. but so is the N-rav. enission fron nearby ACGNS (2)..," The diffuse LSC emission should be correlated with the supergalactic plane, but so is the X-ray emission from nearby AGNs \citep{shaver_pierre89}."
" Receutlv. 7 analyzed the full sky 2.10 keV N-vav map and found evidence for the ""diffuse (Quuresolved) X-ray e1uüission associated with the superealactic plauc."," Recently, \citet{boughn99} analyzed the full sky $2-10$ keV X-ray map and found evidence for the “diffuse” (unresolved) X-ray emission associated with the supergalactic plane."
 The quote anit surface brightuess of the diffuse cluission constitutes about P4 of the 2.)10 keV XRD aud is Zx25«10.MoresstemP3 (see his Table 1).," The quoted maximum surface brightness of the diffuse emission constitutes about $1\%$ of the $2-10$ keV XRB and is $I_{\rm X}\approx
5\times 10^{-10}\ {\rm ergs\ s^{-1} cm^{-2} sr^{-1}}$ (see his Table 1)."
" This correspouds to z0,01keV s+ cni5jrckeV- assundue. that flux djs. constant over 2.10 keV. This flux is consistent with the predicted fiux from the WIT gas in our simulation (sce Fig. ο) ", This corresponds to $\approx 0.04\ {\rm keV}$ ${\rm s^{-1}}$ ${\rm cm^{-2} sr^{-1} keV^{-1}}$ assuming that flux is constant over $2-10$ keV. This flux is consistent with the predicted flux from the WH gas in our simulation (see Fig. \ref{fig:xspec}) ):
107.2IkeVostemJyκο|: at chereies 25 keV (ie. zmOL of the total ARB at hese energies: sec ?)) the fux decreases steeplv at higher cucreics.," $10^{-2}-10^{-1}
{\rm keV\ s^{-1} cm^{-2} sr^{-1} keV^{-1}}$ at energies $2-5$ keV (i.e., $\approx 0.1-1\%$ of the total XRB at these energies; see \citealt{kuntz_etal01}) ); the flux decreases steeply at higher energies."
 As can be secu from the X-ray brightucss nap in Figure 5.. most of the XN-rav cussion is indeed concentrated towards the supergalactic aue.," As can be seen from the X-ray brightness map in Figure \ref{fig:sky2}, most of the X-ray emission is indeed concentrated towards the supergalactic plane."
 However. the cussion is pateliv auc is far roni beime uniforin.," However, the emission is patchy and is far from being uniform."
 Iu contrast. ? modelled the eas distribution using a simple “pillbox” mode or the distribution of gas iu the LSC reeionu: le uniform gas istribution within a disk of radius Πως and thickuess JueczO.25Rac.," In contrast, \citet{boughn99} modelled the gas distribution using a simple “pillbox” model for the distribution of gas in the LSC region: the uniform gas distribution within a disk of radius $R_{\rm SC}$ and thickness $H_{\rm SC}\approx 0.25R_{\rm SC}$."
 For 1e detected diffuse N-ray. flux this model Προς eas density of 2.5«10Gcm? for temperatures of around 10 keV. This temperature is much higher iu the typical tempcratures of the LSC eas iu nr simulation., For the detected diffuse X-ray flux this model implies gas density of $2.5\times 10^{-6}{\ \rm cm^{-3}}$ for temperatures of around $10$ keV. This temperature is much higher than the typical temperatures of the LSC gas in our simulation.
 The difference is due to the fac iat the eas distribution iu the sinulated LSC is not described by the pillbox model., The difference is due to the fact that the gas distribution in the simulated LSC is not described by the pillbox model.
 As can be seen from Fieures Lo and 5.. the distribution of matter in the LSC region is filamentary rather than disk-like aud the N-ray cussion is far frou beius uniform.," As can be seen from Figures \ref{fig:sky1} and \ref{fig:sky2}, the distribution of matter in the LSC region is filamentary rather than disk-like and the X-ray emission is far from being uniform."
 The bulk of the X-ray cinissiou thus comes from the relatively high-deusitv. ης~10an* regions within and around eroups and clusters.," The bulk of the X-ray emission thus comes from the relatively high-density, $n_e\sim
10^{-5}-10^{-3}{\ \rm cm^{-3}}$ , regions within and around groups and clusters."
 The hot aud deuse regions of the LSC should also distort the cosnüc microwave backeround radiation via the inverse Compton or Doppler scattering. the thermal and kinetic Sunvaev-Zeldovich effect. respectively (SZ:7.andrefor- ," The hot and dense regions of the LSC should also distort the cosmic microwave background radiation via the inverse Compton or Doppler scattering, the thermal and kinetic Sunyaev-Zel'dovich effect, respectively \citep[SZ;][and references therein]{SZ80}. ."
The temperature fluctuations of the CMD due to the non-relativistic thermal SZ effect cau be written as: where. —ΠοΤε y=LAT41091fuTd (all quantities are iu ces uuits aud the inteeral is along line-of-sight).," The temperature fluctuations of the CMB due to the non-relativistic thermal SZ effect can be written as: where $x\equiv hv/kT_{\rm CMB}$, $y=1.117\times 10^{-34}\int n_eT_edl$ (all quantities are in cgs units and the integral is along line-of-sight)."
 In the Bavleigli-Jeaus regine (heκαςkTi NEB): AT/Tx2g: the deviation of CMD eniperature along a eiven direction is thus xoportional to the eas pressure integrated. along lis cürection., In the Rayleigh-Jeans regime $hv\ll kT_{\rm CMB}$ ): $\Delta T/T\approx 2y$; the deviation of CMB temperature along a given direction is thus proportional to the gas pressure integrated along this direction.
 Although for teuiperatures aud densities typical or the gas in the LSC the SZ effect is not expected o be very strong. due to the large angular extent ofthe LSC. it should contribute to the anisotropy ou the ou the angular scales of ~1107.," Although for temperatures and densities typical for the gas in the LSC the SZ effect is not expected to be very strong, due to the large angular extent of the LSC, it should contribute to the anisotropy on the on the angular scales of $\sim 1-10^{\circ}$."
 The iecrual SZ imap has patehy appearance simular o that of the X-ray brightness in Figure 5 with AT/T~5.10%10 within groups and clusters. ~10.© on their outskirts. and ~10.* iu he strongest filameuts.," The thermal SZ map has patchy appearance similar to that of the X-ray brightness in Figure \ref{fig:sky2} with $\Delta T/T\sim 5\times 10^{-6}-10^{-5}$ within groups and clusters, $\sim 10^{-7}$ on their outskirts, and $\sim 10^{-8}$ in the strongest filaments."
 The kinetic SZ effect has simular inagnitude., The kinetic SZ effect has similar magnitude.
 These fluctuations are siuall and are below current seusitivity limits of the SZ observations: their coutribution to the large angular scale anisotropies imcasured bv CODE satellite is iusienificant., These fluctuations are small and are below current sensitivity limits of the SZ observations; their contribution to the large angular scale anisotropies measured by COBE satellite is insignificant.
" Finally, we presented (Figure 5)) skv maps of coluun densities of the three ionic species of oxvecn: OVI. OVIL. and OVIIL."," Finally, we presented (Figure \ref{fig:sky2}) ) sky maps of column densities of the three ionic species of oxygen: OVI, OVII, and OVIII."
 The colui densities were caleulated using the densities and temperatures of eas m the high+vesolition region of the simulation and assmuuiung the wniformmetallicitv of 0.3 solu and the observed N- backerouud., The column densities were calculated using the densities and temperatures of gas in the high-resolution region of the simulation and assuming the uniformmetallicity of $0.3$ solar and the observed X-ray background.
 Although. the OVI is the least i»nudaut (ts abundance is more than an order of," Although, the OVI is the least abundant (its abundance is more than an order of"
The Na-O auticorrelatiou is slightly different for the MIR aud ALP components. but in both cases the exteusiou is quite modest. more similar to that in AL 1 than iu NGC 2808.,"The Na-O anticorrelation is slightly different for the MR and MP components, but in both cases the extension is quite modest, more similar to that in M 4 than in NGC 2808."
 This suggests only a modest spread in Ue iu NGC 1851 because a high Y fraction secu to be associated only to very loug tails of very O-poor stars. not preseut in this CC.," This suggests only a modest spread in He in NGC 1851 because a high Y fraction seems to be associated only to very long tails of very O-poor stars, not present in this GC."
 Our finding fom a chemical approach confiniis the claims by Salarisetal.(2008).. based on the absence of a tilt along the IIB aud the lack of a splitting iu the MS.," Our finding from a chemical approach confirms the claims by \cite{sal08}, based on the absence of a tilt along the HB and the lack of a splitting in the MS."
 Iu addition. we observe a slight change of the mean value of ο aud Na abundances at the level of the buup ou the RGB.," In addition, we observe a slight change of the mean value of O and Na abundances at the level of the bump on the RGB."
 We already showed (Carrettaetal.2007):Dragaeliaetal.2010) that this variation is expected roni theoretical models. which predict a change iu the »uup huninosity with We content (Salavisetal.2006).. ic. with elements involved in p-capture reactions.," We already showed \citep{car07b,bra10} that this variation is expected from theoretical models, which predict a change in the bump luminosity with He content \citep{sal06}, i.e. with elements involved in p-capture reactions."
 The xesence of a mix of first and second generation stars results iu a concentration of Na-poor/IIe-poor stars just fore the bump accompanied by an accunulation of Niwexich/Ileadeh stars just above the bump level., The presence of a mix of first and second generation stars results in a concentration of Na-poor/He-poor stars just before the bump accompanied by an accumulation of Na-rich/He-rich stars just above the bump level.
 This accounts for the observed abundance changes at the nup without resurrecting the internal mixing scenario or ο and Na (e.g. Lee 2010)) and then overcomine he unpalatable requirement of basic stellar structure differences between field aud. CC stars., This accounts for the observed abundance changes at the bump without resurrecting the internal mixing scenario for O and Na (e.g. \citealt{lee10}) ) and then overcoming the unpalatable requirement of basic stellar structure differences between field and GC stars.
 For NGC 1851. where we hvpothesize two distinct clusters (see below). we expect a further seamiug of the bump in the RGB huninosity function (Carretta et al.," For NGC 1851, where we hypothesize two distinct clusters (see below), we expect a further smearing of the bump in the RGB luminosity function (Carretta et al.,"
 in prep)., in prep).
 Iu Fig. L.," In Fig. \ref{f:fig4},"
 we use the Stroummeren vin 0οΝΤΟ to test where stars of different components aud populations are located ou the RGD.," we use the Strömmgren $u,u-b$ CMD to test where stars of different components and populations are located on the RGB."
 This plane is optimally suited to separate first and second generation stars. probably because of N (euliauced in O-depleted. second generation stars) via the formation of NIT. CN aud their relevauce ou the «—b (or the Johusou ( B). see Yougetal.(2008):Marinoetal...Carrettaal. (20092).," This plane is optimally suited to separate first and second generation stars, probably because of N (enhanced in O-depleted, second generation stars) via the formation of NH, CN and their relevance on the $u-b$ (or the Johnson $U-B$ ), see \cite{yon008,mar08,car09a}."
.. Stars of the first ecueration (P. Carrettaetal. 200923) in NGC 1851 lie alone a narrow strip to the blue of the RGB (Fig. L.," Stars of the first generation (P, \citealt{car09a}) ) in NGC 1851 lie along a narrow strip to the blue of the RGB (Fig. \ref{f:fig4},"
 bottom pancl). as expected from their uiprocessed Chemical abiunudauces.," bottom panel), as expected from their unprocessed chemical abundances."
 On the contrary. the second generation stars are spread out to the red. as in NGC 6752 (Carrettaetal. 20092): the I stars are in the muddle aud the extreme E component. with the lowest O abuudauces. is located at the reddest edge.," On the contrary, the second generation stars are spread out to the red, as in NGC 6752 \citealt{car09a}) ): the I stars are in the middle and the extreme E component, with the lowest O abundances, is located at the reddest edge."
" This scerceation is followed also within cach metallicity conrponeut. aud it is ""orthogonal to the separation of MIR and MP stars (Fie. {νι"," This segregation is followed also within each metallicity component, and it is “orthogonal"" to the separation of MR and MP stars (Fig. \ref{f:fig4},"
 top panel). which are well intermingled across all the RGB in this color.," top panel), which are well intermingled across all the RGB in this color."
 The same holds if we separate the RGD stars using the average value for Ca ([Ca/II|2.— 0.853): stars with low and high Ca are spread across the entiro RGB (Fie. L.," The same holds if we separate the RGB stars using the average value for Ca $=-0.83$ ): stars with low and high Ca are spread across the entire RGB (Fig. \ref{f:fig4},"
 middle panel)., middle panel).
 Therefore the spread of Ca does not track the abundances of p-capture elements., Therefore the spread of Ca does not track the abundances of p-capture elements.
 A ERK-S test on the σπανο distributious of |Ca/II]| for stars of the first and second eeueratious in NGC 1851 iudicates that they are indistinguishable., A K-S test on the cumulative distributions of [Ca/H] for stars of the first and second generations in NGC 1851 indicates that they are indistinguishable.
 Instead. the Ca abunudauices closely track those of Fe.," Instead, the Ca abundances closely track those of Fe."
 The cumulative distribution of [Ca/II| values for the MB. and MP. coniponeuts on the RGB are definitively ciffereut., The cumulative distribution of [Ca/H] values for the MR and MP components on the RGB are definitively different.
 Moreover. also the radial distributions of Ca-ricli and Ca-poor stars coufirui the close correspondence with metallicity: the Ca-poor eqjauts are iore concentrated. while Ca-rich stars slow a tendency toward iore exterual regions.," Moreover, also the radial distributions of Ca-rich and Ca-poor stars confirm the close correspondence with metallicity: the Ca-poor giants are more concentrated, while Ca-rich stars show a tendency toward more external regions."
 Is there a comprehensive scenario able to account for all the evidence found here aud iu previous works in NGC' 18517, Is there a comprehensive scenario able to account for all the evidence found here and in previous works in NGC 1851?
 In our view. the answer is affirmative if we consider NGC 1851 as the result of a chain of eveuts that started with two distinct clusters;," In our view, the answer is affirmative if we consider NGC 1851 as the result of a chain of events that started with two distinct clusters."
 Several sugeestious of duplicitv come from the bimodal distribution of IID stars. the double SCD. and hints of double sequences ou the RGB.," Several suggestions of duplicity come from the bimodal distribution of HB stars, the double SGB, and hints of double sequences on the RGB."
 Up to now. the main objection was the absence of a metallicity spread. owing to the lack of precise abuudances for a statistically significant nunber of stars.," Up to now, the main objection was the absence of a metallicity spread, owing to the lack of precise abundances for a statistically significant number of stars."
 This linütatiou has finally been overcome by our study., This limitation has finally been overcome by our study.
 As a tentative working livpothesis we can think of two different clusters. born in a much larger syste. perhaps a dSphl.," As a tentative working hypothesis we can think of two different clusters, born in a much larger system, perhaps a dSph."
 Being distinct. cach one might have formed with asheltly differeut metallicity and with a differeut level of oclemceutst!.," Being distinct, each one might have formed with a slightly different metallicity and with a different level of $\alpha-$."
", Each object is rightfully a GC. since cach component show the Na-O auticorrelation. the classical signature of the processes endiug ina GC (Carrettaotal.201053)."," Each object is rightfully a GC, since each component show the Na-O anticorrelation, the classical signature of the processes ending in a GC \citealt{car10b}) )."
 After a while. the two clusters uuderwenut ainerecr. likely because both were dragged to the cceuter of the dSph by dynamical friction (see Bellazzinictal.2008)) and the result is NGC 1851.," After a while, the two clusters underwent a merger, likely because both were dragged to the center of the dSph by dynamical friction (see \citealt{bel08}) ) and the result is NGC 1851."
 Finally. the dSpl moreed with the Milkv. Way.," Finally, the dSph merged with the Milky Way."
 We think that this is the siuplest scenario. that with a minima of hypothesis lav account for many observational constraints.," We think that this is the simplest scenario, that with a minimum of hypothesis may account for many observational constraints."
 The two MB and MP. coimiponeuts do not show any significaut difference iu kinematics. the velocity dispersion being the same for both componcuts.," The two MR and MP components do not show any significant difference in kinematics, the velocity dispersion being the same for both components."
 A comprehensive dynamical model would be very welcome. although we do not know when the merging occurred.," A comprehensive dynamical model would be very welcome, although we do not know when the merging occurred."
 We will rely on the observed chemistry., We will rely on the observed chemistry.
 The observables include: a double SCD. where the faint SGD (fSGD) includes of the stars aud the bright SGB (bSCGCD) the remaimime (Milonectal. 2008)). aud with coutroversial evidence of different concentration: the MB aud ALP components on the ROB. with a clear cdiffercuce in radial concentration: a bimodal distribution on the IB. with ~10 of the stars ou the BIIB aud ~605( on the ROB (Miloneetal. 2008)): the observed luminosity of ΠΟ stars aud the mocerate exteusion of the Na-O auticorrelation in both the AIR and AIP ROB components. which both sugeests small He abundance variations.," The observables include: a double SGB, where the faint SGB (fSGB) includes of the stars and the bright SGB (bSGB) the remaining \citealt{mil08}) ), and with controversial evidence of different concentration; the MR and MP components on the RGB, with a clear difference in radial concentration; a bimodal distribution on the HB, with $\sim 40\%$ of the stars on the BHB and $\sim 60\%$ on the RHB \citealt{mil08}) ); the observed luminosity of HB stars and the moderate extension of the Na-O anticorrelation in both the MR and MP RGB components, which both suggests small He abundance variations."
 We may explain these observables in different wavs: (i) À single CC with two populations having a cdiffereut total CNO abundance (but a similar Πο abuudauce)., We may explain these observables in different ways: (i) A single GC with two populations having a different total CNO abundance (but a similar He abundance).
 This may explain the SCD but fails to reproduce the uunuberratios on the WB (if the same efficicucy for mass loss on the RGB is assumed for both sub-populations). because the iinor fSGD component should be associated with the major RNB one. G," This may explain the SGB but fails to reproduce the number ratios on the HB (if the same efficiency for mass loss on the RGB is assumed for both sub-populations), because the minor fSGB component should be associated with the major RHB one. ("
i) Α siugle CC with two populations having a differeut Ho and total CNO abundance.,ii) A single GC with two populations having a different He and total CNO abundance.
 In this scenario there is wich were He in CNO-vich than in CNO normal, In this scenario there is much more He in CNO-rich than in CNO normal
(TDB) using the position and proper motion of Geminea listed in Table 2..,(TDB) using the position and proper motion of Geminga listed in Table \ref{tbl-2}.
" This was done using the ""timeconv program. which is part of the FTOOLS software package."," This was done using the “timeconv” program, which is part of the FTOOLS software package."
 We determined optimal radii for the source extraction and background annulus in the GIS by maximizing the signal-to-noise ratio in the resulting light curve., We determined optimal radii for the source extraction and background annulus in the GIS by maximizing the signal-to-noise ratio in the resulting light curve.
 As shown in Figure 2.. (he radius of the source circle was chosen to be 3% and a concentric annulus of inner and outer radii of 5 and 6/25. respectively gave a good estimate of the background.," As shown in Figure \ref{point1}, the radius of the source circle was chosen to be $3^{\prime}$, and a concentric annulus of inner and outer radii of $5^{\prime}$ and $6.\!^{\prime}25$, respectively gave a good estimate of the background."
 We confirm that. as found by Beckeretal.(1999) using previous aand limages. there is no evidence lor diffuse emission (svnchrotron nebulosity) associated. with Geninga.," We confirm that, as found by \cite{be99} using previous and images, there is no evidence for diffuse emission (synchrotron nebulosity) associated with Geminga."
 In Figure 3. we compare the resulting 0.5—4.0 keV light curve with the only previous hard X-ray. light. curve of Geminga made with the same instrument in 1994 March. and described by Halpern&Wang(1997)., In Figure \ref{pulse1} we compare the resulting $0.5-4.0$ keV light curve with the only previous hard X-ray light curve of Geminga made with the same instrument in 1994 March and described by \cite{hw97}.
 The curves have been aligned. according to the EGRET ephemeris. and (he resulting agreement in phase confirms the drifine EGRET ephemeris shown in Figure 1..," The curves have been aligned according to the EGRET ephemeris, and the resulting agreement in phase confirms the drifting EGRET ephemeris shown in Figure \ref{phaseplot}."
 To evaluate whether (he pulse shape experienced anv change between the 1994 and 1999 oobservations. the (vo light curves in Figure 3. were compared using \7 test. after accounting for the difference in the exposure times.," To evaluate whether the pulse shape experienced any change between the 1994 and 1999 observations, the two light curves in Figure \ref{pulse1} were compared using $\chi^2$ test, after accounting for the difference in the exposure times."
 It was determined that the light curves do not differ significantly., It was determined that the light curves do not differ significantly.
 The stability of the light curve. and its large pulsed [raction and strong main peak allow the possibility of continuing (lie rotational ephemeris of Geminga using hard observations. e.g.. wilh andChandra. during the current epoch in which there are no high-energv >-rav instruments in orbit.," The stability of the light curve, and its large pulsed fraction and strong main peak allow the possibility of continuing the rotational ephemeris of Geminga using hard X-ray observations, e.g., with and, during the current epoch in which there are no high-energy $\gamma$ -ray instruments in orbit."
 Figure + shows the lolded light curves for the 1999 GIS data divided into three energy bands. 0.7—1.5 keV. 1.5—3.5 keV. and 3.5—7.0 keV. For comparison. we also reproduce soft. N-rav. pulse profiles from a 1993 September observation with the PPSPC that was published by Halpern&Wane(1997).. in energv bands 0.08—0.28 keV. 0.28—0.53 keV. and 0.53—1.50 keV. The summed EGRET light. eurve above LOO MeV is also shown.," Figure \ref{pulse2} shows the folded light curves for the 1999 GIS data divided into three energy bands, $0.7-1.5$ keV, $1.5-3.5$ keV, and $3.5-7.0$ keV. For comparison, we also reproduce soft X-ray pulse profiles from a 1993 September observation with the PSPC that was published by \cite{hw97}, in energy bands $0.08-0.28$ keV, $0.28-0.53$ keV, and $0.53-1.50$ keV. The summed EGRET light curve above 100 MeV is also shown."
 For CGRO viewing period 1. comparing the number of events selected with I100 MeV to the likelihood estimate of Geminga flux (Mattox. of the events selected are estimated to be from Geminga.," For CGRO viewing period 1, comparing the number of events selected with $>$ 100 MeV to the likelihood estimate of Geminga flux \citep{ma96} of the events selected are estimated to be from Geminga."
 The remaining would be primarily diffuse Galactic ganmia-ray enission., The remaining would be primarily diffuse Galactic gamma-ray emission.
 This is the estimated background for the EGRET lighteurve., This is the estimated background for the EGRET lightcurve.
 It is statistically consistent with the claim in Maver-Iasselwander(1994) that there is neeligible unpulsed emission from Geminea in that energv band., It is statistically consistent with the claim in \cite{mh94} that there is negligible unpulsed emission from Geminga in that energy band.
 These light curves resemble closely those given in Figure 9 of Halpern&Wane(1997)., These light curves resemble closely those given in Figure 9 of \cite{hw97}.
. Because the 1999 GIS observation had 2.5 (mes more exposure time than the 1994 observation. the GIS light curves could be split into smaller energv. bands. vielding more information about," Because the 1999 GIS observation had 2.5 times more exposure time than the 1994 observation, the GIS light curves could be split into smaller energy bands, yielding more information about"
Our analysis confirms that 1138213 is a marginal Am star - Fe is overabundant and O is underabundant.,Our analysis confirms that 138213 is a marginal Am star - Fe is overabundant and O is underabundant.
 Ca is almost solar abundant., Ca is almost solar abundant.
 For C and Li we give only the upper limits., For C and Li we give only the upper limits.
 75(LIIUmi66385. {12 8161. 5840536 1102¢2632 is a spectroscopic binary star.," 6385, +12 3161, 84036, 102632, A1m) is a spectroscopic binary star."
 Osawa determined. the spectral class of the star. as ALfA3/A5 from W/ll/metallic lines and Cowleyοἱal.(1969) - as Al., \citet{osawa58} determined the spectral class of the star as A1/A3/A5 from K/H/metallic lines and \citet{ccjj69} - as A1.
 Later Abt&Alorrell(1995) specified it as AQUI class and Paunzenctal.(2001) - as VY class., Later \citet{am95} specified it as III class and \citet{pdhkw01} - as V class.
 Abt&Morrell(1995). also determined the projected rotational velocity as (sin;=.25 kmss1 but scaling' this: value to the system of Roveretal.(2002). the latter ‘hanged the velocity to esinf=33kmss |., \citet{am95} also determined the projected rotational velocity as $\vsini=25~$ $^{-1}$ but scaling this value to the system of \citet{rgbgz02} the latter changed the velocity to $\vsini=33~$ $^{-1}$.
 Our value of MM=3l is in good agreement with Roveretal.02) result., Our value of $\vsini=31~$ $^{-1}$ is in good agreement with \citet{rgbgz02} result.
 We used the orbital elements. given by Ima(2002): Pay= 23.25 Wo=31.42kms e——0.22. Vy=103kms tow=114.586.," We used the orbital elements given by \citet{Debernardi02}: $P_{\rm orb}=23.25^{d}$ , $K=31.42~{\rm km\,s^{-1}}$, =0.422, $V_0=1.03~{\rm km\,s^{-1}}$, $\omega={\rm 114^{\circ}.86}$."
 The radial velocities measured. from our spectra are. in. excelent agreement with the racial velocities. calculated by using these elements (see Figure 2))., The radial velocities measured from our spectra are in excelent agreement with the radial velocities calculated by using these elements (see Figure \ref{fnew}) ).
 There are many Fe lines in the spectral region of. so the abundance of Fe is very well determined.," There are many Fe lines in the spectral region of, so the abundance of Fe is very well determined."
 Ca is uncerabundant., Ca is underabundant.
 Phe abundances given at Table 3. for € and O are upper limits., The abundances given at Table \ref{t5} for C and O are upper limits.
 Phe situation with Li is the same as in the case of 1116657 - the Li line is very weak and the obtained Li abundance is only upper limit., The situation with Li is the same as in the case of 116657 - the Li line is very weak and the obtained Li abundance is only upper limit.
 Alore than one spectrum have been obtained in order to check the possible variability of some lines in the spectrum of 115537, More than one spectrum have been obtained in order to check the possible variability of some lines in the spectrum of 155375.
 As it is seen. a few lines have changed their profiles(see Figure 4))," As it is seen, a few lines have changed their profiles (see Figure \ref{fhd155}) )."
 The relative changes of two lines. bel A6419.95 and Fel A6421.35AA. have been most obvious.," The relative changes of two lines, FeI $\lambda$ and FeI $\lambda$, have been most obvious."
 Two calcium lines. Cal A6439.08 and Cal A6462.57AA.. have shown changes. too.," Two calcium lines, CaI $\lambda$ and CaI $\lambda$, have shown changes, too."
 The center of the lines has been changed. and also there could. be seen some features emerging from the blue side of the lines., The center of the lines has been changed and also there could be seen some features emerging from the blue side of the lines.
 All these observable clues forced us to suspect 1155375 as a new SB2 star., All these observable clues forced us to suspect 155375 as a new SB2 star.
 1159560. (G7. Dra. 66555. οὗ 1945. ss5s829. 2200. 110628. ENA. Adm) is a member of the visual binary system.," 159560 $\nu^2$ Dra, 6555, +55 1945, 85829, 30450, 10628 A, A4m) is a member of the visual binary system."
 The angular separation between the components is 61.97., The angular separation between the components is ${\rm 61.9}\arcsec$.
 Both components of the binary system are Am stars., Both components of the binary system are Am stars.
 “Phere have been many determinationsof the stellar spectral class in the literature., There have been many determinationsof the stellar spectral class in the literature.
 Phe first evaluation was given bv Slettebak(1949). - A2/E0/E51LIV from Ix/HL/moetallic lines., The first evaluation was given by \citet{Slettebak49} - IV from K/H/metallic lines.
 Later the author specified the spectral class from. H-lines as AT (Slettebak 1963))., Later the author specified the spectral class from H-lines as A7 \citealt{Slettebak63}) ).
 According to Abt&Carcona the star was ΕΣ. Cowl, According to \citet{ac84} the star was A4/F2V/F3.
eyctal.(1969) also determined. the spectral class of 1159560. as At [rom L-lines., \citet{ccjj69} also determined the spectral class of 159560 as A4 from H-lines.
 The evaluations of the projected: rotational velocity of the star have been very different., The evaluations of the projected rotational velocity of the star have been very different.
 Moved(1973) eave esin/=35 kmss 3ohm-Vitense&Dettmann(1980). - esini=47kmss +. Abt&Levy - esiné=50kmss 5. Abt&Morrell(1095). - rsiné=58 kimss L," \citet{am73} gave $\vsini=35~$ $^{-1}$, \citet{bvd80} - $\vsini=47~$ $^{-1}$ , \citet{al85} - $\vsini=50~$ $^{-1}$ , \citet{am95} - $\vsini=58~$ $^{-1}$ ."
opinally. Roveretal.(2002). determined the rotational velocity of the star as esing=68 kmss ," Finally, \citet{rgbgz02} determined the rotational velocity of the star as $\vsini=68~$ $^{-1}$ "
" l. . where uland arethe imposed velocity patternsat the top and bottomu? planes. Forrun G,we exciteall wavenumbers 3 €","As shown in Figures \ref{fig1}, , \ref{fig2} and \ref{fig3} initially the system until time $t \sim 79\, \tau_A$ follows the linear curves \ref{eq:lin2}) ) and \ref{eq:diff1}) )."
" n,<4, whilefor run we excite onlyone Fouriercompo- nentas uv= —u?= sin (8ra+1)éy,ie.", Up to this point the shear velocity at the top boundary induces a sheared magnetic field in the volume.
" wwe are injecting energy inthe system only atnin = 4-2πς, the wavenumber4 alongx. This can", As discussed in \ref{sec:runa} we have introduced a perturbation mimicking those naturally present in the corona.
" be noticedalsoin Figure[I0],wherethe kinetic spectrum for run Gat n= 3 is higher thanforrun F,aspart ofthe energy is injected alsoatn— 3 in thevorticalcase."," With no perturbation the system would relax over the resistive diffusive timescale $\tau_R$ $\sim 25\, \tau_A$ for run A) in a saturated diffusive equilibrium as described in \ref{eq:lin2}) ) and \ref{eq:diff1}) )."
 'The lowerlevel forthe kinetic spectrumis dueto the boundary conditionsthat roughlyset the value or the velocityat the injection wavenumbersinside the volume.," While the simulation presented here used a very small amplitude for the perturbation $\epsilon = 10^{-16}$ ), we have performed shortest simulations with different values for the amplitude."
" Inthe simple linear casethis is givenby eq. (12)).Inthe shearcase we wouldhave = V- (u?)=2.5 =10isthe in Ek(4)thelinear1/2-regime,andfrom (V Figurewe noticevolume)thatalso"," As expected for higher values of $\epsilon$ the instability develops sooner and for smaller values later, always following the linear curves until the instability transitions to the nonlinear stage."
 Ey(4)~2.5.Ontheother hand themagnetic fieldgrows linearlyin time [eq. (11))]until a balance is r," The more complete and systematic analysis of \cite{rom04,rom09} in 2D confirms this behavior."
eached betweentheenergy flux that is injectedat this scaleand the flux of energy flowing towardssmaller scales through a turbulent cascade. The magnetic energyspectra of thetwosimulations, \cite{dlkn09} have performed a similar simulation with a lower resolution and with a fixed value for the perturbation and for a time interval that covers only the initial stage of our simulations.
" areslightly different atthelarge scales withn, < 5.The largescale dynamics isin fact slightly diff"," They in fact stop right after the first big dissipative peak, that in our Figures \ref{fig1}, \ref{fig2} and \ref{fig3} corresponds at $t \sim 100\, \tau_A$."
"erentinthetwo cases. Inthe vortical case (run energyis injected in allmodes with wavenumbers G)3 <n,4 thatthen cascades toward smaller scales.In the shear (runF) energy isinjected onlyat one wavenumber n, =(4,0). We"," Continuing the simulation, and using a higher numerical resolution, the system reaches a statistically steady state where magnetic energy consistently fluctuates around a mean value and the shear is not recreated in the topology of the orthogonal magnetic field."
 have already noticed in § wecontinue shearing thefootpoints ofthe field-lines with our," Their analysis is then limited to a event taking place only during the early stages of the dynamics, and that afterward does not repeat."
" 1Dforcing [eq. (1))] inthenonlinear stage theorthogonal magnetic field isorganized in magnetic islands (Figure [3)), sothatit isnolonger amapping of the boundary velocity. wavenumber 4along x,energy is", As shown in \cite{rved08} during the linear stage the system is able to accumulate energy well beyond the average value maintained in the nonlinear stage only if the boundary forcing velocity satisfies the condition that its.
"then redistributed bythe nonlinear terms also tomodes with wavenumbers alongyatthe large scales, andasmall inversecascade is presentas inrun G.", The sheared profiles used in this paper satisfy this condition as well the profile used by \cite{dlkn09} (a linear combination of 6 sheared profiles).
 This is thebasic mechanism by which magnetic islands aresustained throughout thesimulation inthenonlinear stage. AND DISCUSSION 6.CONCLUSIONS In thispaper wehave investigated the dynamics of the Parkerproblem for theheating of coronal loops whenthefootpoints bya1Dshear velocity pattern at," These profiles are a very small subset of all the possible forcing profiles, and while they are very useful to get insight into the origin of turbulence in coronal loops they are not representative of the disordered photospheric motions, for which the strong stress buildup required for secondary instability to develop does not take place."
" the photosphere- mimicking boundary, and", The significance of their conclusions is then strongly diminished.
" compared these results withthose previously obtained whenamore complex “vortex- like"" velocitypattern was imposed 2008)..", Furthermore the Parker angle for this system cannot be defined as the relative angle between magnetic field-lines at which the system becomes unstable.
 This verysimple, This is not a definition.
 forcing is ideal to investigate the origin ofturbulence in coronal loops andtheinfluence of theboundary velocity forcing onthe dynamics ofthe system. Wewill also compare our resultsof (1992) and ofthemore, In fact for given initial conditions the angle or equivalently the time (as the linear equation \ref{eq:lin2}) ) and \ref{eq:diff1}) ) imply) at which the instability develops depend on the value of the amplitude of the perturbation that we add to the system.
" recent simulations of ∙∙ In summary, the main results presented", Depending on the value of the perturbation the Parker angle so defined.
 inthis paper are the following: 1. Initiallythe shearedvelocity forcing induces sheared perpendicular , On the other hand in the fully nonlinear stage the average magnetic field line magnitude fluctuates around a mean value.
"insidethe volume. aThe currentmagneticfield resultingtearingmodes ⋅layers are knowntoIn befactunstablewhen theto system transitions nonlinear it is to frommultiple thelinear to instability,the"," It is then possible to give a unique value for the Parker angle, defined now as the average inclination of the magnetic field-lines respect to the axial direction as done in \cite{rved07,rved08}, and as originally introduced by \cite{park88}."
 stageshownin due aBut thetearing systemhas becomeas fully nonlinearFigurethe| dynamicsonce fundamentally dif- A," As summarized in \ref{sec:ed} the one-point closure model developed by \cite{hp92}
 splits the domain into large and small scales."
sthe nonlinear termsare longer vanishthey ferent.do transport from theno the small nowscales where in correspondenceenerg," They conjecture that the large-scale fields evolve into a stationary laminar regime, the field magnitudes determined by the effective diffusion coefficients."
y ofthelarge X-pointsto linear magnetic takes withoutnon- , These laminar regimes correspond to our linear saturated diffusive regimes computed in \ref{par3}. .
"ing througha seriesreconnectionof equilibria disruptedplace,by tearing- go- like", In Figure \ref{fig8} the dotted lines show such diffusive curves for different values of the Reynolds numbers.
 instabilities. the with disordered Similarlyforcingto velocitiescase vortical boundarythefully nonlinear stage system highly in and chaotic (and increasingl," In their model the large-scale fields computed in this way are used to obtain $S$, the energy flowing into the system for unit time at the boundary (the power) due do the work done by photospheric motions on the magnetic field-lines footpoints."
ythe atis dynamicalReynolds numbers). For this do notsoob- higher secondary tearingof the current sheetswe in2D serve," They also calculate, through an EDQNM approximation, the value of the spectral energy flux $\epsilon$ flowing along the inertial range at the small scales."
" high-resolution simulations of decaying MHDasturbu- lence (Biskamp Welter||1989), thesmall scales fast & dynamics takeasplace.now at turbulent 2. The dynamics of theParker modeldonot depend onthe ofthe velocity forcingthat mimics strongly photosphericpatternmotions,as far stantin time (wedefer the of time- dependent fo"," Both $S$ and $\epsilon$ are functions of the effective diffusion coefficients, and the solution of the problem results requiring balance between the two powers $S = \epsilon$ $S$ and $\epsilon$ have both the dimension of a power, energy over time, as $S$ is the Poynting flux integrated over the boundary surface and $\epsilon$ is integrated over the whole volume as in \cite{rved07}] ]."
rcingtoa future investigationwork).The shear forcing[eq. boundary(1)] onlyatthetoppl," As shown in our simulations the large-scale fields are not laminar, and they are stationary only statistically."
ateisa simple and orderedapplied one-dimensional forcing. We have, Nevertheless itis useful to use \cite{hp92} model in order tounderstand why it is not applicable.
 shown thatthe resulting are sim- tothose developed whena more dynamics veryanddis- ordered “vortex-type” forcing velocity is complexapplied, From Figure \ref{fig8} we can estimate that the Reynolds number for which the diffusive regime dissipation matches the dissipation of the simulated
 shown thatthe resulting are sim- tothose developed whena more dynamics veryanddis- ordered “vortex-type” forcing velocity is complexapplied., From Figure \ref{fig8} we can estimate that the Reynolds number for which the diffusive regime dissipation matches the dissipation of the simulated
 shown thatthe resulting are sim- tothose developed whena more dynamics veryanddis- ordered “vortex-type” forcing velocity is complexapplied.W, From Figure \ref{fig8} we can estimate that the Reynolds number for which the diffusive regime dissipation matches the dissipation of the simulated
 shown thatthe resulting are sim- tothose developed whena more dynamics veryanddis- ordered “vortex-type” forcing velocity is complexapplied.We, From Figure \ref{fig8} we can estimate that the Reynolds number for which the diffusive regime dissipation matches the dissipation of the simulated
Thus if we take Py=107 we infer that the perturbed inary cecentricity needs to be greater than 0.75. in. order hat we may view the ring features at their inferred age.,"Thus if we take $P_{\rm min}=10^5$ we infer that the perturbed binary eccentricity needs to be greater than $0.75$, in order that we may view the ring features at their inferred age."
 In Fig., In Fig.
 6 we show a trajectory with e=0.75., $6$ we show a trajectory with $e=0.75$.
 The similarity. of such high eccentricty bound orbits to. the xwabolic trajectory suggests that the cise response should uu gaimilar., The similarity of such high eccentricty bound orbits to the parabolic trajectory suggests that the disc response should be similar.
 In Fig., In Fig.
 16 we give the apocentre distance versus longitude of pericentre plots for a coplanar encounter simulation with e=OS. jp;=0.3. and q=2.6.," $16$ we give the apocentre distance versus longitude of pericentre plots for a coplanar encounter simulation with $e=0.8$, $\mu=0.3$, and $q=2.6$."
 The »erturber starts at apocentre anc we give results. for he first and second completed orbital periods., The perturber starts at apocentre and we give results for the first and second completed orbital periods.
 Clearly he length asvmumetry is completely disrupted: after. the second encounter., Clearly the length asymmetry is completely disrupted after the second encounter.
 In. Fig., In Fig.
 17 we give the corresponding xwticle positions. which show that the ring structure is also destroved after the second encounter.," $17$ we give the corresponding particle positions, which show that the ring structure is also destroyed after the second encounter."
 Thus as suspected only a single Lv-by can be allowed., Thus as suspected only a single fly-by can be allowed.
 Possible scenarios in which a single pericentre passage could have occurred in the recent past are: 1) the perturber was originally a bound companion to 2 Pie with a sulliciently large pericentre distance that no significant interaction with the disc occurred. over its lifetime. and sullicently large semimajor axis and eccentricity that it could. be perturbed into the required trajectory by the close Lv-by of a massive field star SLO’ in the past: 2) 3 Pic was a hierarchical multiple svstem that was dynamically unstable on a time-scale comparable to its present age. resulting in the ejection of one or more stars and the development of a. disc-intercepting orbit for another. or an ejected star passed by ? Pic on its way out of the system.," Possible scenarios in which a single pericentre passage could have occurred in the recent past are: 1) the perturber was originally a bound companion to $\beta$ Pic with a sufficiently large pericentre distance that no significant interaction with the disc occurred over its lifetime, and sufficently large semimajor axis and eccentricity that it could be perturbed into the required trajectory by the close fly-by of a massive field star $\sim$ $^5$ in the past; 2) $\beta$ Pic was a hierarchical multiple system that was dynamically unstable on a time-scale comparable to its present age, resulting in the ejection of one or more stars and the development of a disc-intercepting orbit for another, or an ejected star passed by $\beta$ Pic on its way out of the system."
 In the first scenario the three-body interaction should rave occurred at least LO ago., In the first scenario the three-body interaction should have occurred at least $10^5$ ago.
 The further in the past we out this event. the more eccentric the resultant binary orbit must be in order to delay closest approach until LO’ ago (if the perturbed star has an initially out-going trajectory)," The further in the past we put this event, the more eccentric the resultant binary orbit must be in order to delay closest approach until $10^5$ ago (if the perturbed star has an initially out-going trajectory)."
 Although perhaps more attractive than the lower probability close binary encounter model used by INLSS. the problem of ow relative velocity is shifted from the star-disc interaction o the interaction between the second and third bodies.," Although perhaps more attractive than the lower probability close binary encounter model used by KLSS, the problem of low relative velocity is shifted from the star-disc interaction to the interaction between the second and third bodies."
 In respect of the second scenario we point out that Weinberger et al. (, In respect of the second scenario we point out that Weinberger et al. (
2000) report an age of 5 for the triple svstenm LID 141569.,2000) report an age of $5$ for the triple system HD 141569.
 Phis consists of a Vega-like DVe primary with a 500 disc and two M stars each at ~ projected separation., This consists of a Vega-like BVe primary with a $500$ disc and two M stars each at $\sim$ projected separation.
 Weinberger et al. (, Weinberger et al. (
2000) deduce that the relative separations involved may make the svstem unstable. eiving the possibility that HD. 141569 represents an earlier stage in the evolution of the 3 Pie svstem (with an age of 20 Myr).,"2000) deduce that the relative separations involved may make the system unstable, giving the possibility that HD 141569 represents an earlier stage in the evolution of the $\beta$ Pic system (with an age of $20$ )."
 Studies of the dynamies involved suggest that voung stars with multiple companions may not lose some of them. via dynamical instability until main sequence ages (IZeeleton Wiscleva 1995). which is supported by observations of high binary [frequency amongst voung stars compared. to field stars (Mathieu 1994).," Studies of the dynamics involved suggest that young stars with multiple companions may not lose some of them via dynamical instability until main sequence ages (Eggleton Kiseleva 1995), which is supported by observations of high binary frequency amongst young stars compared to field stars (Mathieu 1994)."
 Lere we have. considered. the stellar. Ilv-by./— hypothesis that may. account for the asvmimetrical structure in. the scattered. light. dust. cise of «3 Pie., Here we have considered the stellar fly-by hypothesis that may account for the asymmetrical structure in the scattered light dust disc of $\beta$ Pic.
 We investigated. the dvnamies of the perturbed planetesimal disc in an attempt to learn. more about the possible parameters of a IHlyv-by encounter., We investigated the dynamics of the perturbed planetesimal disc in an attempt to learn more about the possible parameters of a fly-by encounter.
 Additionally. we have described the origin of transient circumstellar eccentric ring structures as. being a general outcome of an encounter. due to the rellex motion of the primary star as the perturber passes through closest approach.," Additionally, we have described the origin of transient circumstellar eccentric ring structures as being a general outcome of an encounter, due to the reflex motion of the primary star as the perturber passes through closest approach."
 The apparent ring system: is) actually an eccentric tighth-wound one-armed. spiral pattern tha eraclually disappears owing to phase-mixing in the particles orbital motion., The apparent ring system is actually an eccentric tightly-wound one-armed spiral pattern that gradually disappears owing to phase-mixing in the particles' orbital motion.
 In inclined-orbit simulations we find that the vertical disc response also has analogues in the scattere ight disc of 3 Pie., In inclined-orbit simulations we find that the vertical disc response also has analogues in the scattered light disc of $\beta$ Pic.
 We summarise the coplanar results in Fig. 18.," We summarise the coplanar results in Fig. $18$,"
 and the inclined-orbit results in Fig. 19., and the inclined-orbit results in Fig. $19$.
 For the perturber we conclude that the mass should no » very. dilferent. from 0.5 MAL. (corresponding to spectra vpe MOV)., For the perturber we conclude that the mass should not be very different from $0.5$ $_\odot$ (corresponding to spectral type M0V).
 In this case we favour a low inclination progracde encounter with relative velocity of a few +. which makes a bound perturber seem the most natural choice.," In this case we favour a low inclination prograde encounter with relative velocity of a few $^{-1}$, which makes a bound perturber seem the most natural choice."
 One xoblem with this is that asvmmetrics would be destroved in the second. pericentre passage., One problem with this is that asymmetries would be destroyed in the second pericentre passage.
" Thus. given the deduced vouth of the asvmmetries. the perturber in this scenario is required to make a single close approach in the recent past (Le. 10"" ago)."," Thus, given the deduced youth of the asymmetries, the perturber in this scenario is required to make a single close approach in the recent past (i.e. $\sim$ $^5$ ago)."
" This can be accounted for by making the perturber-:2 Pie svstem an initial wide binary that is perturbed by a massive. passing field star 10"" in the past. such that the perturber is sent into a close approach orbit about 3 Pic."," This can be accounted for by making the $\beta$ Pic system an initial wide binary that is perturbed by a massive passing field star $\sim$ $^5$ in the past, such that the perturber is sent into a close approach orbit about $\beta$ Pic."
 A more probable solution is to arrange [or 3 Pie to have originally been a hierachical multiple system (cf, A more probable solution is to arrange for $\beta$ Pic to have originally been a hierachical multiple system (cf.
 LID 141569: Weinberecr et al., HD 141569; Weinberger et al.
 2000) that became, 2000) that became
 , 
dusty disk and envelope around a central illuminating star by means of a Monte Carlo algorithm.,dusty disk and envelope around a central illuminating star by means of a Monte Carlo algorithm.
 After the self-consistent computation of the dust temperatures throughout the disk. the SED and images at different wavelengths are computed with a ray-tracing procedure.," After the self-consistent computation of the dust temperatures throughout the disk, the SED and images at different wavelengths are computed with a ray-tracing procedure."
 The density model for the circumstellar disk and envelope we have used in our modeling is defined in cylindrical coordinates (7.z) by the following The density of the disk is given by where ro is set to the stellar radius and {η is the vertical pressure scale height.," The density model for the circumstellar disk and envelope we have used in our modeling is defined in cylindrical coordinates $(r,z)$ by the following The density of the disk is given by where $r_0$ is set to the stellar radius and $h$ is the vertical pressure scale height."
 The parameter {1 can increase with the radial distance from the center. resulting in a flaring of the disk: This disk is embedded in a spherical envelope with density gradients: and For our computations. we used a grid size ofR25000 AU. 130 radial grid points. 50 angular grid points (with. 10 extra angular points to refine the grid near the equator). a gas-to-dust mass ratio of 100:1. and 10? photon packets for the Monte Carlo simulations of each model.," The parameter $h$ can increase with the radial distance from the center, resulting in a flaring of the disk: This disk is embedded in a spherical envelope with power-law density gradients: and For our computations, we used a grid size of $R = 5000$ AU, 130 radial grid points, 50 angular grid points (with 10 extra angular points to refine the grid near the equator), a gas-to-dust mass ratio of 100:1, and $10^5$ photon packets for the Monte Carlo simulations of each model."
 For the dust opacities we used the updated model of Draine&Lee(1984) with silicate and graphite., For the dust opacities we used the updated model of \citet{Draine84} with silicate and graphite.
" We assumed a canonical grain size distribution nta)e«77? with a minimum grain size of 0.005 jm. For the maximum grain size we used a density dependent value ranging from ej,= 0.254mm for p<lolg/cm? UP to das=LO gem forp>107g/eny."," We assumed a canonical grain size distribution $n(a) \propto a^{-3.5}$ with a minimum grain size of $0.005\,\mu$ m. For the maximum grain size we used a density dependent value ranging from $a_{\rm max} = 0.25\,\mu$ m for $\rho < 10^{-17}\,{\rm g/cm^2}$ up to $a_{\rm max} = 10\,\mu$ m for $\rho > 10^{-13}\,{\rm g/cm^2}$."
" Our density model is described by eleven free parameters: (1) Stellar luminosity Z,; (2) Stellar temperature 7.: (3) Disk density power law a: (4) Disk flaring parameter6: (5) Disk vertical scale height /ig: (6) Disk density paio: (7) Inner envelope density power law exponent y: (8) Outer envelope density power law exponent 6: (9) Envelope characteristic radius Rey: (10) Envelope density pen: (11) Inclination i."," Our density model is described by eleven free parameters: (1) Stellar luminosity $L_\ast$; (2) Stellar temperature $T_{\ast}$ ; (3) Disk density power law $\alpha$; (4) Disk flaring parameter $\beta$; (5) Disk vertical scale height $h_0$ ; (6) Disk density $\rho_{\rm disk, 0}$ ; (7) Inner envelope density power law exponent $\gamma$; (8) Outer envelope density power law exponent $\delta$; (9) Envelope characteristic radius $R_{\rm env}$; (10) Envelope density $\rho_{\rm env, 0}$; (11) Inclination $i$."
 Furthermore. as the object is obviously located inside a dark cloud. we also considered the effects of foreground extinction.," Furthermore, as the object is obviously located inside a dark cloud, we also considered the effects of foreground extinction."
" A general problem of such a radiative transfer modeling is that the high dimensionality and the complicated topology of the parameter space make a search for the ""best model"" very difficult in practice.", A general problem of such a radiative transfer modeling is that the high dimensionality and the complicated topology of the parameter space make a search for the “best model” very difficult in practice.
 Simple scanning of the parameter space 1s not feasible: for example. even an extremely coarse discretization of only 5 different values for each of the I] parameters would already require the computation (and evaluation) of about 50 million different models.," Simple scanning of the parameter space is not feasible: for example, even an extremely coarse discretization of only 5 different values for each of the 11 parameters would already require the computation (and evaluation) of about 50 million different models."
 Another problem is how to evaluate the fit-quality of a specific model., Another problem is how to evaluate the fit-quality of a specific model.
 For the comparison of the model SED to the observed SED. a simple y analysis is easy to implement.," For the comparison of the model SED to the observed SED, a simple $\chi^2$ analysis is easy to implement."
 A quantitative evaluation of the model images. however. is not so straightforward.," A quantitative evaluation of the model images, however, is not so straightforward."
 First attempts to compare the individual pixel values in the model images to those in the observed images were not successful., First attempts to compare the individual pixel values in the model images to those in the observed images were not successful.
 Instead. we performed a quantitative assessment of the most important morphological features in the images.," Instead, we performed a quantitative assessment of the most important morphological features in the images."
 This was implemented by computing 2D discrete cosme transformations (DCTs) of the model images and comparing 192 DCT coefficients per model image to those derived from the observed images., This was implemented by computing 2D discrete cosine transformations (DCTs) of the model images and comparing 192 DCT coefficients per model image to those derived from the observed images.
" In this way we made sure that nodel images classified as ""good fits"" showed a central dark lane and a roundish nebulosity above and below this lane.", In this way we made sure that model images classified as “good fits” showed a central dark lane and a roundish nebulosity above and below this lane.
 —1 the first part of our modeling. we computed several thousand models to explore the main effects of the individual parameters on the fit quality.," In the first part of our modeling, we computed several thousand models to explore the main effects of the individual parameters on the fit quality."
" Based on these initial results. we then implemented a ""genetic algorithm"". in which small random changes of some of the parameters of good models are introduced in order to find an even better model."," Based on these initial results, we then implemented a “genetic algorithm”, in which small random changes of some of the parameters of good models are introduced in order to find an even better model."
 In total. more than 0000 models were computed.," In total, more than 000 models were computed."
 It was very easy to find models that reproduce the SED very well., It was very easy to find models that reproduce the SED very well.
 However. the images of these models deviate considerably from the observed images: they generally produced a much thicker dark lane. and often the shape of the upper and lower reflection nebulosity was more cone-like than hemisphere-like (as observed).," However, the images of these models deviate considerably from the observed images: they generally produced a much thicker dark lane, and often the shape of the upper and lower reflection nebulosity was more cone-like than hemisphere-like (as observed)."
 On the otherhand. we alsofoundseveral models that reproduced the near-infrared," On the otherhand, we alsofoundseveral models that reproduced the near-infrared"
be coustaut.,be constant.
 Here we clemmoustrate this analytically aud derive conditious for the infrared opacity [uuction such that the atinosphere cau be convective., Here we demonstrate this analytically and derive conditions for the infrared opacity function such that the atmosphere can be convective.
 If both the infrared and optical absorption coefficients are constant with pressure. the temperat profiles will never include a radiative-couvective boundary. regardless of how deep the bottom boundary is set or how high of an internal heat [lux is used.," If both the infrared and optical absorption coefficients are constant with pressure, the temperature-pressure profiles will never include a radiative-convective boundary, regardless of how deep the bottom boundary is set or how high of an internal heat flux is used."
 From Equation 27 of we can calculate that the night side temperature-pressure [or an atinosphere with constant. absorption coelficients will follow: and as the optical thickness goes to infinity. d(lnT)/d(lnP) reaches a maximum: value of 0.25.," From Equation 27 of \citet{Guillot2010} we can calculate that the night side temperature-pressure for an atmosphere with constant absorption coefficients will follow: and as the optical thickness goes to infinity, $d (\ln T)/ d (\ln P)$ reaches a maximum value of 0.25."
 Convection occurs when d(lnT)/d(lnP)zορ: iuthe case of a diatomic gas {ο=0.286 and an atiuosphere with constant absorption coellicieuts will never be convective., Convection occurs when $d (\ln T)/ d (\ln P) \geq R/c_p$; inthe case of a diatomic gas $R/c_p=0.286$ and an atmosphere with constant absorption coefficients will never be convective.
 For au atiuosphere in which the infrared absorption coellicieut scales exponentially with pressure (as per Equation 2)). the formalisin of Cuillot(2010) can be expanded (see Equation A2)) to lind that the uight side profile will follow: where the optical depth is uo longer linear with pressure: 7=(μεο)ία).," For an atmosphere in which the infrared absorption coefficient scales exponentially with pressure (as per Equation \ref{eqn:kir}) ), the formalism of \citet{Guillot2010} can be expanded (see Equation \ref{eqn:tprof}) ) to find that the night side profile will follow: where the optical depth is no longer linear with pressure: $\tau=(k_{\mathrm{IR,0}}/g)(P/P_{\mathrm{ref}})^\alpha$."
 As the optical depth goes to infinity. d(luT)/d(lnP) reaches a ruaximum value of (o+1)/I.," As the optical depth goes to infinity, $d (\ln T)/ d (\ln P)$ reaches a maximum value of $(\alpha+1)/4$."
" The atinosphere will be convective at depth when: (a+1)/1>οd/e,: lor the case of a diatomic gas this requires a>1/7.", The atmosphere will be convective at depth when: $(\alpha+1)/4 \geq R/c_p$; for the case of a diatomic gas this requires $\alpha \geq 1/7$.
 Our code trausitious to using fluxes [rom the diffusion approximation in the deep. optically thick attnosphere.," Our code transitions to using fluxes from the diffusion approximation in the deep, optically thick atmosphere."
 Arras&Bildsten(2006) solve for analytic pressure-temperature proliles deep in a gas glant atmosphere. assumiug flux-linited diffusion aud an absorption coellicient that scales as a powerlaw with pressure (and temperature).," \citet{Arras2006} solve for analytic pressure-temperature profiles deep in a gas giant atmosphere, assuming flux-limited diffusion and an absorption coefficient that scales as a powerlaw with pressure (and temperature)."
 Their requirement [or a couvective zoue to exist (Vs2 Vad). when converted into our notation. is: (a+1)/1 ορ.," Their requirement for a convective zone to exist $\nabla_\infty \geq \nabla_{\mathrm{ad}}$ ), when converted into our notation, is: $(\alpha+1)/4 \geq R/c_p$ ."
 This is consistent with the result above aud sets a robust requirement for our mocels., This is consistent with the result above and sets a robust requirement for our models.
For comparison. the upper πιά on the£emporal a-variation obtained frou high-redshift quasar absorbers is Δια]2 ppm (Sect.,"For comparison, the upper limit on the $\alpha$ -variation obtained from high-redshift quasar absorbers is $|\Delta \alpha/\alpha| < 2$ ppm (Sect."
 1)., 1).
 Tf dependence of coustauts ou the ambicnt matter deusitv cdomduates over teniporal (costuological). as sugeested in clhameleonu-like scalar field models. then one mav expect that [Aafal<0.2 ppl at high redshifts as well. since quasar absorbers have gas densitics simular to those in the interstellar clouds.," If dependence of constants on the ambient matter density dominates over temporal (cosmological), as suggested in chameleon-like scalar field models, then one may expect that $|\Delta\alpha/\alpha| < 0.2$ ppm at high redshifts as well, since quasar absorbers have gas densities similar to those in the interstellar clouds."
 Takine info accotut that the predicted cjuges iu e aud fr are not independent aud that p-variations may exceed variations ina (e.c. Calmet Fritzsch 2002: Laugacker 22002: Dine 22003: Flambamn 9950011. even a lower bouid of [Aofa]0.03 pp is conceivable within the ienuework of the chameleon models.," Taking into account that the predicted changes in $\alpha$ and $\mu$ are not independent and that $\mu$ -variations may exceed variations in $\alpha$ (e.g., Calmet Fritzsch 2002; Langacker 2002; Dine 2003; Flambaum 2004), even a lower bound of $|\Delta\alpha/\alpha| \leq 0.03$ ppm is conceivable within the framework of the chameleon models."
 We note that if a theoretical prediction [Aafal<[Apfp] is valid. then ΔΕΕςApp. iud. hence. the F-estimate with a further order of magnitude inprovemoenut in sensitivity will provide an indepeudent test of the tentative cliauge of ji.," We note that if a theoretical prediction $|\Delta\alpha/\alpha| \ll |\Delta\mu/\mu|$ is valid, then $\Delta F/F \approx -\Delta\mu/\mu$, and, hence, the $F$ -estimate with a further order of magnitude improvement in sensitivity will provide an independent test of the tentative change of $\mu$."
 The factors limiting accuracy of the curent estimate of aat ;=0 are a relatively low spectral resolution of the available observatious in subnuu- and nuu-wave bands. a rather large uncertainty of the rest frequencies of the 1]] FS lines. aud a sinall ος of objects observed in both 1]] aud CO transitions.," The factors limiting accuracy of the current estimate of at $z = 0$ are a relatively low spectral resolution of the available observations in submm- and mm-wave bands, a rather large uncertainty of the rest frequencies of the ] FS lines, and a small number of objects observed in both ] and $^{13}$ CO transitions."
 Modern telescopes like the recently launched Ierschel Space Observatory can provide for Calactic objects the cAoectral resolution as Ligh as 3hans! (e.g.Yo the Heterodvue Tustrmment for the Far Infrared. HIFI has resolving power A7? 10°).," Modern telescopes like the recently launched Herschel Space Observatory can provide for Galactic objects the spectral resolution as high as 30 (e.g., the Heterodyne Instrument for the Far Infrared, HIFI, has resolving power $R = 10^7$ )."
 This means trat the positions of the [C1] FS lires can be measured with t1ο uncertainty of ~35., This means that the positions of the ] FS lines can be measured with the uncertainty of $\sim$ 3.
 In the near future. high precision measurements will be also available with the Atacama Large Millimeter/ubinilliuieter Array (ALAA). the Stratospheric Observatory For Tufrared AsTOLOLUV (SOFIA). the Cornell Caltech Atacama Telescope (CCAT) aud others.," In the near future, high precision measurements will be also available with the Atacama Large Millimeter/submillimeter Array (ALMA), the Stratospheric Observatory For Infrared Astronomy (SOFIA), the Cornell Caltech Atacama Telescope (CCAT) and others."
 Thus. any firther advances in exploring AF’F depend crucially on new laoratory nieasureniens of he ICY] FS frequencies.," Thus, any further advances in exploring $\Delta F/F$ depend crucially on new laboratory measurements of the ] FS frequencies."
 IE these frequencies will be known with uncertaintics of a fowi. then the paraieter AF/F can be probed at the level of 10.5 which would be comparable with tle non-zero signal iun the spatial variation of the clectrou-to-protou lass ratio fr.," If these frequencies will be known with uncertainties of a few, then the parameter $\Delta F/F$ can be probed at the level of $10^{-8}$ which would be comparable with the non-zero signal in the spatial variation of the electron-to-proton mass ratio $\mu$ ."
of 30. taken within a 2 week time span. we can reconstruct most variables present in tje data (with those periods and amplitudes in the range ¢etermined in Section 3.2)).,"of 30, taken within a 2 week time span, we can reconstruct most variables present in the data (with those periods and amplitudes in the range determined in Section \ref{an:fmp}) )."
 OL course. this argument assitmes that the lighteurves ο ‘the variables are close to sinusoidal.," Of course, this argument assumes that the lightcurves of the variables are close to sinusoidal."
 Fig., Fig.
 v7. presents the cistribution of variables we lind in the FSVS according to he period and zunplitude o “the variability as well as their cumulative period. clistribution., \ref{res:timeamp:fig2} presents the distribution of variables we find in the FSVS according to the period and amplitude of the variability as well as their cumulative period distribution.
 The top panels consider the total number of variables in the trusted range of periods ancl amplitudes. (689). where the error on the periods and amplitudes is less than 30. per cent., The top panels consider the total number of variables in the trusted range of periods and amplitudes (689) where the error on the periods and amplitudes is less than 30 per cent.
 “Phe bottom panels. present. the distributions when we only take the svstems where the period and. amplitude determined. has a maximum error of LO per cent., The bottom panels present the distributions when we only take the systems where the period and amplitude determined has a maximum error of 10 per cent.
 In both cases we find that most systems lie at short. periods and low amplitudes. with only a few systems showing larger amplitudes ancl periods.," In both cases we find that most systems lie at short periods and low amplitudes, with only a few systems showing larger amplitudes and periods."
 We find that 50 per cent of the objects show periods below 6 hours with peaks in the 30 per cent error distribution at ~24mmin. 70.03 delays (~43 mimin) 0.12 ddavys (~2.9 hhours). 70.179 delays (19 hhours)). ~1.8clelavs and ~4ededays. and in the 19 per cent distribution at ~0.12 ddavs (~2-9bhours).," We find that 50 per cent of the objects show periods below 6 hours with peaks in the 30 per cent error distribution at $\sim$ min, $\sim$ days $\sim$ min), $\sim$ days $\sim$ hours), $\sim$ days $\sim$ hours), $\sim$ days and $\sim$ days, and in the 10 per cent distribution at $\sim$ days $\sim$ hours)."
 In the 30 per cent period. distribution. the clump of sources between 24 and 36mmin L778«fogP< 1.6) contains 67 sources.," In the 30 per cent period distribution, the clump of sources between $\sim$ 24 and min $-$ $<log P<$$-$ 1.6) contains 67 sources."
 To confirm that these are short. period. variables. anc not just à systematic problem caused. by the sampling (after all the minimum period. we are sensitive to is mmin) they were inspected. by eve resulting in SO per cent being bona-fide short period. variables with the remaining 20 per cent showing only one point olf the average brightness of the target. and thus giving the short. period. result. based only on one point variability.," To confirm that these are short period variables, and not just a systematic problem caused by the sampling (after all the minimum period we are sensitive to is min) they were inspected by eye resulting in 80 per cent being bona-fide short period variables with the remaining 20 per cent showing only one point off the average brightness of the target and thus giving the short period result based only on one point variability."
 These one point olf sources are not present in the LO per cent sample., These one point off sources are not present in the 10 per cent sample.
 Regarding the amplitude distribution. 50 per cent of the objects show amplitudes lower than 70.07 mmag in the 30 »er Cent error sample and lower than ~0.12 mmag in the 10 per cent sample.," Regarding the amplitude distribution, 50 per cent of the objects show amplitudes lower than $\sim$ mag in the 30 per cent error sample and lower than $\sim$ mag in the 10 per cent sample."
 When we combine the number of sources we find. per »eriod and amplitude bin with the sensitivity of the Moating mean periocogram search. plotted in Fig. 4..," When we combine the number of sources we find per period and amplitude bin with the sensitivity of the floating mean periodogram search, plotted in Fig. \ref{an:fmp:fig3},"
 we obtain ower limits for the space density of variables. ie. number of variables per square degree. versus. period.," we obtain lower limits for the space density of variables, i.e. number of variables per square degree, versus period."
 These are xesented in the form of a histogram in Fig. S.., These are presented in the form of a histogram in Fig. \ref{res:timeamp:fig5}.
 We see four distinct peaks in the distribution centred at 6 hours. 1 clay. 3.75 davs and 12.75 davs with a somewhat less significant »eak at Gdaws.," We see four distinct peaks in the distribution centred at 6 hours, 1 day, 3.75 days and 12.75 days with a somewhat less significant peak at 6 days."
 The highest density of variables show periods xlow 12 hours., The highest density of variables show periods below 12 hours.
 Phese periods include CVs. RR Lyr stars. and other short period pulsators such as 0 SScuti stars.," These periods include CVs, RR Lyr stars, and other short period pulsators such as $\delta$ Scuti stars."
 The »riod range centred at 1 day includes also possible CVs. th Lyr and other pulsators like + DDoradus stars and Pop LL Cepheids.," The period range centred at 1 day includes also possible CVs, RR Lyr and other pulsators like $\gamma$ Doradus stars and Pop II Cepheids."
 At 3.75 days we would still find some longer »eriod. CVs. 5 DDoradus stars. Pop LE Cepheids and longer »eriod pulsators such as subedwarf D stars.," At 3.75 days we would still find some longer period CVs, $\gamma$ Doradus stars, Pop II Cepheids and longer period pulsators such as subdwarf B stars."
 At periods around 12.75 days. we expect to find. apart from binaries with those orbital periods. Pop IE Cepheids contributing to the space density of variables.," At periods around 12.75 days, we expect to find, apart from binaries with those orbital periods, Pop II Cepheids contributing to the space density of variables."
 Ixeeping in mind the uncertainty of the variability timescales determined. when we combine the variability. information with the colour information available for the FSVS we obtain the colour-colour diagrams shown in Fig. 9..," Keeping in mind the uncertainty of the variability timescales determined, when we combine the variability information with the colour information available for the FSVS we obtain the colour-colour diagrams shown in Fig. \ref{res:ccdiag:fig1}."
 We find that. if we take only the sources with less than 30 per cent errors in their timescales and. the amplitudes. 344 point sources show variabilities shorter than hh ddavs).," We find that, if we take only the sources with less than 30 per cent errors in their timescales and the amplitudes, 344 point sources show variabilities shorter than h days)."
 These short imescale variables are found along the main sequence in the colour-colour diagram (see first panel of Fig. 9)).," These short timescale variables are found along the main sequence in the colour-colour diagram (see first panel of Fig. \ref{res:ccdiag:fig1}) ),"
 where we expect to find for example ὁ SScuti stars. as well as above he main sequence and in more extreme colour areas usually illeck by binary systems. where one of the components is Xue and the other red. c.g. detached red dwarlwhite dwacl xnaries.," where we expect to find for example $\delta$ Scuti stars, as well as above the main sequence and in more extreme colour areas usually filled by binary systems where one of the components is blue and the other red, e.g. detached red dwarf-white dwarf binaries."
 At these short timescale variabilities we also fin a clump of objects above the. so called. blue. cut-olf a V)«0.38.," At these short timescale variabilities we also find a clump of objects above the, so called, blue cut-off at $-$ $<$ 0.38."
 The blue cut-oll of the main sequence results rom the combination of the number density of cülferen spectral tvpes and the scale. height of the Galaxy., The blue cut-off of the main sequence results from the combination of the number density of different spectral types and the scale height of the Galaxy.
 The colours and the short variability timescales. of the order of characteristic close binary orbital periods. suggest that these sources above the blue cut-olf are possibly interacting binary systems of the CW type or detached. binary systems such as subchwarl D. binaries.," The colours and the short variability timescales, of the order of characteristic close binary orbital periods, suggest that these sources above the blue cut-off are possibly interacting binary systems of the CV type or detached binary systems such as subdwarf B binaries."
 Longer coverage. better sampled lighteurves combined with spectroscopy are necessary in order to identify the sources.," Longer coverage, better sampled lightcurves combined with spectroscopy are necessary in order to identify the sources."
 The variable sources with timescales shorter than hh represent 50 per cent of the total number of short timescale variables in the survey., The variable sources with timescales shorter than h represent 50 per cent of the total number of short timescale variables in the survey.
 Fig., Fig.
 10 shows how most of those short. period. sources subdivide in smaller. variability, \ref{res:ccdiag:fig1b} shows how most of those short period sources subdivide in smaller variability
infrared.,infrared.
 The diffuse UV radiation is about of the total radiation emitted [rom the LMC' (Parkeretal.1998). and understanding its distribution is important to models of galactic evolution., The diffuse UV radiation is about of the total radiation emitted from the LMC \citep{Parker98} and understanding its distribution is important to models of galactic evolution.
 More recently. Coleetal.(1999a) used the rocket-borne to map the scattered light in the near ultraviolet (2150 A)). fincling (hat the scattered light is actually a complex combination of the relative geometry of the dust and the stags.," More recently, \citet{Cole99a} used the rocket-borne to map the scattered light in the near ultraviolet (2150 ), finding that the scattered light is actually a complex combination of the relative geometry of the dust and the stars."
 It is not sufficient to merely have bright stars or to have dust: both must be present to show the scattered light., It is not sufficient to merely have bright stars or to have dust: both must be present to show the scattered light.
 In (his work. we use serendipitous observations made with the ιο report. for the first time. measurements of the diffuse FUV (1000 - 1150 À)) emission in an external galaxy.," In this work, we use serendipitous observations made with the to report, for the first time, measurements of the diffuse FUV (1000 - 1150 ) emission in an external galaxy."
 TheFUSE spacecralt ancl ils mission have been described by Moosοἱal.(2000) aid sahnowetal.(2000)., The spacecraft and its mission have been described by \citet{Moos00} and \citet{Sahnow00}.
. The primary purpose of the nmüssion was (o take high resolution spectra (A/ AAez 20.000) of galactic and extragalactic sources.," The primary purpose of the mission was to take high resolution spectra $\lambda/\Delta \lambda$ $\approx$ 20,000) of galactic and extragalactic sources."
"FUSE included 3 apertures: the LIRS (11.25""x 2207): the MDRS x 2207): and the LWRS x 3307)). all of which obtained data simultaneouslv."," included 3 apertures: the HIRS x ); the MDRS x ); and the LWRS x ), all of which obtained data simultaneously."
" Although only the LWRS with its relatively large field of view was useful for diffuse observations. there were many fields in which the primary aperture was either the HIRS or the MDRS leaving the LWRS aperture (separated [rom the other two apertures by and2200"".. respectively) to observe a nominally blank region of the skv."," Although only the LWRS with its relatively large field of view was useful for diffuse observations, there were many fields in which the primary aperture was either the HIRS or the MDRS leaving the LWRS aperture (separated from the other two apertures by and, respectively) to observe a nominally blank region of the sky."
 Murthy.&Sahnow(2004) have described the analvsis of these serendipitous backeround observations aud we have followed their extraction of diffuse surface brishtnesses from theFUSE spectra. except Chat we have used the latest version of the data pipeline software (CalFUSE v3.2: Dixonetal. (2007))).," \citet{Murthy04} have described the analysis of these serendipitous background observations and we have followed their extraction of diffuse surface brightnesses from the spectra, except that we have used the latest version of the data pipeline software (CalFUSE v3.2; \citet{Dixon07}) )."
 This involved treatingFUSE as a broad band photometer and collapsing the spectra into (wo wavelength: bands. per detector. excluding the terrestrial airglow lines (primarily Ly).," This involved treating as a broad band photometer and collapsing the spectra into two wavelength bands per detector, excluding the terrestrial airglow lines (primarily $\beta$ )."
 This resulted in seven wavelength bands with elective wavelengths of 1004 (LAL). 1058 (142). L117 (1DI). 1157 192). 1159 (PAL). L112 (242). 1056 (2DI).," This resulted in seven wavelength bands with effective wavelengths of 1004 (1A1), 1058 (1A2), 1117 (1B1), 1157 (1B2), 1159 (2A1), 1112 (2A2), 1056 (2B1)."
 The instrumental background. was derived from strips off the spectrum and subtracted from the band f[Iuxes., The instrumental background was derived from strips off the spectrum and subtracted from the band fluxes.
 Our wavelength: bands are shown in Fig., Our wavelength bands are shown in Fig.
 1. on top of a spectrum of the diffuse emission in N11. one of the brighter diffuse regions of the LAIC.," \ref{Fig1} on top of a spectrum of the diffuse emission in N11, one of the brighter diffuse regions of the LMC."
 The integrated [αν of each band is marked by a solid circle at ils effective wavelength., The integrated flux of each band is marked by a solid circle at its effective wavelength.
 There are more than 600FUSE pointings in and around the LMC ancl we have examined all for suitability [or dilfiise measurements., There are more than 600 pointings in and around the LMC and we have examined all for suitability for diffuse measurements.
 We immecdiately rejected all observations which specifically observed a bright star in (he LAWS aperture but there were others where a star was coincidentallv in the aperture., We immediately rejected all observations which specifically observed a bright star in the LWRS aperture but there were others where a star was coincidentally in the aperture.
 These were identified and rejected through (heir FEWIIM.," These were identified and rejected through their FWHM,"
ancl the above arguments don't apply to them.,and the above arguments don't apply to them.
 Therefore. when considering the overall burst duration distribution we expect a flat section for durations significantly lower (han 50 sec down to the duration where these non-Collapsars dominate.," Therefore, when considering the overall burst duration distribution we expect a flat section for durations significantly lower than 50 sec down to the duration where these non-Collapsars dominate."
 The observed duration of a GRB is characterized using Τομ2£5. during which 90% of the (Inence is acciunulated.," The observed duration of a GRB is characterized using $T_{90}\approx t_\gamma$, during which $90\%$ of the fluence is accumulated."
 We use the data from the three major GRD detectors: DATSE. and Fermi GDM.," We use the data from the three major GRB detectors: BATSE, and Fermi GBM."
 For BATSE we use the current. catalog (04/21/91-05/26/00: containing 2041 bursts)., For BATSE we use the current catalog (04/21/91-05/26/00; containing 2041 bursts).
 The data ofSwifl is taken [rom its online (12/17/0408/27/11: containing 582 bursts)., The data of is taken from its online (12/17/04-08/27/11; containing 582 bursts).
 Fermi data is extracted from GCNs using the GRBox (08/12/08-07/21/11: containing 194 bursts)., Fermi data is extracted from GCNs using the GRBox (08/12/08-07/21/11; containing 194 bursts).
 Each data set is binnecl into equally spaced logarithmic bins. where (he minimal number of events per bin is limitel to five (Pressetal.1989).," Each data set is binned into equally spaced logarithmic bins, where the minimal number of events per bin is limited to five \citep{NumRes}."
.. A bin with less than five events is merged with its neighbor., A bin with less than five events is merged with its neighbor.
 We use a 7 minimization to look for the longest logarithmic time interval that is consistent with a flat line within 1 σ. where the only [ree parameter is the normalization.," We use a $\chi^2$ minimization to look for the longest logarithmic time interval that is consistent with a flat line within 1 $\sigma$, where the only free parameter is the normalization."
 We verily that varying the bin size doesnt change the length of the plateau by much., We verify that varying the bin size doesn't change the length of the plateau by much.
 Fie., Fig.
" 1 depicts the observed distribution of Loy. p,(Zoo). for the three major GRB satellites."," 1 depicts the observed distribution of $T_{90}$, $p_\gamma(T_{90})$, for the three major GRB satellites."
" Note that we show here the quantitw p,(Z54)=dN/dI and not dN/dloeT traditionally shown in such plots (e.g..IXouveliotouοἱal.1993)."," Note that we show here the quantity $p_\gamma(T_{90})=dN/dT$ and not ${dN}/{d\log
T}$ traditionally shown in such plots \citep[e.g.,][]{Kouveliotou93}."
. The best fitted flat regions are hiehlighted in a solid bold line on top of each distribution., The best fitted flat regions are highlighted in a solid bold line on top of each distribution.
 In all satellites these plateaus range about an order of magnitude in durations (BATSE 5-25 sec. 3.6/6 \7/dol: sec. 8.9/7 A? /dol: Fermi 1.2-31 sec. 4.1/6 42/dof).," In all satellites these plateaus range about an order of magnitude in durations (BATSE 5-25 sec, 3.6/6 $\chi^2$ /dof; 0.7-21 sec, 8.9/7 $\chi^2$ /dof; Fermi 1.2-31 sec, 4.1/6 $\chi^2$ /dof)."
 The extent of the platean varies slightly [rom one detector to another., The extent of the plateau varies slightly from one detector to another.
 This is expected given the different detection threshold sensitivities in dillerent energy. windows (see below)., This is expected given the different detection threshold sensitivities in different energy windows (see below).
 At the high end of the plateau the τομ distribution decreases rapidly aud can be fitted at long cdurations (100 8) by a power law with an index. a. in the range —4<a<—3.," At the high end of the plateau the $T_{90}$ distribution decreases rapidly and can be fitted at long durations $>$ 100 s) by a power law with an index, $\alpha$, in the range $-4<\alpha<-3$."
 The existence of the plateaus and (heir duration range (ου2—25 s) agrees well with the expectation of the Collapsar model., The existence of the plateaus and their duration range $\sim 2-25$ s) agrees well with the expectation of the Collapsar model.
 ILowever. one cannot exclude the possibility Chat the origin of the observed flat sections is unrelated to the effect. of the jet breakout time that we discuss above.," However, one cannot exclude the possibility that the origin of the observed flat sections is unrelated to the effect of the jet breakout time that we discuss above."
 For example. these plateaus may somehow arise coincidentallv [rou the combination of two distributions: one increasing (LGRBs) and one decreasing (SGRBs).," For example, these plateaus may somehow arise coincidentally from the combination of two distributions: one increasing (LGRBs) and one decreasing (SGRBs)."
relativistic flows could possibly be detectable and. therefore. provide a new probe οἱ. e.g. GRB central engine physics and expansion dynamics.,"relativistic flows could possibly be detectable and, therefore, provide a new probe of, e.g., GRB central engine physics and expansion dynamics."
 We acknowledge discussions with Ix. Abazajian. D. INirkman. J. OMeara. and ο. Woosley.," We acknowledge discussions with K. Abazajian, D. Kirkman, J. O'Meara, and S. Woosley."
 This work was supported in part by NSF Grant PILY-0099499 at UCSD and DOE Sci-Dac supernova grants at LLNL and UCSD., This work was supported in part by NSF Grant PHY-0099499 at UCSD and DOE Sci-Dac supernova grants at LLNL and UCSD.
oxvgen-neon-WD.,oxygen-neon-WD.
 This Μα] example. and (he one presented by IIurley et al. (," This lurid example, and the one presented by Hurley et al. ("
2001). highlight an important point.,"2001), highlight an important point."
 Though (he possibility of anv particular star becoming a DS as a result of a dynamical encounter is quite random. They also show that if a binary emerees from (he indiscriminate and short-lived relationship (hat is an exchange interaction il will more often than not comprise the (vo most massive stars involved in the interaction. i.e. in this society it is desirable to be heavy.," Though the possibility of any particular star becoming a BS as a result of a dynamical encounter is quite random, They also show that if a binary emerges from the indiscriminate and short-lived relationship that is an exchange interaction it will more often than not comprise the two most massive stars involved in the interaction, i.e. in this society it is desirable to be heavy."
 In a tvpical GRAPE-6 simulation with V=20000. initially comprised of 18000 single stars and 2000 binaries. where the evolution was followed for 5 Gyr. the number of exchange interactions observed was 500.," In a typical GRAPE-6 simulation with $N = 20\,000$, initially comprised of $18\,000$ single stars and $2\,000$ binaries, where the evolution was followed for $5\,$ Gyr, the number of exchange interactions observed was $\sim 500$."
 These involved. 730 different stars with some stars taking part in multiple interactions., These involved 730 different stars with some stars taking part in multiple interactions.
 The munber of stars that swapped partner once was 494. twice was 105. three (times was 48. four times was 27. and 14 stars swapped partner on five occasions.," The number of stars that swapped partner once was 494, twice was 105, three times was 48, four times was 27, and 14 stars swapped partner on five occasions."
 Amonest (hese were a number of re-amarriages where a “star changed its mind” and returned (o its original partner., Amongst these were a number of re-marriages where a “star changed its mind” and returned to its original partner.
 The component stars of one particularly [irtatious binary were actually involved in 22 exchange interactions. including 10 re-marriages.," The component stars of one particularly flirtatious binary were actually involved in 22 exchange interactions, including 10 re-marriages."
 The total number of merger events observed curing the entire simulation was 104., The total number of merger events observed during the entire simulation was 104.
 Of these. mass (ransler in a primordial binary accounted lor 46 cases while 13 mergers came from mass transfer in a binary formed via an exchange interaction.," Of these, mass transfer in a primordial binary accounted for 46 cases while 13 mergers came from mass transfer in a binary formed via an exchange interaction."
 The remainder were the result of collisions in eccentric binaries: 21 in primordial svstems where (he orbit was strongly perturbed by nearby stars., The remainder were the result of collisions in eccentric binaries: 21 in primordial systems where the orbit was strongly perturbed by nearby stars.
 Bs formation as a result of a hyperbolic collision is rare in simulations of open clusters but does occur., BS formation as a result of a hyperbolic collision is rare in simulations of open clusters but does occur.
 One example involved a 0.32.4. star Chat began life in (he core of the cluster and slowly αντος out to the hall-mass racius. owing to mass-segregation. where at 1300 Myr it collided with a 0.922. star.," One example involved a $0.32 M_\odot$ star that began life in the core of the cluster and slowly drifted out to the half-mass radius, owing to mass-segregation, where at $1\,300\,$ Myr it collided with a $0.93 M_\odot$ star."
 The relative velocity of the two stars at infinity was 2.8kms.! and the collison product was assumed to be a fullv-mixed 1.254/. MS star.," The relative velocity of the two stars at infinity was $2.8 {\rm km} \, {\rm s}^{-1}$ and the collison product was assumed to be a fully-mixed $1.25 M_\odot$ MS star."
 When the cluster was 3850 Myr old (his star first appeared as a BS and by 4500 Myr. when (he simulation ended. it had sunk inside the cluster core.," When the cluster was $3\,850\,$ Myr old this star first appeared as a BS and by $4\,500\,$ Myr, when the simulation ended, it had sunk inside the cluster core."
 The incidence of direct collisions will be greater in the higher density. conditions of a globular cluster simulation., The incidence of direct collisions will be greater in the higher density conditions of a globular cluster simulation.
 However. these are abrupt encounters and much less interesting than the sociable nature of exchange interactions ancl the binary systems (μον produce.," However, these are abrupt encounters and much less interesting than the sociable nature of exchange interactions and the binary systems they produce."
 The presence of binaries also acts to magnify the chance Lf collisions because in the ease of a binary it is the semi-major axis that sets the relevant cross-section for collision. rather than the stellar radius which is used in the case of single stars.," The presence of binaries also acts to magnify the chance of collisions because in the case of a binary it is the semi-major axis that sets the relevant cross-section for collision, rather than the stellar radius which is used in the case of single stars."
 Here we have to be careful with the terminology used., Here we have to be careful with the terminology used.
 If a binary is involved in a hyperbolic anv that results will occur in a two-step process and will nol bedirech firstly a (resonant) capture may produce a hierarchical svstem and subsequently a, If a binary is involved in a hyperbolic any that results will occur in a two-step process and will not be: firstly a (resonant) capture may produce a hierarchical system and subsequently a
(e.g.. Chae2005:οἱal. 2006)) and galaxy evolutions (e.g.. al.2003)).,"(e.g., \citealt{Cha05,Cha06}) ) and galaxy evolutions (e.g., \citealt{CM03,Ofe03}) )."
 The sample from the completed CLASS. in particular its subsample of 13 lenses strictly salislving well-defined selection criteria (the CLASS statistical sample: Browneetal.2003:Chae 2003)). was first extensively analvzed by Chaeetal.(2002) ancl Chae(2003).. who found μις&0.3 assuming; a flat cosmology and adopting non-evolving galaxy. populations.," The sample from the completed CLASS, in particular its subsample of 13 lenses strictly satisfying well-defined selection criteria (the CLASS statistical sample; \citealt{Bro03,Cha03}) ), was first extensively analyzed by \citet{Cha02} and \citet{Cha03}, who found $\Om \approx 0.3$ assuming a flat cosmology and adopting non-evolving galaxy populations."
 Mitchelletal.(2005) re-analvzed the CLASS statistical sample based on the velocity dispersion Iunction (WDE) of early-type galaxies directly derived from (he SDSS Data Release 1 (DRI: Stoughtonetal. 2002)) galaxies (Shethetal.2003))., \citet{Mit05} re-analyzed the CLASS statistical sample based on the velocity dispersion function (VDF) of early-type galaxies directly derived from the SDSS Data Release 1 (DR1; \citealt{Sto02}) ) galaxies \citealt{She03}) ).
 However. Chae(2005) fines that the Shethetal.(2003) VDE of early-type galaxies would imply a significantly underestimated abundance of earh-type galaxies based on the Wilkinson Microwave Anisotropy Probe (WAIAP) Ist vear cosmology (Spergeletal. 2003)) and the CLASS statistical sample.," However, \citet{Cha05} finds that the \citet{She03} VDF of early-type galaxies would imply a significantly underestimated abundance of early-type galaxies based on the Wilkinson Microwave Anisotropy Probe (WMAP) 1st year cosmology \citealt{Spe03}) ) and the CLASS statistical sample."
 Just recently. Choietal.(2007) have made a new measurement of the VDF of earlv-tvpe ealaxies based on the much larger SDSS Data Release 5 (DR5: Adelman-MeCarthy galaxies emploving a new and more reliable method of classilving galaxies 2005)).," Just recently, \citet{Cho06} have made a new measurement of the VDF of early-type galaxies based on the much larger SDSS Data Release 5 (DR5; \citealt{Ade07}) galaxies employing a new and more reliable method of classifying galaxies \citealt{PC05}) )."
 The Choietal.(2007) VDF has a much higher comoving number density of early-type galaxies and a dillerent shape for the lower velocity part compared with the VDF., The \citet{Cho06} VDF has a much higher comoving number density of early-type galaxies and a different shape for the lower velocity part compared with the \citet{She03} VDF.
 The Choietal.(2007). earlv-tvpe nmunber density is in favor of the results., The \citet{Cho06} early-type number density is in favor of the \citet{Cha05} results.
 The goal of this work is to improve strong lensing statistics using the SDSS DR5 VDF of earlv-tvpe galaxies., The goal of this work is to improve strong lensing statistics using the SDSS DR5 VDF of early-type galaxies.
" Our focus shall be to put independent constraints on O,,, and i, assuming a flat cosmology.", Our focus shall be to put independent constraints on $\Om$ and $w_x$ assuming a flat cosmology.
 We shall consider both no evolution and a evolution of galaxies based on the prediction by a semi-analvtical model of galaxy. formation etal. 2006)).," We shall consider both no evolution and a evolution of galaxies based on the prediction by a semi-analytical model of galaxy formation \citealt{Kan05,Cha06}) )."
 In 82. we brielly describe the data and the analvsis method.," In 2, we briefly describe the data and the analysis method."
 We present and discuss the results in 33., We present and discuss the results in 3.
 The comoving number density of galaxies as a function of velocity dispersion (0) can be described by the modified Schechter lunction @(7) given by 2005))," The comoving number density of galaxies as a function of velocity dispersion $\sigma$ ) can be described by the modified Schechter function $\phi(\sigma)$ given by \citealt{She03, Mit05}) ) dn = ) = _*"
"where (,4,25:54 are the limb-darkening passbancl-specilic coefficients. Z(1) is the passband-specifie intensity at. the center. of the stellar disc. and ye=coss. where 5 ds the angle between the line of sight and the local surface normal.","where $a_{1,2,3,4}$ are the limb-darkening passband-specific coefficients, $I(1)$ is the passband-specific intensity at the center of the stellar disc, and $\mu=\cos \gamma$, where $\gamma$ is the angle between the line of sight and the local surface normal."
 Phe central intensity is calculated for the cllective wavelengths of the observations (600 nm for the ColtoT light curves and 550 nm for the V-band light curve). using a simple blackhocky approximation.," The central intensity is calculated for the effective wavelengths of the observations (600 nm for the CoRoT light curves and 550 nm for the V-band light curve), using a simple blackbody approximation."
 For a given metallicity. the values of the passbancl-specilic limb-darkening coellicients are derived [rom the current values of the stellar elective temperature Tar and surface gravity logg in cach iteration. bybi-lincar interpolation (Pressetal. 1992)... of both «quantities from tables of Claret.(2000). for the V. light curves and of Sine(2010) for the CoRoT light. curves.," For a given metallicity, the values of the passband-specific limb-darkening coefficients are derived from the current values of the stellar effective temperature ${\rm T_{eff}}$ and surface gravity ${\rm log} \ g$ in each iteration, by interpolation \citep{press}, , of both quantities from tables of \citet{claret} for the V light curves and of \citet{sing} for the CoRoT light curves."
 The procedure is described in more detail by Djuraseviéetal. (2004)., The procedure is described in more detail by \citet{djura04}.
. The limb-darkening was applied to the cdisk in the same way. with logg corresponding to the middle of the disk radius.," The limb-darkening was applied to the disk in the same way, with $\log g$ corresponding to the middle of the disk radius."
 The results of the light-curve analysis based on the described model of AU Alon are given in Table 1.., The results of the light-curve analysis based on the described model of AU Mon are given in Table \ref{TabAUMon}.
 The first column contains parameter designations. and the following five columns give the values derived from each of the five Colo lisht-curves. with the mean values and their. estimated uncertainties.," The first column contains parameter designations, and the following five columns give the values derived from each of the five CoRoT light-curves, with the mean values and their estimated uncertainties."
 Phe uncertainties were estimated [rom a set of solutions obtained for the five observed light curves ancl three dillerent. values of the mass ratio: q;= 0.14. qo=0.17 and gq;=0.20 (chosen according to the error assigned. to the mass ratio by Desmetctal. 2010... q¢=0.17 0.03). resulting in a total of 15 values for cach parameter.," The uncertainties were estimated from a set of solutions obtained for the five observed light curves and three different values of the mass ratio: $q_1=0.14$ , $q_2=0.17$ and $q_3=0.20$ (chosen according to the error assigned to the mass ratio by \citealt{des10}, , $q=0.17\pm 0.03$ ), resulting in a total of 15 values for each parameter."
 The uncertainties given in Table 1 are the maximal deviations of these values from the mean., The uncertainties given in Table \ref{TabAUMon} are the maximal deviations of these values from the mean.
 Table 1. also lists the results ol applving a simple Roche mocel (seee.g.Djurasevié1992) to the CoRoT light curves and the results of applving the accretion disk mocel to the ground-based V-band light curves (discussed in detail in Section 4.1))., Table \ref{TabAUMon} also lists the results of applying a simple Roche model \citep[see e.g.][]{djura92} to the CoRoT light curves and the results of applying the accretion disk model to the ground-based V-band light curves (discussed in detail in Section \ref{longterm}) ).
 The first three rows of Table 1. present the number of points in the light curve (n)) the final sum of the squares of the. residuals between the observed (LOCO) and. the synthetic (LCE) licht curves. (0CY. and the root-mean-square of the resicluals Toys. respectively.," The first three rows of Table \ref{TabAUMon} present the number of points in the light curve $n$ ), the final sum of the squares of the residuals between the observed (LCO) and the synthetic (LCF) light curves, $\sum (O-C)^2$, and the root-mean-square of the residuals $\sigma_{rms}$, respectively."
 The best fit model of AU Mon contains an optically and ecometrically thick accretion cise around the hotter. more massive gainer star.," The best fit model of AU Mon contains an optically and geometrically thick accretion disc around the hotter, more massive gainer star."
 With a radius of Z2;21342. the disk is more than twice as large as the central star (4252542. ).," With a radius of $R_d\approx13 R_{\odot}$, the disk is more than twice as large as the central star $R_h\approx5 R_{\odot}$ )."
" The isk has a moderately concave shape. with central thickness old.=OAR. and the thickness at the edge of d,= 1.672..."," The disk has a moderately concave shape, with central thickness of $d_c\approx 0.4 R_{\odot}$ and the thickness at the edge of $d_e\approx1.6 R_{\odot}$ ."
 The temperature of the disk increases from at s edge. to 2;=158T0A. at the inner radius (where it is in —rermal ancl physical contact with the eainer). according to Eq. 1..," The temperature of the disk increases from at its edge, to $T_h=15870 K$ at the inner radius (where it is in thermal and physical contact with the gainer), according to Eq. \ref{eq1},"
 with the temperature profile exponent ap=6.5., with the temperature profile exponent $a_T=6.5$.
 The ellective temperature of the disk is significantly: higher than 10 temperature at its edge., The effective temperature of the disk is significantly higher than the temperature at its edge.
 We were able to model the asvounetry of the light curve very precisely by incorporating three regions of enhanced radiation on the aceretion disk: the hot spot (hs). and two bright spots (bsl and bs2).," We were able to model the asymmetry of the light curve very precisely by incorporating three regions of enhanced radiation on the accretion disk: the hot spot (hs), and two bright spots (bs1 and bs2)."
 The hot spot (hs) is situated at longitude Aj;z330. roughly. between the components of the system. at the place where the gas stream falls onto the disk.," The hot spot (hs) is situated at longitude $\lambda_{hs}\approx 330^\circ$, roughly between the components of the system, at the place where the gas stream falls onto the disk."
 Phe longitude A is measured clockwise (as viewed from the direction of the |Z-axis. which is orthogonal to the orbital plane) with respect to the line connecting the star centers (|X-axis). in the range 0°360°.," The longitude $\lambda$ is measured clockwise (as viewed from the direction of the +Z-axis, which is orthogonal to the orbital plane) with respect to the line connecting the star centers (+X-axis), in the range $0^\circ-360^\circ$."
 The temperature of the hot spot isapproximately higher then the disk edge temperature. Le. Z5;8STOOA.," The temperature of the hot spot isapproximately higher then the disk edge temperature, i.e. $T_{hs}\approx 8700 K$."
" The hot spot can be interpreted. as a rough approximation of the ""hot line? which forms at the edge of the gas stream between the components."," The hot spot can be interpreted as a rough approximation of the ""hot line"" which forms at the edge of the gas stream between the components."
" According to Atwood-Stoneetal...(2010)... a gas-stream with a temperature near SOOQOA. can explain the excess recd- and blue-shifted £4, emission near phases 0.2 ane 0.7. respectively."," According to \citet{atwood10}, a gas-stream with a temperature near $8000 K$ can explain the excess red- and blue-shifted $H_\alpha$ emission near phases 0.2 and 0.7, respectively."
 We note that this spectroscopic result 1s an independent. confirmation of our photometrically estimatec temperature of the gas stream., We note that this spectroscopic result is an independent confirmation of our photometrically estimated temperature of the gas stream.
 Although including the hot spot region into the moce significantly improves the fit. it cannot explain the light-curve asymmetry completely.," Although including the hot spot region into the model significantly improves the fit, it cannot explain the light-curve asymmetry completely."
" By introducing two adcditiona bright spots (bsl and bs2). larger than the hot spot aux located on the disk edge at Ans,2170 and As,»750. the fit becomes much better."," By introducing two additional bright spots (bs1 and bs2), larger than the hot spot and located on the disk edge at $\lambda_{bs1}\approx 170^\circ$ and $\lambda_{bs2}\approx 50^\circ$, the fit becomes much better."
 The bright spots can be related to the spiral shocks tha result from radiative cooling and form at the outer boundary of the disk., The bright spots can be related to the spiral shocks that result from radiative cooling and form at the outer boundary of the disk.
 Since the disk is larec. filling about of the eainer’s critical Roche surface. the tidal forces exerted by 10 donor can cause a spiral-shaped tidal shock in the disk - see e.g. Llecmskerk(1994).," Since the disk is large, filling about of the gainer's critical Roche surface, the tidal forces exerted by the donor can cause a spiral-shaped tidal shock in the disk - see e.g. \citet{Heemskerk}."
.. Such a shock wave can produce one or two extended: spiral arms in the outer parts of the isk., Such a shock wave can produce one or two extended spiral arms in the outer parts of the disk.
 The first arm. represented in our model by a bright spot (bs1l). is located on the disk edge at longitude As;τεnu ," The first arm, represented in our model by a bright spot (bs1), is located on the disk edge at longitude $\lambda_{bs1}\approx 170^\circ$ ."
This is also a region where we can expect loss of matter from the gas stream and the disk through the Lagrangian point L5. forming some kind of a circumbinary shell.," This is also a region where we can expect loss of matter from the gas stream and the disk through the Lagrangian point ${\rm L_3}$, forming some kind of a circumbinary shell."
 The fit was additionally improved by introducing the second bright spot (bs2). located at longitude Apsoe50°," The fit was additionally improved by introducing the second bright spot (bs2), located at longitude $\lambda_{bs2}\approx 50^\circ$."
 This spot is the largest active region. with a temperature about higher than the disk edge temperature.," This spot is the largest active region, with a temperature about higher than the disk edge temperature."
 It can be interpreted as the second spiral arm in the disk., It can be interpreted as the second spiral arm in the disk.
 We note that the svstem can also be modeled: with active regions (dark spots) on the donor. and. without the active regions on the accretion disk.," We note that the system can also be modeled with active regions (dark spots) on the donor, and without the active regions on the accretion disk."
 Such a model would explain the period-to-period variations in the light curves by the presence. development anc migration of spots over the surface of the donor.," Such a model would explain the period-to-period variations in the light curves by the presence, development and migration of spots over the surface of the donor."
 However. the model with active regions on the accretion disk seems to be more appropriate.," However, the model with active regions on the accretion disk seems to be more appropriate."
 NameA the relatively fast variations of the light curves are more likely to originate from the changes in the clisk structure. produced: by variable mass outllow from the donor. than from the motion of stellar spots. not expected to take place on these timescales.," Namely, the relatively fast variations of the light curves are more likely to originate from the changes in the disk structure, produced by variable mass outflow from the donor, than from the motion of stellar spots, not expected to take place on these timescales."
 In order to make a comparison between the results of our study and that of Desmetetal. (2010).. we made several trial runs with a simple semicectached Roche mocel of AU Aon.," In order to make a comparison between the results of our study and that of \citet{des10}, , we made several trial runs with a simple semidetached Roche model of AU Mon."
 The semicletached model cannot fit theobservations as well as the model with an accretion disk., The semidetached model cannot fit theobservations as well as the model with an accretion disk.
 The fit can be improved by using an anomalously [largegravity clarkeningexponent of the donor.as doneby Desmetetal. (2010)...," The fit can be improved by using an anomalously largegravity darkeningexponent of the donor,as doneby \citet{des10}. ."
 Our previous investigationsof anomalously high values of gravity, Our previous investigationsof anomalously high values of gravity
roughly the same spatial scale as those described here have been reported. (NeumanneSparks.DirettaandAlacehetto1996:Perlmanefal. 1998).,"roughly the same spatial scale as those described here have been reported \citep{neu97,spa96,per98}."
. In particular. recent observations of the M87 jet have detected morphological differences remarkably similar to those seen here (Perlman.Marshall.and.Biretta9001:Marshallefa£.2001b).," In particular, recent observations of the M87 jet have detected morphological differences remarkably similar to those seen here \citep{per01,mar01b}."
. IST imaging and polarimetric observations of MIST have provided evidence (hat (he more energetic particles responsible for the optical and X-ray emission [rom the M37 jet are located closer to the axis of the jet which is surrounded by a sheath or cocoon of lower-energy particles responsible lor the radio emission., HST imaging and polarimetric observations of M87 have provided evidence that the more energetic particles responsible for the optical and X-ray emission from the M87 jet are located closer to the axis of the jet which is surrounded by a sheath or cocoon of lower-energy particles responsible for the radio emission.
 The knots in (he M87 jet are identilied as (he sites of shocks in the flow where (he magnetic fields are compressed and (he particles accelerated (Sparks.andMacchetto1996:Perlinane£a£. 1998).," The knots in the M87 jet are identified as the sites of shocks in the flow where the magnetic fields are compressed and the particles accelerated \citep{spa96,per98}."
. The optical and X-ray emission is then due to svuchrotron radiation at (he sites of the shocks., The optical and X-ray emission is then due to synchrotron radiation at the sites of the shocks.
 Differences between the optical and radio emission are caused by particle diffusion and aging., Differences between the optical and radio emission are caused by particle diffusion and aging.
 A similar model can be used to qualitatively describe all the features of the bright N-rav/radio knots of the Cen A inner jet., A similar model can be used to qualitatively describe all the features of the bright X-ray/radio knots of the Cen A inner jet.
 The differences in the positions of the knots in (he X-ray ancl radio emission are naturally explained by particle aging., The differences in the positions of the knots in the X-ray and radio emission are naturally explained by particle aging.
 This hypothesis could be greatly strengthened if optical emission could be detected between the various radio and X-ray knots. but the dark dust lane makes such a detection unlikely (Marconiefa£2000).," This hypothesis could be greatly strengthened if optical emission could be detected between the various radio and X-ray knots, but the dark dust lane makes such a detection unlikely \citep{mar00}."
.. In [act. it is now becoming clear that such morphological differences are a common feature of N-rav. and. radio emission from FR. I ealaxies (Iarcleastle.Birkinshaw.anclWorrall2001).," In fact, it is now becoming clear that such morphological differences are a common feature of X-ray and radio emission from FR I galaxies \citep{hrd01}."
. The existence of shock sites and the eeneralion of X-rav emitting plasma may be a fundamental feature of jets in FR I galaxies., The existence of shock sites and the generation of X-ray emitting plasma may be a fundamental feature of jets in FR I galaxies.
 We have presented high-resolution. Chandra/ACIS-I X-ray images and spectra of the X-ray jet in Centaurus A and have found the following:, We have presented high-resolution /ACIS-I X-ray images and spectra of the X-ray jet in Centaurus A and have found the following:
he WR star. but is significantly offset in radius. as determined by he location of the stagnation point on the line of centers between he two stars and the width of the compressed WR wind.,"the WR star, but is significantly offset in radius, as determined by the location of the stagnation point on the line of centers between the two stars and the width of the compressed WR wind."
 This means there is a minimum radius to the CWIR. interior to which a spherical WR wind makes a simple flat-top contribution to the ine profile.," This means there is a minimum radius to the CWIR, interior to which a spherical WR wind makes a simple flat-top contribution to the line profile."
" Example profiles for the conical bow shock approximation are shown in Figure 3. for different binary separations relative ο the critical radius afro. cavity opening angles 2.and viewing inclinations 7,"," Example profiles for the conical bow shock approximation are shown in Figure \ref{fig3} for different binary separations relative to the critical radius $a/r_{\rm c}$, cavity opening angles $\beta$,and viewing inclinations $i$."
 Model parameters for the different panels are oovided in Table 2.., Model parameters for the different panels are provided in Table \ref{tab2}.
 The principle conclusions are that: (2) only a pole-on view to the orbit produces a symmetric profile. with a double-horned appearance. (b) an edge-on view produces one that is maximally lopsided. and (c) generally an asymmetric double-horned profile shape results whose appearance relates to the viewing perspective and orbital parameters.," The principle conclusions are that: (a) only a pole-on view to the orbit produces a symmetric profile, with a double-horned appearance, (b) an edge-on view produces one that is maximally lopsided, and (c) generally an asymmetric double-horned profile shape results whose appearance relates to the viewing perspective and orbital parameters."
 Note that these profiles have been gaussian smoothed to simulate limited spectral resolution., Note that these profiles have been gaussian smoothed to simulate limited spectral resolution.
 Given that typical WR winds have ος=1000.3000L.. smoothing with a gaussian of HWHM defer.=0.1 was adopted to match roughly the resolution of /50's SWS06 instrument (de Graauw 11996).Finally. all of the examples in Figure 3. have emission relative to a purely spherical wind by factors of10-204c.," Given that typical WR winds have $\vinf \approx
1000-3000$, smoothing with a gaussian of HWHM $\delta v/\vinf =0.1$ was adopted to match roughly the resolution of 's SWS06 instrument (de Graauw 1996).Finally, all of the examples in Figure \ref{fig3} have emission relative to a purely spherical wind by factors of."
. It happens that the total line flux as a function of the opening angle and binary separation is derivable analytically., It happens that the total line flux as a function of the opening angle and binary separation is derivable analytically.
 There are four basic zones., There are four basic zones.
 As previously noted for radii r«rii. the WR wind is spherical and contributes a flat-top contribution to the profile.," As previously noted for radii $r<r_{\rm orb}$, the WR wind is spherical and contributes a flat-top contribution to the profile."
 In the cavity sector. there is no contribution.," In the cavity sector, there is no contribution."
 Then there is the spherical zone and the compressed layer for rrog," Then there is the spherical zone and the compressed layer for $r\ge r_{\rm
orb}$."
 Accounting for these zones. the solution for the line luminosity. relative to a pure spherical wind. is denoted by X and given by In the limit that «ή=0. equation (30)) reduces to unity. as it mustbecause there is no CWIR.," Accounting for these zones, the solution for the line luminosity, relative to a pure spherical wind, is denoted by $\Lambda$ and given by where In the limit that $\beta=0$, equation \ref{eq:lineflux}) ) reduces to unity, as it mustbecause there is no CWIR."
 For two identical winds. the opening angle is 3?=907. and the maximum line flux for a given value of THyfue=orofro becomes: If the binary orbit is exceedingly large. \=1 is again recovered because the CWIR is displaced to a location of irrelevance with FasD re.," For two identical winds, the opening angle is $\beta= 90^\circ$, and the maximum line flux for a given value of $\sigma u_0/u_{\rm c} =
\sigma r_{\rm c}/r_0$ becomes: If the binary orbit is exceedingly large, $\Lambda =1$ is again recovered because the CWIR is displaced to a location of irrelevance with $r_{\rm orb} \gg r_{\rm c}$ ."
 On the other hand. ifr Mey. all the arctangent factors reduce to 7/2. thus," On the other hand, if $r_{\rm c} \gg r_{\rm orb}$ , all the arctangent factors reduce to $\pi/2$ , thus"
Brightness variation of V1154CCyg was discovered by Strohmeieretal.(1963).,Brightness variation of Cyg was discovered by \citet{stretal63}.
. Cepheid type variation and a periodicity somewhat shorter than 5dd were obvious from the photographic magnitudes leading to the discovery., Cepheid type variation and a periodicity somewhat shorter than d were obvious from the photographic magnitudes leading to the discovery.
 The first reliable light curve based on photoelectric UBV observations was published by Wachmann(1976)., The first reliable light curve based on photoelectric $UBV$ observations was published by \citet{wac76}.
". Further multicolour photoelectric and CCD photometric data were published by Szabados (1977),, ArellanoFerroetal.(1998),, Ignatova&Vozyakova (2000), Berdnikov(2008) and Pigulskietal.(2009)."," Further multicolour photoelectric and CCD photometric data were published by \citet{sza77}, , \citet{areetal98}, \citet{igvo00}, \citet{ber08} and \citet{asas09}."
. This latter paper contains the data of a dedicated photometric survey of the whole Kepler field., This latter paper contains the data of a dedicated photometric survey of the whole Kepler field.
 Space photometric data of CCyg are also available from the Hipparcos satellite ESA(1997) and the Scientific Archive of the Optical Monitoring Camera (OMC) on board INTEGRAL., Space photometric data of Cyg are also available from the Hipparcos satellite \citet{esa97} and the Scientific Archive of the Optical Monitoring Camera (OMC) on board INTEGRAL.
 None of these previous data can compete withKepler in photometric quality., None of these previous data can compete with in photometric quality.
" In addition, a large number of radial velocity data have been collected on CCyg by the Moscow CORAVEL team (Gorynyaetal.1998)."," In addition, a large number of radial velocity data have been collected on Cyg by the Moscow CORAVEL team \citep{goretal98}."
". These data obtained between 1990—1996 show a slight change in the mean velocity averaged over the pulsation cycle, thus Gorynyaetal.(1996) suspected spectroscopic binarity of this Cepheid."," These data obtained between $1990-1996$ show a slight change in the mean velocity averaged over the pulsation cycle, thus \citet{goretal96} suspected spectroscopic binarity of this Cepheid."
" However, radial velocity data obtained by Imbert(1999) much earlier than the Moscow data and covering a reasonably long time interval do not indicate binarity."," However, radial velocity data obtained by \citet{imb99} much earlier than the Moscow data and covering a reasonably long time interval do not indicate binarity."
" Molenda-Zakowiczetal.(2008) determined basic parameters for our target: [Fe/H]=0.06+0.07, spectral type: G2Ib, Tet=5370+ 118KK, logg=1.49+0.34, and vsini= 12.34+1.6kkmss~‘."," \citet{mfl08} determined basic parameters for our target: ${\rm [Fe/H]} = 0.06 \pm 0.07$, spectral type: G2Ib, $T_{\rm eff}=5370\pm118$ K, $\log{g}=1.49\pm0.34$, and $v \sin i = 12.3\pm 1.6$ $^{-1}$."
" These parameters are all consistent with a Cepheid, and place CCyg inside the theoretical instability strip presented in refKIC.."," These parameters are all consistent with a Cepheid, and place Cyg inside the theoretical instability strip presented in \\ref{KIC}. ."
 The chemical composition of V1154CCyg was determined independentlyby Lucketal.(2006) in a major project ofCepheid spectroscopy., The chemical composition of Cyg was determined independentlyby \citet{lucketal06} in a major project ofCepheid spectroscopy.
 They published a value of [Fe/H]= —0.10., They published a value of ${\rm [Fe/H]} = - 0.10$ .
FOC46192 bands.,$^{12}{\rmn C}^{12}{\rmn C}~\lambda6192$ bands.
 Another is the complete. absorption oetween andA... whatever its precise cause.," Another is the complete absorption between and, whatever its precise cause."
 The equivalent width of this (Wzzo». 6202) was measured relative o the pseudo-continuum., The equivalent width of this $_{5722-6202}$ ) was measured relative to the pseudo-continuum.
 Fig., Fig.
 I1. shows the relationship »etween these two measures., \ref{c_index} shows the relationship between these two measures.
 The correlation is obvious and here is no lateral dependence on J-ndex., The correlation is obvious and there is no lateral dependence on j-index.
 Hence. à c-index was set as the position along this sequence as indicated in he figure.," Hence, a c-index was set as the position along this sequence as indicated in the figure."
 Lt transpires that the spectral groups are dependent on roth the c-index and the j-index as can be seen in Fig., It transpires that the spectral groups are dependent on both the c-index and the j-index as can be seen in Fig.
 12 which is Fie., \ref{slopes_c} which is Fig.
 4. redrawn this time with svmbols showing 10 c-index., \ref{slopes} redrawn this time with symbols showing the c-index.
 In this case. the c-index increases. clockwise in sectors of a circle centred. near the point 11.0).," In this case, the c-index increases clockwise in sectors of a circle centred near the point 1.0)."
 By comparing Figs LO and 12. it is clear that the two spectral gaopes (4) eive a two-dimensional classification that can be related to the j- and c-indices and are more easily. measured ian the latter., By comparing Figs \ref{slopes_j} and \ref{slopes_c} it is clear that the two spectral slopes $\Phi$ ) give a two-dimensional classification that can be related to the j- and c-indices and are more easily measured than the latter.
 lig., Fig.
 13 shows the c-index polted against (JJfy) emperature), \ref{c_t} shows the c-index plotted against $(J-K)$ (temperature).
 nie (rend is as ex»eeted: the cooler stars have the higher c-indices., The trend is as expected; the cooler stars have the higher c-indices.
 It should be noted that the index here is purely a nxsure of the observed carbon strength and not a measure of hat strength relaive to a standard of the same temperature., It should be noted that the index here is purely a measure of the observed carbon strength and not a measure of that strength relative to a standard of the same temperature.
 Phere is considerable spread. but that is significantly recluc'ed when the brig ht.) stars are excluded.," There is considerable spread, but that is significantly reduced when the bright J stars are excluded."
 Phere is a distinc tendeney for the brighter J stars to have weaker carbon bands (smaller c-indices) than the fainter (normal) JJ stars of the same (JA) colours., There is a distinct tendency for the brighter J stars to have weaker carbon bands (smaller c-indices) than the fainter (normal) J stars of the same $(J-K)$ colours.
 Fig., Fig.
 1d shows the expected correlation between the depths of the C7C (1.33) and (0.22) bands at and respectively., \ref{c_c} shows the expected correlation between the depths of the $^{12}{\rmn C}^{12}{\rmn C}$ 3) and 2) bands at and respectively.
 Phe svmbols mark the (JA) colour., The symbols mark the $(J-K)$ colour.
 Lt is clear hat the relationship between Doj2» ancl Dojo» is dilferent or cach colour range in that it appears progressively to he right as colour increases (i6. empoerature. decreases)., It is clear that the relationship between $_{6122}$ and $_{6192}$ is different for each colour range in that it appears progressively to the right as colour increases (i.e. temperature decreases).
 This would. give rise to a correlaticon between the ratio of he band depths. Dijo» DDojos. and CS ÁA)inthe sense hat the 33) band is stronger wih respect to the 22) xuxd for the bluer (hotter) stars.," This would give rise to a correlation between the ratio of the band depths, $_{6122}$ $_{6192}$, and $(J-K)$ in the sense that the 3) band is stronger with respect to the 2) band for the bluer (hotter) stars."
 This ratio changes from Wat (FIN)— 110 0.6 at CSA)~2 with a spread, This ratio changes from 0.9 at $(J-K)\sim 1$ to 0.6 at $(J-K)\sim 2$ with a spread
(see AAVSO website).,(see AAVSO website).
 Considering that the shock wave is the consequence of the pulsation of the star. as the luminosity variation is. then it is no wonder that the polarizing mechanism linked to the shock wave does not behave exactly the same as before.," Considering that the shock wave is the consequence of the pulsation of the star, as the luminosity variation is, then it is no wonder that the polarizing mechanism linked to the shock wave does not behave exactly the same as before."
 However. as we noticed for the third cycle. the signatures’ widths do not change from one cycle to another for a given Balmer line.," However, as we noticed for the third cycle, the signatures' widths do not change from one cycle to another for a given Balmer line."
 appearance of an. emission lineD. indicates Thethat the specific intensity of the back-lighting radiation (creating the continuum) is less than the source funetion inside the shock., The appearance of an emission line indicates that the specific intensity of the back-lighting radiation (creating the continuum) is less than the source function inside the shock.
 This is linked to the intrinsic. photon emission behind the shock’s front., This is linked to the intrinsic photon emission behind the shock's front.
 Emission lines can thus help in determining the properties of this area., Emission lines can thus help in determining the properties of this area.
 Polarization also happens to be much stronger in the Balmer lines than in the pseudo-continuum around them., Polarization also happens to be much stronger in the Balmer lines than in the pseudo-continuum around them.
 Its temporal evolution follows the evolution of, Its temporal evolution follows the evolution of
curve’ of the emission lines.,curve' of the emission lines.
 LW the observed feature is indeed a rotating disk. the morphology of the emission. would require it to be nearly egde-on.," If the observed feature is indeed a rotating disk, the morphology of the emission would require it to be nearly egde-on."
 Given this fact and the mass discrepaney. we thus rule out a fully sampled rotating disk in dvnamical equilibrium with the host potential being the source of the highest. surface. brightness. aligned. [eatures.," Given this fact and the mass discrepancy, we thus rule out a fully sampled rotating disk in dynamical equilibrium with the host potential being the source of the highest surface brightness aligned features."
 Llowever. it is possible that we might be observing a disk feature which is not fully sampled. ancl may only. be observing the part of it which lies within the lonization cone of the quasar.," However, it is possible that we might be observing a disk feature which is not fully sampled, and may only be observing the part of it which lies within the ionization cone of the quasar."
 The observations might instead. be better explained ow eas strewn along the elliptical orbit of the companion. which is seen almost edge-on (see Fie.," The observations might instead be better explained by gas strewn along the elliptical orbit of the companion, which is seen almost edge-on (see Fig."
 1200)., \ref{cartoons}b b).
 In this case. he gas observed. close to the nucleus of PINS2250-41. could » moving almost perpencicularly to the line of sight. xwticularlv if this σας doesn't sample the perigee of the companion orbit.," In this case, the gas observed close to the nucleus of PKS2250-41 could be moving almost perpendicularly to the line of sight, particularly if this gas doesn't sample the perigee of the companion's orbit."
 The higher surface brightness. of. the emission closers to the nucleus would be due to the gas being it up by the quasar: the ollset of the emission from the racio axis is naturally explained in this scenario., The higher surface brightness of the emission closer to the nucleus would be due to the gas being lit up by the quasar; the offset of the emission from the radio axis is naturally explained in this scenario.
 The exact orm of the velocity curve would depend on the distrubution of gas around the orbit. ancl could. easily account for the observations.," The exact form of the velocity curve would depend on the distrubution of gas around the orbit, and could easily account for the observations."
 Another option is a tidal streamer pulled. off from the disk of the radio source precursor object by the companion ealaxy (see. e.g.. the numerical simulations of di Matteo et al (20053).," Another option is a tidal streamer pulled off from the disk of the radio source precursor object by the companion galaxy (see, e.g., the numerical simulations of di Matteo et al (2005))."
 At some stages following the first pass of the two merging galaxies the gas will stream back (almost raciallv) towards the nuclei of the galaxies (see Fig., At some stages following the first pass of the two merging galaxies the gas will stream back (almost radially) towards the nuclei of the galaxies (see Fig.
 13ec)., \ref{cartoons}c c).
 In this case. the velocity curve could rellect the streaming infall motions on either side of the nucleus.," In this case, the velocity curve could reflect the streaming infall motions on either side of the nucleus."
 Such infalling material could be responsible for triggering the AGN/jet activity within PINS2250-41., Such infalling material could be responsible for triggering the AGN/jet activity within PKS2250-41.
 Spitzer AILPS photometry of a complete. sanmiple of 2Jvy sources (Cladbunter οἱ al 2007) bas been used. to address the issue of the dominant dust heating mechanism in AGN., Spitzer MIPS photometry of a complete sample of 2Jy sources (Tadhunter et al 2007) has been used to address the issue of the dominant dust heating mechanism in AGN.
 The eemission line Luminosity (which can be used as a tracer of AGN power. e.g. Rawlings Saunders 1991: Tadhunter et al 19958: Simpson 1998) is strongly. correlated with the mic-LR emission (see Fig.," The emission line luminosity (which can be used as a tracer of AGN power, e.g. Rawlings Saunders 1991; Tadhunter et al 1998; Simpson 1998) is strongly correlated with the mid-IR emission (see Fig."
 1 in Tadhunter et al 2007). indicating that AGN illumination is the principal dust-heating mechanism.," 1 in Tadhunter et al 2007), indicating that AGN illumination is the principal dust-heating mechanism."
 Phe mid- to far-LR uses of PINS2250-41 (22mv. ane 11.6 η]ν at τομ and 244/20 respectively) imply an overall far-infrared luminosity of Lig~2.I0L. (Sanders Mirabel 1996).," The mid- to far-IR fluxes of PKS2250-41 (22mJy and 11.6 mJy at $\mu m$ and $\mu m$ respectively) imply an overall far-infrared luminosity of $L_{IR}
\sim 2\times 10^{11}L_{\odot}$ (Sanders Mirabel 1996)."
 Although this is consistent with a LlliCi-class source. this source falls firmly on the correlation between mid- to far-IR emission and the cemission line luminosity. suggesting that in the case of this source the mid- to far-L luminosity is derived. [rom AGN rather than starburst illumination.," Although this is consistent with a LIRG-class source, this source falls firmly on the correlation between mid- to far-IR emission and the emission line luminosity, suggesting that in the case of this source the mid- to far-IR luminosity is derived from AGN rather than starburst illumination."
 Further to this. the mid- to far-LR. emission is well centered. on the centroid of the radio source host galaxy. with no displacement towards the companion such as might be expected if the companion galaxy was actively forming stars at a high rate.," Further to this, the mid- to far-IR emission is well centered on the centroid of the radio source host galaxy, with no displacement towards the companion such as might be expected if the companion galaxy was actively forming stars at a high rate."
 lt is clear from these results that PIXS250-41 is not à case of a source trigeered in the final stages of a major galaxy merger., It is clear from these results that PKS2250-41 is not a case of a source triggered in the final stages of a major galaxy merger.
 However. the scenario wherein a luminous AGN has been triggered by an encounter with another galaxy. some ime after the point of closest approach. is far more plausible.," However, the scenario wherein a luminous AGN has been triggered by an encounter with another galaxy, some time after the point of closest approach, is far more plausible."
 either the radio galaxy nor its companion can be classed as a very massive elliptical., Neither the radio galaxy nor its companion can be classed as a very massive elliptical.
 Acdcditionallv. the lack of very ο far-LRo emission. from either source is consistent with low levels of starburst activity. with the AGN being the dominant cause of dust heating.," Additionally, the lack of very powerful far-IR emission from either source is consistent with low levels of starburst activity, with the AGN being the dominant cause of dust heating."
 Lf this is a case of triggering ollowing a (first pass) galaxy encounter. the companion galaxy would. most. likely be on a highly elliptical. orbit. and eventually merge with the radio source host at a much ater point in time.," If this is a case of triggering following a (first pass) galaxy encounter, the companion galaxy would most likely be on a highly elliptical orbit, and eventually merge with the radio source host at a much later point in time."
 Many. of the different merger. models edict relatively low levels of star formation until the nuclei inallv merge (e.g. Llernquist Mihos 1995. Springel et al 2005).," Many of the different merger models predict relatively low levels of star formation until the nuclei finally merge (e.g. Hernquist Mihos 1995, Springel et al 2005)."
 The tidal forces associated with 1e initial encounter could be expected to cause the formation of the filamentary spur. while gas settling onto the central black hole for some time after the time of closest approach could. trigger the AGN activity.," The tidal forces associated with the initial encounter could be expected to cause the formation of the filamentary spur, while gas settling onto the central black hole for some time after the time of closest approach could trigger the AGN activity."
 Similar features are observed in the case of PINSO349-27 (Danziger et al 1984). which displays a tidal bridge linking it to a companion galaxy.," Similar features are observed in the case of PKS0349-27 (Danziger et al 1984), which displays a tidal bridge linking it to a companion galaxy."
 lt is also interesting to contrast the results [or PINS2250-41 with those of the z=0.23 ETUL radio galaxy PINS1932-464 (Inskip et al 2007). which. like. PINS2250-41. exists within an interacting eroup environment and also displays a spectacular ELL.," It is also interesting to contrast the results for PKS2250-41 with those of the z=0.23 FRII radio galaxy PKS1932-464 (Inskip et al 2007), which, like PKS2250-41, exists within an interacting group environment and also displays a spectacular EELR."
 However. the similarities end there.," However, the similarities end there."
 Despite the materially enriched IM. the racio source of PINS1932-464 does not interact stronely with its environment.," Despite the materially enriched IGM, the radio source of PKS1932-464 does not interact strongly with its environment."
 Xdeditionallv. while we observe little in the way of star formation associated with PIS2250-41. evidence for significant star formation is seen in the IG. surrounding PINS1932-464. as well as within the neighbouring companion ealaxv.," Additionally, while we observe little in the way of star formation associated with PKS2250-41, evidence for significant star formation is seen in the IGM surrounding PKS1932-464, as well as within the neighbouring companion galaxy."
 These results illustrate that even when ACN are trigecred under apparently similar circumstances. a diversity of diferent outcomes are possible.," These results illustrate that even when AGN are triggered under apparently similar circumstances, a diversity of different outcomes are possible."
 We have carried out a multiwavelength study. of the racio source PINS2250-41. for the first time combining imaging. long-slit and LIEU spectroscopic observations of this source.," We have carried out a multiwavelength study of the radio source PKS2250-41, for the first time combining imaging, long-slit and IFU spectroscopic observations of this source."
 In addition to analysing the ELLR in greater. detail. we have now also been able to constrain the nature of the companion objects surrounding the host galaxy. as well as the nature of the host galaxy. itself.," In addition to analysing the EELR in greater detail, we have now also been able to constrain the nature of the companion objects surrounding the host galaxy, as well as the nature of the host galaxy itself."
 This has allowed us to produce a far cleaver picture of the cillerent astrophysical processes ongoing within this source. and how the activity was Lriggered.," This has allowed us to produce a far clearer picture of the different astrophysical processes ongoing within this source, and how the activity was triggered."
 The key results are as follows:, The key results are as follows:
therein).,.
 Today we know that the narrow lines (Forbidden and permitted) are emitted in the Weigelt: blobs (Weigelt&Ibersberger.LOSG).. at ~0.3 aresee from. the central star (Davidsonetal.1995) and the broad emission ines are formed. in the wind of the central object. (μουandAllen1992:Davidsonetal. 1995).," Today we know that the narrow lines (forbidden and permitted) are emitted in the Weigelt blobs \citep{weigelt86}, at $\sim0.3$ arcsec from the central star \citep{b6} and the broad emission lines are formed in the wind of the central object \citep{HA92_eta,b6}."
. The combination of vigh and low excitation lines in the same object. however. was paracdoxical.," The combination of high and low excitation lines in the same object, however, was paradoxical."
 A Κον to understanding this interesting object was ound recently through the study of the variability of the veh excitation lines., A key to understanding this interesting object was found recently through the study of the variability of the high excitation lines.
 Vhe high excitation forbidden lines disappeared in 1948. and again in 1965. 1981. 1987 ane 1992.," The high excitation forbidden lines disappeared in 1948, and again in 1965, 1981, 1987 and 1992."
" These ""spectroscopic events! (Gaviola1953:Rodgers&Searle1967:ThackerayZanellactal.1984). or ‘low excitation events’ (Darminelietal.1998). were believec to be part of S Doracus eveles. commonly. seen in other LBY stars similar to eta Car."," These `spectroscopic events' \citep{b9,b17,b21,b25} or `low excitation events' \citep{b29} were believed to be part of S Doradus cycles, commonly seen in other LBV stars similar to eta Car."
 This interpretation seemec to be supported. by the ALOS30 line which went to minimum (Damineli1096) when the near-infrared lih curve went to maximum (Whitelocketal., This interpretation seemed to be supported by the $\lambda$ 10830 line which went to minimum \citep{b3} when the near-infrared light curve went to maximum \citep{b23}.
1904)... The maxima in the near-infrared light curves were not truly periodic ancl the length. of the quasi-period was dilleren for dilferent. pass-bands., The maxima in the near-infrared light curves were not truly periodic and the length of the quasi-period was different for different pass-bands.
 However. the spectroscopic events were demonstrated to be periodic (Daminelietal. 2000).. in contrast to the incoherent character. of the S Dor oscillations.," However, the spectroscopic events were demonstrated to be periodic \citep{b5}, , in contrast to the incoherent character of the S Dor oscillations."
 Damincli.Contiand.Lopes(1997). and others proposed a binary model with a highly eccentric. orbit. a hotter secondary component and a strong. wind-wine collision (WNCO.," \citet{b4} and others proposed a binary model with a highly eccentric orbit, a hotter secondary component and a strong wind-wind collision (WWC)."
 Binarity is interesting as it. potentially. allows the clirect measurement of the mass of the stars. heir most fundamental parameter.," Binarity is interesting as it potentially allows the direct measurement of the mass of the stars, their most fundamental parameter."
 “Phe binary scenario las provided a framework for understanding the star and »ovided. euidelines for fruitful observations. although some owefer a model in which there are periodic shell ejections (Martinctal. ," The binary scenario has provided a framework for understanding the star and provided guidelines for fruitful observations, although some prefer a model in which there are periodic shell ejections \citep{b15}. ."
2006).., In Fig.
 In Fig. 1 we present examples of high and low excitation state spectra of η Carinae., \ref{highlow} we present examples of high and low excitation state spectra of $\eta$ Carinae.
 The observation of an event in 1997.95. as was wedieted. brought more confidence to the true. periodic nature of the variation (Daminelietal.," The observation of an event in 1997.95, as was predicted, brought more confidence to the true periodic nature of the variation \citep{b5}."
2PO00).. Feastetal.(2001) used archival spectra to identify three previously unreported events. in 1953. 1959 anc 1970. which also it the 5.5-vr. period.," \citet{b8} used archival spectra to identify three previously unreported events, in 1953, 1959 and 1970, which also fit the 5.5-yr period."
 Moreover. those authors discovered hat the dips on top of the broad quasi-periodic near-infrared maxima were truly periodic ancl correlated with he behavior of the high excitation lines.," Moreover, those authors discovered that the dips on top of the broad quasi-periodic near-infrared maxima were truly periodic and correlated with the behavior of the high excitation lines."
 An extensive. X-rav monitoring campaign was started in 1996 with theRNTE satellite anc revealed deep minima in 1997.95 and 2003.49 which coincided with the minima seen at other wavelengths (Corcoran2005)., An extensive X-ray monitoring campaign was started in 1996 with the satellite and revealed deep minima in 1997.95 and 2003.49 which coincided with the minima seen at other wavelengths \citep{b2}.
. X-ray observations inside and outside the minimum performed. with andNALAL furnished details on the column density CVg). temperature and. chemical composition of the colliding wind. shock (Lamaguchietal. 2007).," X-ray observations inside and outside the minimum performed with and furnished details on the column density $N_{\rm{H}}$ ), temperature and chemical composition of the colliding wind shock \citep{b11}. ."
. vanGenderenetal.(2006). showed that the optical light curve displays periodic clips like those inthe near-infrarecl ancl Lajüsetal.(2003)0 reported a very detailed light curve in the D. V. # and £ bands for the 2003.49 event.," \citet{b22} showed that the optical light curve displays periodic dips like those inthe near-infrared and \citet{b14} reported a very detailed light curve in the $B$, $V$, $R$ and $I$ bands for the 2003.49 event."
 The events were recorded. also at cm (DuncanandWhite2003). and racio-mm (Abrahametal. 2005).. but no specific value to the period length was reported for those wavelengths.," The events were recorded also at radio-cm \citep{b7} and radio-mm \citep{b1}, but no specific value to the period length was reported for those wavelengths."
 Alany other features vary periodically. in intensity and. radial velocity. like the broad. emission and D? Cvgni absorption components. and can also be used. to. derive he period. length.," Many other features vary periodically in intensity and radial velocity, like the broad emission and P Cygni absorption components, and can also be used to derive the period length."
 One of them is A686 discovered » SteinerandDamineli(2004)... which raises anc drops just before. minimum faster than any other cature over he entire spectrum.," One of them is $\lambda$ 4686 discovered by \citet{b20}, which raises and drops just before minimum faster than any other eature over the entire spectrum."
 Although faint ολλ2 A)) it was requenthy monitored. with high signal/noise along the last event., Although faint $<2$ ) it was frequently monitored with high signal/noise along the last event.
 Unfortunately. it was observed. only occasionally in he previous events. preclucling its use to measure the period.," Unfortunately, it was observed only occasionally in the previous events, precluding its use to measure the period."
 This spectral line deserves. better monitoring in. future events. not only to improve the accuracy of the derived »eriod. but also because it is the highest excitation feature observed. at optical wavelengths. and its origin remains a nivstery.," This spectral line deserves better monitoring in future events, not only to improve the accuracy of the derived period, but also because it is the highest excitation feature observed at optical wavelengths, and its origin remains a mystery."
 To facilitate discussion we label the events by numbers as described by GrohaneDamineli(2005):: number one (£11) is assigned to the event observed in 1948 by Caviola. so that the event of 2003.49. is #111.," To facilitate discussion we label the events by numbers as described by \citet{b10}: number one 1) is assigned to the event observed in 1948 by Gaviola, so that the event of 2003.49 is 11."
 We celine as he time interval between the starting of two consecutive minima. so that evele #99 started at the 1992.42 minimum and finished. when cevele #110 was starting in 1997.95.," We define as the time interval between the starting of two consecutive minima, so that cycle 9 started at the 1992.42 minimum and finished when cycle 10 was starting in 1997.95."
 Because of observational reasons. that will become clear ater in this paper. the starting point of a evcle is defined by he clisappearance of the A6678 narrow line component.," Because of observational reasons, that will become clear later in this paper, the starting point of a cycle is defined by the disappearance of the $\lambda6678$ narrow line component."
 With this definition. phases along the evele are defined in a unique way for all measured quantities.," With this definition, phases along the cycle are defined in a unique way for all measured quantities."
 The paper is organized as follows., The paper is organized as follows.
 We present in section 77 the observations: in §?? the definition of the phase 0 of the minimum: in §?? the determination of the period length: in 877. the stability of the period: in 877. the relation between the sharp peaks during the giant eruption ancl periastron passages: and in gv the discussion. ancl conclusions., We present in section \ref{observations} the observations; in \ref{zero} the definition of the phase 0 of the minimum; in \ref{period} the determination of the period length; in \ref{stability} the stability of the period; in \ref{eruption} the relation between the sharp peaks during the giant eruption and periastron passages; and in \ref{discus} the discussion and conclusions.
 The majority of the ground-based observations presented in this paper came from a monitoring campaign started in 1989 at the Coudé focus of the 1.6-m telescope of Pico clos Dias Observatory (OPD-LNA/Brazil)., The majority of the ground-based observations presented in this paper came from a monitoring campaign started in 1989 at the Coudé focus of the 1.6-m telescope of Pico dos Dias Observatory (OPD-LNA/Brazil).
 The observational setup at OPD was kept essentially unchanged through the campaign: a dispersion grating with 600 L/mm. entrance slit width ~1.3 aresec. exposure time ~5 s in llo increasing to 15 min at 3500 ancl 10800. A.," The observational setup at OPD was kept essentially unchanged through the campaign: a dispersion grating with 600 l/mm, entrance slit width $\sim1.3$ arcsec, exposure time $\sim 5$ s in $\alpha$ increasing to $\sim15$ min at 3500 and 10800 ."
. Spectra were extracted along 2 aresec in the spatial direction and no measurable dillerences in line intensities were seen when changing the extraction size by a [factor of 2., Spectra were extracted along $\sim2$ arcsec in the spatial direction and no measurable differences in line intensities were seen when changing the extraction size by a factor of 2.
 Three different CCDs have been used. with resolving powers 1t. = 25 km + (0.25 ty at La in 2003 and HR. = 50 km | (0.39 1) in the preceding vears.," Three different CCDs have been used, with resolving powers R = 25 km $^{-1}$ (0.25 $^{-1}$ ) at $\alpha$ in 2003 and R = 50 km $^{-1}$ (0.39 $^{-1}$ ) in the preceding years."
 On some occasions. a 1024.1024 Hawaii detector was used to observe he ALOS30 line. delivering a spectral resolution It = 40 uns + (0.65 pixel +).," On some occasions, a $1024\times1024$ Hawaii detector was used to observe the $\lambda$ 10830 line, delivering a spectral resolution R = 40 km $^{-1}$ (0.65 $^{-1}$ )."
 On other occasions. spectra of this ine were taken at It — 15 kms ‘witha thinned CCD.," On other occasions, spectra of this line were taken at R = 15 km $^{-1}$ with a thinned CCD."
 After correcting For fringes ancl degrading the spectral resolution.hese spectra were almostidentical to those collected. with he infrared array at the same date.," After correcting for fringes and degrading the spectral resolution,these spectra were almostidentical to those collected with the infrared array at the same date."
 Observationsin. 1992 and 1907/8 were done with a thick CCD that was almost ree of fringes. but had a lowsensitivity in the blue. which explains the poor coverage of important lines in that spectral range.," Observationsin 1992 and 1997/8 were done with a thick CCD that was almost free of fringes, but had a lowsensitivity in the blue, which explains the poor coverage of important lines in that spectral range."
 For wavelengths longer than 6500 A.. telluric," For wavelengths longer than 6500 , telluric"
0«z<3.,$0<z<3$.
" The simulation box length was ! 1500h-Mpc, resulting in a large simulated cluster sample."," The simulation box length was $1500h^{-1}$ Mpc, resulting in a large simulated cluster sample."
" The force softening length was 17h ‘kpc, so the spatial resolution was more than sufficient to resolve the cluster halos."," The force softening length was $17h^{-1}$ kpc, so the spatial resolution was more than sufficient to resolve the cluster halos."
 Only masses above 2x10?A-!Mo were considered as clusters.," Only masses above $2 \times
10^{13}h^{-1}M_{\sun}$ were considered as clusters."
 The observations include many systematic errors that are not in the simulation., The observations include many systematic errors that are not in the simulation.
" The most important ones are (1) the line-ofsight redshift selection and complications due to photometric redshift error, (2) cluster centroiding error (misidentification of the BCG due to some algorithmic error), (3) noise in the determination of the cluster position angle due the small number of cluster member galaxies (> 5) used to estimate the cluster shape, and (4) the need to exclude clusters that appeared nearly round, due to the difficulty in estimating a position angle."," The most important ones are (1) the line-of-sight redshift selection and complications due to photometric redshift error, (2) cluster centroiding error (misidentification of the BCG due to some algorithmic error), (3) noise in the determination of the cluster position angle due the small number of cluster member galaxies $\ge 5$ ) used to estimate the cluster shape, and (4) the need to exclude clusters that appeared nearly round, due to the difficulty in estimating a position angle."
 The first three of these systematic errors will reduce our ability to detect intrinsic alignments and weaken the observed signal., The first three of these systematic errors will reduce our ability to detect intrinsic alignments and weaken the observed signal.
" Thus, we do not comparewith the direct predictions of ? (for their 0«z0.5 sample), but rather with reduced predictions as described below."," Thus, we do not comparewith the direct predictions of \cite{2005ApJ...618....1H} (for their $0<z<0.5$ sample), but rather with reduced predictions as described below."
" The line-of-sight pair selection criterion, |Azpnot|< 0.015, was chosen to balance competing considerations."," The line-of-sight pair selection criterion, $|\Delta
 z_\mathrm{phot}| < 0.015$ , was chosen to balance competing considerations."
" In the absence of photometric redshift error, we would ideally attempt to mimic the ? selection of pairs within 100h-!Mpc; even with spectroscopic redshifts this would be complicated by redshift-space distortions, but we could at least hope to come fairly close to what was done in the simulations."," In the absence of photometric redshift error, we would ideally attempt to mimic the \cite{2005ApJ...618....1H}
 selection of pairs within $100h^{-1}$ Mpc; even with spectroscopic redshifts this would be complicated by redshift-space distortions, but we could at least hope to come fairly close to what was done in the simulations."
" However, the photometric redshift errors correspond to typical separations of ~50h~!Mpc.,, making itimpossible to imitate a strict line-of-sight separation."," However, the photometric redshift errors correspond to typical separations of $\sim 50$, making itimpossible to imitate a strict line-of-sight separation."
" We could simply choose a Az corresponding roughly to 100h~'Mpc,, but empirical tests showed that the contamination from completely unassociated clusters along the line-of-sight became unacceptably large."," We could simply choose a $\Delta z$ corresponding roughly to $100$, but empirical tests showed that the contamination from completely unassociated clusters along the line-of-sight became unacceptably large."
" These unassociated clusters dilute the expected signal and therefore thedetection significance since their orientations are purely random, and indeed, even with our chosen AZphot, we must impose a correction for it (to be described below)."," These unassociated clusters dilute the expected signal and therefore thedetection significance since their orientations are purely random, and indeed, even with our chosen $\Delta z_\mathrm{phot}$, we must impose a correction for it (to be described below)."
 We therefore err on the conservative side and use a Azpnot corresponding to roughly the 1o photometric redshift error in order to be able to measure the signal., We therefore err on the conservative side and use a $\Delta z_\mathrm{phot}$ corresponding to roughly the $1\sigma$ photometric redshift error in order to be able to measure the signal.
 The contamination due to accidental inclusion of unassociated clusters because of photometric redshift error can be estimated via simulation., The contamination due to accidental inclusion of unassociated clusters because of photometric redshift error can be estimated via simulation.
" We simulate galaxy clusters with constant comoving number density, assign a z assuming σ{ζρμοι)=0.015 (Gaussian), and estimate what fraction of the clusters within |Azpnot|=0.015 are actually >100! Mpc apart on the line-of-sight (with this separation chosen because it is the criterion used by ?))."," We simulate galaxy clusters with constant comoving number density, assign a $z$ assuming $\sigma(z_\mathrm{phot}) =
0.015$ (Gaussian), and estimate what fraction of the clusters within $|\Delta z_\mathrm{phot}| = 0.015$ are actually $>100h^{-1}$ Mpc apart on the line-of-sight (with this separation chosen because it is the criterion used by \citealt{2005ApJ...618....1H}) )."
" Given a contamination fraction 0<I«1 defined as the fraction of cluster pair candidates that satisfy our |AZpnot| cut but that are more than 100h~'Mpc aapart along the line-of-sight, and theoretical predictions (cos?0¢,»)ACDM, we compare our measured signals with We cannot estimate [ from the random, simulated clusters alone, because that simulation only tells us the relative contamination when the intrinsic (3d) cluster correlation function €«1."," Given a contamination fraction $0<\Gamma<1$ defined as the fraction of cluster pair candidates that satisfy our $|\Delta z_\mathrm{phot}|$ cut but that are more than $100$ apart along the line-of-sight, and theoretical predictions $\langle
\cos^2\theta_{c,p}\rangle_{\Lambda\mathrm{CDM}}$, we compare our measured signals with We cannot estimate $\Gamma$ from the random, simulated clusters alone, because that simulation only tells us the relative contamination when the intrinsic (3d) cluster correlation function $\xi\ll 1$."
" When clustering is significant, the relative contamination becomes smaller."," When clustering is significant, the relative contamination becomes smaller."
" We quantify this effect by using the simulated contamination fraction B=0.35 in the absence of clustering, and the observed (projected) cluster correlation function."," We quantify this effect by using the simulated contamination fraction $\beta=0.35$ in the absence of clustering, and the observed (projected) cluster correlation function."
" The simulated contamination fraction can be defined 8=Ns/(Ns+Nn) where Ns represents spurious cluster pairs that appear within our Azpnot due only to photometric redshift error, and Nr represents those real pairs that are expected due toa purely random distribution of clusters in the survey volume."," The simulated contamination fraction can be defined $\beta = N_S/(N_S+N_R)$ where $N_S$ represents spurious cluster pairs that appear within our $\Delta z_\mathrm{phot}$ due only to photometric redshift error, and $N_R$ represents those real pairs that are expected due to a purely random distribution of clusters in the survey volume."
 What we really care about is T=Ns/(Ns+NnΝε) where Ng represents the excess cluster pairs that are there in reality due to a non-zero cluster correlation function., What we really care about is $\Gamma = N_S/(N_S+N_R+N_E)$ where $N_E$ represents the excess cluster pairs that are there in reality due to a non-zero cluster correlation function.
" Fortunately, we also measure the projected correlation function, w(R) +1 = catalogue —"," Fortunately, we also measure the projected correlation function, w(R) +1 = = ."
" NS) To estimate I we then use LR) = In the case that zig(3)clustering is insignificant, as on large scales, w£z0 and ['(R)=8 0.35."," To estimate $\Gamma$ we then use (R) = In the case that clustering is insignificant, as on large scales, $w\approx 0$ and $\Gamma(R)=\beta=0.35$ ."
" On the smallest scales,"," On the smallest scales,"
",, and we select this for the ensemble flux.",", and we select this for the ensemble flux."
" This flux yields about 100 source photons for the 1-year integration period, allowing for asymptotic calibration."," This flux yields about 100 source photons for the 1-year integration period, allowing for asymptotic calibration."
" Since the background strongly affects the spectral analysis, independent realizations of the diffuse background are important and accordingly we also simulated 20 realizations of the diffuse sources."," Since the background strongly affects the spectral analysis, independent realizations of the diffuse background are important and accordingly we also simulated 20 realizations of the diffuse sources."
" To the point source data we added phase from a single-Gaussian light curve with σ=0.03, while we generated uniform random phases for the diffuse photons."," To the point source data we added phase from a single-Gaussian light curve with $\sigma=0.03$, while we generated uniform random phases for the diffuse photons."
" To assess the impact of uncertain parameters, we first calculate the probability weights using the known model parameters for the pulsar and the diffuse background, i.e., the “ideal” case."," To assess the impact of uncertain parameters, we first calculate the probability weights using the known model parameters for the pulsar and the diffuse background, i.e., the “ideal” case."
 We then perform a maximum likelihood spectral fit to estimate the spectral parameters., We then perform a maximum likelihood spectral fit to estimate the spectral parameters.
" Since the simulated point source is very dim relative to the background, it is impossible to fit all three parameters."," Since the simulated point source is very dim relative to the background, it is impossible to fit all three parameters."
" We therefore fix the cutoff energy to 100 GeV, effectively a power law spectrum."," We therefore fix the cutoff energy to $100$ GeV, effectively a power law spectrum."
 This approach is since we are now using an incorrect spectral model., This approach is since we are now using an incorrect spectral model.
" Using the best-fit values for the flux density and the photon index, we calculate a new set of probability weights."," Using the best-fit values for the flux density and the photon index, we calculate a new set of probability weights."
" Finally, we compute Hoo, to determine the significance (a) using the “ideal” weights and (b) using the *measured"" weights."," Finally, we compute $H_{20w}$ to determine the significance (a) using the “ideal” weights and (b) using the “measured” weights."
 We compare the results for the two cases—in o units—in Figure 8.., We compare the results for the two cases—in $\sigma$ units—in Figure \ref{ch5_plot14}.
" In general, the detection significance obtained with the *measured"" weights is slightly lower than that obtained with the “ideal” weights, although in a few cases statistical fluctuations lead to the opposite outcome."," In general, the detection significance obtained with the “measured” weights is slightly lower than that obtained with the “ideal” weights, although in a few cases statistical fluctuations lead to the opposite outcome."
" Comparing the population means, the overall significance using measured weights is decreased to 92% of the ideal case, corresponding to an increase in flux threshold of ~10%.."," Comparing the population means, the overall significance using measured weights is decreased to $92\%$ of the ideal case, corresponding to an increase in flux threshold of $\sim$."
 This effect is small compared to the factor of 1.5—2.0 increase in flux threshold seen between the weighted and unweighted versions of the H-test., This effect is small compared to the factor of $1.5-2.0$ increase in flux threshold seen between the weighted and unweighted versions of the $H$ -test.
 We therefore conclude that—even accounting for uncertainties in the spectral parameters used to calculate the probability weights—weighted statistics offer a significant improvement in sensitivity., We therefore conclude that—even accounting for uncertainties in the spectral parameters used to calculate the probability weights—weighted statistics offer a significant improvement in sensitivity.
 The machinery established above—performing spectral fits on an ensemble to compute the weighted statistics using probabilities estimated from an ML fit—also provides for directly comparing the DC (unpulsed) source significance with the pulsed significance., The machinery established above—performing spectral fits on an ensemble to compute the weighted statistics using probabilities estimated from an ML fit—also provides for directly comparing the DC (unpulsed) source significance with the pulsed significance.
" To estimate the DC significance, we use the result of ? for the asymptotic calibration for the likelihood ratio for a single parameter whose null value lies on a boundary."," To estimate the DC significance, we use the result of \citet{chernoff_miracle} for the asymptotic calibration for the likelihood ratio for a single parameter whose null value lies on a boundary."
" To apply this calibration, we take the free parameter to be the flux density (which has a boundary at 0) and fix the photon index l'=2.0."," To apply this calibration, we take the free parameter to be the flux density (which has a boundary at 0) and fix the photon index $\Gamma=2.0$."
" The cutoff remains fixed at 100 GeV. With this convention, the DC significance is then given by with £5p; the likelihood value obtained with the best-fit flux density and £ the likelihood value obtained with the flux density set to zero."," The cutoff remains fixed at 100 GeV. With this convention, the DC significance is then given by with $\mathcal{L}_{opt}$ the likelihood value obtained with the best-fit flux density and $\mathcal{L}_0$ the likelihood value obtained with the flux density set to zero."
" It should be noted that opc is a significance, so the same value of apc corresponds to a slightly higher value the “o” for the pulsed detection."," It should be noted that $\sigma_{DC}$ is a significance, so the same value of $\sigma_{DC}$ corresponds to a slightly higher value the $\sigma$ ” for the pulsed detection."
 This discrepancy is small compared to the observed difference in significance., This discrepancy is small compared to the observed difference in significance.
" As in the previous section, we calculate the probability weights using both the Monte Carlo truth values of the parameter and with the best-fit spectrum—in this case, with cutoff energy photon index fixed—to estimate a pulsed significance with the Haoy test."," As in the previous section, we calculate the probability weights using both the Monte Carlo truth values of the parameter and with the best-fit spectrum—in this case, with cutoff energy photon index fixed—to estimate a pulsed significance with the $H_{20w}$ test."
 We compare the unpulsed and pulsed significances in Figure 9.., We compare the unpulsed and pulsed significances in Figure \ref{ch5_plot15}.
 The majority of ensemble members are detected more significantly through pulsations than through unpulsed emission., The majority of ensemble members are detected more significantly through pulsations than through unpulsed emission.
 The measured significance for pulsed detections is a factor of 2.2 greater than for DC detection., The measured significance for pulsed detections is a factor of $2.2$ greater than for DC detection.
" If we assume that the detection flux threshold is linear in gpg as it is in the pulsed significance, this means we require sources to be about twice as bright on average to detect them through DC emission rather than pulsed emission, though the factor for a particular pulsar depends strongly on the light curve morphology, especially the sharpness of the peak(s)."," If we assume that the detection flux threshold is linear in $\sigma_{DC}$ as it is in the pulsed significance, this means we require sources to be about twice as bright on average to detect them through DC emission rather than pulsed emission, though the factor for a particular pulsar depends strongly on the light curve morphology, especially the sharpness of the peak(s)."
 This result also suggests a computationally efficient method for calculating pulsed flux upper limits., This result also suggests a computationally efficient method for calculating pulsed flux upper limits.
" The average ""efficiency"" factor (2.2 can be determined for some family of light curve above)templates, and pulsed flux upper limits can be determined directly from DC upper limits, which may themselves be straightforwardly calculated from the data with maximum likelihood methods."," The average “efficiency” factor (2.2 above) can be determined for some family of light curve templates, and pulsed flux upper limits can be determined directly from DC upper limits, which may themselves be straightforwardly calculated from the data with maximum likelihood methods."
" Finally, we note that Hoo, calculated with weights obtained from a spectral fit using only one degree of freedom—the flux density—is comparable to that obtained when calculated with both the flux density and the photon index allowed to vary."," Finally, we note that $H_{20w}$ calculated with weights obtained from a spectral fit using only one degree of freedom—the flux density—is comparable to that obtained when calculated with both the flux density and the photon index allowed to vary."
 This suggests an insensitivity to the precise spectral shape assumed for calculation of the weights., This suggests an insensitivity to the precise spectral shape assumed for calculation of the weights.
" This result has an important practical implication since the uncertainties on power law parameters can be quite large for dim sources (?),, in which case assuming fixed, canonical spectral shape is likely to yield better resultsa than fitting both flux density and power law slope."," This result has an important practical implication since the uncertainties on power law parameters can be quite large for dim sources \citep{1fgl}, in which case assuming a fixed, canonical spectral shape is likely to yield better results than fitting both flux density and power law slope."
values οςusistent with the errors.,values consistent with the errors.
 As suggested by Rovithis et al. (, As suggested by Rovithis et al. (
1996. 1909). we examined whether the orbital period could be represented by a quadratic eplemeris via a coutinuous period increase but [alled to give a satisfactory result.,"1996, 1999), we examined whether the orbital period could be represented by a quadratic ephemeris via a continuous period increase but failed to give a satisfactory result."
 Lusteacl we found that all times of minimnun light are best fitted by the combination of an upward parabolic variation and a liglit-travel time (LTT) ellect caused by the presence of a third body in the system. uamely a quadraticplus LTT ephemeris: where 74 is the LTT due to a third body (Irwin 1952. 1959) and introduces additioal live j)araimeters (ayo sit’. €. αἱ.n. T).," Instead, we found that all times of minimum light are best fitted by the combination of an upward parabolic variation and a light-travel time (LTT) effect caused by the presence of a third body in the system, namely a quadratic LTT ephemeris: where $\tau_{3}$ is the LTT due to a third body (Irwin 1952, 1959) and introduces additional five parameters $a_{12}\sin i_3$ , $e$ , $\omega$ ,$n$, $T$ )."
 Here. a4»sinis. e. and w are the orbital parameters of the eclipsing pair around the mass center oft1e triple system.," Here, $a_{12}\sin i_3$, $e$, and $\omega$ are the orbital parameters of the eclipsing pair around the mass center of the triple system."
 The parameters n aid T denote Ixedleriau neal mWolion of the mass center of the eclipsing pair aud the epoch of its periastron passage. 'espectisrely.," The parameters $n$ and $T$ denote Keplerian mean motion of the mass center of the eclipsing pair and the epoch of its periastron passage, respectively."
 The Leveubere-Marquart aeorithiu (Press et al., The Levenberg-Marquart algorithm (Press et al.
 1992) was applied ο solve or the uikuowii parameters of the eghemeris., 1992) was applied to solve for the unknown parameters of the ephemeris.
 The results are summarized in Tabe 5. together with the hird-bodsy mnasses (Ma) caleulated for tiree different. inclinations of 73.," The results are summarized in Table 5, together with the third-body masses $M_3$ ) calculated for three different inclinations of $i_3$ ."
 Οιr absolute cimeusions lave been usec or these aud subsequent caleulaticLIS., Our absolute dimensions have been used for these and subsequent calculations.
 The Q-C' diagranm constructed witli he liuear terms of the equation is plotted i he top pauel ol Figure D+) where the contiuuous curve aud the dashed: parabola represent. the fLr contribution ALU the quadratic teri. respectively.," The $O$ $C$ diagram constructed with the linear terms of the equation is plotted in the top panel of Figure 3, where the continuous curve and the dashed parabola represent the full contribution and the quadratic term, respectively."
 ΤΙ e1uidele panel displays the LTT orbit. a the bottom patel the residuals rou the complee eplieueris.," The middle panel displays the LTT orbit, and the bottom panel the residuals from the complete ephemeris."
 These appear as O-Ciyy in the edbh column of Taldle 1., These appear as $O$ $C_{\rm full}$ in the fifth column of Table 4.
 As disdlaved in Figure D3.+) the €qua|‘aticplus LTT epelmeris currently provides a good rep‘esentation of al the O-C' residuals.," As displayed in Figure 3, the quadratic LTT ephemeris currently provides a good representation of all the $O$ $C$ residuals."
" If he third companion is ou the main sectelce alcL its orbit is ¢'oplauar with the eclipsing biuary (Le.. £4258Y.2). the jass of the object is Aa42:0.2( DML. all its radius and temperature are calculatec to be Ry=0.27 R.. €and 75230Ls kx. respectively, Slug the mass-racliIs alc driass-LemperatureU relatious f‘om wellbstucied eclipsing binaries (Southworth 2000)."," If the third companion is on the main sequence and its orbit is coplanar with the eclipsing binary (i.e., $i_3$ $^\circ$ .2), the mass of the object is $M_3$ =0.26 $_\odot$ and its radius and temperature are calculated to be $R_3$ =0.27 $_\odot$ and $T_3$ =3048 K, respectively, using the mass-radius and mass-temperature relations from well-studied eclipsing binaries (Southworth 2009)."
 These correspond to a spectral type of abou NO V and a »olometric luminosity ο ‘£3 =0.006 L.. andl ¢ontribute abou to tle otal light of he triple system., These correspond to a spectral type of about M6 V and a bolometric luminosity of $L_3$ =0.006 $_\odot$ and contribute about to the total light of the triple system.
 So. it will be cdifficu1 to detect Sch a αςmpanuion frou. the light-cirve analysis aud specroscopi¢ observatiolis.," So, it will be difficult to detect such a companion from the light-curve analysis and spectroscopic observations."
 The positive coellicient of the quadratic term iu Table 5 indicates a contiuuous period increase jth a ale of 3.78 x10 d which cau be expaimed by a mass transfer [from the cool concary μίας to its mo'e massive primary compolent because WZ Cye is a contact binary syste.," The positive coefficient of the quadratic term in Table 5 indicates a continuous period increase with a rate of 3.78 $\times$ $^{-8}$ d $^{-1}$, which can be explained by a mass transfer from the cool secondary star to its more massive primary component because WZ Cyg is a contact binary system."
 xler te asstumptlion of conservative nass transer. ile trausfer rate is SOxIO 5 M. 1 1H. Of the secodary star transfers its present mass to 1le priuary component on a thermal time scale Τι ETE13)/(RoLs). then Ath 1.05x10* vr ane ass is trauslerrecd to the primary at a rate given roehly by Λούτσα M.I.," Under the assumption of conservative mass transfer, the transfer rate is $\times$ $^{-8}$ $_\odot$ $^{-1}$ If the secondary star transfers its present mass to the primary component on a thermal time scale $\tau_{\rm th}$ $(GM_{2}^2)/(R_{2}L_{2})$, then $\tau_{\rm th}$ = $\times 10^{7}$ yr and mass is transferred to the primary at a rate given roughly by $M_{\rm 2}/ \tau_{\rm th}$$\times 10^{-8}$ $_{\sun}$$^{-1}$."
 TUs vaue is very close to themass trausfer rate calculated from our quadratic term. which means that tLe mass (rausfer between the two stars can explain tie secular period tucrease of the system satisfactorily.," This value is very close to themass transfer rate calculated from our quadratic term, which means that the mass transfer between the two stars can explain the secular period increase of the system satisfactorily."
We have written the source term of the wave equation (5) in terius of the equivalent density pp.,We have written the source term of the wave equation (5) in terms of the equivalent density $\rho_1^\ast$.
 To derive this density. we begin with df. the potential due to the binary.," To derive this density, we begin with $\Phi_1^\ast$, the potential due to the binary."
 For any spatially compact source. the potential is conveniently expressed as a stun over immultipoles.," For any spatially compact source, the potential is conveniently expressed as a sum over multipoles."
 Equation (1.2) ol Jackson(1902). gives the electrostatic expression: this is readily altered to the gravitational oue: Here. the multipole moments are found from the source cleusity (Jackson1962.eq.CL3)): Iu these two expressious. prinied coordinates refer to source points. aud uuprimed ones to the field points where «P1 is measured.," Equation (4.2) of \citet{j62} gives the electrostatic expression; this is readily altered to the gravitational one: Here, the multipole moments are found from the source density \citep[][eq. (4.3)]{j62}: In these two expressions, primed coordinates refer to source points, and unprimed ones to the field points where $\Phi_1^\ast$ is measured."
 Figure 1 shows the binary lying in the 2-7 plane of our coordiuate system. with its center of Inass situated at the origin.," Figure 1 shows the binary lying in the $x$ $y$ plane of our coordinate system, with its center of mass situated at the origin."
 Also indicated are two of the three coordinates (1.0.6) of a liekl point.," Also indicated are two of the three coordinates $(r,\theta,\phi)$ of a field point."
 Ai time /. mass Af; is a distance from the origin. and has rotated [rom the a-axis by the azimuthal angle/£.. where w is the binarys angular speed.," At time $t$, mass $M_1$ is a distance from the origin, and has rotated from the $x$ -axis by the azimuthal angle, where $\omega$ is the binary's angular speed."
 The other mass Alsis located at anddo., The other mass $M_2$is located at and.
 Both masses have polar angles7/2., Both masses have polar angles.
 Evaluating the qu; from equation (8). we fiud that the monopole term is Iu coutrast. all dipole terms vanish.," Evaluating the $q_{lm}$ from equation (8), we find that the monopole term is In contrast, all dipole terms vanish."
 For example. by clefinition of the center of mass.," For example, by definition of the center of mass."
 The quadrupole terius are the most interesting for our purpose., The quadrupole terms are the most interesting for our purpose.
 For 0.. we have," For , we have"
xuxl characteristic of low Ctadlactic latitudes.,band characteristic of low Galactic latitudes.
 Future observations in laree-aperture optical telescopes could establish a definitive spectroscopical identification of thi blazars., Future observations in large-aperture optical telescopes could establish a definitive spectroscopical identification of both blazars.
 This work has been made possible for the first time jecause of the continuous. well sampled light curves iu he eanuna-ray and radio bands obtained over 2.5 vears wFerm? aud the fast-cadeuce OVRO blazar monitoring xoerau.," This work has been made possible for the first time because of the continuous, well sampled light curves in the gamma-ray and radio bands obtained over 2.5 years by and the fast-cadence OVRO blazar monitoring program."
 I completing this analysis. we establish fiii associations for the blazars B2013|370 and B2023|336 with previously unidentified EGRET and Fern-LAT sources.," In completing this analysis, we establish firm associations for the blazars B2013+370 and B2023+336 with previously unidentified EGRET and -LAT sources."
 This confirms the eamuna-ray baud as likely the vest tool to ideutifv blazars located behind the Calactic ralane. which ave heavily absorbed at other wavelengths. and supports the hypothesis that a umuber uuideutified eanmu-ray sources at low Calactic latitudes are indeed of extragalactic origin.," This confirms the gamma-ray band as likely the best tool to identify blazars located behind the Galactic plane, which are heavily absorbed at other wavelengths, and supports the hypothesis that a number unidentified gamma-ray sources at low Galactic latitudes are indeed of extragalactic origin."
"weak lensing is the reduced shear, and the location of a vanishing determinant of the Jacobian in the case of strong lensing.","weak lensing is the reduced shear, and the location of a vanishing determinant of the Jacobian in the case of strong lensing."
" In order to minimise the resulting y?-function with respect to the lensing potential, we have to write the convergence and the shear in terms of the discrete potential values."," In order to minimise the resulting $\chi^{2}$ -function with respect to the lensing potential, we have to write the convergence and the shear in terms of the discrete potential values."
 The relation between these quantities is given by Eqs. (2)). (3))," The relation between these quantities is given by Eqs. \ref{convergence}) ), \ref{shear1}) )"
 and (4)) with the second derivatives replaced by finite-differencing schemes., and \ref{shear2}) ) with the second derivatives replaced by finite-differencing schemes.
" Given a certain grid resolution, we want to obtain the lensing potential at each grid position (x,y) with 1€xM, 1€yN."," Given a certain grid resolution, we want to obtain the lensing potential at each grid position $(x,y)$ with $1\le x\le M$, $1\le y\le N$."
 We enumerate these grid positions sequentially line by line and use the central difference quotients for discrete representations of the second derivatives at a given grid position., We enumerate these grid positions sequentially line by line and use the central difference quotients for discrete representations of the second derivatives at a given grid position.
" Our finite differencing scheme is identical to that chosen by ?,, and the corresponding coefficients are given in Fig. H]."," Our finite differencing scheme is identical to that chosen by \citet{Cacciato2006}, and the corresponding coefficients are given in Fig. \ref{findifschemes}."
 At the edges and borders of the grid different finite-differencing schemes need to be used., At the edges and borders of the grid different (one-sided) finite-differencing schemes need to be used.
" With our way of enumerating the pixels, we can write the discrete potential y on the NxM grid in the form of a vector VεΕΝΜ, "," With our way of enumerating the pixels, we can write the discrete potential $\vec{\psi}$ on the $N\times M$ grid in the form of a vector $\vec{\psi}\in \mathbb{R}^{NM}$."
This only slightly complicates addressing the correct positions in the vector., This only slightly complicates addressing the correct positions in the vector.
" Direct left and right neighbours are just separated by +1 positions in the vector, while top and bottom neighbours are separated by +N positions."," Direct left and right neighbours are just separated by $\pm1$ positions in the vector, while top and bottom neighbours are separated by $\pm N$ positions."
" The advantage of this enumeration is that the finite differencing becomes a simple matrix multiplication, Here, Kj;,σι,σι are sparse band matrices encoding the information on the finite-differencing scheme."," The advantage of this enumeration is that the finite differencing becomes a simple matrix multiplication, Here, $\mathcal{K}_{ij}, \enspace \mathcal{G}^{1}_{ij}, \enspace \mathcal{G}^{2}_{ij}$ are sparse band matrices encoding the information on the finite-differencing scheme."
" The fact that we know these matrices perfectly is the key point to increase the speed of our reconstruction algorithms, which plays an important role with respect to runtime, besides the parallelisation of the code."," The fact that we know these matrices perfectly is the key point to increase the speed of our reconstruction algorithms, which plays an important role with respect to runtime, besides the parallelisation of the code."
 Without several numerical tricks the runtime of a complete reconstruction would not stay at an acceptable level., Without several numerical tricks the runtime of a complete reconstruction would not stay at an acceptable level.
" Using these finite-differencing schemes, y? can be minimised by solving the linear system of equations where the coefficient matrix 8, and the result vector V; contain on the one hand observed ellipticity and critical curve information and on the other hand the convergence and shear matrices."," Using these finite-differencing schemes, $\chi^{2}$ can be minimised by solving the linear system of equations where the coefficient matrix $\mathcal{B}_{lk}$ and the result vector $\mathcal{V}_{l}$ contain on the one hand observed ellipticity and critical curve information and on the other hand the convergence and shear matrices."
 The calculation is detailed in the Appendix., The calculation is detailed in the Appendix.
 We now introduce a regularisation term R(i) into the y? function Eq. (10)), We now introduce a regularisation term $R(\psi)$ into the $\chi^2$ function Eq. \ref{chi2}) )
 to obtain The regularisation term depends on the potential and is defined such as to disfavour unwanted solutions., to obtain The regularisation term depends on the potential and is defined such as to disfavour unwanted solutions.
" This is necessary to prevent the reconstruction from following intrinsic noise patterns, which do not reflect intrinsic features of the true underlying potential."," This is necessary to prevent the reconstruction from following intrinsic noise patterns, which do not reflect intrinsic features of the true underlying potential."
" Here, we use a simple regularisation term scheme simply comparing the convergence obtained with that found in the preceding iteration, Its amplitude 7; controls the relative strength of the regularisation and is crucial for the reconstruction."," Here, we use a simple regularisation term scheme simply comparing the convergence obtained with that found in the preceding iteration, Its amplitude $\eta_{i}$ controls the relative strength of the regularisation and is crucial for the reconstruction."
 It should be chosen such that the overall y? is of order unity and of course low enough for changes in the reconstruction to take effect., It should be chosen such that the overall $\chi^{2}$ is of order unity and of course low enough for changes in the reconstruction to take effect.
 Minimising the regularised y? of Eq. Qo» , Minimising the regularised $\chi^2$ of Eq. \ref{eta}) )
again leads to a linear system which contributes one additional term to the coefficient matrix and the result vector., again leads to a linear system which contributes one additional term to the coefficient matrix and the result vector.
 The complete calculation is presented in the Appendix., The complete calculation is presented in the Appendix.
 In our reconstruction we also regularize on the two shear components which give similar terms., In our reconstruction we also regularize on the two shear components which give similar terms.
" Note that it is not useful to directly regularize the potential since it looks very different in each iteration step, as described in Sect. [3.3]."," Note that it is not useful to directly regularize the potential since it looks very different in each iteration step, as described in Sect. \ref{lensingpotentialandconvergencemaps}."
" As the full expressions for the Y. minimisation show, it is not precisely a linear system, but includes non-linear prefactor terms in the coefficient matrix and the result vector."," As the full expressions for the $\chi^{2}$ minimisation show, it is not precisely a linear system, but includes non-linear prefactor terms in the coefficient matrix and the result vector."
 We solve this problem in the same way as ? did by means of an iterative approach in which the non-linear terms are computed from the preceding iteration., We solve this problem in the same way as \citet{Bradav2005} did by means of an iterative approach in which the non-linear terms are computed from the preceding iteration.
" Starting from a first guess for the convergence, we express the corresponding non-linear terms by constant factors, as can be seen in the Appendix."," Starting from a first guess for the convergence, we express the corresponding non-linear terms by constant factors, as can be seen in the Appendix."
" ? showed that the initial guess of the convergence does not affect the final reconstruction, but at most the number of iterations."," \citet{Bradav2005} showed that the initial guess of the convergence does not affect the final reconstruction, but at most the number of iterations."
" We confirm their result, which implies that the initial guess of a flat convergence is appropriate."," We confirm their result, which implies that the initial guess of a flat convergence is appropriate."
 We minimise X) and obtain a solution for the lensing potential., We minimise $\chi^2$ and obtain a solution for the lensing potential.
" From it, we calculate convergence and shear, which we insert as new guesses into the non-linear terms in the next step of the reconstruction."," From it, we calculate convergence and shear, which we insert as new guesses into the non-linear terms in the next step of the reconstruction."
 This yields a convergent iterative process., This yields a convergent iterative process.
 We control the convergence of the procedure by comparing its results between subsequent iterations., We control the convergence of the procedure by comparing its results between subsequent iterations.
" If the change falls below a given threshold, we stop the iteration."," If the change falls below a given threshold, we stop the iteration."
" The drawback of this method is that we now need several iterations (3-5 in practice) for a complete reconstruction, which takes some time at high resolution."," The drawback of this method is that we now need several iterations (3-5 in practice) for a complete reconstruction, which takes some time at high resolution."
" The regularisation term also helps in a second type of iteration, which we call outer-level iteration."," The regularisation term also helps in a second type of iteration, which we call outer-level iteration."
 The background-galaxy density of today's observations allows a resolution of only 10x uncorrelated pixels., The background-galaxy density of today's observations allows a resolution of only $\sim10\times10$ uncorrelated pixels.
" Higher resolution is desirable at the expense of correlated pixels, which will affect the reconstruction."," Higher resolution is desirable at the expense of correlated pixels, which will affect the reconstruction."
" Also the convergence values in some individual pixels will increase at higher resolution, which renders the initial guess of a vanishing convergence increasingly less accurate, giving rise to many inner-level iterations."," Also the convergence values in some individual pixels will increase at higher resolution, which renders the initial guess of a vanishing convergence increasingly less accurate, giving rise to many inner-level iterations."
algorithms should. be emploved.,algorithms should be employed.
 uses a singular value decomposition., uses a singular value decomposition.
 Equations (8) and (9) are the same. so it does. not matter which form one uses.," Equations (8) and (9) are the same, so it does not matter which form one uses."
 However one must. in most cases. first estimate the covariance matrix and confirm that ib is correct.," However one must, in most cases, first estimate the covariance matrix and confirm that it is correct."
 This requires the Cholesky (or equivalent) transformation., This requires the Cholesky (or equivalent) transformation.
 So the transformation must. be found and applied in any case., So the transformation must be found and applied in any case.
 In practice we find it cllicient to use an OLS solver for al cases so we first [orm My and Ae then use equation (9)., In practice we find it efficient to use an OLS solver for all cases so we first form ${\bf M}_{\rm w}$ and $\vec{R}_{\rm w}$ then use equation (9).
 Spectral analysis is intimately connected to the analysis of pulsar timing observations in (wo wavs., Spectral analysis is intimately connected to the analysis of pulsar timing observations in two ways.
 First. the timing mocel includes a number of parameters which control the amplitude and phase of sine waves. so [fitting for those parameters amounts to a spectral analysis.," First, the timing model includes a number of parameters which control the amplitude and phase of sine waves, so fitting for those parameters amounts to a spectral analysis."
 Second. optimal least squares fitting requires estimation of the covariance matrix and this will require some form of spectral analysis.," Second, optimal least squares fitting requires estimation of the covariance matrix and this will require some form of spectral analysis."
 llere we provide a very brief. discussion of the aspects of spectral analysis important to the analvsis of pulsar timing observations., Here we provide a very brief discussion of the aspects of spectral analysis important to the analysis of pulsar timing observations.
 We will not attempt. to provide a complete mathematical ceseription because there are many books devoted to the subject (e.g... Blackman anc Pukey. 1959. Jenkins anc Watts. 1969: Manolakis. Ingle and Ixogon.2456.," We will not attempt to provide a complete mathematical description because there are many books devoted to the subject (e.g., Blackman and Tukey, 1959, Jenkins and Watts, 1969; Manolakis, Ingle and Kogon,."
. Spectral analwsis of a series. of regularly spaced observations of length T is normally performed via. the discrete. Fourier transform. (DITE)., Spectral analysis of a series of regularly spaced observations of length T is normally performed via the discrete Fourier transform (DFT).
 The squared magnitude of the DET. properly. scaled. known as the periodogram πο. is often used. as the spectral estimator.," The squared magnitude of the DFT, properly scaled, known as the periodogram $P(f)$, is often used as the spectral estimator."
 Here. the sample interval is ὁ and we normalize (f) to have the units ofa power spectral densitv., Here the sample interval is $\delta$ and we normalize $P(f)$ to have the units of a power spectral density.
 Llowever. Pf) will be biased. if the spectrum. being estimated. is not white.," However, $P(f)$ will be biased if the spectrum being estimated is not white."
 The bias is caused. by the finite length. of data., The bias is caused by the finite length of data.
 Ehe data are effectively multiplied by a time window w(f). which is often rectangular (i-c.. it is unity during the observations and. zero otherwise).," The data are effectively multiplied by a time window $w(t)$, which is often rectangular (i.e., it is unity during the observations and zero otherwise)."
 Thus the Fourier transform of r(/)e(/) is the convolution of Af) with Wf)., Thus the Fourier transform of $r(t)w(t)$ is the convolution of $R(f)$ with $W(f)$.
 The measured. power spectrum PCf) is the convolution of the spectral density SC) with VCO., The measured power spectrum $P(f)$ is the convolution of the spectral density $S(f)$ with $|W(f)|^2$.
 Fora rectangular window of length T the spectral window is given This window is most widely used because it provides the jighest frequency resolution., For a rectangular window of length T the spectral window is given by This window is most widely used because it provides the highest frequency resolution.
 This is particularly importan in pulsar timing analvsis., This is particularly important in pulsar timing analysis.
 Phe sidelobes of this window all olf as f7., The sidelobes of this window fall off as $f^{-2}$.
 This makes power law spectra which [fal aster than f7 heavily biased: because window. sidelobes will ominate all frequencies higher. than f=1/7., This makes power law spectra which fall faster than $f^{-2}$ heavily biased because window sidelobes will dominate all frequencies higher than $f=1/T$.
" Steceper spectra must be analysed with ""prewhitening"" anc “postcarkenine” to minimize spectral leakage (c.g.. Jenkins and Watts. 1969)."," Steeper spectra must be analysed with “prewhitening” and “postdarkening” to minimize spectral leakage (e.g., Jenkins and Watts, 1969)."
 One applies a linear pre-whitening filter which is implemented in the time domain., One applies a linear pre-whitening filter which is implemented in the time domain.
 Lhe purpose of his filter is to make the spectrum close enough to white hat leakage is insignificant., The purpose of this filter is to make the spectrum close enough to white that leakage is insignificant.
 For example. if (A) is the input and gy(À) is the output. a first. dillerence filter. is yh)ξ-orth1).," For example, if $x(k)$ is the input and $y(k)$ is the output, a first difference filter is $y(k) = x(k) - x(k-1)$ ."
 The ellect of this filter is to multiply he Fourier transform of the input bv the transfer function H(f)=25sin(xfo)., The effect of this filter is to multiply the Fourier transform of the input by the transfer function $H(f) = 2 \sin(\pi f \delta)$.
 This multiplies the power spectrum by HO»=(Qn[5)Y. which has an F behavior at low requencies.," This multiplies the power spectrum by $|H(f)|^2 = (2 \sin( f \delta))^2$, which has an $f^2$ behavior at low frequencies."
 Ehus it will whiten a power law spectrum with exponent 2., Thus it will whiten a power law spectrum with exponent $-$ 2.
 The whitened data can be analyzed with little spectral leakage if its spectral exponent is between -0.5 and -3.5., The whitened data can be analyzed with little spectral leakage if its spectral exponent is between -0.5 and -3.5.
 We then correct the estimated spectrum of the whitened data by dividing it by |J(fr., We then correct the estimated spectrum of the whitened data by dividing it by $|H(f)|^2$.
 This is called post-clarkening., This is called post-darkening.
 For steeper spectra a second clilfercnee (two applications of the first dillerence) will be required., For steeper spectra a second difference (two applications of the first difference) will be required.
 Phe transfer function of the second. dillerence is. HCf)7, The transfer function of the second difference is $H(f)^2$.
 Unfortunately: spectra often are power law at low frequencies. descending into white noise at high frequencies.," Unfortunately spectra often are power law at low frequencies, descending into white noise at high frequencies."
 The white noise will be ransformed⋅ to f↝⊥ behavior: and spectral leakage can occur vackwards from. high frequencies to. low frequencies., The white noise will be transformed to $f^4$ behavior and spectral leakage can occur backwards from high frequencies to low frequencies.
 To event this one must also filter the original observations with a low-pass filter. implemented. in the time domain. which has a transfer function of the form L(f)=1/(1(f/f 7).," To prevent this one must also filter the original observations with a low-pass filter, implemented in the time domain, which has a transfer function of the form $L(f) = 1/(1 + (f/f_c )^2)$ ."
 Here the corner [requeney. fo is chosen to be near he frequeney at which the power law component reaches he white noise level., Here the corner frequency $f_c$ is chosen to be near the frequency at which the power law component reaches the white noise level.
 Ehe residuals after second dilferencing ancl low pass filtering are then spectrum analyzed., The residuals after second differencing and low pass filtering are then spectrum analyzed.
 Phe post-darkening is done by dividing the estimated spectrum of the ilterecl data by μι”., The post-darkening is done by dividing the estimated spectrum of the filtered data by $|L(f)|^2 |H(f)|^4$.
 Since we do not know the spectral exponent a priori. the analysis must be done in stages.," Since we do not know the spectral exponent a priori, the analysis must be done in stages."
 A first estimate can be done with a periodogram., A first estimate can be done with a periodogram.
 Lit appears to be limited. hy leakage then another analysis must be done using first. dilference pre-whitening and. post-darkening., If it appears to be limited by leakage then another analysis must be done using first difference pre-whitening and post-darkening.
 Lit still appears to be limited by leakage one must do another analysis using second differencing and low pass filtering., If it still appears to be limited by leakage one must do another analysis using second differencing and low pass filtering.
 In some cases à more complex pre-whitening filter may. be required. but we have not seen any such cases in pulsar timing analyses.," In some cases a more complex pre-whitening filter may be required, but we have not seen any such cases in pulsar timing analyses."
 We can apply pre-whitening ancl post-darkening. as described. above. only to regularly sampled: measurements because the temporal filtering operations requires equally spaced. data.," We can apply pre-whitening and post-darkening, as described above, only to regularly sampled measurements because the temporal filtering operations requires equally spaced data."
 Fortunately this is possible for pulsar timing applications because the low frequency. spectra are often power-law and the high [frequencies are white., Fortunately this is possible for pulsar timing applications because the low frequency spectra are often power-law and the high frequencies are white.
 We can use a low pass filter to separate the low and hieh frequencies., We can use a low pass filter to separate the low and high frequencies.
 The low frequencies are heavily over-sampled. and. can be interpolatecl onto a regular erid. without much bias., The low frequencies are heavily over-sampled and can be interpolated onto a regular grid without much bias.
 The high frequencies cannot be interpolated. but as they are white they will not suffer from spectral leakage and we can use existing least-squares methods.," The high frequencies cannot be interpolated, but as they are white they will not suffer from spectral leakage and we can use existing least-squares methods."
 Εθν we can assemble a Composite spectrum. by splitting the spectrum. with a low pass filter ancl analyzing the two halves with dilleren methods The DET is equivalent to a least. squares fitting of complex exponentials. but since the complex exponentials are orthogonal in the regularly sampled DET. they can be fit either individually or simultaneously ancl the result. wil be the same.," Thus we can assemble a composite spectrum by splitting the spectrum with a low pass filter and analyzing the two halves with different methods The DFT is equivalent to a least squares fitting of complex exponentials, but since the complex exponentials are orthogonal in the regularly sampled DFT, they can be fit either individually or simultaneously and the result will be the same."
 Ifthe sample spacing is not regular the complex exponentials are not orthogonal., If the sample spacing is not regular the complex exponentials are not orthogonal.
 His not. possible to fi them all simultaneously because they are highly covariant., It is not possible to fit them all simultaneously because they are highly covariant.
 However we can obtain unbiased estimates of the variance a eachfrequeney using the Lomb-Scargle L-S (Scarele. or Z-Ix (Zechmoelster and Ixurster. methods.," However we can obtain unbiased estimates of the variance at eachfrequency using the Lomb-Scargle L-S (Scargle, or Z-K (Zechmeister and Kurster, methods."
 This is, This is
As already. iueutioned in 11. a comparison of radial velocities from inversion lines of anuuonia (NIT4) with rotational lines of either NIT; or other species can provide sensitive coustraints on variations m the protou- ass ratio f= my fimi.,"As already mentioned in 1, a comparison of radial velocities from inversion lines of ammonia $_3$ ) with rotational lines of either $_3$ or other species can provide sensitive constraints on variations in the proton-to-electron mass ratio $\mu$ = $m_{\rm p}$ $m_{\rm e}$ ."
 Following Fliannbiiun RKozlov (2007). with su and n4; denoting the appareut redshifts of the inversion and rotational lines. 2 represcuting the real redshift of the source. and Ay providing the deviation relative to the current laboratory value of µ. defined such that a negative value indicates a smaller jj in the absorbing cloud.," Following Flambaum Kozlov (2007), with $z_{\rm inv}$ and $z_{\rm rot}$ denoting the apparent redshifts of the inversion and rotational lines, $z$ representing the real redshift of the source, and $\Delta$$\mu$ providing the deviation relative to the current laboratory value of $\mu$, defined such that a negative value indicates a smaller $\mu$ in the absorbing cloud."
 Tn the case where both the NIT; aud rotational trausitious show absorption iu just a single velocity component. defining aud measuring the velocity difference AV is straightforward.," In the case where both the $_3$ and rotational transitions show absorption in just a single velocity component, defining and measuring the velocity difference $\Delta V$ is straightforward."
 The more general case where many velocity colmpoucuts exist must be treated differently by varving all the free parameters. includiug AV. to minimize 47 iu a simmultaucous fit to all NIT; and rotational profiles (c.¢.. Murphliv et al.," The more general case where many velocity components exist must be treated differently, by varying all the free parameters, including $\Delta V$, to minimize $\chi^2$ in a simultaneous fit to all $_3$ and rotational profiles (e.g., Murphy et al."
 2008a)., 2008a).
 Anunonia and evanoacetvlene (IIC4N). are both uolecular lieh density tracers., Ammonia and cyanoacetylene $_3$ N) are both molecular high density tracers.
 Tow 11830211. he NIL3 inversion lines and the rotational Που lines are optically thin aud observed at similar frequencies.," Toward 1830–211, the $_3$ inversion lines and the rotational $_3$ N lines are optically thin and observed at similar frequencies."
 Therefore. in this paper we derive Aqr/p dy comparing he NID; and spectra.," Therefore, in this paper we derive $\Delta\mu/\mu$ by comparing the $_3$ and $_3$ N spectra."
 Civeu the fairly low sigual-to-roise ratio of these spectra. there is little obvious evideuce or iore than a single absorbine velocity component across the NIT; and 10 profiles.," Given the fairly low signal-to-noise ratio of these spectra, there is little obvious evidence for more than a single absorbing velocity component across the $_3$ and $_3$ N profiles."
 We therefore proceed with a simple analysis assume that oulv one absorbing velocity component exists. deferriue a simmltancous 4? nüniuization analysis which incorporates other weak coniponeuts and other rotational transitions to a later paper.," We therefore proceed with a simple analysis assuming that only one absorbing velocity component exists, deferring a simultaneous $\chi^2$ minimization analysis which incorporates other weak components and other rotational transitions to a later paper."
 Despite this assumption of a single velocity coniponent. judicious treatineut of the uncertainties on its velocity im the NID; aud rotational transitions will eusure our constraint on μμ ds robust. if not optimally precise.," Despite this assumption of a single velocity component, judicious treatment of the uncertainties on its velocity in the $_3$ and rotational transitions will ensure our constraint on $\Delta\mu/\mu$ is robust, if not optimally precise."
 The NII; inversion lines were measured with the Green Dauk Telescope (GBT). while the Που data were obtained at Effelsbere.," The $_3$ inversion lines were measured with the Green Bank Telescope (GBT), while the $_3$ N data were obtained at Effelsberg."
 The Local Staudard of Rest velocity scales at both sites are. however. well adjusted (τισ Iluchtineier 1990) and have not been siguificautlv modified in recent vears.," The Local Standard of Rest velocity scales at both sites are, however, well adjusted (Tifft Huchtmeier 1990) and have not been significantly modified in recent years."
" They refer to a peculiar solar velocity of + toward ayouu = LS! an 1900 = 130°. correspouding to VospVary, = 1 in the case of 11830211."," They refer to a peculiar solar velocity of $^{-1}$ toward $\alpha_{1900}$ = $^{\rm h}$ and $\delta_{1900}$ = $^{\circ}$, corresponding to $V_{\rm LSR}-V_{\rm HEL}$ = $^{-1}$ in the case of 1830–211."
 Therefore. the velocity measurements listed in Table 6 can be directly compared to estimate App and. its uncertainty. without further velocity corrections.," Therefore, the velocity measurements listed in Table 6 can be directly compared to estimate $\Delta\mu/\mu$ and its uncertainty, without further velocity corrections."
" For the sinele-CGaussian fits to the NITI, transitions. Table 6 reveals somewhat inconsistent centroid velocities."," For the single-Gaussian fits to the $_3$ transitions, Table 6 reveals somewhat inconsistent centroid velocities."
 Thev scatter around he weighted mean velocity over all transitions of lowdth a 4? value of 117 which. for just Ll degrees of freedom. is far more than expected based on the individual velocity uncertainties.," They scatter around the weighted mean velocity over all transitions of $^{-1}$ with a $\chi^2$ value of 117 which, for just 14 degrees of freedom, is far more than expected based on the individual velocity uncertainties."
 This may dudicate the presence of more than one velocity compoucut., This may indicate the presence of more than one velocity component.
 Despite the large range of NIL; excitation states observed by Tenkel et al. (, Despite the large range of $_3$ excitation states observed by Henkel et al. (
2008) ποιο his we do not observe a trend in velocity with excitation state: For the three lowest NIT; inversion lines. the wuweiehted mean velocity isS.8240.15  and for the higher excite ones we find a mean of 8.912027 + (the errors are standard deviations on the neans).,"2008) – some K – we do not observe a trend in velocity with excitation state: For the three lowest $_3$ inversion lines, the unweighted mean velocity is$\pm$ $^{-1}$ and for the higher excited ones we find a mean of $\pm$ $^{-1}$ (the errors are standard deviations on the means)."
 Iudeed. the lareer than expected scatter in fitted velocities for the different NIT; transitions seenis to be random im sigu aid," Indeed, the larger than expected scatter in fitted velocities for the different $_3$ transitions seems to be random in sign and"
"As we will show in the following section. with too, ~10 s ODo30. Diy~ὃν) ο, 0.5. and ep0.1. most of the internal shocks energy is emitted in the soft N-rav. band and the predicted. [ux matches the observation of the X-ray Hare in GRB 011121.","As we will show in the following section, with $L_{\rm m}\sim
5\times 10^{49}~ {\rm ergs~s^{-1}}$ , $\delta t_\oplus\sim 10$ s, $\Gamma \sim 30$ , $\Gamma_{\rm sh}\sim 2$, $\epsilon_{\rm e} \sim
0.5$ , and $\epsilon_{\rm B} \sim 0.1$, most of the internal shocks energy is emitted in the soft X-ray band and the predicted flux matches the observation of the X-ray flare in GRB 011121."
 Alternatively. as shownin Barrauc οἱ al. (," Alternatively, as shownin Barraud et al. ("
"2005). with ο0.01 (correspondingly. £4,~107—1077CresS+ and Pay~ L06). 64)~ a few seconds. and Eafew hundreds. X-rav dominated. luminosity is expected. too (Please note that their caleulation of the internal shocks emission. is somewhat cillerent from ours: see Barraucl ct al.","2005), with $\epsilon\sim 0.01$ (correspondingly, $L_{\rm m}\sim
10^{51}-10^{52} ~{\rm ergs~s^{-1}}$ and $\Gamma_{\rm sh}\sim
1.06$ ), $\delta t_\oplus\sim $ a few seconds, and $\Gamma\sim {\rm
a~ few}$ hundreds, X-ray dominated luminosity is expected, too (Please note that their calculation of the internal shocks emission is somewhat different from ours; see Barraud et al."
 2005 for cletail)., 2005 for detail).
 Therefore. we believe that with proper parameters. X-ray laves do appear in the late internal shocks scenario.," Therefore, we believe that with proper parameters, X-ray flares do appear in the late internal shocks scenario."
" Following Dai Lu (2002). the comoving number densitv of the unshockecl outllow is estimated. by nozmLu(4x171, me). where my is the rest mass of proton."," Following Dai Lu (2002), the comoving number density of the unshocked outflow is estimated by $n_{\rm e}\approx L_{\rm m}/(4\pi \Gamma^2 R_{\rm int}^2 m_{\rm p}c^3)$ , where $m_{\rm p}$ is the rest mass of proton."
 The thermal cnerey density of the shocked material is calculated by οzΕΕμίαDasme (Blandford Melxee 1971).," The thermal energy density of the shocked material is calculated by $e \approx 4\Gamma_{\rm sh} (\Gamma_{\rm
sh}-1)n_{\rm e}m_{\rm p}c^2$ (Blandford McKee 1977)."
 The intensity of the generated. magnetic. Ποια is estimated by 2Bz(Sircpe)+7610ου...GiCp.1Pan(hoa.1)op2])τοLaliasco∖⊥⊥2 aas ," The intensity of the generated magnetic field is estimated by $B\approx (8\pi \epsilon_{\rm B}e)^{1/2} \approx 6\times 10^3~{\rm G}~
\epsilon_{\rm B,-1}^{1/2}[\Gamma_{\rm sh}(\Gamma_{\rm
sh}-1)/2]^{1/2} L_{\rm m,49.7}^{1/2}\Gamma_{1.5}^{-1}R_{\rm int,14.5}^{-1}$ ."
"As usual. we assume that in. the shock front. the aecelerated electrons distribute as sUforse.> ""caus Where seo=c(lDipημίν 1)m.]is the minimum Lorentz factor of the shocked electrons (Sari ct al."," As usual, we assume that in the shock front, the accelerated electrons distribute as $dn_e/d\gamma_{\rm e}\propto
\gamma_{\rm e}^{\rm -p}~~{\rm for~\gamma_{\rm e}>\gamma_{\rm e,m}}$ , where $\gamma_{\rm e, m}=\epsilon_{\rm e}
(\Gamma_{\rm sh}-1)[(p-2)m_{\rm p}]/
[(p-1)m_{\rm e}]$ is the minimum Lorentz factor of the shocked electrons (Sari et al."
 1998). and m is the rest mass of electron.," 1998), and $m_{\rm e}$ is the rest mass of electron."
 La this section. we take p=2.5.," In this section, we take p=2.5."
"- The observed typical frequency of the svnchrotron radiation reads =a dee aue ie. most of the shock energy is emitted in the soft X-ray band. where q, is the charge of electron."," The observed typical frequency of the synchrotron radiation reads = , i.e., most of the shock energy is emitted in the soft X-ray band, where $q_e$ is the charge of electron."
 The cooling Lorentz factor is estimated by (e.g. Sarl etal," The cooling Lorentz factor is estimated by (e.g., Sari et al."
 1908) 5.zmT10|λαμοι). ancl the corresponding cooling frequency. reacs Lau ," 1998) $\gamma_{\rm e,c}
\approx 7.7\times 10^8 (1+z) /(\Gamma B^2\delta t_\oplus)$, and the corresponding cooling frequency reads = q_e [1.36/(1+z)]^2 ."
The synchirotron selt-absorption frequency is estimated bv (Li Song 2004) , The synchrotron self-absorption frequency is estimated by (Li Song 2004) .
"The maximum spectral ux of the svachrotron radiation is (ec. Wijers Galama 1999). fsm330,01|NomelB(ροπή,Dy). where IN.=La(1|Pre]=8G1077136/0.ΕΕtat) is the number of electrons involved in the emission."," The maximum spectral flux of the synchrotron radiation is (e.g., Wijers Galama 1999) $F_{\rm max} \approx 3\sqrt{3}\Phi_{\rm p}(1+z)N_{\rm e}m_{\rm e}c^2\sigma_{\rm T}\Gamma B/(32\pi^2 q_e D_{\rm L}^2)$ , where $N_{\rm e}= L_{\rm m} \delta t/[(1+z)\Gamma m_{\rm p}c^2]=8\times 10^{51}~[1.36/(1+z)]L_{\rm m,49.7}\Gamma_{1.5}^{-1}
\delta t_{\oplus,1}$ is the number of electrons involved in the emission."
" d, is a function of p. for p=2.5. d,z0.6 Wijers Galama 1999)."," $\Phi_{\rm p}$ is a function of $p$ , for $p=2.5$ , $\Phi_{\rm p}
\approx 0.6$ (Wijers Galama 1999)."
 For of-|Syοι al.," For $\nu_{\rm c, \oplus}<\nu_{\rm a,\oplus}<\nu_{\rm x}<\nu_{\rm m, \oplus} $, the predicted flux is (e.g., Sari et al."
 1998) = 25 , 1998) = 2.5 .
"‘Taking GQ.= land ες=2.4210 Lz. with equation (13)) we have F5,2.5 mJv. which matches the observation of GRB OII121 (~1 mJy)."," Taking $Q_{\rm y}=1$ and $\nu_{\rm x}=2.42\times 10^{17}$ Hz, with equation \ref{Eq:Flux}) ) we have $F_{\nu_{\rm x}}
\approx 2.5$ mJy, which matches the observation of GRB 011121 $\sim 1$ mJy)."
" Phe V band lux can be estimated as (4,4s4A<Ha] ) D.1~PiaVv4|E1/2my240 mv."," The V band flux can be estimated as $\nu_{\rm c, \oplus}<\nu_{\rm v}<\nu_{\rm a,\oplus}$ ) $F_{\nu_{\rm v}}\sim
F_{\nu_{\rm max}} \nu_{\rm a,\oplus}^{-3}\nu_{\rm c, \oplus}^{1/2}\nu_{\rm v}^{5/2}\sim 40$ mJy."
 What happens after the “late internal shocks”?," What happens after the “late internal shocks""?"
 Surely. the refreshed: relativistic outflow. will catch. up with the initial outflow when the latter has swept a large amount of material and. σοι decelerated.," Surely, the refreshed relativistic outflow will catch up with the initial outflow when the latter has swept a large amount of material and got decelerated."
 That energy. injection would eive rise to a llattening (e.g. Rees Alésszárros 1998) or re-brightening signature (e.g. Panaitescu. Mésszárros Rees 1998: Kumar Piran 2000: Zhang Mésszárros 2002). which could potentially account for the late re-brightening of NRE 050406. the late X-ray re-bursting detected in GRB 050502b and the second/weaker X-ray bump observed in GRB 011121.," That energy injection would give rise to a flattening (e.g., Rees Mésszárros 1998) or re-brightening signature (e.g., Panaitescu, Mésszárros Rees 1998; Kumar Piran 2000; Zhang Mésszárros 2002), which could potentially account for the late re-brightening of XRF 050406, the late X-ray re-bursting detected in GRB 050502b and the second/weaker X-ray bump observed in GRB 011121."
 Llowever. the detailed lighteurve modeling is bevond the scope of this Letter.," However, the detailed lightcurve modeling is beyond the scope of this Letter."
 For observers [frequencies above the cooling frequency. the curvature effect. dominates the temporal. behavior of the observed. flux after. the central engine. shuts down ab £540 (PISQO).," For observer's frequencies above the cooling frequency, the curvature effect dominates the temporal behavior of the observed flux after the central engine shuts down at $t_{\rm e, \oplus}$ (PK00)."
 The flux declines as MVἐν)δὲ)NES) (PNOO). where ὁ represents the th pulse. £4. and £4; are the ejection time and the variability timescale of the th. pulse. respectively.," The flux declines as ${\sum\limits_{i}} F_{\nu_{\rm x},i}
[(t_\oplus-t_{{\rm eje},i})/\delta t_{\oplus, i}]^{\rm
-(2+p_i/2)}$ (PK00), where $i$ represents the $i$ th pulse, $t_{{\rm eje},i}$ and $t_{\oplus,i}$ are the ejection time and the variability timescale of the $i$ th pulse, respectively."
 Such a decline is much steeper han (£4fly)1777 for fySmax{olyi}.," Such a decline is much steeper than $(t_\oplus/t_{\rm
e,\oplus})^{-(2+p/2)}$ for $t_{\rm e,\oplus}\gg \max \{\delta
t_{\oplus, i}\}$."
 For example. as shown in Fig. 2..," For example, as shown in Fig. \ref{fig:Illustration},"
" the decline of the [are is dominated by the curvature effect of the last long pulse with 0/4,~0.2445.4."," the decline of the flare is dominated by the curvature effect of the last long pulse with $\delta t_\oplus \sim 0.24 t_{\rm e, \oplus}$."
 A crude power-law fit tothe declinevields fh.x(0ifay) 7s which is steep enough to match the sharpestdecline detected so far.," A crude power-law fit tothe declineyields $F_{\nu_{\rm X}}\propto
(t_\oplus/t_{\rm e,\oplus})^{-8}$ , which is steep enough to match the sharpestdecline detected so far."
 In reality. thecentral enginedoes not turn olf abruptly.," In reality, thecentral enginedoes not turn off abruptly."
 The dimmoer and dimmer emission. poweredby the weaker and weaker 7Iate internal shocks” may dominate over the curvature ellect of the carly pulses. resulting in a shallower decay.," The dimmer and dimmer emission poweredby the weaker and weaker “late internal shocks"" may dominate over the curvature effect of the early pulses, resulting in a shallower decay."
represent those of the entire galaxy.,represent those of the entire galaxy.
" Thus, we do not use the data at mid-infrared and shorter wavelengths, where very small grains and PAHs dominate dust emission (e.g.Draine&Li2001;Dopitaetal.2006) and the spectral energy distribution (SED) is sensitive to grain sizedistribution."," Thus, we do not use the data at mid-infrared and shorter wavelengths, where very small grains and PAHs dominate dust emission \citep[e.g.][]{draine01,dopita06}
 and the spectral energy distribution (SED) is sensitive to grain sizedistribution."
 Figure 1 shows the obtained image., Figure \ref{fig:image} shows the obtained image.
 The source is clearly detected with a peak intensity of 7.4 o (1 o=1.15 mJy beam!)and total flux of 13.6+2.0 mJy., The source is clearly detected with a peak intensity of 7.4 $\sigma$ (1 $\sigma =1.15$ mJy $^{-1}$ ) a total flux of $13.6\pm 2.0$ mJy.
 There is an extended feature in the aeast-west direction with FWHM = 7.1 arcsec (the width of the 2 c contour is ~10 arcsec)., There is an extended feature in the east–west direction with FWHM = 7.1 arcsec (the width of the 2 $\sigma$ contour is $\sim 10$ arcsec).
 Submm emission in galaxies is usually dominated by dust thermal radiation and contaminated by free—free emission., Submm emission in galaxies is usually dominated by dust thermal radiation and contaminated by free–free emission.
" In order to estimate the contribution from free-free emission, we use the VLA 2 cm continuum data, which is sensitive to spatial scales up to 4 arcsec (Becketal.2002)."," In order to estimate the contribution from free–free emission, we use the VLA 2 cm continuum data, which is sensitive to spatial scales up to 4 arcsec \citep{beck02}."
". We adopt the emission on this scale as the contribution from the central star-forming region At this wavelength, the single-dish flux is 18.5 mJy, while the VLA flux is 14+1.5 mJy2002); thus, 76 per cent of the total radiation comes from the central star-forming region at 2 cm."," We adopt the emission on this scale as the contribution from the central star-forming region At this wavelength, the single-dish flux is 18.5 mJy, while the VLA flux is $14\pm 1.5$ mJy; thus, 76 per cent of the total radiation comes from the central star-forming region at 2 cm."
" Since the emission at 2 cm (14+1.5 mJy) is dominated by free-free emission (Beck 2002), we extrapolate the flux to 880 um by assuming a frequency dependence of «v?! (Osterbrock1989) and estimate the contribution of free-free emission to the 880 um flux to be 10x1 mjy."," Since the emission at 2 cm $14\pm 1.5$ mJy) is dominated by free–free emission \citep{beck02}, we extrapolate the flux to 880 $\micron$ by assuming a frequency dependence of $\propto\nu^{-0.1}$ \citep{osterbrock89}
 and estimate the contribution of free–free emission to the 880 $\micron$ flux to be $10\pm 1$ mJy."
" Thus, 7A*1 per cent of the flux detected at 880 uum by SMA is free-free emission from the central star-forming region."," Thus, $74^{+12}_{-10}$ per cent of the flux detected at 880 $\micron$ by SMA is free-free emission from the central star-forming region."
 The contribution from free—free emission is further investigated by theoretical models in Section ??.., The contribution from free–free emission is further investigated by theoretical models in Section \ref{sec:model}.
" The residual (4+2.5 mJy) is possibly dust emission or diffuse (more extended than 4"" but less than ~ 10"") free—free emission.", The residual $4\pm 2.5$ mJy) is possibly dust emission or diffuse (more extended than $4''$ but less than $\sim 10''$ ) free–free emission.
" In order to the radio continuum radiation from a star-forming region, thermal free—free radiation from regions, non-thermal synchrotron emission from supernova remnants (SNRs), and dust thermal emission should generally be taken into account."," In order to the radio continuum radiation from a star-forming region, thermal free–free radiation from regions, non-thermal synchrotron emission from supernova remnants (SNRs), and dust thermal emission should generally be taken into account."
" Here, we adopt theoretical models for thermal free—free and non-thermal synchrotron components from Hirashita&Hunt(2006) to interpret the emission from the central star- region in II Zw 40."," Here, we adopt theoretical models for thermal free–free and non-thermal synchrotron components from \citet{hirashita06} to interpret the emission from the central star-forming region in II Zw 40."
 Below we briefly summarize their models., Below we briefly summarize their models.
 We adopt the same values for the physical parameters as those in Hirashita&Hunt(2006) unless otherwise stated., We adopt the same values for the physical parameters as those in \citet{hirashita06} unless otherwise stated.
 We also add a simple calculation for dust emission in Section ??.., We also add a simple calculation for dust emission in Section \ref{sec:firsed}.
" The star formation is assumed to occur in a uniform star-forming region, whose initial number density of hydrogen nuclei is nyo."," The star formation is assumed to occur in a uniform star-forming region, whose initial number density of hydrogen nuclei is $n_\mathrm{H0}$."
" The mass finally converted into stars, Mo, is also given."," The mass finally converted into stars, $M_0$, is also given."
" The star formation rate (SFR) as a function of time, w(t), is written as w(t)=(esrMo/ta)e-*SF!/'tt, where t is the time measured from the onset of the star formation (thus, we assume y=0 for t<0), and £g is the free-fall time-scale estimated by nyo tthe SFR is regulated by the dynamical time-scale; citealtelmegreen00))."," The star formation rate (SFR) as a function of time, $\psi (t)$, is written as $\psi (t)=(\epsilon_\mathrm{SF}M_0/{\tff})
\, e^{-\epsilon_\mathrm{SF}t/\tff}$, where $t$ is the time measured from the onset of the star formation (thus, we assume $\psi =0$ for $t<0$ ), and $\tff$ is the free-fall time-scale estimated by $n_\mathrm{H0}$ the SFR is regulated by the dynamical time-scale; \\citealt{elmegreen00}) )."
" We simply introduce an efficiency factor esr, which means that the gas is converted into stars on a time-scale of tg/esr=1.4x109(esp/0.1) (nno/10?cm?) yr."," We simply introduce an efficiency factor $\epsilon_\mathrm{SF}$, which means that the gas is converted into stars on a time-scale of $\tff /\epsilon_\mathrm{SF}=1.4\times 10^6
(\epsilon_\mathrm{SF}/0.1)^{-1}
(n_\mathrm{H0}/10^5~\mathrm{cm}^{-3})^{-1/2}$ yr."
 We treat ngo and Mo as independent parameters., We treat $n_\mathrm{H0}$ and $M_0$ as independent parameters.
" Throughout3/2. this paper, we assume a Salpeter initial mass function (IMF) with a stellar mass range of 0.1-100 Mo."," Throughout this paper, we assume a Salpeter initial mass function (IMF) with a stellar mass range of 0.1–100 $\mathrm{M}_{\sun}$."
 Thermal free-free radiation is modeled by considering the evolution of the size and density of rregion., Thermal free–free radiation is modeled by considering the evolution of the size and density of region.
" In this paper, the recombination coefficient excluding captures to the ground level is assumed to be a?)=1.73x (e.g.Spitzer 1978).."," In this paper, the recombination coefficient excluding captures to the ground level is assumed to be $\alpha^{(2)}=1.73\times 10^{-13}~\mathrm{cm^3~s^{-1}}$ \citep[e.g.][]{spitzer78}. ."
" We assume the density in the rregion to be uniform, and neglect the effects of dust in the"," We assume the density in the region to be uniform, and neglect the effects of dust in the"
of complex eight-bit numbers and having a sample time interval equal to 6.4 jtsec were recorded on the subband with central frequency fy=205.781 MHz. bandwidth Af=156.25KHz.,"of complex eight-bit numbers and having a sample time interval equal to 6.4 $\mu sec$ were recorded on the subband with central frequency $f_{0}=205.781MHz$ , bandwidth $\Delta f=156.25KHz$."
 This time-frequency presentation corresponds to 32 frequency channels. 1.e.. the frequency resolution is Of.=4.88KHz. and the integration time ¢=1.7.," This time-frequency presentation corresponds to 32 frequency channels, i.e., the frequency resolution is $\delta f\approx 4.88KHz$, and the integration time $t\approx 1.7s$."
 Fig., Fig.
 11 demonstrates the auto-spectrum calculated from the same data for 1024 frequency channels., 11 demonstrates the auto-spectrum calculated from the same data for 1024 frequency channels.
 The high spectral resolution. 0f.=153Hz. permits the separation of RFI spikes. Le.. not smoothing them. as would have been in the case of 32 channels. see Fig.," The high spectral resolution, $\delta f\approx 153Hz$, permits the separation of RFI spikes, i.e., not smoothing them, as would have been in the case of 32 channels, see Fig."
 10., 10.
 Only half of the 3-hour duration auto-spectrum is shown in Fig., Only half of the 3-hour duration auto-spectrum is shown in Fig.
 11 (due to the constraints of computer memory)., 11 (due to the constraints of computer memory).
 The auto-spectrum in Fig., The auto-spectrum in Fig.
 11 is caleulated without any censoring of outliers., 11 is calculated without any censoring of outliers.
 Fig., Fig.
" 12 shows the ""cleaned"" auto-spectrum where Winsorization was applied.", 12 shows the “cleaned” auto-spectrum where Winsorization was applied.
 The Winsorized spectrum in Fig., The Winsorized spectrum in Fig.
 12 is smoothed so as to get a final frequency resolutionofx488KH- ης in Fig., 12 is smoothed so as to get a final frequency resolution$\delta f\approx 4.88KHz$ as in Fig.
Bailey type-c variable.,Bailey type-c variable.
 This is confirmed. also by the WOT photometry.V74:, This is confirmed also by the K97 photometry.:
 our best. period is slightly longer than S65. in agreement with the mildly positive parabolic shape of the O-C curve.," our best period is slightly longer than S65, in agreement with the mildly positive parabolic shape of the O-C curve."
 The exact times of minimum and maximum light have been missed. in our observations. but the rest. of the curve is very well defined with Little scatter.," The exact times of minimum and maximum light have been missed in our observations, but the rest of the curve is very well defined with little scatter."
 Therefore the period is quite accurate. but the average magnitudes. and amplitudes are not.το:," Therefore the period is quite accurate, but the average magnitudes and amplitudes are not.:"
 this star is among those reanalvsed by NCSO9., this star is among those reanalysed by NC89.
 We find that the period. derived by NCSO fits well also to our observations., We find that the period derived by NC89 fits well also to our observations.
ealaxies of similar luminosities and stellar masses.,galaxies of similar luminosities and stellar masses.
 The olfset from the gas phase metallicity-luminosity for non-interacting galaxies appears to depend on the projected separation between the interacting galaxies and the burst strength (Ixewlevctal..., The offset from the gas phase metallicity-luminosity for non-interacting galaxies appears to depend on the projected separation between the interacting galaxies and the burst strength \citep{kewley06}.
2006).. Peeplesctal.(2009). have identified. a sample of metal-poor high-mass outliers from 1¢ local gas-phase mass-metallicity relation., \citet{peeples09} have identified a sample of metal-poor high-mass outliers from the local gas-phase mass-metallicity relation.
 Phese galaxies we found to exhibit. disturbed: morphologies. bluc colours. and high star formation rates for their masses. implving that rev are undergoing tidal interactions.," These galaxies are found to exhibit disturbed morphologies, blue colours, and high star formation rates for their masses, implying that they are undergoing tidal interactions."
 Similarly. Rupkeetal.(2008) have found that the interstellar gas of local uminous infrared galaxies is underabuncant.," Similarly, \citet{rupke08} have found that the interstellar gas of local luminous infrared galaxies is underabundant."
 Phe ollsct in abundance from the gas phase oxvecn abundance vs. uminositv/mass relation is found to correlate with the total infrared luminosity. with a larger olfset for ultra-Iuminous infrared. galaxies than for luminous infrared. galaxies.," The offset in abundance from the gas phase oxygen abundance vs. luminosity/mass relation is found to correlate with the total infrared luminosity, with a larger offset for ultra-luminous infrared galaxies than for luminous infrared galaxies."
 The olfsets from the mass vs. gas phase metallicity relation in optically selected: mergers is lower than for luminous and ultra-Iuminous infrared. galaxies. suggesting that the olfset from the properties of non-interacting galaxies depends on the strength and the stage of the interaction.," The offsets from the mass vs. gas phase metallicity relation in optically selected mergers is lower than for luminous and ultra-luminous infrared galaxies, suggesting that the offset from the properties of non-interacting galaxies depends on the strength and the stage of the interaction."
 The modest olfset from the stellar mass-metallicitv relation we observe for disturbed: galaxies in the blue cloud. sequence of Abell 1367 could reflect a continuation of the trends observed. for equal-mass mergers to lower strength of galaxy interaction., The modest offset from the stellar mass-metallicity relation we observe for disturbed galaxies in the blue cloud sequence of Abell 1367 could reflect a continuation of the trends observed for equal-mass mergers to lower strength of galaxy interaction.
 Measuring gas phase metallicities of star-forming galaxies in our sample to compare with those of both morphologically. unelisturbed star-forming galaxies. and stronely cisturbed galaxies would test this scenario.," Measuring gas phase metallicities of star-forming galaxies in our sample to compare with those of both morphologically undisturbed star-forming galaxies, and strongly disturbed galaxies would test this scenario."
 , 
"the coefficients 0.345, 0.107, and 0.064 should be 0.7, 0.19, and 0.11, correspondently.","the coefficients 0.345, 0.107, and 0.064 should be 0.7, 0.19, and 0.11, correspondently."
" Exactly as in Paper I, the birthrate of NSs is proportional to the square of the local ISM density."," Exactly as in Paper I, the birthrate of NSs is proportional to the square of the local ISM density."
 With this model we address two main questions., With this model we address two main questions.
 How does the account for the subsonic Propeller stage influence the fraction of Accretors?, How does the account for the subsonic Propeller stage influence the fraction of Accretors?
 How does the fraction of accretors depend on the magnetic field?, How does the fraction of accretors depend on the magnetic field?
" In Paper I the subsonic Propeller stage was not used, and only Crab-like fields, ~10’? G, were considered."," In Paper I the subsonic Propeller stage was not used, and only Crab-like fields, $\sim 10^{12}$ G, were considered."
 Here we include this stage and take into account higher magnetic fields (up to the values typical for decayed fields of magnetars)., Here we include this stage and take into account higher magnetic fields (up to the values typical for decayed fields of magnetars).
" Obviously, the first effect reduces the number of Accretors, while the second — increases."," Obviously, the first effect reduces the number of Accretors, while the second -- increases."
 It is interesting to understand with a simple model the interplay between them before addressing the same questions with a more advanced one., It is interesting to understand with a simple model the interplay between them before addressing the same questions with a more advanced one.
 In this model we use only the function magnetic field distribution., In this model we use only the delta-function magnetic field distribution.
 In Fig., In Fig.
 1 we show the dependence of the fraction of INSs at the Accretor stage on the magnetic field and the ISM number density., 1 we show the dependence of the fraction of INSs at the Accretor stage on the magnetic field and the ISM number density.
 Here both of these parameters are constant during a NS lifetime., Here both of these parameters are constant during a NS lifetime.
 For Crab-like fields and the ISM density ~0.1 — 1 cm? it is below ~1%.., For Crab-like fields and the ISM density $\sim 0.1$ – 1 $^{-3}$ it is below $\sim$.
" The number of Accretors for a given realistic ISM density steadily grows with increasing initial magnetic fields up to ~5x10!? G. This is related to the fact that a NS with larger magnetic field spin-downs faster, and so quicker reaches the stage of accretion."," The number of Accretors for a given realistic ISM density steadily grows with increasing initial magnetic fields up to $\sim 5 \times 10^{13}$ G. This is related to the fact that a NS with larger magnetic field spin-downs faster, and so quicker reaches the stage of accretion."
" However, for larger fields it happens at large periods, but the increase of the spin-down rate is more important than the increase of critical periods for transitions."," However, for larger fields it happens at large periods, but the increase of the spin-down rate is more important than the increase of critical periods for transitions."
" For large fields the fraction of Accretors starts to decrease, because more and more sources appear as Georotators due to their huge magnetospheres."," For large fields the fraction of Accretors starts to decrease, because more and more sources appear as Georotators due to their huge magnetospheres."
" However, it is important to repeat, that here we assume constant fields, but it is normally accepted that large fields gradualy decay while a NS is aging."," However, it is important to repeat, that here we assume constant fields, but it is normally accepted that large fields gradualy decay while a NS is aging."
" The difference with Paper I, where for µ= G cm?, nz1 cm? and comparable velocities we obtained ~ several percents of Accretors, is explained by the influence of the subsonic Propeller stage."," The difference with Paper I, where for $\mu\approx 10^{30}$ G $^3$, $n\approx 1$ $^{-3}$ and comparable velocities we obtained $\sim$ several percents of Accretors, is explained by the influence of the subsonic Propeller stage."
 It is visible in the right panel., It is visible in the right panel.
 The fraction of accretors for n=1 cm? and =109? G cm? is about one percent if the subsonic Propeller stage is not taken into account., The fraction of accretors for $n=1$ $^{-3}$ and $\mu = 10^{30}$ G $^3$ is about one percent if the subsonic Propeller stage is not taken into account.
" For the initial field in the M7 range the fraction of Accretors in our new model (with the subsonic Propeller stage included, see the left panel) goes up to ~10 —30%.."," For the initial field in the M7 range the fraction of Accretors in our new model (with the subsonic Propeller stage included, see the left panel) goes up to $\sim 10$ –."
" Clearly, the main conclusion here is that the subsonic Propeller is not a strong barrier for NSs with realistically large fields in contrast with conclusion by?."," Clearly, the main conclusion here is that the subsonic Propeller is not a strong barrier for NSs with realistically large fields in contrast with conclusion by."
". As now we know that the fraction of such objects (SGRs, AXPs, M7, RRATs) is not low — tens of percent, — one can expect that significant fraction of old INS can start to accrete."," As now we know that the fraction of such objects (SGRs, AXPs, M7, RRATs) is not low – tens of percent, – one can expect that significant fraction of old INS can start to accrete."
" Thus, one has to study the distribution of INSs in different evolutionary stages in more details."," Thus, one has to study the distribution of INSs in different evolutionary stages in more details."
 Some results of this subsection are summarized in the Table., Some results of this subsection are summarized in the Table.
 1., 1.
" We demonstrate twelve tracks for different µ,π,υ."," We demonstrate twelve tracks for different $\mu, n, v$."
" NSs following tracks II, III, and IX always stay at the Ejector stage."," NSs following tracks II, III, and IX always stay at the Ejector stage."
" High velocity NSs with fields larger than the Crab-like value, become Georotators after the Propeller stage (tracks VI and XII)."," High velocity NSs with fields larger than the Crab-like value, become Georotators after the Propeller stage (tracks VI and XII)."
 The NS following track VIII never becomes an Accretor as it stays for a long time at the subsonic Propeller stage., The NS following track VIII never becomes an Accretor as it stays for a long time at the subsonic Propeller stage.
" All the rest NSs (tracks I, IV, V, VII, X, XI) finally start to accrete."," All the rest NSs (tracks I, IV, V, VII, X, XI) finally start to accrete."
 In this subsection we present our main results obtained with the complete model for 10° evolutionary tracks., In this subsection we present our main results obtained with the complete model for $10^5$ evolutionary tracks.
 Our main aim is to calculate the distribution in evolutionary stages for two different distributions of the initial magnetic field., Our main aim is to calculate the distribution in evolutionary stages for two different distributions of the initial magnetic field.
100%.,.
. Similar conclusions have already been reached by other groups performing theoretical DTD studies (see Sect., Similar conclusions have already been reached by other groups performing theoretical DTD studies (see Sect.
 4). e.g. Ruiteretal.(2009) who found a rate almost ten times smaller than the SNuM given by Mannueeietal.(2005).," 4), e.g. \citet{ruiter2009} who found a rate almost ten times smaller than the SNuM given by \citet{mannucci2005}."
. A possible explanation is that the constant conversion of the observational results from SNuk into SNuM may introduce large systematic errors. which ts visible from the uncertainties on the conversion factor as well as the difference in its values given by Mannuecietal.(2005) and Totanietal.(2008).," A possible explanation is that the constant conversion of the observational results from SNuK into SNuM may introduce large systematic errors, which is visible from the uncertainties on the conversion factor as well as the difference in its values given by \citet{mannucci2005} and \citet{totani2008}."
. As a result. the remainder of this section will mainly concentrate on a comparison between the morphological shapes of theoretical and observational DTDs. not on their absolute values.," As a result, the remainder of this section will mainly concentrate on a comparison between the morphological shapes of theoretical and observational DTDs, not on their absolute values."
 e We find that nearly all SD events are created through the WD-MS channel deseribed above., $\bullet$ We find that nearly all SD events are created through the WD+MS channel described above.
 e A great majority of DD SNe Ia (about 80%)) are formed through a RLOF phase followed by a CE phase. as opposed to two successive CE phases.," $\bullet$ A great majority of DD SNe Ia (about ) are formed through a RLOF phase followed by a CE phase, as opposed to two successive CE phases."
 This will be further discussed in the next subsection., This will be further discussed in the next subsection.
 In Figure 3. the influence of mass transfer efficiency. ts investigated.," In Figure 3, the influence of mass transfer efficiency is investigated."
 It shows the DD DTD for different values of accreted mass fraction f., It shows the DD DTD for different values of accreted mass fraction $\beta$.
 The figure shows that in the case of totally non-conservative RLOF (8= 0). the DTD drops dramatically after 0.2 Gyr. leaving virtually no events beyond 0.5 Gyr.," The figure shows that in the case of totally non-conservative RLOF $\beta = 0$ ), the DTD drops dramatically after 0.2 Gyr, leaving virtually no events beyond 0.5 Gyr."
 This directly. implies that in the case of a non-conservative first mass transfer phase. there will be no DD SNe la with a sizable delay time.," This directly implies that in the case of a non-conservative first mass transfer phase, there will be no DD SNe Ia with a sizable delay time."
 Since non-conservative RLOF excludes systems going through such a phase as SN la progenitor. the events which do remain in this case (all with short delay time) must be produced through the double CE channel.," Since non-conservative RLOF excludes systems going through such a phase as SN Ia progenitor, the events which do remain in this case (all with short delay time) must be produced through the double CE channel."
 On the other hand. the events which are absent in the case of B.=0. but account for about of all DD SNe la in the case of conservative RLOF. are those created through a RLOF phase followed by a CE evolution.," On the other hand, the events which are absent in the case of $\beta = 0$, but account for about of all DD SNe Ia in the case of conservative RLOF, are those created through a RLOF phase followed by a CE evolution."
 These results can be understood in terms of the two typical evolution channels leading to a DD SN la discussed before., These results can be understood in terms of the two typical evolution channels leading to a DD SN Ia discussed before.
 It was shown that the RLOF channel (one RLOF followed by one CE phase) produces events with a delay time of multiple Gyr. but ts unable to produce any SN Ia if RLOF ts non-conservative.," It was shown that the RLOF channel (one RLOF followed by one CE phase) produces events with a delay time of multiple Gyr, but is unable to produce any SN Ia if RLOF is non-conservative."
 The CE channel (two successive CE phases) on the other hand ts independent of B. but car only produce events during the first few hundreds of millions of years after starburst.," The CE channel (two successive CE phases) on the other hand is independent of $\beta$, but can only produce events during the first few hundreds of millions of years after starburst."
 In order for the model to be consistent in DTD shape with observations. a nearly conservative mass transfer is necessary. which means that 8 needs to be close to one overall.," In order for the model to be consistent in DTD shape with observations, a nearly conservative mass transfer is necessary, which means that $\beta$ needs to be close to one overall."
 Further investigation has pointed out that. on average. B>0.9 is required for the theoretical model to remain marginally consistent in DTD shape with the observations under the current assumptions.," Further investigation has pointed out that, on average, $\beta \ge 0.9$ is required for the theoretical model to remain marginally consistent in DTD shape with the observations under the current assumptions."
 This 1s also visible from Figure 3., This is also visible from Figure 3.
 An important consequence is that taking the formation timescale of a double WD binary equal to the MS lifetime of the least massive component is unjustified., An important consequence is that taking the formation timescale of a double WD binary equal to the MS lifetime of the least massive component is unjustified.
 This practice. commonly adopted in studies which make use of analytical formalisms instead of detailed evolution. is namely based on the unfounded assumption that the accretor is not affected by the mass transfer process. either because it is totally non-conservative or a CE evolution (which ts by definition taken to be too fast for the aceretor to gain any mass).," This practice, commonly adopted in studies which make use of analytical formalisms instead of detailed evolution, is namely based on the unfounded assumption that the accretor is not affected by the mass transfer process, either because it is totally non-conservative or a CE evolution (which is by definition taken to be too fast for the accretor to gain any mass)."
 The description of angular momentum loss does not affect this quasi-conservative model much. since there is only very little mass which leaves the system.," The description of angular momentum loss does not affect this quasi-conservative model much, since there is only very little mass which leaves the system."
 However. when calculations are made with lower values of 6. the DTD is critically affected by the choice for this description.," However, when calculations are made with lower values of $\beta$, the DTD is critically affected by the choice for this description."
 Mass loss through the corotating ring causes a DTD which descends much steeper than with the assumption of specific accretor angular momentum loss., Mass loss through the corotating ring causes a DTD which descends much steeper than with the assumption of specific accretor angular momentum loss.
 This implies that the latter description, This implies that the latter description
Since our svstem is located very close to the barveenter of the LAIC. we can assume that its distance represents very closely the distance to the LAIC.,"Since our system is located very close to the barycenter of the LMC, we can assume that its distance represents very closely the distance to the LMC."
 However. we cannot exclude that the svstem is located slightly in front or behind the LAIC barveenter.," However, we cannot exclude that the system is located slightly in front or behind the LMC barycenter."
 We will investigate such a possible depth: effect once we have analyzed more Iate-tvpe eclipsing binaries for which we are currently acquiring (he necessary data., We will investigate such a possible depth effect once we have analyzed more late-type eclipsing binaries for which we are currently acquiring the necessary data.
 We have stuclied (he double-linecl late-(wpe LAIC eclipsing binary svstem 635812.3 which we discovered in the database of the OGLLE-II Project., We have studied the double-lined late-type LMC eclipsing binary system OGLE-051019.64-685812.3 which we discovered in the database of the OGLE-II Project.
 From a detailed analvsis of its I-band lisht curve. the radial velocity curves for the (wo giant components of the svstem. and near-intrarecl photometry outside the eclipses we have derived the distance ol the svstem [rom the V.—A surface brightness-color relation for giant stars given by Di Benedetto (2005).," From a detailed analysis of its I-band light curve, the radial velocity curves for the two giant components of the system, and near-infrared photometry outside the eclipses we have derived the distance of the system from the $V-K$ surface brightness-color relation for giant stars given by Di Benedetto (2005)."
 We obtain a true distance modulus of 18.50 mag. with a total estimated uncertaintv of3%.. or 0.06 mag.," We obtain a true distance modulus of 18.50 mag, with a total estimated uncertainty of, or 0.06 mag."
 This result constitutes a significant. improvement on the distance measurements using observations of earlv-tvpe eclipsing binaries in (he LAIC (e.g. Guinan et al., This result constitutes a significant improvement on the distance measurements using observations of early-type eclipsing binaries in the LMC (e.g. Guinan et al.
 1993. Ribas et al.," 1998, Ribas et al."
 2002)., 2002).
 Our derived distance for OGLE-051019.64-685812.3 agrees very. well with most recent determinations of the LMC distance from different methods (e.g. Schaeler 2008. Beneclict et al.," Our derived distance for OGLE-051019.64-685812.3 agrees very well with most recent determinations of the LMC distance from different methods (e.g. Schaefer 2008, Benedict et al."
 2007. Fouqué et al.," 2007, Fouqué et al."
 2007. Guinan οἱ al.," 2007, Guinan et al."
 2004)., 2004).
 This is the first of a number of similar eclipsing binary systems we have discovered in the LMC. whose study in forthcoming papers will allow us to reduce the uncertainty on the curent distance determination to the LAIC from OGLE-051019.64-G85812.3.," This is the first of a number of similar eclipsing binary systems we have discovered in the LMC, whose study in forthcoming papers will allow us to reduce the uncertainty on the current distance determination to the LMC from OGLE-051019.64-685812.3."
 The long orbital periods ancl faint magnitudes of these svstenis make the collection of the required photometric and spectroscopic data cifficult. but our present results demonstrate the high potential offered by Iate-tvpe eclipsing binaries to achieve a breakthrough in the reduction ol the uncertainty of the distance to the LAIC. and to an understanding of the svstematics affecting other methods of distance determination.," The long orbital periods and faint magnitudes of these systems make the collection of the required photometric and spectroscopic data difficult, but our present results demonstrate the high potential offered by late-type eclipsing binaries to achieve a breakthrough in the reduction of the uncertainty of the distance to the LMC, and to an understanding of the systematics affecting other methods of distance determination."
 WG. GP and DM gratefully. acknowledge financial support for this work [rom the Chilean Center lor Astrophysics FONDAP 15010003. and from the BASAL Centro de Astrolisica. v Tecnologias Afines (CATA) PFB-06/2007.," WG, GP and DM gratefully acknowledge financial support for this work from the Chilean Center for Astrophysics FONDAP 15010003, and from the BASAL Centro de Astrofisica y Tecnologias Afines (CATA) PFB-06/2007."
 Support from the Polish grants N203 002 31/046. and N20303032/4215. and the FOCUS subsidy of the Fundation for Polish Science (FNP) is also acknowledged.," Support from the Polish grants N203 002 31/046, and N20303032/4275, and the FOCUS subsidy of the Fundation for Polish Science (FNP) is also acknowledged."
 IBT acknowledges the support of NSF grant. AST-0507325., IBT acknowledges the support of NSF grant AST-0507325.
 It is a pleasure to thank Willie Torres for sharing software wilh us. and to thank the support astronomers al ESO-La Silla and at Las Campanas Observatory for (heir expert help in," It is a pleasure to thank Willie Torres for sharing software with us, and to thank the support astronomers at ESO-La Silla and at Las Campanas Observatory for their expert help in"
"Inspired by the differences between BH+X and BHX in their overall star-formation histories we revisit the topic of and examine the (Figurelow-mass6)), slope of the galaxy mass spectrum, as parametrized by the a of(1976),, which we will call ag to distinguish it from the Bondi accretion parameter ag.","Inspired by the differences between BH+X and BHX in their overall star-formation histories (Figure \ref{fig:sfr-big}) ), we revisit the topic of and examine the low-mass slope of the galaxy mass spectrum, as parametrized by the $\alpha$ of, which we will call $\alpha_S$ to distinguish it from the Bondi accretion parameter $\alpha_B$."
" As in ag is defined for our purposes by n(M)dMος(M/M,)9*5*e-M/M-qM."," As in $\alpha_S$ is defined for our purposes by $n(M)dM\propto (M/M_*)^{\alpha_S}
e^{-M/M_*}dM$ ."
" The results are given in Table 2:: while the No BH and BH models agree within error with the values corresponding to their background from (and BH+X has a modestly steeper slope likely related to the suppression of cooling flows), the BHX and BHXRP models indeed have an extreme effect on ας, flattening it well beyond the observed value of ~—1 to —0.44 and —0.55, respectively."," The results are given in Table \ref{tab3:sch}: while the No BH and BH models agree within error with the values corresponding to their background from (and BH+X has a modestly steeper slope likely related to the suppression of cooling flows), the BHX and BHXRP models indeed have an extreme effect on $\alpha_S$ , flattening it well beyond the observed value of $\sim\! -1$ to $-0.44$ and $-0.55$, respectively."
" In we found that o increases as any heating is added: from —1.6 with no heating, to —1.3 with a UV background, to —1.0 with UV and SN feedback, to —0.75 with UV, feedback and an X-ray background."," In we found that $\alpha$ increases as any heating is added: from $-1.6$ with no heating, to $-1.3$ with a UV background, to $-1.0$ with UV and SN feedback, to $-0.75$ with UV, feedback and an X-ray background."
" As mentioned above, the timing of the X-rays seems to explain the difference between BH+X and BHX, since the BHX X-rays peak just before the epoch of primary star formation in low-mass galaxies (found in to be 3Zz=2), while the X-ray background peaks toward the end of it, and doesn't significantly suppress star formation until z1.5."," As mentioned above, the timing of the X-rays seems to explain the difference between BH+X and BHX, since the BHX X-rays peak just before the epoch of primary star formation in low-mass galaxies (found in to be $3\gtrsim z \gtrsim 2$ ), while the X-ray background peaks toward the end of it, and doesn't significantly suppress star formation until $z\approx 1.5$."
" 'This result suggests that our X-ray feedback might be too strong: in fact, our formula converting X-ray flux to heating rate has a term (representing the photoionization heating and line and recombination cooling, Eq."," This result suggests that our X-ray feedback might be too strong: in fact, our formula converting X-ray flux to heating rate has a term (representing the photoionization heating and line and recombination cooling, Eq."
" A35 in 2005)) which is linear in Z/Zo—i.e. the metallicity as a fraction of solar—where solar metallicity is assumed (Sergey Sazonov, private "," A35 in ) which is linear in $Z/Z_{\odot}$ —i.e. the metallicity as a fraction of solar---where solar metallicity is assumed (Sergey Sazonov, private communication)."
"Thus, since the metallicity of the ISM communication).could be 0.179 at early times, our heating rate could be too high by that factor."," Thus, since the metallicity of the ISM could be $0.1 Z_{\odot}$ at early times, our heating rate could be too high by that factor."
" Moreover, the work of suggests that after reionization the equilibrium temperature of the IGM is roughly independent of the intensity of the radiation field, depending only on its spectrum, so our r~? attenuation factor may be less significant than we would vely think."," Moreover, the work of suggests that after reionization the equilibrium temperature of the IGM is roughly independent of the intensity of the radiation field, depending only on its spectrum, so our $r^{-2}$ attenuation factor may be less significant than we would vely think."
 The results for fMit are more modest., The results for $f M_{\text{crit}}$ are more modest.
" The quantity fMerit is defined inII:: in short, it represents the largest halo mass at which halos have an average star:DM mass ratio of less than half the global value (0.08)."," The quantity $f
M_{\text{crit}}$ is defined in: in short, it represents the largest halo mass at which halos have an average star:DM mass ratio of less than half the global value $0.08$ )."
 (Merit is the theoretically-calculated virial mass whose escape velocity is equal to the sound speed of its gas at the epoch when it should be forming stars; the effective correction factor fz 0.75.)," $
M_{\text{crit}}$ is the theoretically-calculated virial mass whose escape velocity is equal to the sound speed of its gas at the epoch when it should be forming stars; the effective correction factor $f\approx 0.75$ .)"
" Here again, No BH, BH, and BH+X agree well with the values that gives from their ionizing background models."," Here again, No BH, BH, and BH+X agree well with the values that gives from their ionizing background models."
" The X-ray background models have somewhat higher values, as we would expect from their flatter low-mass slopes."," The X-ray background models have somewhat higher values, as we would expect from their flatter low-mass slopes."
 The effects of X-ray feedback from AGN are manifold., The effects of X-ray feedback from AGN are manifold.
" We find that X-ray heating and radiation pressure are only moderately effective at self-regulation: they reduce the black hole's mass far less than increasing the thermal feedback efficiency does, primarily by suppressing bursts of Eddington-limited accretion at early times."," We find that X-ray heating and radiation pressure are only moderately effective at self-regulation: they reduce the black hole's mass far less than increasing the thermal feedback efficiency does, primarily by suppressing bursts of Eddington-limited accretion at early times."
" The model with radiation pressure also accretes significantly less gas at the time of a major merger, instead accreting it more smoothly over the following several Gyr."," The model with radiation pressure also accretes significantly less gas at the time of a major merger, instead accreting it more smoothly over the following several Gyr."
" TheX-ray feedback produces a significant reduction in the host galaxy's baryon- efficiency compared to a traditional feedback model,but only slightly more than a model with"," TheX-ray feedback produces a significant reduction in the host galaxy's baryon-conversion efficiency compared to a traditional feedback model,but only slightly more than a model with"
In our previous paper. we studied dynamical. bar-moce stabilities of clillerentially rotating stars in. Newtonian eravity (Shibataetal.2002).,"In our previous paper, we studied dynamical bar-mode stabilities of differentially rotating stars in Newtonian gravity \cite{SKE}."
. In. that. studs: we adopted a polvtropic equation of state with the polvtropic: index n—|] and the so-called. 7j-constant-like angular. velocity profile in which the magnitude of the angular velocity decreases as ce7 at laree values of ze. where ze denotes the cvlindrical radius.," In that study, we adopted a polytropic equation of state with the polytropic index $n=1$ and the so-called $j$ -constant-like” angular velocity profile in which the magnitude of the angular velocity decreases as $\varpi^{-2}$ at large values of $\varpi$, where $\varpi$ denotes the cylindrical radius."
 We found that rotating stars of a hieh degree of dillerential rotation are dynamically unstable even or 3=[LAW|~0.03. where T. and Wo are rotational and eravitational potential energies.," We found that rotating stars of a high degree of differential rotation are dynamically unstable even for $\beta \equiv |T/W| \sim 0.03$, where $T$ and $W$ are rotational and gravitational potential energies."
 This value is much smaller han the long-believed. criterion of ο)=0.27 lor onset. οἱ he bar-mode dvnamical instability of rotating stars. (see Shibata et al., This value is much smaller than the long-believed criterion of $\beta \approx 0.27$ for onset of the bar-mode dynamical instability of rotating stars (see Shibata et al.
 2002 for a review)., 2002 for a review).
 We also found that after he instability. sets in. such unstable rotating stars with 1.03z;2£z:0.15 eventually settle down to nonaxisymmetric ellipsoidal quasi-stationary states.," We also found that after the instability sets in, such unstable rotating stars with $0.03 \alt \beta \alt 0.15$ eventually settle down to nonaxisymmetric ellipsoidal quasi-stationary states."
 However. there are two questions which have not been answered in the previous work (Shibataetal.2002).," However, there are two questions which have not been answered in the previous work \cite{SKE}."
. One is associated. with our choice of the j-constant-like angular velocity. profile., One is associated with our choice of the $j$ -constant-like angular velocity profile.
 It is. well known that accretion isks around a central body with constant specific angular momenttun are unstable against the Papaloizou-Pringle instability (Papaloizou&Pringle1984)., It is well known that accretion disks around a central body with constant specific angular momentum are unstable against the Papaloizou-Pringle instability \cite{PP}.
. On the other hand. rc aceretion disks are stable if the velocity profile is Ixepler-ike.," On the other hand, the accretion disks are stable if the velocity profile is Kepler-like."
 One could claim that the bar-mocde instability which we found would not be universal ancl might set in only [or 16 very special rotational profile such as the j-constant-like aw as in the Papaloizou-Prinele instability., One could claim that the bar-mode instability which we found would not be universal and might set in only for the very special rotational profile such as the $j$ -constant-like law as in the Papaloizou-Pringle instability.
 In addition. we ocused only on a still equation of state with »=1 in the oevious paper. so that one could also ask if the instability sets in for softer equations of state.," In addition, we focused only on a stiff equation of state with $n=1$ in the previous paper, so that one could also ask if the instability sets in for softer equations of state."
 To answer these questions. we have performed numerical simulations of differentially rotating stars varving the polvtropic index and the angular velocity. profile.," To answer these questions, we have performed numerical simulations of differentially rotating stars varying the polytropic index and the angular velocity profile."
 In this paper. we report the numerical results.," In this paper, we report the numerical results."
 We will show that irrespective of the polvtropic index m and the angular velocity profile. a rotating star ofa high degree of dillerential rotation is dynamically unstable against the bar-moce deformation even for the case that 3 is of order 0.01.," We will show that irrespective of the polytropic index $n$ and the angular velocity profile, a rotating star of a high degree of differential rotation is dynamically unstable against the bar-mode deformation even for the case that $\beta$ is of order 0.01."
"f=0.25, and chemical abundances typical of LBVs, with a He mass fraction of Y=0.62, N/No=7, C/Co=0.08, O/Og= 0.025, and solar abundance for the iron group elements.","$f=0.25$, and chemical abundances typical of LBVs, with a He mass fraction of Y=0.62, $_\odot$ =7, $_\odot=0.08$, $_\odot=0.025$ , and solar abundance for the iron group elements."
 These parameters are roughly characteristic for CCar., These parameters are roughly characteristic for Car.
 Here we briefly describe our implementation of time-dependent density and velocity structures., Here we briefly describe our implementation of time-dependent density and velocity structures.
" The reader is referred to Grohetal.(2008) and Groh Hillier (2011, in prep) for further details."," The reader is referred to \citet{groh08_clumping} and Groh Hillier (2011, in prep) for further details."
" At the moment, no time dependence is accounted for in the solution of the radiative transfer and rate equations, since the recombination timescale is much shorter than the flow timescale for the wind regions studied here."," At the moment, no time dependence is accounted for in the solution of the radiative transfer and rate equations, since the recombination timescale is much shorter than the flow timescale for the wind regions studied here."
" As the wind hydrodynamics is currently not yet solved for, in order to account for time-dependent outflows, we allow for arbitrary v(r) and p(r) stratifications as inputs intocMFGEN."," As the wind hydrodynamics is currently not yet solved for, in order to account for time-dependent outflows, we allow for arbitrary $v(r)$ and $\rho(r)$ stratifications as inputs into."
" We consider cases when the star crosses the bi-stability limit, evolving between two epochs: the initial epoch 1, when Teg=23000K and a fast wind is present with γουι. and epoch 2, when Tor decreases and a slow wind is present, according to voo»«vos."," We consider cases when the star crosses the bi-stability limit, evolving between two epochs: the initial epoch 1, when $\teff=23\,000~K$ and a fast wind is present with $\vinfi$, and epoch 2, when $\teff$ decreases and a slow wind is present, according to $\vinff < \vinfi$."
" For simplicity, we keep M fixed, so that p(r) is computed following the equation of mass conservation."," For simplicity, we keep $\mdot$ fixed, so that $\rho(r)$ is computed following the equation of mass conservation."
" We explore two scenarios as to whether the double P-Cygni profiles can be caused by the time variability of wind parameters: a gradual change in νοο, or an abrupt instantaneous variability."," We explore two scenarios as to whether the double P-Cygni profiles can be caused by the time variability of wind parameters: a gradual change in $\vinf$, or an abrupt instantaneous variability."
" For models with gradual changes in v,..(t), for simplicity we assume that, starting at t= 0, νο decreases linearly with time."," For models with gradual changes in $\vinf (t)$, for simplicity we assume that, starting at $t=0$ , $\vinf$ decreases linearly with time."
" For a given time Af after t=0, we compute the distance traveled by material ejected between t=0 and At."," For a given time $\Delta t$ after $t=0$, we compute the distance traveled by material ejected between $t=0$ and $\Delta t$."
 The result is a velocity law that increases linearly with the distance up to a point where the initial wind is reached., The result is a velocity law that increases linearly with the distance up to a point where the initial wind is reached.
" From that point onwards, v(r) follows the standard B-type, asymptotically reaching γοοι (Fig."," From that point onwards, $v(r)$ follows the standard $\beta$ -type, asymptotically reaching $\vinfi$ (Fig."
 2aa)., \ref{fig2}a a).
" For models with an abrupt change, we assume that at ¢=0, Vo, SWitches instantaneously from γοοΙ to Voo.2."," For models with an abrupt change, we assume that at $t=0$, $\vinf$ switches instantaneously from $\vinfi$ to $\vinff$."
" To compute a model for a given At, we evaluate the distance traveled by the last particles of the initial wind (rj), as well as the distance traveled by the first particles of the final wind (r2)."," To compute a model for a given $\Delta t$, we evaluate the distance traveled by the last particles of the initial wind $r_1$ ), as well as the distance traveled by the first particles of the final wind $r_2$ )."
" Following a suggestion by S.P. Owocki (2010, priv."," Following a suggestion by S.P. Owocki (2010, priv."
" comm.),"," comm.),"
" we assume that the initial wind characterizes the region from r, up to the outer boundary of the computational domain, whilst a 8-type law characterized by γοοι and βι is assumed."," we assume that the initial wind characterizes the region from $r_1$ up to the outer boundary of the computational domain, whilst a $\beta$ -type law characterized by $\vinfi$ and $\beta_1$ is assumed."
" From the star up to r?, we assume that v(r) is given by the velocity law of the final wind according to γου2 and 62."," From the star up to $r_2$, we assume that $v(r)$ is given by the velocity law of the final wind according to $\vinff$ and $\beta_2$."
" We further assume that the region in between the two winds, i.e. between r» and ri, v(r) and p(r) are linear, and join the two winds smoothly."," We further assume that the region in between the two winds, i.e. between $r_2$ and $r_1$, $v(r)$ and $\rho(r)$ are linear, and join the two winds smoothly."
" We defer to future hydrodynamical computations for a more accurate treatment of this region, but we note that initial tests using somewhat different assumptions may lead to small changes in the resulting P-Cygni line profiles."," We defer to future hydrodynamical computations for a more accurate treatment of this region, but we note that initial tests using somewhat different assumptions may lead to small changes in the resulting P-Cygni line profiles."
 These changes are not anticipated to affect the conclusions reached here., These changes are not anticipated to affect the conclusions reached here.
" For both gradual and abrupt models, we assume that the stellar wind is at t=0 in steady-state with νοι=150kms""! and B,=1 and, as t—co, evolves to a final state with and f»=1."," For both gradual and abrupt models, we assume that the stellar wind is at $t=0$ in steady-state with $\vinfi=150~\kms$ and $\beta_1=1$ and, as $t\to \infty$, evolves to a final state with $\vinff=70~\kms$ and $\beta_2=1$."
" The choice of νο values is based on LBV spectral data, whilst the factor of two change in v,, is adopted from theoretical predictions (Vinketal.1999)."," The choice of $\vinf$ values is based on LBV spectral data, whilst the factor of two change in $\vinf$ is adopted from theoretical predictions \citep{vink99}."
". Figure 2bb presents the P-Cygni line profiles calculated for both a gradual and an abrupt variation of the wind parameters as a function of time, here computed at At=120 days."," Figure \ref{fig2}b b presents the P-Cygni line profiles calculated for both a gradual and an abrupt variation of the wind parameters as a function of time, here computed at $\Delta t=120$ days."
 A gradual change in v.(f) causes a linear variationof the velocity with distance (Fig., A gradual change in $\vinf(t)$ causes a linear variationof the velocity with distance (Fig.
 2aa)., \ref{fig2}a a).
" At At—120 days, the P-Cygni profile deviates only slightly from a steady-state model with the same stellar parameters."," At $\Delta t=120$ days, the P-Cygni profile deviates only slightly from a steady-state model with the same stellar parameters."
" To test the robustness of our results, we computed models for different At, as well as stellar and wind"," To test the robustness of our results, we computed models for different $\Delta t$ , as well as stellar and wind"
between tabulatecl points are obtained via. cubic spline interpolation.,between tabulated points are obtained via cubic spline interpolation.
 We have considered. two cilferent models for the small infalling satellite galaxy., We have considered two different models for the small infalling satellite galaxy.
 One is a rigid Plummer sphere with density distribution given bv and the other is a rigid Llernquist sphere with density clistribution given bv For both of these cases we caleulate the mass and velocity dispersion as for the host galaxy. except with a eric spaced logarithmically between 0.1kpe and LOkpe.," One is a rigid Plummer sphere with density distribution given by and the other is a rigid Hernquist sphere with density distribution given by For both of these cases we calculate the mass and velocity dispersion as for the host galaxy, except with a grid spaced logarithmically between $0.1 \kpc$ and $10 \kpc$."
" The total mass (AJ=10"" M.) and scale lengths (b=O.4kpe and @=0.23 kpc) were chosen to match the estimated mass and half-mass racius (0.55 kpe) of Sagittarius (Llelmi 2000).. as a typical cwarf galaxy."," The total mass $M = 10^9 \ \msun$ ) and scale lengths $b = 0.4\kpc$ and $a = 0.23\kpc$ ) were chosen to match the estimated mass and half-mass radius $0.55 \kpc$ ) of Sagittarius \cite{helmi00}, as a typical dwarf galaxy."
 We will see in Section that the choice of model for the satellite does not make a significant dillerence to our results., We will see in Section \ref{tidalstripping} that the choice of model for the satellite does not make a significant difference to our results.
 ln computing5 the orbit of the infallingg satellite. we neglectο changes which might be induced by the satellite in the rost galaxy.," In computing the orbit of the infalling satellite, we neglect changes which might be induced by the satellite in the host galaxy."
 This approximation has been the subject. of debate in the literature (Zaritsky&WhiteLOSS:Llern-quist&Weinberg 1989).. with the conclusion being reached hat it is reasonable to ignore the galaxy response for low mass satellites (Velázquez&White1999).," This approximation has been the subject of debate in the literature \cite{zaritsky88,hernquist89}, with the conclusion being reached that it is reasonable to ignore the galaxy response for low mass satellites \cite{velazquez99}."
". Physically. this is what is expected. as the response induced in the galaxy is ooportional to the mass ratio of the satellite to the galaxy AZM, and the elect of this response Is again proportional o the mass ratio. so the overall elfect will be proportional o CALMY."," Physically, this is what is expected, as the response induced in the galaxy is proportional to the mass ratio of the satellite to the galaxy ${M_s}/{M_g}$, and the effect of this response is again proportional to the mass ratio, so the overall effect will be proportional to $\left({M_s}/{M_g}\right)^2$."
 We model. the motion of the satellite using Chancrasekhar’s cyvnamical friction formula &Tremaine 1987)., We model the motion of the satellite using Chandrasekhar's dynamical friction formula \cite{binney87}.
. Thus where us is the velocity ofthe satellite. p is the density of the medium. & is given by C ds. Newton's gravitational constant. ἂν the dimensionless speed of the satellite. is given hy and iX is the Coulomb logarithin discussed below.," Thus where $\bmath{v_s}$ is the velocity of the satellite, $\rho$ is the density of the medium, $k$ is given by $G$ is Newton's gravitational constant, $X$, the dimensionless speed of the satellite, is given by and $\Lambda$ is the Coulomb logarithm discussed below."
 Equation (7)) applies to the case of a point mass moving through an infinite homogeneous medium. with a Maxwellian velocity cistribution., Equation \ref{chandra}) ) applies to the case of a point mass moving through an infinite homogeneous medium with a Maxwellian velocity distribution.
 However. numerica experiments (Velázquez&White1999). show that it also provides a reasonable approximation to the eravitationa drag experienced. in more realistic situations. proviclec one chooses the Columb logarithm In.X appropriately.," However, numerical experiments \cite{velazquez99} show that it also provides a reasonable approximation to the gravitational drag experienced in more realistic situations, provided one chooses the Columb logarithm $\ln\Lambda$ appropriately."
 To determine whether it is adequate for the situation we are considering. we use it to reproduce the results from the infal o a satellite into a dark halo when the galaxy was mocelle using a full N-body treatment (vandenBoschetal.1999).," To determine whether it is adequate for the situation we are considering, we use it to reproduce the results from the infall of a satellite into a dark halo when the galaxy was modelled using a full $N$ -body treatment \cite{vandenbosch99}."
. We note that this procedure has been investigated in some etail by Tavlor Babul (2001). who demonstrate that i oes provide a reasonable approximation.," We note that this procedure has been investigated in some detail by Taylor Babul (2001), who demonstrate that it does provide a reasonable approximation."
" ‘Yo perform the integration.o we used the algorithmὃν from 1e fifth-order Cash-Ixarp Itunge-Ixutta integrator 'odeint"" (Pressctal."," To perform the integration, we used the algorithm from the fifth-order Cash-Karp Runge-Kutta integrator ' \cite{numrec}."
1992)... The timestep used. is automatically adjusted by the routine to control errors. and in no case did 10 timestep reach a significant fraction of the dynamical time of the satellite.," The timestep used is automatically adjusted by the routine to control errors, and in no case did the timestep reach a significant fraction of the dynamical time of the satellite."
 Asa test of our procedure. we use our method to mocel van den Bosch et al," As a test of our procedure, we use our method to model van den Bosch et al."
s (1999). Plummer sphere sinking into a truncated: isothermal halo.,'s \shortcite{vandenbosch99} Plummer sphere sinking into a truncated isothermal halo.
 We find that we are able to reproduce their results over a wide range of conditions if we choose InXe2.," We find that we are able to reproduce their results over a wide range of conditions if we choose $\ln
\Lambda\approx2$."
 In Figure 1. we show a comparison between he full N-bocly calculations of van cen Bosch et al (1999) and our approximate procedure for a typical case., In Figure \ref{fig-vandenbosch} we show a comparison between the full $N$ -body calculations of van den Bosch et al (1999) and our approximate procedure for a typical case.
 As can be seen. there is reasonable agreement between the two.," As can be seen, there is reasonable agreement between the two."
 Using hese results we shall adopt the same choice In.X=2 for he orbital computations contained in this paper., Using these results we shall adopt the same choice $\ln\Lambda=2$ for the orbital computations contained in this paper.
 We should note that. for the spherical case. our treatment of dynamical riction also ignores the clleets of possible galaxy rotation ancl velocity. anisotropy. but. unless these are extreme we expect that this approximation will be adequate in. the Current case.," We should note that, for the spherical case, our treatment of dynamical friction also ignores the effects of possible galaxy rotation and velocity anisotropy, but unless these are extreme we expect that this approximation will be adequate in the current case."
"estimate second derivatives of those quantities. so V77"" and Wy"" suller much more than VI and WZ/""*.","estimate second derivatives of those quantities, so $W_{xz}^{visc}$ and $W_{xy}^{visc}$ suffer much more than $W_{xz}^{turb}$ and $W_{xy}^{turb}$."
 Some representative fitted curves [or each series of rins are shown in figures 5.. G and 7..," Some representative fitted curves for each series of runs are shown in figures \ref{fig: XZ_W_fit}, \ref{fig: XY_W_fit} and \ref{fig: Amp_W_fit}."
 We see thal. except for one case. the two curves match closely. which shows that the effective viscosity assumption captures the effects of turbulent dissipation to a good degree.," We see that, except for one case, the two curves match closely, which shows that the effective viscosity assumption captures the effects of turbulent dissipation to a good degree."
 The 7=10 curves for the strong z dependent x forcing case do not match well at all., The $T=10$ curves for the strong $z$ dependent $x$ forcing case do not match well at all.
 The bad fit is due to the [act that if there is any dissipation present it is undetectable from within the turbulent noise., The bad fit is due to the fact that if there is any dissipation present it is undetectable from within the turbulent noise.
 The reason for the stronglv suppressed dissipation is Chat for such short periods the corresponding eddy sizes are too small to be reliably simulated at our current resolution., The reason for the strongly suppressed dissipation is that for such short periods the corresponding eddy sizes are too small to be reliably simulated at our current resolution.
 This fact points to the Goldreich et., This fact points to the Goldreich et.
 al., al.
 picture. that the dissipation is dominated by ecdcdies with periods close to close that of the external forcing (7). as opposed to the Ζ picture in which the largest eddies are always the most important. since in the latter case we would expect the linear scaling to continue to arbitrarily short periods.," picture, that the dissipation is dominated by eddies with periods close to close that of the external forcing $T$ ), as opposed to the Zahn picture in which the largest eddies are always the most important, since in the latter case we would expect the linear scaling to continue to arbitrarily short periods."
" An alternative way to get a value for the scaling1313 constants AT,1212 and AV, is to equate the overall power deposited into the box by the external forcing:", An alternative way to get a value for the scaling constants $K_{1313}^0$ and $K_{1212}^0$ is to equate the overall power deposited into the box by the external forcing:
To the extent that acliabatic compression occurs. but the Blumenthaletal.(1986). model exaegerates its degree (Gnedinetal.2004).. reality is intermediate to these two extremes.,"To the extent that adiabatic compression occurs, but the \citet{blumenthal86} model exaggerates its degree \citep{gnedin04}, reality is intermediate to these two extremes."
" We also fit the mass-to-light ratio as log(M,/L)=loga+bz. where «=(AL,/L)q is the mass-to-lisht ratio al z=0 and 5b=d(log(M,/L))/dz is ils evolution with redshilft."," We also fit the mass-to-light ratio as $\log(M_{\ast}/L)=\log a+bz$, where $a=(M_{\ast}/L)_{0}$ is the mass-to-light ratio at $z=0$ and $b=d(\log(M_{\ast}/L))/dz$ is its evolution with redshift."
 Fig. 13..," Fig. \ref{fig:figure8},"
 shows the likelihood contours for these two parameters for both: compressed ancl unconmpressed models., shows the likelihood contours for these two parameters for both compressed and uncompressed models.
 For the compressed isotropic models. we find α=(7.220.5)M.fL. and b=—0.72z 0.08.," For the compressed isotropic models, we find $a=(7.2 \pm 0.5) M_{\sun}/L_{\sun}$ and $b=-0.72\pm 0.08$ ."
 This agrees with the local value of «=(7.342.1M./L. from Gerhardοἱal.(2001) (that was used by Treuetal.(2006)., This agrees with the local value of $a = (7.3\pm 2.1) M_{\sun}/L_{\sun}$ from \citet{gerhard01} that was used by \citet{tkbbm06}.
. It also agrees with the estimate for the rate of the evolution b=—0.69x0.08., It also agrees with the \citet{tkbbm06} estimate for the rate of the evolution $b=-0.69\pm 0.08$.
 Changing the isotropy over the range ?=—1/3.0.1/3 has little effect. while the model without adiabatic compression requires higher normalizations for the mass-to-lieht ratio (10.0£0.3) ancl slehtly slower rates of evolution.," Changing the isotropy over the range $\beta=-1/3, 0, 1/3$ has little effect, while the model without adiabatic compression requires higher normalizations for the mass-to-light ratio $10.0\pm0.3$ ) and slightly slower rates of evolution."
 Our analvsis includes 2 lenses (PG1115+080 and 1115434535) that were not used by Treuetal.(2006).. but excluding them from (he analysis has little effect on the nmass-bto-light ratios.," Our analysis includes 2 lenses (PG1115+080 and H1543+535) that were not used by \cite{tkbbm06}, but excluding them from the analysis has little effect on the mass-to-light ratios."
" There are no significant changes in (M,/L),; and dlog(M,/L)/dz i we neglect these (wo lenses.", There are no significant changes in $(M_{\ast}/L)_0$ and $d\log(M_{\ast}/L)/dz$ if we neglect these two lenses.
 We reanalvzed the data from the SLACS and LSD surveys of gravitational lenses with velocity dispersion measurements., We reanalyzed the data from the SLACS and LSD surveys of gravitational lenses with velocity dispersion measurements.
 Our mass distribution consists olfa Leruquist model for the luminous galaxy embedded in a theoretically constrained NEW halo model., Our mass distribution consists of a Hernquist model for the luminous galaxy embedded in a theoretically constrained NFW halo model.
" We investigated the homogeneity of the sample. the stellar mass fraction f. the local (2= 0) stellar ratio (AL,/L) and its evolution d(log(M./L))/dz."," We investigated the homogeneity of the sample, the stellar mass fraction $f_*$, the local $z=0$ ) stellar mass-to-light ratio $(M_{\ast}/L)_0$ and its evolution $d(\log(M_{\ast}/L))/dz$."
 As in the earlier study by (2006).. we found that the effects of orbital anisotropy on both (he stellar mass Iraction and the mass-to-light ratio are small.," As in the earlier study by \citet{ktbbm06}, we found that the effects of orbital anisotropy on both the stellar mass fraction and the mass-to-light ratio are small."
 In most cases. a central velocity dispersion measurement provides only weak a constraint on halo structure in the physically interesting region.," In most cases, a central velocity dispersion measurement provides only weak a constraint on halo structure in the physically interesting region."
 Typical limits on the mass fraction represented by the stars have logarithmic errors of order 0.5 dex., Typical limits on the mass fraction represented by the stars have logarithmic errors of order 0.5 dex.
 While (his appears to contradict the conclusions of (for example) ουράςοἱal.(2006).. this is not (he case.," While this appears to contradict the conclusions of (for example) \citet{ktbbm06}, this is not the case."
 Koopmansetal.(2006) fit mass models where por. and find values in the range 1-8«5<2.3.," \citet{ktbbm06} fit mass models where $\rho\varpropto r^{-\gamma}$, and find values in the range $1.8<\gamma<2.3$."
 Fie., Fig.
" 14. shows the expected range of this slope For our models of SDSS where we estimated the slope by fitting the pocr power law model to the projected mass distribution over a radial baseline of 2,./8 to Hp that approximates (he leverage in using stellar dvnamies combined with gravitational lensing to determine halo structure."," \ref{fig:figure9} shows the expected range of this slope for our models of SDSS J0037--0942, where we estimated the slope by fitting the $\rho\varpropto r^{-\gamma}$ power law model to the projected mass distribution over a radial baseline of $R_e/8$ to $R_E$ that approximates the leverage in using stellar dynamics combined with gravitational lensing to determine halo structure."
 For (his twpical lens. the variations in 5 of 1.6<5S2.06 are comparable tothe spread in 5 observed for the SLACS svstems (INoopmansetal. (2006))).," For this typical lens, the variations in $\gamma$ of $1.6 \la \gamma \la 2.06$ are comparable tothe system-to-system spread in $\gamma$ observed for the SLACS systems \citet{ktbbm06}) )."
 Thus. the," Thus, the"
Umeda. IL. Nomoto. IX... Yamaoka. Wanajo. 11999. ApJ. 513. 861.,"Umeda, H., Nomoto, K., Yamaoka, Wanajo, 1999, ApJ, 513, 861."
 van den Bergh. S.. Li. Filippenko. 22005. PASP. 117. 773.," van den Bergh, S., Li, Filippenko, 2005, PASP, 117, 773."
 Wang. L.. et 22003. ApJ. 591. 1110.," Wang, L., et 2003, ApJ, 591, 1110."
 Wang. X.. et 22008a. ApJ. 675. 626.," Wang, X., et 2008a, ApJ, 675, 626."
 Wang. X.. et 22003b. ApJ. in press (arXiv:0ill.25," Wang, X., et 2008b, ApJ, in press (arXiv:0711.2570). ("
,W+08)
scales <3+ Alpe.,scales $\la 3\;h^{-1}\;$ Mpc.
 In this analysis we are binning the data onto a Ingrid with bin size 12h.+ Alpe. and only consideringin] power on scales 2325.+ Alpe. so we expect this to have little ellect on our final results.," In this analysis we are binning the data onto a grid with bin size $\sim 12\;h^{-1}\;$ Mpc, and only considering power on scales $\ga 32\;h^{-1}\;$ Mpc, so we expect this to have little effect on our final results."
 The power spectrum estimation is carried out as described in Outram. Hovle Shanks (2001). using the method outlined in Llovle et al. (," The power spectrum estimation is carried out as described in Outram, Hoyle Shanks (2001), using the method outlined in Hoyle et al. ("
1999) ane Fadros Efstathiou (1996).,1999) and Tadros Efstathiou (1996).
 1n order to apply the distant observer approximation (see section ??)) the mock QSO survey is. divided. into four regions., In order to apply the distant observer approximation (see section \ref{datageom}) ) the mock QSO survey is divided into four regions.
 Each region is embedded: into a larger cubica volume. which is rotated such that the central line of sigh lies along the same axis of the cube in each case.," Each region is embedded into a larger cubical volume, which is rotated such that the central line of sight lies along the same axis of the cube in each case."
 Ehe density field is binned onto a 256 mesh. using nearest erid-poin assignment.," The density field is binned onto a $^3$ mesh, using nearest grid-point assignment."
 Phe power spectrum of each region is estimate using a Fast Fourier Transform (EET). aad the average of the resulting power spectra is taken.," The power spectrum of each region is estimated using a Fast Fourier Transform (FFT), and the average of the resulting power spectra is taken."
 The result. binnec logarithmically into Ay and ky. is plotted in figure I..," The result, binned logarithmically into $k_{\parallel}$ and $\mathbf{k}_{\perp}$, is plotted in figure \ref{fig2}."
 The 2072 covers two declination strips of approximately 5757. and. we have three such mock declination strips available.," The 2QZ covers two declination strips of approximately $5^{\circ} \times 75^{\circ}$, and we have three such mock declination strips available."
 When calculating Pryκι) information about the line of sight must be retained.," When calculating $P^S(k_{\parallel},\mathbf{k}_{\perp})$ information about the line of sight must be retained."
 To achieve this the data are divided into subsamples that subtend a small solid angle on the sky. and the distant-observer approximation is then applied.," To achieve this the data are divided into subsamples that subtend a small solid angle on the sky, and the distant-observer approximation is then applied."
 The data from each declination strip are split into 4 regions each approximately 590 during this analysis. and the average of the resulting power spectra is taken.," The data from each declination strip are split into 4 regions each approximately $5^{\circ} \times 20^{\circ}$ during this analysis, and the average of the resulting power spectra is taken."
 This approach introduces a small svstematic error: the measured. P7(Ej.κι} is in fact a convolution between the rue power spectrum and a function related to the sample window introduced due to the assumption of parallel lines of sight.," This approach introduces a small systematic error; the measured $P^S(k_{\parallel},\mathbf{k}_{\perp})$ is in fact a convolution between the true power spectrum and a function related to the sample window introduced due to the assumption of parallel lines of sight."
 The resulting power spectrum is more isotropic. aud his leads to a small. vel svstematic uncder-estimate of 3.," The resulting power spectrum is more isotropic, and this leads to a small, yet systematic under-estimate of $\beta$."
" Fora 5"".20 window this error is <14 and considerably smaller than the statistical errors obtained.", For a $5^{\circ} \times 20^{\circ}$ window this error is $< 1\%$ and considerably smaller than the statistical errors obtained.
 Although it is »ossible to correct for this clleet (see Cole et al., Although it is possible to correct for this effect (see Cole et al.
 1994). for he purposes of this analysis we have chosen to ignore it.," 1994), for the purposes of this analysis we have chosen to ignore it."
 The measured. power spectrum is also convolved with he power spectrum of the window function., The measured power spectrum is also convolved with the power spectrum of the window function.
 As the QSO co-moving number density is almost constant out to very large scales (Llovle ct al., As the QSO co-moving number density is almost constant out to very large scales (Hoyle et al.
 2001a). we are effectively. considering a volume limited sample.," 2001a), we are effectively considering a volume limited sample."
 Hence cach QSO ceuries. equal weight. and the survey window function. VW(x). simply takes a value of unity in the volume of the universe included in the survey. and zero elsewhere.," Hence each QSO carries equal weight, and the survey window function, $W({\mathbf x})$, simply takes a value of unity in the volume of the universe included in the survey, and zero elsewhere."
 This is approximated. using a catalogue containing a large number of unclustered points with the same racial and angular distribution as the survey. and its power spectrum. is calculated. in a similar manner to that of the data.," This is approximated using a catalogue containing a large number of unclustered points with the same radial and angular distribution as the survey, and its power spectrum is calculated in a similar manner to that of the data."
 Figure 2. shows the power spectrum of, Figure \ref{fig1} shows the power spectrum of
Either of the above two possibilities might provide a better approximation depending on whether a vortex unpins as a whole along its length. or only small segments of it are unpinned randomly.,"Either of the above two possibilities might provide a better approximation depending on whether a vortex unpins as a whole along its length, or only small segments of it are unpinned randomly."
 For the superfIuid in the crust of a neutron star. simultaneous unpinning of a vortex as a whole must be ruled out. given the huge iunber of the pinning centers (ie.," For the superfluid in the crust of a neutron star, simultaneous unpinning of a vortex as a whole must be ruled out, given the huge number of the pinning centers (ie."
 (he nuclei of the solid crust) along each vortex (having a length of a km or so): hence case should be more probable., the nuclei of the solid crust) along each vortex (having a length of a km or so); hence case should be more probable.
 In contrast. case might be the proper choice for the laboratory experiments in which a vortex pins only al ils end points (Illledge Glaberson 1980: SSchwarz 1981).," In contrast, case might be the proper choice for the laboratory experiments in which a vortex pins only at its end points \markcite{hed80}H Hedge Glaberson 1980; \markcite{sch81}S Schwarz 1981)."
 Further theoretical work max indicate the extent to which the (statistically averaged) motion of individual vortices could deviate from a uniform local densitv. and distinguish between the above alternative possibilities lor (he initial conditions of the rotation rate of the vortices upon unpinning.," Further theoretical work may indicate the extent to which the (statistically averaged) motion of individual vortices could deviate from a uniform local density, and distinguish between the above alternative possibilities for the initial conditions of the rotation rate of the vortices upon unpinning."
 The relaxation timescale of the vortex array to (he new conditions. in each case. would be likewise relevant for making a decision.," The relaxation timescale of the vortex array to the new conditions, in each case, would be likewise relevant for making a decision."
 Also. the distinct behavior of the superfIuid spin-down. for Vo=0. in the (wo cases might be possible to be Lested. experimentally.," Also, the distinct behavior of the superfluid spin-down, for $N=0$, in the two cases might be possible to be tested, experimentally."
" An assumed general model of à normal (non-superfIuid) component plus the ""erust. with moments of inertia. J, aud. /.. and rotation Irequencies Q, and Q.. respectively. under (he influence of an external negative torque. —N acting primarily on (he crust-component. would obey the following dvnamical relations Bavin et."," An assumed general model of a normal (non-superfluid) component plus the “crust”, with moments of inertia $I_{\rm n}$ and $I_{\rm c}$ , and rotation frequencies $\Omega_{\rm n}$ and $\Omega_{\rm c}$, respectively, under the influence of an external negative torque $- N$ acting primarily on the crust-component, would obey the following dynamical relations \markcite{Bayet69}
 Baym et."
 al., al.
" 1969) where /=/.+/,. ancl 7, is (he velocitv-relaxation time lor (he dissipation of microscopic relative motion between the constituent particles of the two components."," 1969) where $I = I_{\rm c} + I_{\rm n}$, and $\tau_v$ is the velocity-relaxation time for the dissipation of microscopic relative motion between the constituent particles of the two components."
 A solution of the two coupled equations indicate exponential relaxations of the rotation frequencies Q.(/) and Qu). with time /.," A solution of the two coupled equations indicate exponential relaxations of the rotation frequencies $\Omega_{\rm
c}(t)$ and $ \Omega_{\rm n}(t)$, with time $t$."
" The exponential time constant το. referred to as the dynamical coupling timescale of the svstem. is given as In (he case of a superfIuid component. (he relation between 7, and 7, would be in general different (hanthat in Eq."," The exponential time constant $\tau_{\rm D}$, referred to as the dynamical coupling timescale of the system, is given as In the case of a superfluid component, the relation between $\tau_{\rm D}$ and $\tau_v$ would be in general different thanthat in Eq."
 12. as discussed below.," 12, as discussed below."
 Further. the steady-state behavior inferred," Further, the steady-state behavior inferred"
B379 by fitting its deep photometry extending below the niain-sequence turn-off to the isochrones of the Padova eroup (Ciürardietal.2002.2008:Marigo2008).,"B379 by fitting its deep photometry extending below the main-sequence turn-off to the isochrones of the Padova group \citep{Girardi02,Girardi08,Marigo08}."
. The age obtained in this paper is consistent with the cleternunation of Brownctal.(200la).. aud confirms the couclusion of Brownetal.(200Lla) that. D379 is 23 Cyr voungcr than the oldest Galactic GC.," The age obtained in this paper is consistent with the determination of \citet{brown04a}, and confirms the conclusion of \citet{brown04a} that, B379 is 2–3 Gyr younger than the oldest Galactic GC."
 This paper also coufirius the consistence of the age scale of D379 between the Padova group isochrones used im this paper and the 2006 VaudenBere isochrones used by Brownal.(200 12)., This paper also confirms the consistence of the age scale of B379 between the Padova group isochrones used in this paper and the 2006 VandenBerg isochrones used by \citet{brown04a}.
. So. the results’ comparison between Brownetal.(200la) and Maetal.(2007) is imeanineful.," So, the results' comparison between \citet{brown04a} and \citet{Ma07} is meaningful."
 Iu addition. the metal abundance. reddening value. aud distance modulus obtained in this paper are consistent with the previous deteruinuations.," In addition, the metal abundance, reddening value, and distance modulus obtained in this paper are consistent with the previous determinations."
 We are indebted to the referee for thoughtful coments and insightful sugeestious that improved this paper sienificauth., We are indebted to the referee for thoughtful comments and insightful suggestions that improved this paper significantly.
 We are thankful to Dr. Brown for providing the HST/ACS WEC data for D379., We are thankful to Dr. Brown for providing the /ACS WFC data for B379.
 This work was supported by the Chinese National Natural Science Foundation grands No., This work was supported by the Chinese National Natural Science Foundation grands No.
 10873016. 106323020. LOGO3006. and 10803007. aud by National Basic Research Program of China (973 Program). No.," 10873016, 10633020, 10603006, and 10803007, and by National Basic Research Program of China (973 Program), No."
 2007CD815103., 2007CB815403.
several times.,several times.
 Finally for 25l out of 397 sources improved coordinates were deterinined., Finally for 254 out of 397 sources improved coordinates were determined.
 Iu cases where positional correction was possible the relmaiing svsteniatic error consists of the error iu former optical measurements and the statistical error of the identified source., In cases where positional correction was possible the remaining systematic error consists of the error in former optical measurements and the statistical error of the identified source.
" For not corrected sources the systematic error was set to7"".", For not corrected sources the systematic error was set to.
. The positional error was finally computed as a composite of the statistical uncertainty with 90 confidence aud the svstematic error., The positional error was finally computed as a composite of the statistical uncertainty with 90 confidence and the systematic error.
 It is used throughout the paper for he error circle., It is used throughout the paper for the error circle.
 After the source detection procedure the mean positional error was 8733, After the source detection procedure the mean positional error was 3.
", The coordinate correction reduced the mean positional error of all sources to LL.", The coordinate correction reduced the mean positional error of all sources to 4.
 For position corrected sources the mean positional error is 57/11., For position corrected sources the mean positional error is 1.
 The catalogue was cross-correlated with the SINBAD data base and the TYCIIO catalogue in order to identify IRI sources., The catalogue was cross-correlated with the SIMBAD data base and the TYCHO catalogue in order to identify HRI sources.
 The IIRI catalogue coutaius samples of known SSSs. N-rav binaries. SNRs. Galactic foreground stars. and backeround ACN.," The HRI catalogue contains samples of known SSSs, X-ray binaries, SNRs, Galactic foreground stars, and background AGN."
 The catalogue was also cross-correlated with the source list from the pointed PSPC observations (ΠΡΟ)., The catalogue was also cross-correlated with the source list from the pointed PSPC observations (HP99b).
 138 IIRI sources are identical with sources Which were detected i1 PSPC data aud thus for most of thei the harduess ratios (IRI. ITR2) are known.," 138 HRI sources are identical with sources which were detected in PSPC data and thus for most of them the hardness ratios (HR1, HR2) are known."
 Since the TRI had no spectra resolution no information ou the N-ray spectrmau could be obtaiued for ITRI sources which are completely new detections., Since the HRI had no spectral resolution no information on the X-ray spectrum could be obtained for HRI sources which are completely new detections.
 A total of 91 IIRI, A total of 94 HRI
pixel noise but produces a slight variation of the wavelength bin along the spectra.,pixel noise but produces a slight variation of the wavelength bin along the spectra.
 The calibrated spectra were corrected for atmospheric extinc(üon adopting the airmass values previously menüoned., The calibrated spectra were corrected for atmospheric extinction adopting the airmass values previously mentioned.
 Finally they have been corrected also lor interstellar extinction using the E(B-V) values of 0.015 and 0.05 obtained from the NASA/IPAC Extragalactic Database for the galaxies D6 ancl C2. respectively aud adopting (he standard Seaton (1979) exlinclion curve.," Finally they have been corrected also for interstellar extinction using the E(B-V) values of 0.015 and 0.05 obtained from the NASA/IPAC Extragalactic Database for the galaxies D6 and C2, respectively and adopting the standard Seaton (1979) extinction curve."
 The final extracted and calibrated spectra are shown in Fig.l and Fig.2 (upper panels) with the resulting noise level., The final extracted and calibrated spectra are shown in Fig.1 and Fig.2 (upper panels) with the resulting noise level.
 Strong SilI1260. OIT303. CIIL334. S3IV1393. SilI1526. CIV1548 absorption lines are clearly visible in (he spectra and give the same absorption redshifts [or the (wo galaxies as provided bv Pettini et al. (," Strong SiII1260, OI1303, CII1334, SiIV1398, SiII1526, CIV1548 absorption lines are clearly visible in the spectra and give the same absorption redshifts for the two galaxies as provided by Pettini et al. ("
1993).,1998).
 Their rest frame equivalent widths are of the order of 3., Their rest frame equivalent widths are of the order of 3.
Α. At the same time. the Lyman-a emissions are different being stronger in D6 with W=41/(14-:)10.5 thanin C2. IV=14.5/(1+2)3.4 and correspondingly. the UV dust absorption estimated by Petini et al. (," At the same time, the $\alpha$ emissions are different being stronger in D6 with $W=41/(1+z)=10.5$ thanin C2, $W=14.5/(1+z)=3.4$ and correspondingly the UV dust absorption estimated by Pettini et al. ("
1998) in the (wo galaxies is stronger in C2 (A4544—1.14) than in DG (A450070.8).,1998) in the two galaxies is stronger in C2 $_{1500}$ =1.14) than in D6 $_{1500}$ =0.8).
 It is clear from Fig.l and Fig.2 (lower panels) that there is no significant emission shortwarcd of the Lyman limit (912 A) in both galaxies., It is clear from Fig.1 and Fig.2 (lower panels) that there is no significant emission shortward of the Lyman limit (912 ) in both galaxies.
 The average (lux is zero and the noise level in (he spectra can provide a useful upper limit to the fizoo/Joon ratio aud hence io the UV escape [raction., The average flux is zero and the noise level in the spectra can provide a useful upper limit to the $f_{1500}/f_{900}$ ratio and hence to the UV escape fraction.
 As in SPA we define a spectral region shortward of the Lyman Limit sufficiently large to have enough statistics bul keeping al (he same time a low noise level., As in SPA we define a spectral region shortward of the Lyman Limit sufficiently large to have enough statistics but keeping at the same time a low noise level.
 The selected rest [rame region is between 880 and 910A., The selected rest frame region is between 880 and 910.
" The pixel-to-pixel noise rms obtained in (he selected region is e(880—910)=3.3 (relative [lux densities in !) for the galaxy C2 which corresponds (o an average upper limit [lux in the region of foy,«3.3/V54Ap.r!=0.45."," The pixel-to-pixel noise rms obtained in the selected region is $\sigma
(880-910)=3.3$ (relative flux densities in $^{-1}$ ) for the galaxy C2 which corresponds to an average upper limit flux in the region of $f_{900}<3.3/\sqrt{54 pxl}=0.45$."
 The corresponding average flux level at 1500 ls 33cx2.5/y21Iprl= 20.6.," The corresponding average flux level at 1500 is $33\pm 2.5/\sqrt{21 pxl} =33\pm
0.6$ ."
" The resulting lower limit to the observed. fi305/foi ratio is Jison/ου>73 lor the galaxy CY,", The resulting lower limit to the observed $f_{1500}/f_{900}$ ratio is $f_{1500}/f_{900}>73$ for the galaxy C2.
 Sinularly for the galaxy DG the pixel-to-pixel noise rms obtained in (he selected region is c(880—910)=1.1 corresponding to an average upper limit flux in (he regionof 0.16.," Similarly for the galaxy D6 the pixel-to-pixel noise rms obtained in the selected region is $\sigma (880-910)=1.1$ corresponding to an average upper limit flux in the regionof $f_{900}<1.1/\sqrt{50
pxl}=0.16$ ."
 The corresponding average fhix level nearby 1500 rest [rame is 11dL5/430prf=11zx0.3., The corresponding average flux level nearby 1500 rest frame is $11\pm 1.5/\sqrt{30 pxl} =11\pm 0.3$.
 The resulting lower limit to the observed f1500/foo) ratio is toon.Py2»69.ET) , The resulting lower limit to the observed $f_{1500}/f_{900}$ ratio is $f_{1500}/f_{900}>69$ .
The escape fraction defined as the fraction of emitted900 A photons that escapes the, The escape fraction defined as the fraction of emitted900 photons that escapes the
second Fourier transform we would get YP= (n | 119) whichi tends to a constant as we add more fields.,second Fourier transform we would get )|^2= (n + ) which tends to a constant as we add more fields.
 We rave to take the second Fourier transform of each field averaging to see the non-Gaussianity., We have to take the second Fourier transform of each field averaging to see the non-Gaussianity.
 This is because the atures in pz(oa) are repeated at dillerent à in each field and heir sum will tend to a constant for many fields., This is because the features in $\rho^2_k(\alpha)$ are repeated at different $\alpha$ in each field and their sum will tend to a constant for many fields.
 Taking the second Fourier transform takes the feature [rom pz(a) and ransfornis it to a feature at the same 2 in £463) for each ield., Taking the second Fourier transform takes the feature from $\rho^2_k(\alpha)$ and transforms it to a feature at the same $\beta$ in $F_k(\beta)$ for each field.
 The idea of looking at small fields in this way is an extension of the idea in (Ferreira& MagueijoLOOT) of analvsing several small fields separately., The idea of looking at small fields in this way is an extension of the idea in \cite{fermag} of analysing several small fields separately.
 For the example ofa cosmic string on a Gaussian background. used in Alagueijo1997) the shape spectrum shows the feature well. but. the string was always considere to have the same orientation.," For the example of a cosmic string on a Gaussian background used in \cite{fermag} the shape spectrum shows the feature well, but the string was always considered to have the same orientation."
 Wo the shape spectrum were averaged over dillerent. directions the result. woulc disappear., If the shape spectrum were averaged over different directions the result would disappear.
 Llere we can look at the sane fields with differen orientations and the will still be apparent., Here we can look at the same fields with different orientations and the will still be apparent.
 However. if we randomly choose small fields from a larger one containing many shapes we would have to take translations into accoun as well and the οσοι will not be so clear.," However, if we randomly choose small fields from a larger one containing many shapes we would have to take translations into account as well and the effect will not be so clear."
 As a simple demonstration of this statistic on an obviously non-Gaussian field. we will consider applying it to some regular shapes.," As a simple demonstration of this statistic on an obviously non-Gaussian field, we will consider applying it to some regular shapes."
 Phese are intended as mock small scale dust maps., These are intended as mock small scale dust maps.
 We expect our geometric shapes to present the same tvpe of non-Ciaussianity. ancl obstacles for its detection. as the actual dust maps.," We expect our geometric shapes to present the same type of non-Gaussianity, and obstacles for its detection, as the actual dust maps."
 First to demonstrate. the ideas in Section. ?? we will look at fields cach containing one square at. dilferent orientations., First to demonstrate the ideas in Section \ref{small} we will look at fields each containing one square at different orientations.
 We first Fourier transform the whole field and take the amplitude: squared. (p» (&))., We first Fourier transform the whole field and take the amplitude squared $\rho^2({\bmath k})$ ).
 For one square the amplitude is not alfected. by translation of the square., For one square the amplitude is not affected by translation of the square.
 We multiply the field by a Gaussian window to get rid of edge ellects (Llobson&Magueijo 1996)., We multiply the field by a Gaussian window to get rid of edge effects \cite{hobmag}. .
. In Figure 1. we have plotted [or one field. ρα). and ο) at hk=60 from a 200° array., In Figure \ref{1square1} we have plotted for one field $\rho^2(\alpha)$ and $|F(\beta)|$ at $k = 60$ from a $200^2$ array.
 This corresponds to a physical scale. of &=21600/L where L is the size of the real field in degrees., This corresponds to a physical scale of $k= 21600/L$ where $L$ is the size of the real field in degrees.
 The v-axis is in arbitrary units., The y-axis is in arbitrary units.
 Figure 2 shows the same functions averaged over 100 fields., Figure \ref{1square100} shows the same functions averaged over 100 fields.
 The averaged. [/762)] is not in fact identical to 1) for one field. because of numerical ellects due to the discrete Fourier transform being on a lattice.," The averaged $|F(\beta)|$ is not in fact identical to $|F(\beta)|$ for one field, because of numerical effects due to the discrete Fourier transform being on a lattice."
 The average settles down quickly. however: |£62)| looks the same averaged over 10 fieles as it does averaged over 100 fields.," The average settles down quickly, however: $|F(\beta)|$ looks the same averaged over 10 fields as it does averaged over 100 fields."
 Phe first Fourier transform. poe). tends to à constant as we average over more fields as we would expect. but is also alfected by the lattice.," The first Fourier transform, $\rho^2(\alpha)$ , tends to a constant as we average over more fields as we would expect, but is also affected by the lattice."
 Figures 3. and 4 show the equivalent graphs for a Gaussian random field., Figures \ref{gauss1} and \ref{gauss100} show the equivalent graphs for a Gaussian random field.
 Here [£(2)| converges rapidly. and is distinct. [rom the result. lor squares., Here $|F(\beta)|$ converges rapidly and is distinct from the result for squares.
 The dips in. p-(e) are due to the window function superimposed in real space to avoid edge ellects from the first Fourier transform., The dips in $\rho^2(\alpha)$ are due to the window function superimposed in real space to avoid edge effects from the first Fourier transform.
 The Gaussian is treated in this way even though we generate the original map in Fourier space so as to model as closely as possible theway in which the map is treated., The Gaussian is treated in this way even though we generate the original map in Fourier space so as to model as closely as possible theway in which the map is treated.
 We would of course expect (20.7)| totend to a delta function, We would of course expect $|F(\beta)|$ totend to a delta function
—wadüauüedt The extragalactic radio source population ranges from normal galaxies with hunuinosities ~107? W + to ealaxics whose radio cussion is as iucl as 10!10° times ereater owing to regious of massive star formation or to an active galactic nucleus (ACN).,0 The extragalactic radio source population ranges from normal galaxies with luminosities $\sim 10^{20}$ W $^{-1}$ to galaxies whose radio emission is as much as $10^4 -10^7$ times greater owing to regions of massive star formation or to an active galactic nucleus (AGN).
 The population of radio sources iu the sky with fiux densities 2 Lundy is dominated by ACN driven ciission eenecrated from the eravitational poteutial associated with a supermassive black-hole iu the nucleus., The population of radio sources in the sky with flux densities $> 1$ mJy is dominated by AGN driven emission generated from the gravitational potential associated with a supermassive black-hole in the nucleus.
 For these sources. the observed radio emission includes the classical extended jet and double lobe radio sources as well as compact radio components more directly associated with the cucrey generation and collimation near the central eugme.," For these sources, the observed radio emission includes the classical extended jet and double lobe radio sources as well as compact radio components more directly associated with the energy generation and collimation near the central engine."
 Below 1 iimJw there Is anu increasing coutribution to the radio source population fron svuchrotron cussion resulting from relativistic plasma ejected from supernovae associated with massive star formation in galaxies or groups of galaxies. often associated with ierecrs or iuteractious (e.g.Wind-Ww 102).," Below 1 mJy there is an increasing contribution to the radio source population from synchrotron emission resulting from relativistic plasma ejected from supernovae associated with massive star formation in galaxies or groups of galaxies, often associated with mergers or interactions \citep[e.g.,][]{win95,ric98,fom02}."
.. However. the mix of star forming galaxies (SFC) xb ACN and their dependence on epoch is not well deteruiued.," However, the mix of star forming galaxies (SFG) and AGN and their dependence on epoch is not well determined."
" For example. Couppioietal.(2003).. bv using pau Infrared Space Observatory (ISO) aud 1.1 GIIz data have sugeested that starburst galaxies start to become a significant fraction of the radio source population for S,m05OS indy. exceeding ~GO only di S,«50 ud."," For example, \cite{gru03}, by using $\mu$ m Infrared Space Observatory (ISO) and 1.4 GHz data have suggested that starburst galaxies start to become a significant fraction of the radio source population for $S_{\rm r}
\la 0.5 - 0.8$ mJy, exceeding $\sim 60\%$ only at $S_{\rm r} < 50~
\mu$ Jy."
 Ou the other haud. Πανetal.beep(2005).. bv fitting radio uunber counts of the Πιο] Field South aud other surveys with evolutionary starburst models. have estimated that the wmuber of SFC exceeds 50% of the total already at ~0.25 indy.," On the other hand, \cite{hu05}, by fitting radio number counts of the Hubble Deep Field South and other surveys with evolutionary starburst models, have estimated that the number of SFG exceeds $50\%$ of the total already at $\sim 0.25$ mJy."
 The most commonly accepted paradiem has been that the subauJw population is larecly made up of SFC., The most commonly accepted paradigm has been that the sub-mJy population is largely made up of SFG.
 A (small but erowiue) number of rescarchers. however," A (small but growing) number of researchers, however"
Once we have determined (or assumed) the cluster formation rates and tlie age-metallicity relation. we know the mass of clusters formed at each age and metallicity.,"Once we have determined (or assumed) the cluster formation rates and the age-metallicity relation, we know the mass of clusters formed at each age and metallicity."
 To produce a svnthetic CXIRD. however. we need the actual number of clusters formed.," To produce a synthetic CMRD, however, we need the actual number of clusters formed."
 We first divide the mass of clusters at one age and metallicity according to an assumed cluster initial mass function., We first divide the mass of clusters at one age and metallicity according to an assumed cluster initial mass function.
 The function can be anything: we will assume a power law of the form Or As (his function diverges at low mass for a<—1. we arbitrarily choose a low-mass cutoff of 1 solar mass for now. and we hope that we can constrain the cluster initial mass function at low mass from the closer galaxies in our sample.," The function can be anything; we will assume a power law of the form or As this function diverges at low mass for $\alpha \le -1$, we arbitrarily choose a low-mass cutoff of 1 solar mass for now, and we hope that we can constrain the cluster initial mass function at low mass from the closer galaxies in our sample."
 The second. division of clusters will be according to cluster radius., The second division of clusters will be according to cluster radius.
 This is extremely difficult. as our observed data have very. strong selection effects in radius: small clusters will be classified as stars while large clusters will be too difficult to photometer or may be identified as small associations and thus rejected.," This is extremely difficult, as our observed data have very strong selection effects in radius: small clusters will be classified as stars while large clusters will be too difficult to photometer or may be identified as small associations and thus rejected."
 Because we are sensitive to only a limited range of core radii (0.2—4 pe in our NGC 3627 data). we will initially assume a [Lal distribution of cluster radii. We will be able to constrain this huther by comparing the sharpness distributions of the observed and synthetic CMIBRDS.," Because we are sensitive to only a limited range of core radii $0.2 - 4$ pc in our NGC 3627 data), we will initially assume a flat distribution of cluster radii, We will be able to constrain this further by comparing the sharpness distributions of the observed and synthetic CMRDs."
 With the clusters now divided by age. metallicity. initial cluster mass. and core radius. one other factor affects (he physical properties of the cluster population: lifetimes.," With the clusters now divided by age, metallicity, initial cluster mass, and core radius, one other factor affects the physical properties of the cluster population: lifetimes."
 However. in terms of generating a synthetic CMBRD. we consider this effect unimportant due to the very short timescales over which clusters are visible above our photometric eutoffs.," However, in terms of generating a synthetic CMRD, we consider this effect unimportant due to the very short timescales over which clusters are visible above our photometric cutoffs."
 With our limiting magnitude of V.=24. we only require that a cluster with an initial mass of 105M. lives 10—15 Mir and that one with a mass of LO?AL. lives 30—40 Myr. (," With our limiting magnitude of $V = 24$, we only require that a cluster with an initial mass of $10^4 M_{\odot}$ lives $10-15$ Myr and that one with a mass of $10^5 M_{\odot}$ lives $30-40$ Myr. ("
That the upper end of the mass function is unalfected by disruption is demonstrated by the analvtical models of Fall Zhang 2001.),That the upper end of the mass function is unaffected by disruption is demonstrated by the analytical models of Fall Zhang 2001.)
 As cluster lifetimes are expected to be at least an order of magnitude bevond these values (IX&roupa.Aarseth.&ILIurlev2001:FallZhang2001).. we can thus ignore the effect of cluster disruption.," As cluster lifetimes are expected to be at least an order of magnitude beyond these values \citep{kro01b,fal01}, we can thus ignore the effect of cluster disruption."
"Although the asymptotic calibration for Hoo has been previously characterized by Monte Carlo (??),, the analytic solution extends the calibration to an arbitrary collection of harmonics. ","Although the asymptotic calibration for $H_{20}$ has been previously characterized by Monte Carlo \citep{dejager_1,dejager_2}, the analytic solution extends the calibration to an arbitrary collection of harmonics. ("
"method used here to develop the asymptotic distribution(The can be easily extended to cases where the harmonics are not sequential, e.g. a test incorporating only the 2nd, 10th, and 20th harmonics.","The method used here to develop the asymptotic distribution can be easily extended to cases where the harmonics are not sequential, e.g. a test incorporating only the 2nd, 10th, and 20th harmonics."
" However, for simplicity's sake we retain the original formulation of an inclusive set of harmonics.)"," However, for simplicity's sake we retain the original formulation of an inclusive set of harmonics.)"
" It is apparent from Figure A10 that the distributions approach a limiting distribution as m, the maximum harmonic, becomes large."," It is apparent from Figure \ref{ch5_plot6} that the distributions approach a limiting distribution as $m$, the maximum harmonic, becomes large."
" Indeed, as we see in Figure A12,, for m>10, there is very little difference in the tail probability for *practical"" values of H."," Indeed, as we see in Figure \ref{ch5_plot18}, for $m\geq10$, there is very little difference in the tail probability for “practical” values of $H$."
" For instance, the chance probabilities of observing H,,>50 for any harmonic m>10 are all within a factor of 2 of each other, a negligible distinction at this significance level, and for H,,«50, the discrepancy is even smaller."," For instance, the chance probabilities of observing $H_m>50$ for any harmonic $m\geq10$ are all within a factor of 2 of each other, a negligible distinction at this significance level, and for $H_m<50$, the discrepancy is even smaller."
" We conclude that, once one has made the decision to choose m large enough to allow for sharply-peaked light curves—and indeed, this is the whole point of the H-test—that there is no penalty for making m as large as feasible."," We conclude that, once one has made the decision to choose $m$ large enough to allow for sharply-peaked light curves---and indeed, this is the whole point of the $H$ -test—that there is no penalty for making $m$ as large as feasible."
" In this sense,omnibus, as m can be chosen large enough to make the test sensitive to light curves with arbitrarily sharp features."," In this sense, as $m$ can be chosen large enough to make the test sensitive to light curves with arbitrarily sharp features."
" Finally, we note a practical formula for the cdf applicable for m> 10: 1—F(A)&exp(—0.398405H)."," Finally, we note a practical formula for the cdf applicable for $m\geq10$ : $1-F(H)\approx\exp(-0.398405\,H)$."
" This is, naturally, quite close to the formula reported by ?.."," This is, naturally, quite close to the formula reported by \citet{dejager_2}."
" For m« 10, quadratic terms in the exponential are important for an accurate evaluation of the tail probability and the fullexpression Eq."," For $m<10$ , quadratic terms in the exponential are important for an accurate evaluation of the tail probability and the fullexpression Eq."
 A10 should be used., \ref{f3} should be used.
data the W ratio is 1.4 at its higher-resolution.,data the W ratio is 1.4 at its higher-resolution.
 When the McDonald. data were artificially degraded. for comparison. the W ratio increased to 1.5.," When the McDonald data were artificially degraded for comparison, the W ratio increased to 1.5."
 Phis suggests that in the lower resolution APO data the bluc-shifted emission is allecting the W ratio and that the material seen against the star is not optically thin. but not entirely. optically thick.," This suggests that in the lower resolution APO data the blue-shifted emission is affecting the W ratio and that the material seen against the star is not optically thin, but not entirely optically thick."
 Taking the AleDonald data into consideration. it seems that we are seeing the same outflow of material in the cavity that is flowing into the outer nebula.," Taking the McDonald data into consideration, it seems that we are seeing the same outflow of material in the cavity that is flowing into the outer nebula."
 Schmidt&Witt(1991) observed that the D-line emission. was strongest. along the bi-cone walls or ‘whiskers’., \citet{schmidt1991} observed that the D-line emission was strongest along the bi-cone walls or `whiskers'.
 At large distances [rom the centre of he nebula we directly. see the emission. which mav be domünated bv emission from. walls of the outflow: cavity.," At large distances from the centre of the nebula we directly see the emission, which may be dominated by emission from walls of the outflow cavity."
 The optical depth. through the edge of outflow: would. be argest., The optical depth through the edge of outflow would be largest.
" These observations are consistent with the work of Warren-Smith.Searrott.&Murdin(1981). who found that he sodium emission was optically thin on either side of the ""whiskers! and optically-thick on the whiskers’.", These observations are consistent with the work of \citet{warren-smith1981} who found that the sodium emission was optically thin on either side of the `whiskers' and optically-thick on the `whiskers'.
 Similarly. the ratio of red-shifted. emission. height. 1t-wight 5890 to R-height 5896 is also constant with hase: however. the ratio is about 0.84.," Similarly, the ratio of red-shifted emission height R-height 5890 to R-height 5896 is also constant with phase; however, the ratio is about 0.84."
 The 5896A DI-line consistently has a stronger red component than the 5890 D2-line., The 5896 D1-line consistently has a stronger red component than the 5890 D2-line.
 This remains a puzzle: it was expected that the recd-shifted. emission component max be more highly scattered and thus weaker than the blue-shifted: component., This remains a puzzle; it was expected that the red-shifted emission component may be more highly scattered and thus weaker than the blue-shifted component.
 This ellect should apply equally to both lines of the doublet., This effect should apply equally to both lines of the doublet.
 Since the ratio for the blue-shifted: emission. is unity. implving near optically-thick emission. the expectation. would. be that the red emission ratio would also be unity.," Since the ratio for the blue-shifted emission is unity, implying near optically-thick emission, the expectation would be that the red emission ratio would also be unity."
 However. the smaller ratio of the red. component is. unexpected for the optically thick case.," However, the smaller ratio of the red component is unexpected for the optically thick case."
 This could indicate that the photospheric subtraction technique of 5.5. is insullicicnt. but without the means to spatially separate the components at high-resolution we cannot resolve the issue in this paper.," This could indicate that the photospheric subtraction technique of \ref{technique} is insufficient, but without the means to spatially separate the components at high-resolution we cannot resolve the issue in this paper."
 The bluc-shifted: emission ancl blue-shifted. absorption in the APO spectra are seen to be strongest near phase ó= 0.48 (just after periastron) coinciding with the period of the fastest mass outllow from the primary star. based on the asvnimetry of the photospheric absorption lines. see. panel c of Fig. 10..," The blue-shifted emission and blue-shifted absorption in the APO spectra are seen to be strongest near phase $\phi =$ 0.48 (just after periastron) coinciding with the period of the fastest mass outflow from the primary star, based on the asymmetry of the photospheric absorption lines, see panel c of Fig. \ref{fig10}."
 During these phases. mass is being added to the outflow resulting in increased. absorption and emission.," During these phases, mass is being added to the outflow resulting in increased absorption and emission."
 The close proximity in time of the strongest. emission aud absorption is suggestive that the outllows seen in absorption and emission are one and the same., The close proximity in time of the strongest emission and absorption is suggestive that the outflows seen in absorption and emission are one and the same.
 This argument provides urther support to the discussion in 6.2. relating the velocities of these components., This argument provides further support to the discussion in \ref{semissioncomp} relating the velocities of these components.
 Note that the gas seen in D absorption is significantly slower than that of lle. seen in absorption near superior conjunction. and attributed o the jet from the secondary (Wittetal.2009).," Note that the gas seen in D absorption is significantly slower than that of $\alpha$, seen in absorption near superior conjunction, and attributed to the jet from the secondary \citep{witt2009}."
 Ehis also supports the idea that the outflow seen in D emission and absorption originates from. the wimary star. rather than from the secondary stars accretion disc jet.," This also supports the idea that the outflow seen in D emission and absorption originates from the primary star, rather than from the secondary star's accretion disc jet."
 Hlowever. why the emission components. which extend. to large distances from the central source. are correlated with the motion of the primary star is unknown.," However, why the emission components, which extend to large distances from the central source, are correlated with the motion of the primary star is unknown."
 A comparison of single-peaked emission. (for cxaniple La (Wittctal. 2009))) and clouble-peakecl emission (às seen in D-line). reveals that they are representative of, A comparison of single-peaked emission (for example $\alpha$ \citep{witt2009}) ) and double-peaked emission (as seen in D-line) reveals that they are representative of
"suggesting that while the line peak may be moderately optically thick (tTu¢q~1.4—1.6 at most), the integrated !9ΟΟ(1-0) emission is optically thin over the central region.","suggesting that while the line peak may be moderately optically thick $\tau_{\rm ^{13}CO} \sim 1.4-1.6$ at most), the integrated $^{13}$ CO(1-0) emission is optically thin over the central region."
" This is also consistent with the observed lack of variation of the integrated (2-1)/(1-0) line ratios, as a function of the integrated line intensities (Fig. 6))."," This is also consistent with the observed lack of variation of the integrated (2-1)/(1-0) line ratios, as a function of the integrated line intensities (Fig. \ref{fig:opt-thin}) )."
 The points in the plots correspond to the values of the line ratios taken over 46” beams., The points in the plots correspond to the values of the line ratios taken over $^{\prime\prime}$ beams.
" Given that the !?CO emission at the peak of the line is close to be optically thick, as an estimate of the excitation temperature of the molecular gas we assumed 20 K, the peak line brightness temperature as measured in the clump with the highest optical depth."," Given that the $^{13}$ CO emission at the peak of the line is close to be optically thick, as an estimate of the excitation temperature of the molecular gas we assumed 20 K, the peak line brightness temperature as measured in the clump with the highest optical depth."
 In the following we have assumed this excitation temperature for all the clumps., In the following we have assumed this excitation temperature for all the clumps.
" The total column density of the CO molecule for each identified clump was derived using the J=1—0 transition, under the assumption of local thermodynamic equilibrium (LTE) at an excitation temperature T,,."," The total column density of the $^{13}$ CO molecule for each identified clump was derived using the $J=1-0$ transition, under the assumption of local thermodynamic equilibrium (LTE) at an excitation temperature $T_{ex}$."
" The column density of this molecule, in the optically thin limit, is given by where: (see Eq. ["," The column density of this molecule, in the optically thin limit, is given by where: (see Eq. ["
A1] and [A4] of 1,A1] and [A4] of ).
"986)). fTadv is the integrated line brightness temperature in K km s! of the transition with frequency visco (Hz), Bisco is the rotational constant of the molecule and asco is CO's permanent electric dipole moment, which is taken to be 0.1101 Debye."," $\int T_{\rm B} {\rm dv}$ is the integrated line brightness temperature in K km $^{-1}$ of the transition with frequency $\nu_{\rm ^{13}CO}$ (Hz), $B_{\rm ^{13}CO}$ is the rotational constant of the molecule and $\mu_{\rm ^{13}CO}$ is $^{13}$ CO's permanent electric dipole moment, which is taken to be 0.1101 Debye."
" Assuming Τεν=20 K for every clump, we computed the total column density of each clump identified in the FCRAO CO(1-0) emission."," Assuming $T_{ex} = 20$ K for every clump, we computed the total column density of each clump identified in the FCRAO $^{13}$ CO(1-0) emission."
 The results are shown in col., The results are shown in col.
 2 of Table 6.., 2 of Table \ref{mass_nubi}. .
" The total LTE mass of gas in the identified clumps, Mir, can be computed from the CO column density as follows: 1986)), where ]/[ [H5CO] is the abundance ratio of molecular hydrogen to CO, ug=2.72 is the mean molecular weight of the gas, my is the mass of the hydrogen atom, Os is the angular diameter of each clump (deconvolved FWHM, see col."," The total LTE mass of gas in the identified clumps, $M_{\rm LTE}$, can be computed from the $^{13}$ CO column density as follows: ), where $_2$ $^{13}$ CO] is the abundance ratio of molecular hydrogen to $^{13}$ CO, $\mu_{\rm G} = 2.72$ is the mean molecular weight of the gas, $m_{\rm H}$ is the mass of the hydrogen atom, $\Theta_{\rm
S}$ is the angular diameter of each clump (deconvolved FWHM, see col."
 5 of Table 3)) and d is the distance of the source., 5 of Table \ref{param_nubi}) ) and $d$ is the distance of the source.
" Adopting an abundance ratio [H5]/[""CO]=10 (eg. 1986)), [H;]/[CO]- ([H5/[?CO]x([ C]/[PC]) can be computed for each clump, using the values of ['?C]/[?C] presented in col."," Adopting an abundance ratio $_2$ $^{12}$ $] = 10^{4}$ (e.g. ), $_2$ $^{13}$ $_2$ $^{12}$ $\times$ $^{12}$ $^{13}$ C]) can be computed for each clump, using the values of $^{12}$ $^{13}$ C] presented in col."
 13 of Table 3.., 13 of Table \ref{param_nubi}.
 We can thus derive the gas mass of each clump (col., We can thus derive the gas mass of each clump (col.
 3 of Table 6))., 3 of Table \ref{mass_nubi}) ).
" The estimated gas masses can be compared with the masses computed under the assumption of virial equilibrium for an homogeneous sphere, neglecting contributions from the magnetic field and surface pressure: (e.g. MacLaren et al."," The estimated gas masses can be compared with the masses computed under the assumption of virial equilibrium for an homogeneous sphere, neglecting contributions from the magnetic field and surface pressure: (e.g. MacLaren et al."
" 1988), where Δνιρ is the PCco(1- linewidth in kilometers per second and Os is the angular diameter of each clump in arcsec within the intensity contour of the integrated '*CO(1-0) emission (see col."," 1988), where $\Delta{\rm v_{1/2}}$ is the $^{13}$ CO(1-0) linewidth in kilometers per second and $\Theta_{\rm S}$ is the angular diameter of each clump in arcsec within the intensity contour of the integrated $^{13}$ CO(1-0) emission (see col."
 5 of, 5 of
the shock oscillations) due to the strong magnetic field which prevents mass and energy exchange across magnetic field lines.,the shock oscillations) due to the strong magnetic field which prevents mass and energy exchange across magnetic field lines.
" We expect therefore that accretion streams consisting of many (200— 300) different fibrils with different instability periods and random phases would produce no periodic variations of the X-ray emission from shocked plasma, as recently found in the case of TW Hya (?))."," We expect therefore that accretion streams consisting of many $200-300$ ) different fibrils with different instability periods and random phases would produce no periodic variations of the X-ray emission from shocked plasma, as recently found in the case of TW Hya \citealt{Drake2009ApJ}) )."
" As discussed in Sect. 2.4,"," As discussed in Sect. \ref{sect:absorption},"
 the key parameters for understanding the effects of the absorption by the stellar chromosphere on the X-ray emission from shocked accreted plasma are the thickness of the shocked slab πια and the sinking of the slab in the chromosphere to the position κκ at which the ram pressure of the post-shock plasma equals the thermal pressure of the chromosphere., the key parameters for understanding the effects of the absorption by the stellar chromosphere on the X-ray emission from shocked accreted plasma are the thickness of the shocked slab $l_{\rm max}$ and the sinking of the slab in the chromosphere to the position $h_{\rm sink}$ at which the ram pressure of the post-shock plasma equals the thermal pressure of the chromosphere.
" Figure 4 shows that [max is anti-correlated with the density of the accretion stream and with the metal abundance, and is correlated with the stream velocity (see also Eq. 9))."," Figure \ref{fig:absval} shows that $l_{\rm max}$ is anti-correlated with the density of the accretion stream and with the metal abundance, and is correlated with the stream velocity (see also Eq. \ref{eqn:lslab}) )."
" On the other hand, /gink is correlated with both the stream density and velocity, the ram pressure of the post-shock plasma being Pram©Pacei. and the thermal pressure of the chromosphere decreasing with height."," On the other hand, $h_{\rm sink}$ is correlated with both the stream density and velocity, the ram pressure of the post-shock plasma being $P_{\rm ram} \propto \rho_{\rm acc} u_{\rm acc}^2$ and the thermal pressure of the chromosphere decreasing with height."
" Consequently, the shocked slab from high density streams with high metal abundances is expected to hardly emerge from the dense chromospheric layers, where the absorption is strong, because the post-shock zone is rooted deeply in the chromosphere (hsink is large) and the slab is rather thin (/max is small)."," Consequently, the shocked slab from high density streams with high metal abundances is expected to hardly emerge from the dense chromospheric layers, where the absorption is strong, because the post-shock zone is rooted deeply in the chromosphere $h_{\rm sink}$ is large) and the slab is rather thin $l_{\rm max}$ is small)."
" On the other hand, streams with high velocity form extended post-shock zones (being /maxοςu4,,) and the shocked slab may easily emerge above the chromosphere even if it is deeply rooted in the chromosphere, the ram pressure being proportional to the square of velocity."," On the other hand, streams with high velocity form extended post-shock zones (being $l_{\rm
max}\propto u_{\rm acc}^4$ ) and the shocked slab may easily emerge above the chromosphere even if it is deeply rooted in the chromosphere, the ram pressure being proportional to the square of velocity."
" The distribution of emission measure versus temperature EM(T) of the shock-heated plasma is a useful source of information of the plasma components contributing to the X-ray emission and is directly comparable with EM(T) distributions derived from X-ray observations (see, for instance, ?))."," The distribution of emission measure versus temperature $(T)$ of the shock-heated plasma is a useful source of information of the plasma components contributing to the X-ray emission and is directly comparable with $(T)$ distributions derived from X-ray observations (see, for instance, \citealt{Argiroffi2009A&A}) )."
" Figure 5 shows how the time-averaged EM(T) of the post-shock plasma varies as a function of velocity (upper panel), density (middle panel) and metal abundance (lower panel) of the accretion stream."," Figure \ref{fig:emdis}
 shows how the time-averaged $(T)$ of the post-shock plasma varies as a function of velocity (upper panel), density (middle panel), and metal abundance (lower panel) of the accretion stream."
" The case with density nac=107? cm, velocity Uacc=400 km s! and solar abundance Z—1 (red dotted line in Fig. 5))"," The case with density $n_{\rm acc}=10^{12}$ $^{-3}$, velocity $u_{\rm acc}=400$ km $^{-1}$ and solar abundance $\zeta=1$ (red dotted line in Fig. \ref{fig:emdis}) )"
 is the reference in all the panels., is the reference in all the panels.
" In all cases, our model predicts a monolithic distribution of emission measure of the post-shock plasma that covers the entire range of values below the temperature of the shock front."," In all cases, our model predicts a monolithic distribution of emission measure of the post-shock plasma that covers the entire range of values below the temperature of the shock front."
 The ascending part of the EM(T) distribution corresponds to cooled post-shock plasma of the slab and to the transition region between the shocked chromosphere and the hot slab., The ascending part of the $(T)$ distribution corresponds to cooled post-shock plasma of the slab and to the transition region between the shocked chromosphere and the hot slab.
 This characteristic of the EM(T) distribution cannot be reproduced by heuristic models (e.g. ???)) and has been proved to be in agreement with observations (e.g. ?)).," This characteristic of the $(T)$ distribution cannot be reproduced by heuristic models (e.g. \citealt{Lamzin1998ARep, Calvet1998ApJ, Argiroffi2007A&A}) ) and has been proved to be in agreement with observations (e.g. \citealt{Argiroffi2009A&A}) )."
 The total emission measure of plasma above 1 MK and the maximum temperature of the emission measure distribution depend on the velocity of the accretion stream (upper panel in Fig. 5))., The total emission measure of plasma above 1 MK and the maximum temperature of the emission measure distribution depend on the velocity of the accretion stream (upper panel in Fig. \ref{fig:emdis}) ).
" In fact, the temperature of the post-shock plasma depends on the square of the stream velocity (see Eq. 8)),"," In fact, the temperature of the post-shock plasma depends on the square of the stream velocity (see Eq. \ref{eqn:posttemp}) ),"
 leading to a shift of the EM(T) profile towards higher temperature for higher value of acc.," leading to a shift of the $T$ ) profile towards higher temperature for higher value of $u_{\rm
acc}$."
 The thickness of the post- zone strongly increases with the stream velocity (see Fig., The thickness of the post-shock zone strongly increases with the stream velocity (see Fig.
" 4 and Eq. 9)),"," \ref{fig:absval} and Eq. \ref{eqn:lslab}) ),"
 so that its volume and therefore the overall emission measure of the shock-heated plasma increases with Uacc-, so that its volume and therefore the overall emission measure of the shock-heated plasma increases with $u_{\rm acc}$ .
is discussed brielly in 5.,is discussed briefly in \ref{sec:variability}.
 However. the variability must be removed. in order to search for the signature of transiting planets in the lightcurve.," However, the variability must be removed in order to search for the signature of transiting planets in the lightcurve."
 Both the shape and the timescale of these sources of variability. οος significantly from. that of a transit. so hey do not induce contamination in the transit search.," Both the shape and the timescale of these sources of variability differ significantly from that of a transit, so they do not induce contamination in the transit search."
 During our observing campaign. three large Dares occured on AU Mic which were detected in all four filters.," During our observing campaign, three large flares occured on AU Mic which were detected in all four filters."
 During the ]ares. the magnitude of AU Mic increased. sharply within minutes and then decavect slowly to its original value over an hour to a few hours.," During the flares, the magnitude of AU Mic increased sharply within minutes and then decayed slowly to its original value over an hour to a few hours."
 We simply remove the regions of he AU Mic lighteurves which include the three large Hares., We simply remove the regions of the AU Mic lightcurves which include the three large flares.
 ‘Table 1. gives the range of Heliocentric Julian Dates (11.9Ds) which are used to exclude points from the lighteurve near he time of the Hares., Table \ref{tab:flarereg} gives the range of Heliocentric Julian Dates (HJDs) which are used to exclude points from the lightcurve near the time of the flares.
 The cquasi-sinusoidal. starspot) variability occurs. on imeseales of ~5 days and varies in amplitude for the different: filters.," The quasi-sinusoidal, starspot variability occurs on timescales of $\sim 5$ days and varies in amplitude for the different filters."
 The Hipparcos period for this object is 4.8902 davs. however we find the variation occurs on a 4.847 clay timescale (see relsee:variability)).," The Hipparcos period for this object is 4.8902 days, however we find the variation occurs on a 4.847 day timescale (see \\ref{sec:variability}) )."
 Lo remove the. modulation. we experimented with fitting a truncated fourier series to the phase-foleecl lighteurves in cach filter as well as applying incar least-squares fits} to the individual nights of data.," To remove the modulation, we experimented with fitting a truncated fourier series to the phase-folded lightcurves in each filter as well as applying linear least-squares fits to the individual nights of data."
 We note that neither approach is a physical interpretation of the data or an attempt to model the starspots., We note that neither approach is a physical interpretation of the data or an attempt to model the starspots.
 Both echniques produce corrected. lighteurves with similar noise ooperties. thus. for this analysis. we subtract a linear fi rom each night of data[to produce the AU Mic lighteurve or cach filter which i4 removed of intrinisic. variability.," Both techniques produce corrected lightcurves with similar noise properties, thus, for this analysis, we subtract a linear fit from each night of data to produce the AU Mic lightcurve for each filter which is removed of intrinisic variability."
 ‘Table 2. gives the number of measurements in the correctec ighteurve for cach filter. as well as the median. sampling. he rms (weighted by the error. bars) ancl the correlatec noise on a 2-hour timescale. the tvpical planetary. transi duration.," Table \ref{tab:lcinfo} gives the number of measurements in the corrected lightcurve for each filter, as well as the median sampling, the rms (weighted by the error bars) and the correlated noise on a 2-hour timescale, the typical planetary transit duration."
 Lo measure the correlated noise. we calculate the rms of the lishteurve where cach point is replaced. by the average of the points in à 2-hour window around that poin (accounting for edge ellects) (Pontetal.2006).," To measure the correlated noise, we calculate the rms of the lightcurve where each point is replaced by the average of the points in a 2-hour window around that point (accounting for edge effects) \citep{Pont06}."
. Phe rms of the smoothed lighteurve is 1.7 mmae which is higher than what is expected if the data were only. white noise., The rms of the smoothed lightcurve is 1.7 mmag which is higher than what is expected if the data were only white noise.
 Since the transit signals for which we are searching will produce dips in brightness of the same depth in all four filters. we combine the four corrected. lighteurves into a combined final AU Mie lighteurve. shown in Figure 3..," Since the transit signals for which we are searching will produce dips in brightness of the same depth in all four filters, we combine the four corrected lightcurves into a combined final AU Mic lightcurve, shown in Figure \ref{fig:finallc}."
 The combined. intrinsic variabilitv-corrected. lighteurve. is input into the periodic transit. searching. algorithm.," The combined, intrinsic variability-corrected lightcurve is input into the periodic transit searching algorithm."
 The lighteurve contains 17593 points. has à median sampling of 0.35 minutes. and has an rms of 6:8 mamas (see Table 2)).," The lightcurve contains 17593 points, has a median sampling of 0.35 minutes, and has an rms of 6.8 mmag (see Table \ref{tab:lcinfo}) )."
 The final. combined. intrinsic variability-correctecl lighteurve was searched for features that could. have been caused by a transiting planet.," The final, combined, intrinsic variability-corrected lightcurve was searched for features that could have been caused by a transiting planet."
 We performed a systematic search for periodic. square-shaped clips in brightness using the Box-Least-Squares algorithm of Ixováesetal.(2002).," We performed a systematic search for periodic, square-shaped dips in brightness using the Box-Least-Squares algorithm of \citet{Kovacs}."
. We apply the algorithm to the final lightcurve. testing periods between 0.5r15 davs.," We apply the algorithm to the final lightcurve, testing periods between $0.5-15$ days."
 The short. period. limit is set by the Roche radius of a Jupiter mass planet. around AU Mic. and the long period limit is approximately one," The short period limit is set by the Roche radius of a Jupiter mass planet around AU Mic, and the long period limit is approximately one"
swept-up circumstellar density behiud the forward shock.,swept-up circumstellar density behind the forward shock.
 If à cool shell forms because the reverse shoek is radiative. that shell absorbs X-rays below 10 keV aud the observed 1-10 keV eiission must be from the forward shock as discussed for 8N1993J. (Franssonetal.1996).," If a cool shell forms because the reverse shock is radiative, that shell absorbs X-rays below 10 keV and the observed 1-10 keV emission must be from the forward shock as discussed for SN1993J \citep{Fran96}."
. Other than SNLOSTA. [ew supernova have been followed iu the X-ray baud.," Other than SN1987A, few supernova have been followed in the X-ray band."
 The total current pool cousists of «20 objects: of those. ouly four supernovae have data sets consistiug of more than four observatious (SNIO9Y7SIx. SN10993J. SN1998S. SN1999en: Schlegel1905:Imiuler&Lewin2002)).," The total current pool consists of $<$ 20 objects; of those, only four supernovae have data sets consisting of more than four observations (SN1978K, SN1993J, SN1998S, SN1999em; \citealt{S95,IL02}) )."
 SNIOSTSIÍX is oue of a sinall number of extremely luminous. N-ray emitting supernovae Lewin 2002).," SN1978K is one of a small number of extremely luminous, X-ray emitting supernovae \citep{S95,IL02}."
. As X-ray observations are difficult to obtain. each bright supernova contributes valuable knowledge to our uuderstaudiugW. of the X-ray. emission [rom superuovae.," As X-ray observations are difficult to obtain, each bright supernova contributes valuable knowledge to our understanding of the X-ray emission from supernovae."
 The X-ray emission of ΝΟΤΙΑΣ was first. uucovered. in au observation obtained. with theROSAT PSPC iu 1992 (Ryderetal.19903). Cierealter. R93).," The X-ray emission of SN1978K was first uncovered in an observation obtained with the PSPC in 1992 \citep{Ryder93} (hereafter, R93)."
 It had been detected optically during a spectropliotometric survey of H IE regions two years earlier aic reported as a nova at that time., It had been detected optically during a spectrophotometric survey of H II regions two years earlier and reported as a nova \citep{DopRyd90} at that time.
 Π had also been detected as a powerful radio source in 1982 but never followed up (1193)., It had also been detected as a powerful radio source in 1982 but never followed up (R93).
 As part of the investigation spurred by the X-ray. detection. examination of archival optical plates uucovered a light curve that while erude was sufficieut to assign au explosiou date near 1 June 1978 (R93).," As part of the investigation spurred by the X-ray detection, examination of archival optical plates uncovered a light curve that while crude was sufficient to assign an explosion date near 1 June 1978 (R93)."
 To further study SNI97931Ix. observations were obtained during the mid- and LIate-19908 withASCA.Telescope.ROSAT (High Resolution Lager).Array (ATCA). and theObservatory (AAO) covering the X-ray. ultraviolet. optical. and radio bands (Schlegeletal.1999) (hereafter. 509).," To further study SN19798K, observations were obtained during the mid- and late-1990s with, (High Resolution Imager), (ATCA), and the (AAO) covering the X-ray, ultraviolet, optical, and radio bands \citep{Schl99} (hereafter, S99)."
 An early spectrum audROSAT light curve were reported in Petreetal.(1991). aud Schlegel.Petre.&Colbert(1996).. respectively.," An early spectrum and light curve were reported in \cite{Petre94} and \cite{Schl96}, respectively."
 The radio light curve showed behavior typical of a circumstellar interaction. with a rapid turn-on followed by a slow decline at all radio wavelengths (e.g.. Weileretal.1996:: S99).," The radio light curve showed behavior typical of a circumstellar interaction, with a rapid turn-on followed by a slow decline at all radio wavelengths (e.g., \citealt{Weiler96}; S99)."
 Optically. the440 spectra of SNLITSIN showed strong emission lines of the Balmer series as well as He L [O III]. [Ne III]. anc [Fe HI] (93).," Optically, the spectra of SN1978K showed strong emission lines of the Balmer series as well as He I, [O III], [Ne III], and [Fe II] (R93)."
 In the ultraviolet. the Me II doublet ancl [Ne IV] lines were detected using theHST Faint Object Spectrograph (S99).," In the ultraviolet, the Mg II doublet and [Ne IV] lines were detected using the Faint Object Spectrograph (S99)."
 The X-ray light curve was consistent with uo decline throughout the decade (S99)., The X-ray light curve was consistent with no decline throughout the decade (S99).
 The spectra were best-flit with absorbed. siugle-componeut continuum models: πο evidence for line emission existed iu the data (Petreetal.199[.. 899).," The spectra were best-fit with absorbed, single-component continuum models; no evidence for line emission existed in the data \citealt{Petre94}, S99)."
" SNIOTSIx is important because it is amoue the first known X-ray. emittiug supernovae with luminositiesn above 10""uc erg +t.", SN1978K is important because it is among the first known X-ray emitting supernovae with luminosities above $^{39}$ erg $^{-1}$.
 It isn also a nearby object. and may be studiedn inn greater detailn than more distaut superuovae., It is also a nearby object and may be studied in greater detail than more distant supernovae.
 lu this paper. we present the ACIS spectrum of SNLOTSIN aud contrast it with the recent. aud earlier spectra.," In this paper, we present the ACIS spectrum of SN1978K and contrast it with the recent and earlier spectra."
 Throughout we adopt a distance to NGC 1313 of L.1320.11 Mpe as cletermined by Méudezetal. (2002).., Throughout we adopt a distance to NGC 1313 of$\pm$ 0.11 Mpc as determined by \cite{Mendez02}. .
Then the perturbed photou distribution is found using eq.(25)).,Then the perturbed photon distribution is found using \ref{photons}) ).
 As already mentioned in 82.1. the οταν conirpactuess cannot become arbitrarily high but instead saturates at a critical value2B.," As already mentioned in 2.1, the $\gamma$ -ray compactness cannot become arbitrarily high but instead saturates at a critical value$\lcr$."
 When B9cmu only a negligible fraction of the hard photon luninositv will be absorbed. Le. (or equivalently 7!sPS ," When $\linj<<\lcr$ only a negligible fraction of the hard photon luminosity will be absorbed, i.e., $l_{\gamma}^{h'} \rightarrow 0$, or equivalently $l_{\gamma}^{h}\propto \linj$."
We call critical the injected huuinositv required for P=P to occurs, We call critical the injected luminosity required for $l_{\gamma}^{h'}=l_{\gamma}^{s'}$ to occur.
 After this poiut any further increase of the injected. huuinosity will be seen as au increase of the lhuuimnositv in soft photons., After this point any further increase of the injected luminosity will be seen as an increase of the luminosity in soft photons.
 Using eq.(21)). (25)). aud (30)) one derives iim1) -- τομέα rnb P9) where suas=beD|.," Using \ref{hard}) ), \ref{photons}) ), and \ref{elec_final}) ) one derives ) = dx x ) , where $\xmax=b\eg^2/4$."
" Solving the above equation with respect to PH, we findbe."," Solving the above equation with respect to $\lcr$, we find."
..(33) To compare our results with the value nEIN1—3 found by Sk. who have used catastrophic svuchrotrou losses. we have to take the following steps: Thus. our result should be compared to i—— ," To compare our results with the value $l_{\rm cr}^{\rm SK}=3$ found by SK, who have used catastrophic synchrotron losses, we have to take the following steps: Thus, our result should be compared to $\frac{4 l_{\rm cr}^{\rm SK}}{3\eg}$."
The critical compactness as a function of e. for two values of the magnetic field is shown in Fig. 1.., The critical compactness as a function of $\eg$ for two values of the magnetic field is shown in Fig. \ref{lcr1}.
 The two crosses mark the values of SIX. which completely agree with ours iu the lit ofcatastrophic losses.," The two crosses mark the values of SK, which completely agree with ours in the limit of catastrophic losses."
 We note also that eq. (1)).," We note also that eq. \ref{lcr0}) ),"
 which was obtained iby making robust calculations. also agrees with the eq. (32))," which was obtained by making robust calculations, also agrees with the eq. \ref{lcrit_delta}) )"
 iu the lit ofcatastrophic losses., in the limit of catastrophic losses.
 Iu this paragraph we analytically find the critical huninosity compactness after adopting a more realistic expression for the annihilation cross section: ασ where the normalization used is oy=1/3., In this paragraph we analytically find the critical luminosity compactness after adopting a more realistic expression for the annihilation cross section: _0 where the normalization used is $\sigma_0=4/3$.
 The initial equations (7)}-(9)) keep the same form., The initial equations \ref{gammaray}) \ref{elec}) ) keep the same form.
 Ouly the operator of photon losses becomes = (Qs)deo T Following the same steps as those described in 82.2.1. we fud for s=0 dsa4Ql08)m:Pte)CLfoem80 (94). where A—Fe5 and tax€./2.," Only the operator of photon losses becomes = n(2 dx _0 Following the same steps as those described in 2.2.1, we find for $s=0$ ^2 ) = } ), where $A=\frac{2\bar{n} \sigma_0}{b^2 \eg}$ and $\gmax=\eg/2$."
 Equation. (38))ae can be solved iteratively when writing Ες) as a sum of approximations. Le. o(5)9i89?!|if s ," Equation \ref{elec_theta1}) ) can be solved iteratively when writing $n_e'(\gamma)$ as a sum of approximations, i.e., $n_e'(\gamma)=n^{(0)}+n^{(1)}+\cdots$ ."
"Assmimine for. the first* approsimation.H the formJ ηο)=(>4272, we findd the other terms of the seres expausion: =: =Cr: 2. where A=iu4) and |A|«1 fortypical values of the parameters."," Assuming for the first approximation the form $n^{(0)}=C\gamma^{-2}$, we find the other terms of the series expansion: = = where $\lambda=\frac{A}{3}\left(\frac{1}{\gmax^3}-\frac{1}{\gcr^3}\right)$ and $|\lambda|<1$ fortypical values of the parameters."
" Thus. the seres of approximuations couverges andοἱ, cau be found iu closed fora:"," Thus, the series of approximations converges and$n_e'$ can be found in closed form: ) = ="
being photo-ionised/photo-excited. by the intrinsic power law continuum. as found by Sako (20003.,"being photo-ionised/photo-excited by the intrinsic power law continuum, as found by Sako (2000)."
 In all these respects the soft. X-ray emitting eas in lis very similar to that seen in another nearby. and still brighter. Sevfert 2 galaxy NOGCTO (Ixinkhabwala 22002).," In all these respects the soft X-ray emitting gas in is very similar to that seen in another nearby, and still brighter, Seyfert 2 galaxy NGC1068 (Kinkhabwala 2002)."
 Xs those authors sente €out. the οίκος outflow seen in emission in NGCLOGS - and in -- is consistent with that tvpically seen in absorption against the power law continuum in Sevfert 1 galaxies.," As those authors pointed out, the ionised outflow seen in emission in NGC1068 - and in - is consistent with that typically seen in absorption against the power law continuum in Seyfert 1 galaxies."
 The soft rav luminositv in iis of order of the intrinsic (absorption corrected) 0.3-10 keV luminosity., The soft X-ray luminosity in is of order of the intrinsic (absorption corrected) 0.3-10 keV luminosity.
" ]t is interesting to compare that figure with the ""non-varying component of the soft. N-rav. excess found. in the Sevfert 1 galaxies (NGC 4051. Pounds 22004a) and 1110419-577 Maas(Pounds ""1010). where soft luminosities of and of dominant power law component were indicated."," It is interesting to compare that figure with the `non-varying' component of the soft X-ray excess found in the Seyfert 1 galaxies (NGC 4051, Pounds 2004a) and 1H0419-577 (Pounds 2004b), where soft X-ray luminosities of and of the dominant power law component were indicated."
 It was suggested. in those papers tha laree part of the non-varving soft flux might be explained by the same outflow usually seen in absorption., It was suggested in those papers that a large part of the non-varying soft flux might be explained by the same outflow usually seen in absorption.
 For that explanation to hold up it would appear that a bright inner region. shielded. from direct. view in a Sevlert 2 such as3. must have a sullicienthy high velocity dispersion to remain uncdetected in high resolution absorption spectra.," For that explanation to hold up it would appear that a bright inner region, shielded from direct view in a Seyfert 2 such as, must have a sufficiently high velocity dispersion to remain undetected in high resolution absorption spectra."
 Interestingly. Gierlinski and Done outlined an extreme case of such a scenario in proposing an absorption-hascd alternative to the strong soft excess in the luminous Sevfert 1 galaxy P€GI211|577 (Cierlinski and. Done. 2004).," Interestingly, Gierlinski and Done outlined an extreme case of such a scenario in proposing an absorption-based alternative to the strong soft excess in the luminous Seyfert 1 galaxy PG1211+577 (Gierlinski and Done, 2004)."
 Phe blue wings seen on some of the stronger emission lines in the RCS spectrum may be an indication of such a trend to higher velocity gas extending into the unobscured outflow in3., The blue wings seen on some of the stronger emission lines in the RGS spectrum may be an indication of such a trend to higher velocity gas extending into the unobscured outflow in.
. However. significant emission could. also come from a second οίκος emission component out of the line-of-sight to the continuum source (and hence not seen in absorption in Sevíert Is) within the cone of the obscuring torus.," However, significant emission could also come from a second ionised emission component out of the line-of-sight to the continuum source (and hence not seen in absorption in Seyfert 1s) within the cone of the obscuring torus."
 Such a component. might be identified with the ‘equatorial scatterer’ required by optical polarisation studies of Sevfert galaxies (Smith 22002. 2004).," Such a component might be identified with the `equatorial scatterer' required by optical polarisation studies of Seyfert galaxies (Smith 2002, 2004)."
 The results. reported. here are. based. on observations obtained with.XAZAZ-Neiwlon.. an ESA science mission with instruments ancl contributions cirectly funded. by ESA Alember States and the USA (NASA).," The results reported here are based on observations obtained with, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA)."
 Phe authors wish to thank Leicester colleagues for valuable input. the SOC anc SSC teams for organising the oobservations and initial data reduction.," The authors wish to thank Leicester colleagues for valuable input, the SOC and SSC teams for organising the observations and initial data reduction."
 WAP acknowledges the support of a Leverhulme Trust Emeritus Fellowship and WPA of a PPABC research grant., KAP acknowledges the support of a Leverhulme Trust Emeritus Fellowship and KPA of a PPARC research grant.
minute curation to reduce noise and all baselines have been usec.,minute duration to reduce noise and all baselines have been used.
 A number of features are visible in the plots., A number of features are visible in the plots.
 Points to note in the data are: To search for any periodic signals in the data. photometric points from all six epochs were included. anc the Fourier spectrum was determined.," Points to note in the data are: To search for any periodic signals in the data, photometric points from all six epochs were included and the Fourier spectrum was determined."
 This is shown in refPower.poedriun.., This is shown in \\ref{Power_spectrum}.
 The main part of the figureὃν is a cirect Fourier transform of the data., The main part of the figure is a direct Fourier transform of the data.
 Vhe insert is the Fourier. transform of the windowing function., The insert is the Fourier transform of the windowing function.
 The main transform of the data is the time-domain equivalent of an interlerometric dirty map. anc the transform of the windowing function is the equivalent of a dirty beam.," The main transform of the data is the time-domain equivalent of an interferometric dirty map, and the transform of the windowing function is the equivalent of a dirty beam."
 In the same way as one can re-create a clean map from a cirty map and dirty beam. the power spectrum can be cleaned by an iterative subtraction of the window function (Clark 1980: Roberts. Lebar Dreher 1987).," In the same way as one can re-create a clean map from a dirty map and dirty beam, the power spectrum can be cleaned by an iterative subtraction of the window function (Clark 1980; Roberts, Lehar Dreher 1987)."
 The result of this procedure is shown in refeleanpectrum.., The result of this procedure is shown in \\ref{clean_spectrum}.
 To check whether the periods shown in refclean;pectrumearespurious.lhefirstlhrecepochsiwereanalgscedscparatelyj Asthe," To check whether the periods shown in \\ref{clean_spectrum} are spurious, the first three epochs were analysed separately from the last three epochs."
rci, The clean spectra are shown in \\ref{clean_spectrum_split}.
snocorretationbelweenthelwohalveso fthed periodicoscillalions.," As there is no correlation between the two halves of the data sets, we can conclude that there are no stable periodicities and any periodicity shown is due to spurious features or quasi-periodic oscillations."
 A brightness temperature of a [lare can be. caleulated associated with the Lux change from the source., A brightness temperature of a flare can be calculated associated with the flux change from the source.
 A’ black body sphere with radius r. Εαν 5. at a distance 2. assuming a Ravicigh-Jeans tail. would have a brightness temperature of where e is the speed of light. Ay is Boltzmann's constant and 7 is the frequency.," A black body sphere with radius $r$, flux $S_{\nu}$, at a distance $D$, assuming a Rayleigh-Jeans tail, would have a brightness temperature of where $c$ is the speed of light, $k_{\rm B}$ is Boltzmann's constant and $\nu$ is the frequency."
 Therefore a tare of duration. At minutes and flux change of AS mJy has a brightness temperature of where quantities are measured in their subscripted units., Therefore a flare of duration $\Delta t$ minutes and flux change of $\Delta S$ mJy has a brightness temperature of where quantities are measured in their subscripted units.
 For a typical fare: ALJD 50419.56. the Lux increase AS=SO mJy over Af=46 minutes at vy=5 Cllz.," For a typical flare: MJD 50419.56, the flux increase $\Delta S = 80$ mJy over $\Delta t = 46$ minutes at $\nu = 5$ GHz."
 I Cre X-3 djs at a distance of 10 kpe. the brightness temperature is dc1.5101 Ix. This value is tvpical for the majority of Lares. including the large Hare at ALJD 50432.," If Cyg X-3 is at a distance of 10 kpc, the brightness temperature is $T_{\rm b} \geq 1.5
\times 10^{10}$ K. This value is typical for the majority of flares, including the large flare at MJD 50432."
 Phe largest brightness temperature. occurred. at. ΑΙΤΟ 50438.6. if the change within | integration bin is believable.," The largest brightness temperature occurred at MJD 50438.6, if the change within 1 integration bin is believable."
" For this flare the Hux change is AS=10 my over Af=4.0 min. and so the brightness temperature is 71,=2Lot K. There is à maximum brightness temperature. above which inverse Compton. losses become catastrophic."," For this flare the flux change is $\Delta S = 10$ mJy over $\Delta t = 4.0$ min, and so the brightness temperature is $T_{\rm b} \geq 2 \times 10^{11}$ K. There is a maximum brightness temperature above which inverse Compton losses become catastrophic."
 This niximum brightness temperature can be written for galactic SOULCOS a where £p ds the upper frequeneyv of the synchrotron emission in Gllz. D is the Doppler boosting factor of the emission and ez; is the a-dependent. parts of the ratio of Pacholezvk constants e; and e; where," This maximum brightness temperature can be written for galactic sources as where $\nu_{U}$ is the upper frequency of the synchrotron emission in GHz, $\cal{D}$ is the Doppler boosting factor of the emission and $c_{56}^{(\alpha)}$ is the $\alpha$ -dependent parts of the ratio of Pacholczyk constants $c_{5}$ and $c_{6}$ where"
without the problems and assmiuptious simroundins coluposite spectra.,without the problems and assumptions surrounding composite spectra.
" Lastly. Main Sequence and Red Cüant Brauch stars areunderstood. at least relative to hot subdwarts,"," Lastly, Main Sequence and Red Giant Branch stars are, at least relative to hot subdwarfs."
 We have therefore searched. for common propor motion pairs and visual binaries coutainime a hot subdwart component., We have therefore searched for common proper motion pairs and visual binaries containing a hot subdwarf component.
 This is the focus of the present work. where we present our most promising candidates and some very. preliminary ⋞⋯↧↕⊔↸↨⋅↸⋯↕∐⊔↕⊔⋞⊔⋅↸⊔∖∏↕↾∖⊺↕⊔⊔⋞⊔↧∪∖↕↓∪∏↕↸↧results.," This is the focus of the present work, where we present our most promising candidates and some very preliminary results."
 The reader s noteote that; thereere are no definiteFite or* clear eleavocconclusions⊳4 in the worl ispresented here: this anis exploratoryan exploratorycases pilot-tvpe study. which is in part intended to inspire others to explore] ‘alternative methodsuethods to deteriuiniusdetermining the fundamental'‘al parameterspari of these enigniaticgluatic objects.," The reader should note that there are no definite or clear conclusions in the work presented here; this is an exploratory pilot-type study, which is in part intended to inspire others to explore alternative methods to determining the fundamental parameters of these enigmatic objects."
object With. separations. hugeOo enoughC» to show no spectral contanunation.⋅⋅ wuderstancding⋅ visual⋅ binary⋅ aud common proper motion⋅ svstenis containiugDe hot subchwart. will⋅ be more straightforward⋅. than other methods.," With separations large enough to show no spectral contamination, understanding visual binary and common proper motion systems containing hot subdwarf components will be more straightforward than other methods."
 It allow us to use the spectroscopic. parallax to determine. the hot subedwart’s. nass and metallicity.∙∙ (, It allow us to use the spectroscopic parallax to determine the hot subdwarf's mass and metallicity. (
Note- that while we use the terms “spectroscopic parallax” and -photometric. parallax” ..in this⋅ paper. they may be considered as more related to a distance modulus than an aneular measurement.),"Note that while we use the terms “spectroscopic parallax” and “photometric parallax” in this paper, they may be considered as more related to a distance modulus than an angular measurement.)"
 The latter has previously been unknown in field hot subdwiufs (uulike their stellar cluster counterparts). as atmospheric diffusion leads to peculiar abundance patterns (e.g.O'Toole&IIeber 2006).," The latter has previously been unknown in field hot subdwarfs (unlike their stellar cluster counterparts), as atmospheric diffusion leads to peculiar abundance patterns \citep[e.g.][]{OH2006}."
. Aud while pulsations detected iu some sdD stars have allowed the estimation of masses of a haudful of objects Charpinetet⋅al.2008)..∖they rely on sinele-star htt∙ evolutionary models where a mass range is assuned.," And while pulsations detected in some sdB stars have allowed the estimation of masses of a handful of objects \citep[e.g.][]{Charpinet2008}, they rely on single-star evolutionary models where a mass range is assumed."
 Both mass and metallicity are vital iu πιοταςας the evolution of both the RGB progenitor aud the hot subdscarf itsclt., Both mass and metallicity are vital in understanding the evolution of both the RGB progenitor and the hot subdwarf itself.
 As a first step in this project. we lave used the Subdwiuf Database (Osteusen2006) as the source of known objects.," As a first step in this project, we have used the Subdwarf Database \citep{Oestensen2006} as the source of known objects."
 Latter work will use large databases such as the Sloan Digital Sky Survey Data Release, Latter work will use large databases such as the Sloan Digital Sky Survey Data Release.
 After ideutifving hot subeawarf candidates. we then cross-correlated their co-ordinates with proper motion and visual binary databases including. but not luuted to: The Washington Double Star Catalogue: SuperCosmos: UCAC2: NOMADI: aud Tipparcos.," After identifying hot subdwarf candidates, we then cross-correlated their co-ordinates with proper motion and visual binary databases including, but not limited to: The Washington Double Star Catalogue; SuperCosmos; UCAC2; NOMAD1; and Hipparcos."
 Ounce pairs are identified. the goal of the project is to determine parameters of companion star aud the hot aubdgart.," Once pairs are identified, the goal of the project is to determine parameters of companion star and the hot subdwarf."
 There are some limitations aud difficulties with this approach to be considered., There are some limitations and difficulties with this approach to be considered.
 The main concern is that rot subclwarts are in seueral quite distaut aud faiut. which means that their proper motious are often very sinall and/or not well coustrained.," The main concern is that hot subdwarfs are in general quite distant and faint, which means that their proper motions are often very small and/or not well constrained."
 Figure 1 shows he proper motions for the known binary sdB pulsator Feiee I8: the star has a moderately hieh proper motion ⋖∿−≽⋅↱⊐∐∐⊔⋜↧↴∖↴∙∏⋅⋝⋅↴⋝∏↑⋯∐∐∪↑↴⋝↸∖↸∖⋜↧↴∖↴∏⋅↖↽≼∐↴∖↴⊓∐∶↴∙⋯↴∖↴∐↸∖≼↧: ∣ ⋅ ⋅⋅⋅ .roni the other stars iu. the field. around it.," Figure \ref{fig:feige48}
 shows the proper motions for the known binary sdB pulsator Feige 48; the star has a moderately high proper motion $\sim -25$ mas/yr), but cannot be easily distinguished from the other stars in the field around it."
. Iu the worst thoueh. no⋅ proper motions have∖↴ been measured at all.," In the worst cases though, no proper motions have been measured at all."
all Once↽ candidate: pairs: are selected. we follow: two possible: aveuues to determine: their: photometric: parallax ixd therefore their distance: empirical relations involvinel colours (Ihwvlevetal.2002:Ivezie2008) or relatious using spectral iudices (e.g.Cruz&ReidWw )02).," Once candidate pairs are selected, we follow two possible avenues to determine their photometric parallax and therefore their distance: empirical relations involving colours \citep{Hawley2002,Ivezic2008} or relations using spectral indices \citep[e.g.][]{CR2002}."
. The combination of two should lead to a more precise mass estimate., The combination of two should lead to a more precise mass estimate.
 The photometric parallax cau then be combined with the spectroscopic parallax to eive the nass of the subdwurf., The photometric parallax can then be combined with the spectroscopic parallax to give the mass of the subdwarf.
 The spectroscopic component:arallax requires. the subdwarfs» magnitude.. effective⋅⋅ teiiperature and surface. gravitv.," The spectroscopic parallax requires the subdwarf's magnitude, effective temperature and surface gravity."
. The mass is. derived. by matchiug. the two parallaxes., The mass is derived by matching the two parallaxes.
 From our initial⋅⋅⋅ search described⋅ above. we have found these pronusing candidates: 1811777: FBS 2251)373: 9932: aud 116181563.," From our initial search described above, we have found these promising candidates: 1777; FBS 2254+373; 932; and 1618+563."
 Thomas Bauch has also ciscussed the possibility that ο1Lsl-2303 is also a member ofa close pair (these proceccines)8, Thomas Rauch has also discussed the possibility that EC11481-2303 is also a member of a close pair (these proceedings).
 Below we cliscuss the candidates in more detail., Below we discuss the candidates in more detail.
The standard. unified scheme for active galactic nuclei (AGN) attempts to explain the observational appearance of dillerent AGN types mainly. as an inclination effect (Antonucci1993:Urey&Padovani1995).,"The standard unified scheme for active galactic nuclei (AGN) attempts to explain the observational appearance of different AGN types mainly as an inclination effect \citep{antonucci1993,urry1995}."
. The continuum emission. being produced. by the accretion disc of the central. supermassive black hole (SAIBLL) irractiates the surrounding broad line region (BL) and gives rise to broad. optical/UV line emission., The continuum emission being produced by the accretion disc of the central supermassive black hole (SMBH) irradiates the surrounding broad line region (BLR) and gives rise to broad optical/UV line emission.
 Obscuring equatorial dust at larger distances plavs a crucial role in the unified mocel as at higher inclinations it covers the BLR. and thereby clisentaneles twpc-l AGN. which spectroscopically show broad line emission. from twpe-2 objects. whieh do not.," Obscuring equatorial dust at larger distances plays a crucial role in the unified model as at higher inclinations it covers the BLR and thereby disentangles type-1 AGN, which spectroscopically show broad line emission, from type-2 objects, which do not."
 The accretion process onto the SALBIL partly leads to re-ejeetion of matter and supposedly. causes ionised. winds in the polar direction., The accretion process onto the SMBH partly leads to re-ejection of matter and supposedly causes ionised winds in the polar direction.
 The presence of these winds has brought strong support to the unified model as it allowed to indirectly detect hidden BLRs in obscurecl type-2 objects by scattering of BLK. lisht around the torus (Antonucci&AlillerLOS5)., The presence of these winds has brought strong support to the unified model as it allowed to indirectly detect hidden BLRs in obscured type-2 objects by scattering of BLR light around the torus \citep{antonucci1985}.
. Systematic studies in optical spectropolarimetry lead. to a classification of Sevlert galaxies according to the direction of the optical polarisation angle (seee.g.Antonucci1984:Smithetal.2002.2004.anclreferencestherein) that can »' explained in the framework of the unified scheme when inclucing equatorial scattering inside the torus funnel.," Systematic studies in optical spectropolarimetry lead to a classification of Seyfert galaxies according to the direction of the optical polarisation angle \citep[see e.g.][and references therein]{antonucci1984,smith2002,smith2004} that can be explained in the framework of the unified scheme when including equatorial scattering inside the torus funnel."
 The simplest approach of the unified. model assumes hat the rotation axis of the accretion disc is aligned to he axis of the torus and of the outllows., The simplest approach of the unified model assumes that the rotation axis of the accretion disc is aligned to the axis of the torus and of the outflows.
 The alignment of hese components is justified by assuming a symmetric mass ransfer from the inner torus towards the outer aceretion disc as well as by the svmmetric collimation elfect that the orus should have on the polar outflows., The alignment of these components is justified by assuming a symmetric mass transfer from the inner torus towards the outer accretion disc as well as by the symmetric collimation effect that the torus should have on the polar outflows.
 But recent work by, But recent work by
 Ifa statistically significant peak is found in the SLLARC- map within a small search raclius (discussed. below) of a SILADIZS position. we infer that this is the counterpart source and can use its Dux. after deboosting as per etal. (2005). to constrain the source SEL).,"If a statistically significant peak is found in the SHARC-II map within a small search radius (discussed below) of a SHADES position, we infer that this is the counterpart source and can use its flux, after deboosting as per \citet{Coppin}, to constrain the source SED."
 There are. several dillerent approaches to estimating the lux density associated with SLLADES sources if a SUARC-LL source is detected. and they are biased in. cliflerent ways. (," There are several different approaches to estimating the flux density associated with SHADES sources if a SHARC-II source is detected, and they are biased in different ways. ("
1) Aleasuring the Bux density. of cach object at. the pposition will be biased low on average because of the uncertainty in the true source location due to the large beam sizes ancl modest. S/N in the SILXDIZS maps. (,1) Measuring the flux density of each object at the position will be biased low on average because of the uncertainty in the true source location due to the large beam sizes and modest S/N in the SHADES maps. (
2) Measuring the flux density of the brightest pixel within given racius of the pposition will be biased sightlv. high on (e.g. Coppinetal.2005: see discussion. below). (,2) Measuring the flux density of the brightest pixel within given radius of the position will be biased sightly high on (e.g. \citealt{Coppin}; see discussion below). (
3) Measuring the 350pm. flux. density of each object at the precise radio position. if one exists. is less biased. on average than methods I and 2 because of the small positional uncertainty.,"3) Measuring the $350\,\mathrm{\mu m}$ flux density of each object at the precise radio position, if one exists, is less biased on average than methods 1 and 2 because of the small positional uncertainty."
 Not all of the S50p/m sources mapped at 350yam have racio counterparts.," Not all of the $850\,\mathrm{\mu m}$ sources mapped at $350\,\mathrm{\mu m}$ have radio counterparts."
 We therefore choose to measure the Dux densities using method 2 and to correct. for flux boosting ellects. since this measurement can be performed for all of our target. SAIGs in a uniform way.," We therefore choose to measure the flux densities using method 2 and to correct for flux boosting effects, since this measurement can be performed for all of our target SMGs in a uniform way."
 Detections. and measurements from methods 1. 2 and 3 are given in Table 3.. without deboosting. for completeness and inter-comparison.," Detections, and measurements from methods 1, 2 and 3 are given in Table \ref{tab:sharc_fluxes}, without deboosting, for completeness and inter-comparison."
 The deboosted Uuxes are listed in Table 4.., The deboosted fluxes are listed in Table \ref{tab:sharc_photom}.
 In fact we deboost the detected. sources and the other sources in the same way., In fact we deboost the detected sources and the other sources in the same way.
 The first step in either source detection or method 2 is to choose a search radius for companion sources., The first step in either source detection or method 2 is to choose a search radius for companion sources.
 Too large a radius increases the false detection rate and the Dux boost factor., Too large a radius increases the false detection rate and the flux boost factor.
 Too small a radiuscauses sources to be missed., Too small a radiuscauses sources to be missed.
 The rows labeled Px(107) in Table 2 show the percentage chance that an arbitrary aarcesec circle contains a peak., The rows labeled $P_\mathrm{N}(10'')$ in Table \ref{tab:peaks} show the percentage chance that an arbitrary arcsec circle contains a peak.
 Using the negative peaks as a reliable measurement of the noise level in the maps. we infer that a 2.56 peak found. within lOaaresce of a SLLADIES location is 5 per cent likely to be a false positive association. i.c. à 2.5@ peak found within a aaresce radius is a 95 per cent confidence detection.," Using the negative peaks as a reliable measurement of the noise level in the maps, we infer that a $2.5\,\sigma$ peak found within arcsec of a SHADES location is 5 per cent likely to be a false positive association, i.e. a $2.5\,\sigma$ peak found within a arcsec radius is a 95 per cent confidence detection."
 We use Monte Carlo techniques on the actual data to test the number of false identifications mace., We use Monte Carlo techniques on the actual data to test the number of false identifications made.
 An area with a aaresce radius around. cach 850pmi source is masked out so that any real counterpart will not contaminate the test.," An area with a arcsec radius around each $850\,\mathrm{\mu m}$ source is masked out so that any real counterpart will not contaminate the test."
 We select a random position on cach map and search for a peak in the map within the given search radius above the designated S/N threshold., We select a random position on each map and search for a peak in the map within the given search radius above the designated S/N threshold.
 his is repeated 10.000 times over the SUARC-L maps. and we find that using a search radius of 2lOaresec and Z2.5 finds a SELARC-LL peak in 5 per cent of the trials. confirming the conclusion above from Table 2..," This is repeated 10,000 times over the SHARC-II maps, and we find that using a search radius of $\simeq\!10\,\mathrm{arcsec}$ and $\gtrsim 2.5$ finds a SHARC-II peak in 5 per cent of the trials, confirming the conclusion above from Table \ref{tab:peaks}."
 We therefore adopt counterpart search criteria of 2.5e and aarcsec.," We therefore adopt counterpart search criteria of $2.5\,\sigma$ and arcsec."
 We can estimate the fraction of sources for. which a lOaresee search misses the true companion due to »oxitional uncertainty in both the SILADES and SHARC.lb data.," We can estimate the fraction of sources for which a $10\,\mathrm{arcsec}$ search misses the true companion due to positional uncertainty in both the SHADES and SHARC-II data."
 Ivisonetal.(2007). find a one-climensional yositional uncertainty of aaresec for SILADES SMCs » comparing the S50ym determined. positions (beam FWIIM-AIAd.Saresec) with more precise radio positions (svnthesised. beam ENWIIMeIES200800).," \citet{paper3} find a one-dimensional positional uncertainty of arcsec for SHADES SMGs by comparing the $850\,\mathrm{\mu m}$ determined positions (beam $14.8\,\mathrm{arcsec}$ ) with more precise radio positions (synthesised beam $1.3\,\mathrm{arcsec}$ )."
 This is consistent with the theoretical expectation (see equation 2 of Xppencdix D in Ivisonetal.20073) of a~0.6(S/N)ENIM.3 for FWIIM-IA.Saresec and S/N=3.," This is consistent with the theoretical expectation (see equation 2 of Appendix B in \citealt{paper3}) ) of $\sigma\simeq
0.6\,\mathrm{(S/N)^{-1}\,FWHM}\simeq\!3$ for $14.8\,\mathrm{arcsec}$ and $\mathrm{S/N}\simeq3$."
 The positional uncertainty of ΗΛΙΟ observations stems from. the clescope pointing uncertainty. typically aaresec.," The positional uncertainty of SHARC-II observations stems from the telescope pointing uncertainty, typically arcsec."
 Adding hese uncertainties in quacrature ancl converting one-dimensional uncertainties to two climensions and then a racial distance. we expect the true ccounterpart to Lic within a aarcsec circle 93 per cent of the time.," Adding these uncertainties in quadrature and converting one-dimensional uncertainties to two dimensions and then a radial distance, we expect the true counterpart to lie within a arcsec circle 93 per cent of the time."
 Using a lareer search radius would increase the level of Lux boosting without including more sources., Using a larger search radius would increase the level of flux boosting without including more sources.
 Peaks in the SILARC-IL map will be displaced [rom the true counterpart. location bv a distance that scales with the SLIARC-LL beam size (Qaaresec) and the S/N (22.56). as for SILADIZS.," Peaks in the SHARC-II map will be displaced from the true counterpart location by a distance that scales with the SHARC-II beam size arcsec) and the S/N $\geq2.5\,\sigma$ ), as for SHADES."
 Adding this term. we find that of apparent peaks associated with a counterpart will lic within our chosen aaresee circle.," Adding this term, we find that of apparent peaks associated with a counterpart will lie within our chosen arcsec circle."
 Proper correction for the [ux boosting that results [rom picking off peaks in low S/N data requires knowledge of the source count distribution (e.g. Coppinetal. 2005))., Proper correction for the flux boosting that results from picking off peaks in low S/N data requires knowledge of the source count distribution (e.g. \citealt{Coppin}) ).
 Llowever. prior information about the ssource counts is not vet. sulliciently well-constrained.," However, prior information about the source counts is not yet sufficiently well-constrained."
 We have estimated the [ux boosting ellect on this sample of low S/N 350jun- detected and non-detected SMCGs Cs 40) by caleulating the error-weiehted: mean excess flux of method 2 compared to method 3 to be a factor of 21.5.," We have estimated the flux boosting effect on this sample of low S/N $350\,\mathrm{\mu m}$ -detected and non-detected SMGs $\lesssim4\,\sigma$ ) by calculating the error-weighted mean excess flux of method 2 compared to method 3 to be a factor of $\simeq1.5$."
 We therefore divide each 350pm Hux measurement. from method 2 (or from mocoest S/N detections). by 1.5 and. present the deboosted [Duxes in ‘Table 4. for use in fitting SEDs.," We therefore divide each $350\,\mathrm{\mu m}$ flux measurement, from method 2 (or from modest S/N detections), by 1.5 and present the deboosted fluxes in Table \ref{tab:sharc_photom} for use in fitting SEDs."
 For the highest S/N 350 sam SMCGs (245m: Ixovacsetal. 2006)). Dux boosting ellects are negligible and so we do not correct these IEuxes.," For the highest S/N $350\,\mathrm{\mu m}$ -detected SMGs $\gtrsim4.5\,\sigma$; \citealt{Kovacs}) ), flux boosting effects are negligible and so we do not correct these fluxes."
" The 350,42 [lux densities have been constrained for 25 per cent of the SILADIZS. N50,/m sources. with the current work providing 21/31 of the SLLADIES 350p/m Lux density constraints."," The $350\,\mathrm{\mu m}$ flux densities have been constrained for 25 per cent of the SHADES $850\,\mathrm{\mu m}$ sources, with the current work providing 21/31 of the SHADES $350\,\mathrm{\mu m}$ flux density constraints."
 Complementary SUARC-LL observations of SAICGs including several SHADES sources have been performed. by Laurentctal.(2006) and IWovaesctal. (2006)., Complementary SHARC-II observations of SMGs including several SHADES sources have been performed by \citet{Laurent06} and \citet{Kovacs}. .
. Since we include their photometry in the SED fits of seven SILADIZS sources. it is worth describing how Iluxes are measured by these groups.," Since we include their photometry in the SED fits of seven SHADES sources, it is worth describing how fluxes are measured by these groups."
so could not apply exactly the same procedure.,so could not apply exactly the same procedure.
 However. we found it perfectly adequate to fit a quadratic function to the fiducial points around the MSTO for these clusters.," However, we found it perfectly adequate to fit a quadratic function to the fiducial points around the MSTO for these clusters."
 Johnson et al., Johnson et al.
 (1999). provide measurements of (Y{νο and τοος for NGC 1466. 2257. Hodge LI. M92. and M3.," \shortcite{johnson:99} provide measurements of $(V-I)_{{\rm TO}}$ and $V_{0.05}$ for NGC 1466, 2257, Hodge 11, M92, and M3."
 We find our calculations (again. see Table 39) to match their results very closely. which leads us to have confidence in our procedure and our measurements for M5.," We find our calculations (again, see Table \ref{t:vages}) ) to match their results very closely, which leads us to have confidence in our procedure and our measurements for M5."
 We estimate that our determinations of Vipo and Vous are accurate to cO.)5 mag. while those for (VI)po and (V.1loos areaccurate ο better jan 30.01. mag.," We estimate that our determinations of $V_{{\rm TO}}$ and $V_{0.05}$ are accurate to $\pm 0.05$ mag, while those for $(V-I)_{{\rm TO}}$ and $(V-I)_{0.05}$ areaccurate to better than $\pm 0.01$ mag."
 Our measured viues for AV“HBpe. the difference in V between ye level of the HB and the MSTO are listed in Table 3..," Our measured values for $\Delta V_{{\rm TO}}^{{\rm HB}}$, the difference in $V$ between the level of the HB and the MSTO are listed in Table \ref{t:vages}."
 Both of 1986 levels are quite uncertain — the HB because of some intrinsic width. as well as scater due to stellar variability for Reticuum. and le significant weighing to the blue for NGC 1928 and 1939: and 1ο MSTO because t1ο MS is vertical in the turn-off region.," Both of these levels are quite uncertain – the HB because of some intrinsic width, as well as scatter due to stellar variability for Reticulum, and the significant weighting to the blue for NGC 1928 and 1939; and the MSTO because the MS is vertical in the turn-off region."
 We estimate our measurement errors in zM“HLTO to be approximately +0.1 mag., We estimate our measurement errors in $\Delta V_{{\rm TO}}^{{\rm HB}}$ to be approximately $\pm 0.1$ mag.
 The resuts in Table 3. show only a small dispersion among the clusters. even ignoring metallicity effects.," The results in Table \ref{t:vages} show only a small dispersion among the clusters, even ignoring metallicity effects."
 Rosenberg et al., Rosenberg et al.
 (1999) used tw«» sets of theoretically calculated isochrones to provide a calibration Or age asa :unction. of: AV-fey-HB and metallicity., \shortcite{rosenberg} used two sets of theoretically calculated isochrones to provide a calibration for age as a function of $\Delta V_{{\rm TO}}^{{\rm HB}}$ and metallicity.
D. This appears in their Figure 3. wlich shows the significant majority of the 34 Galactic gl»bular clusers in their sample to lie within a narrow band of ~2 Gyr width about mean values of 14.3 Gyr and 14.9 Gyr for the two isochrone sets.," This appears in their Figure 3, which shows the significant majority of the $34$ Galactic globular clusters in their sample to lie within a narrow band of $\sim 2$ Gyr width about mean values of $14.3$ Gyr and $14.9$ Gyr for the two isochrone sets."
 Placing our clusters on this diagram (using the metalliciies listed in Table 4)) shows NGC 1928 and Reticulum to lie witin this band. along with all the reference clusters except Hodge L1.," Placing our clusters on this diagram (using the metallicities listed in Table \ref{t:hages}) ) shows NGC 1928 and Reticulum to lie within this band, along with all the reference clusters except Hodge 11."
 This cluster. along with NGC 1939 fall ~2 Gyr oleer than the upper limit of the Rosenberg et al.," This cluster, along with NGC 1939 fall $\sim 2$ Gyr older than the upper limit of the Rosenberg et al."
 2 Gyr band., $2$ Gyr band.
 In order to combat the uncertainty in AWE) introduced by measuring Vere. Buonanno et al.," In order to combat the uncertainty in $\Delta V_{{\rm TO}}^{{\rm HB}}$ introduced by measuring $V_{{\rm TO}}$, Buonanno et al."
" (1998) introduced a calibration for a similar vertical parameter. A\“HLOL he difference in V between the level of he HB and 15,55."," \shortcite{buonanno} introduced a calibration for a similar vertical parameter, $\Delta V_{0.05}^{{\rm HB}}$ – the difference in $V$ between the level of the HB and $V_{0.05}$."
 Because the MS is sloped at Vous. this measurement is in theory more accurate (although formally. our errors are the same).," Because the MS is sloped at $V_{0.05}$, this measurement is in theory more accurate (although formally, our errors are the same)."
 We herefore calculated AVBE for each cluser for comparison with the Buonanno et al., We therefore calculated $\Delta V_{0.05}^{{\rm HB}}$ for each cluster for comparison with the Buonanno et al.
 calibration., calibration.
 These measurements are also lised in Table 3.., These measurements are also listed in Table \ref{t:vages}.
 The relevant calibration appears in Figure 7(a)-(c) of Buonanno et al..," The relevant calibration appears in Figure 7(a)-(c) of Buonanno et al.,"
 for three isochrone sets., for three isochrone sets.
 We note that this caibration is based on (V.D13 CMDs. however it should provide some indicaion of age homogeneity or otherwise in our cluster sample.," We note that this calibration is based on $(V\,,\,B-V)$ CMDs, however it should provide some indication of age homogeneity or otherwise in our cluster sample."
 Ploting the clusters on this Figure again shows them to lie within a band of width ~2 Gyr. along with 14 Galactic globular clusters from the Buonanno et al.," Plotting the clusters on this Figure again shows them to lie within a band of width $\sim 2$ Gyr, along with $14$ Galactic globular clusters from the Buonanno et al."
 sample., sample.
 The one outlier in our sample is GC 1928. which lies ~2 Gyr older than the band.," The one outlier in our sample is NGC 1928, which lies $\sim 2$ Gyr older than the band."
 For the three isochrone sets. the best fitting mean ages are 14 Gyr. 12 Gyr and 15 Gyr. respectively.," For the three isochrone sets, the best fitting mean ages are $14$ Gyr, $12$ Gyr and $15$ Gyr, respectively."
 Rosenberg et al., Rosenberg et al.
 (1999). provide a calibration for a horizontal dating method in addition to their vertical method calibration., \shortcite{rosenberg} provide a calibration for a horizontal dating method in addition to their vertical method calibration.
 The relevant parameter in this case is 3»5. the difference in V between the MSTO and a point on tye RGB 2.5 mag brightert the MSTO.," The relevant parameter in this case is $\delta_{2.5}$, the difference in $V-I$ between the MSTO and a point on the RGB $2.5$ mag brighter than the MSTO."
 Our measurements of this parameter for all the clus are listed in Table 4.., Our measurements of this parameter for all the clusters are listed in Table \ref{t:hages}. .
 Again. these results show a good dea of internal consistency. suggesting sma| relative age differences.," Again, these results show a good deal of internal consistency, suggesting small relative age differences."
 This is confirmed by consideration of Figure 4 in Rosenberg et al., This is confirmed by consideration of Figure 4 in Rosenberg et al.
 which shows their age calibration including metallicity effects. again for," which shows their age calibration including metallicity effects, again for"
relative to the [O LI] lines.,relative to the [O III] lines.
 Second. the adoption of such a small characteristic velocity width for the BLR. results in a verv low black hole mass.," Second, the adoption of such a small characteristic velocity width for the BLR results in a very low black hole mass."
 The lxaspietal.(2000). value based on the rms spectrum is 7.5x105ALS.," The \citet{kaspietal00} value based on the rms spectrum is $7.5 \times 10^6
M\sun$."
 The black hole mass based on the DG92 line width is 2.0x10?A ," The black hole mass based on the BG92 line width is $2.0 \times
10^9 M\sun$."
The small black hole mass would implv an unreasonably small Eddington luminosity. resulting in an Ecldington ratio of 30. while the high black hole mass gives a more sensible Exldington ratio of 0.1.," The small black hole mass would imply an unreasonably small Eddington luminosity, resulting in an Eddington ratio of 30, while the high black hole mass gives a more sensible Eddington ratio of 0.1."
 Finally. we note that a drawback of the use of the rms specirun to identify the changing part of the line is (hat variations in line proliles due (o seeing differences or instrumental differences [rom observation to observation will cause residuals (see footnote 9 in Ixaspietal. (2000))).," Finally, we note that a drawback of the use of the rms spectrum to identify the changing part of the line is that variations in line profiles due to seeing differences or instrumental differences from observation to observation will cause residuals (see footnote 9 in \citet{kaspietal00}) )."
 The rms spectrum of PG 17042-608 shows signilicant residual emission at the [O III] lines. ten times as strong as the narrow IL? spike in that spectrum (Ilxaspi2002).," The rms spectrum of PG 1704+608 shows significant residual emission at the [O III] lines, ten times as strong as the narrow $\beta$ spike in that spectrum \citep{kaspi02}."
. This suggests that in this object. the rms spectrum is misleading as an inclicator of what part of the IL? line is varving as a response to continuum variations.," This suggests that in this object, the rms spectrum is misleading as an indicator of what part of the $\beta$ line is varying as a response to continuum variations."
" Table 1 also lists for each object the black hole mass. derived using the formula M,= /G. where e=V3/2FWIAL,;. and Πριν=32.9(AL240,/10!eress.5!"" light davs (IXaspiοἱal.2000)."," Table 1 also lists for each object the black hole mass, derived using the formula $_{\bullet} = v^2 {\rm R_{BLR}} / G$ , where $v=\sqrt{3}/2 {\rm FWHM_{H\beta}}$, and ${\rm R_{BLR}} = 
32.9(\lambda{\rm L}_{5100}/10^{44} {\rm ergs\ s^{-1}})^{0.7}$ light days \citep{kaspietal00}."
. The values of μμ were derived from the r* photometry in the SDSS database., The values of $_{5100}$ were derived from the ${\rm r^*}$ photometry in the SDSS database.
 Flixes were converted to Iuminosities using the Schlegeletal.(1993) maps for correcting for galactic absorption. Hl τὸ km + 1|. and qu = 0.5.," Fluxes were converted to luminosities using the \citet{schlegeletal98} maps for correcting for galactic absorption, $_0$ = 75 km $^{-1}$ $^{-1}$, and $_0$ = 0.5."
 Finally. Table 1 also lists log R. à measure of radio loudness.," Finally, Table 1 also lists log R, a measure of radio loudness."
 The FIRST catalog (Becker. and the NVSS catalog (Condonetal.1998) were searched at the position of each object., The FIRST catalog \citep{beckeretal95} and the NVSS catalog \citep{condonetal98} were searched at the position of each object.
 All but one of the objects. SDSS J173348.314-5835651.0. are in regions covered bv one or the other radio survey.," All but one of the objects, SDSS J173348.81+585651.0, are in regions covered by one or the other radio survey."
 Àny source within 30 arcseconds was declared a malch. though we note that the positional accuracy of all these survevs is good enough that a match with a radio core ought to lie within 2 arcseconds.," Any source within 30 arcseconds was declared a match, though we note that the positional accuracy of all these surveys is good enough that a match with a radio core ought to lie within 2 arcseconds."
 Thus. we flag those that have differences between the radio and optical centroid of more than 2 arcseconds will a colon in Table 1.," Thus, we flag those that have differences between the radio and optical centroid of more than 2 arcseconds with a colon in Table 1."
 Rois the ratio of flux density at 5 Gllz to flux density at A2500., R is the ratio of flux density at 5 GHz to flux density at $\lambda$ 2500.
 In computing (his. a spectral slope of -0.3 was used to transform the total observed radio flux density at 1.4 GlIIz. and an optical spectral slope of -1.0 was used to translorm the g magnitude.," In computing this, a spectral slope of -0.3 was used to transform the total observed radio flux density at 1.4 GHz, and an optical spectral slope of -1.0 was used to transform the $^*$ magnitude."
 Figure 1 shows the black hole mass plotted against the [O III] line width for the 115 SDSS low-redshift QSOs., Figure 1 shows the black hole mass plotted against the [O III] line width for the 115 SDSS low-redshift QSOs.
 The seven objects with log I. > 1. the usual criterion lor are plotted as open circles.," The seven objects with log R $>$ 1, the usual criterion for radio-loud, are plotted as open circles."
 The one object unobserved in (he radio is plotted as an open triangle., The one object unobserved in the radio is plotted as an open triangle.
 The remaining LOT objects are plotted as solid squares., The remaining 107 objects are plotted as solid squares.
"algorithm, (Gal-Yametal.2004).","algorithm, \citep{gy+04}."
. Template images for subtraction and reference magnitudes for zeropoint computation were taken from the Sloan Digital Sky Survey (Abazajianetal.2009)., Template images for subtraction and reference magnitudes for zeropoint computation were taken from the Sloan Digital Sky Survey \citep{aaa+09}.
. Near-infrared images were obtained with the Peters Automated Infrared Imaging Telescope (PAIRITEL; Bloometal. and reduced by an automated reduction pipeline.," Near-infrared images were obtained with the Peters Automated Infrared Imaging Telescope (PAIRITEL; \citealt{bsb+06}) ), and reduced by an automated reduction pipeline."
" 2006)),We lack sufficiently deep template images, which are free of light from PTF110fqs, to perform reliable image subtraction."," We lack sufficiently deep template images, which are free of light from 10fqs, to perform reliable image subtraction."
" Thus, we measure the flux from the source in a small circular aperture, removing the sky with a nearby background region, and adopt a systematic error of 0.2 mag in theJ and bands and 0.3 mag in band."," Thus, we measure the flux from the source in a small circular aperture, removing the sky with a nearby background region, and adopt a systematic error of 0.2 mag in the and bands and 0.3 mag in band."
 The values reported in Table 2 have been calibrated against the 2MASS system 2003).., The values reported in Table \ref{tab:grijhk} have been calibrated against the 2MASS system \citep{cwm03}. .
and the duration of the survey (Tsurv) which we take as one year.,and the duration of the survey $T_{\mbox{surv}}$ ) which we take as one year.
" 'To check the likelihood for individual detectors, we construct noise power spectra for one detector at a time (Figure 9))."," To check the likelihood for individual detectors, we construct noise power spectra for one detector at a time (Figure \ref{fig:detnoise}) )."
" To gain intuition on the magnitude of the beam impact on parameters, we use a symmetric Gaussian model for the beam function, b?—exp(—constantx 1?), where the constant describes the width of the beam."," To gain intuition on the magnitude of the beam impact on parameters, we use a symmetric Gaussian model for the beam function, $b^2_l = \exp(-\mbox{constant} \times l^2)$ , where the constant describes the width of the beam."
" Deconvolving a mismatched beam yields a ratio of functions which may be parameterized by the fractional error in theFWHM, which we denote A."," Deconvolving a mismatched beam yields a ratio of functions which may be parameterized by the fractional error in theFWHM, which we denote $\Delta$."
" For small errors, the function ratio in this case is which corresponds to τῇ=1+2Aloga at the multipole defined by bL—a, where the beam function has fallen by a factor a."," For small errors, the function ratio in this case is which corresponds to $r^2_{l_a} = 1 + 2\Delta \log a$ at the multipole defined by $b^2_{l_{a}} = a$, where the beam function has fallen by a factor $a$."
" We include data only for |<Io.91, that is, scales larger than where the beam function has fallen to1%."," We include data only for $l<l_{0.01}$, that is, scales larger than where the beam function has fallen to."
". We fix all parameters except Τις, which we expect to be the most sensitive to errors in the beam."," We fix all parameters except $n_s$, which we expect to be the most sensitive to errors in the beam."
" For an incorrect function, the datadistorted, as depicted in Figure 10 for the beam and noise of a 100 GHz detector."," For an incorrect function, the data, as depicted in Figure \ref{fig:simplelikeslice} for the beam and noise of a 100 GHz detector."
" We can quantify the bias in the likelihood by computing the distorted mean,A, normalized by the error in n, along that slice, given by Figure 11 summarizes the likelihood bias for a detector in each of the channels, when the fitted beam is too small."," We can quantify the bias in the likelihood by computing the distorted mean, normalized by the error in $n_s$ along that slice, given by Figure \ref{fig:simplelikebias} summarizes the likelihood bias for a detector in each of the channels, when the fitted beam is too small."
" The highest frequency channels (545 and 857 GHz) show no bias simply because the errors are so large, and are not displayed."," The highest frequency channels (545 and 857 GHz) show no bias simply because the errors are so large, and are not displayed."
" The pivot point for the family of spectra in our slice is 1~ 570, so the channels (30 and 44 GHz) with larger beams, which weight lower | more strongly, show a negative bias in n,, while the others show a positive bias."," The pivot point for the family of spectra in our slice is $l \sim 570$ , so the channels (30 and 44 GHz) with larger beams, which weight lower $l$ more strongly, show a negative bias in $n_s$, while the others show a positive bias."
" For the three channels with the best noise, the beam requirement is strictest, and the fidelity in the beam required is striking."," For the three channels with the best noise, the beam requirement is strictest, and the fidelity in the beam required is striking."
" For example, at 100 GHz for a single detector, to limit the distortion in the likelihood slice to 0.1c, the FWHM must be known to for a symmetric Gaussian."," For example, at 100 GHz for a single detector, to limit the distortion in the likelihood slice to $0.1\sigma$ , the FWHM must be known to for a symmetric Gaussian."
 Note that our parametric fits (Table 3)) are achieving this precision in all nine channels., Note that our parametric fits (Table \ref{tab:vanillaparameters}) ) are achieving this precision in all nine channels.
" For a general beam, this means that the beam function b? must be known to almost where it has fallen to 1 percent, and better at lower multipoles."," For a general beam, this means that the beam function $b_l^2$ must be known to almost where it has fallen to 1 percent, and better at lower multipoles."
" We can marginalize the likelihood over the errors in the function, using the function ratio r; from Eq. 22.."," We can marginalize the likelihood over the errors in the function, using the function ratio $r_l$ from Eq. \ref{eqn:transrat}."
" Assuming a flat prior on the function, we may approximate the posterior probability distribution of functions with our Monte Carlo ensemble of beam functions."," Assuming a flat prior on the function, we may approximate the posterior probability distribution of functions with our Monte Carlo ensemble of beam functions."
 We findthe fidelity of the reconstructionfor the parametric model is very impressive., We findthe fidelity of the reconstructionfor the parametric model is very impressive.
 Inthe bandsmost, Inthe bandsmost
morphologv ἵ often. found in. wide-angle tail sources 2004).,morphology is often found in wide-angle tail sources .
. We have presented deep. high-resolution. multi-configuration/ VLA imaging of four Η radio —sources: 1193. 02061235. 0755[37 and ALSS4. together with lower-resolution observations of 0326|39 anc a reanalysis of our published. images of 2296.," We have presented deep, high-resolution, multi-configuration VLA imaging of four I radio sources: 193, 0206+35, 0755+37 and 84, together with lower-resolution observations of 0326+39 and a reanalysis of our published images of 296."
 These sources are all examples of “lobed” FRAIL radio galaxies., These sources are all examples of “lobed” I radio galaxies.
 Our results. clisplavecdl as images of total intensity. brightness eracient. degree of polarization. apparent magnetic-Lield direction and spectral index. show common features. as follows.," Our results, displayed as images of total intensity, brightness gradient, degree of polarization, apparent magnetic-field direction and spectral index, show common features, as follows."
 We will present quantitative modelling of the inner jets in later papers., We will present quantitative modelling of the inner jets in later papers.
 We thank Anita Richards for a useful perl script., We thank Anita Richards for a useful perl script.
local ULIRGs (e.g. Howell et al.,local ULIRGs (e.g. Howell et al.
 2010) and z>2 submm-luminous galaxies., 2010) and $>$ 2 submm-luminous galaxies.
" On the other hand, it appears that the UV provides a better estimate (closer to the measured Lin) for the averaged BX/BMs."," On the other hand, it appears that the UV provides a better estimate (closer to the measured $L_{\rm IR}$ ) for the averaged $/$ BMs."
" Finally, it is instructive to look at variations of the obscuration of these UV-selected galaxies."," Finally, it is instructive to look at variations of the obscuration of these UV-selected galaxies."
 For this purpose we examine the bolometric luminosity (defined as the sum of the IR and UV luminosities) as a function of obscuration (approximated by the ratio of IR-to-UV luminosity) for LBGs and BX/BMs., For this purpose we examine the bolometric luminosity (defined as the sum of the IR and UV luminosities) as a function of obscuration (approximated by the ratio of IR-to-UV luminosity) for LBGs and $/$ BMs.
 For comparison we include submm-luminous and UGR-selected z~2 galaxies (from Reddy et al., For comparison we include submm-luminous and UGR-selected $z\sim$ 2 galaxies (from Reddy et al.
 2006 and references therein)., 2006 and references therein).
 The resulting plot is shown in Fig., The resulting plot is shown in Fig.
 5., 5.
 The straight line indicates the correlation found by Reddy et al. (, The straight line indicates the correlation found by Reddy et al. (
"2006, 2010) for z~2 UGR-selected galaxies.","2006, 2010) for $z\sim$ 2 UGR-selected galaxies."
 The averaged z ~3 LBG and the two individually detected ones appear to follow the relation defined for the z~2 galaxies., The averaged $z\sim$ 3 LBG and the two individually detected ones appear to follow the relation defined for the $\sim$ 2 galaxies.
" In terms of luminosities, both averaged LBGs and BX/BMs have similar Luy (few x 10!9L5) but LBGs have higher Lig and thus higher Lrrn/Luv ratio."," In terms of luminosities, both averaged LBGs and $/$ BMs have similar $L_{UV}$ (few $\times$ $^{10}L_{\odot}$ ) but LBGs have higher $L_{\rm IR}$ and thus higher $L_{FIR}/L_{UV}$ ratio."
 Since it is well established that obscuration decreases with increasing redshift (Reddy et al., Since it is well established that obscuration decreases with increasing redshift (Reddy et al.
" 2006, 2010, Adelberger and Steidel 2000), the difference in the Lp:n/Luv ratio must be attributed to different causes."," 2006, 2010, Adelberger and Steidel 2000), the difference in the $L_{FIR}/L_{UV}$ ratio must be attributed to different causes."
" While selection effects are likely to play a role (see section 2.2) we argue that possible differences in morphologies, dust distribution and extent of star-forming regions are also likely to contribute."," While selection effects are likely to play a role (see section 2.2) we argue that possible differences in morphologies, dust distribution and extent of star-forming regions are also likely to contribute."
 Morphological studies of UV-selected z ~2 and z~3 galaxies in the GOODS-N field find few differences between the two samples (Law et al., Morphological studies of UV-selected $z\sim$ 2 and $\sim$ 3 galaxies in the GOODS-N field find few differences between the two samples (Law et al.
" 2007) although, dustier galaxies (as evidenced by E(B-V)) were found to show more nebulous UV morphologies."," 2007) although, dustier galaxies (as evidenced by E(B-V)) were found to show more nebulous UV morphologies."
" Finally, since MIPS-detected LBGs have ULIRG-like luminosities (section 3.1) it is possible that their UV and IR emission originates in different regions (as observed in local ULIRGs e.g. Wang et al 2004, see also Huang et al."," Finally, since MIPS-detected LBGs have ULIRG-like luminosities (section 3.1) it is possible that their UV and IR emission originates in different regions (as observed in local ULIRGs e.g. Wang et al 2004, see also Huang et al."
 2007) and thus could account for the higher Lr5/Luv ratio observed., 2007) and thus could account for the higher $L_{IR}/L_{UV}$ ratio observed.
 SPIRE has been developed by a consortium of institutes led by Cardiff Univ. (, SPIRE has been developed by a consortium of institutes led by Cardiff Univ. (
UK) and including Univ.,UK) and including Univ.
" Lethbridge (Canada); NAOC (China); CEA, LAM (France); IFSI, Univ."," Lethbridge (Canada); NAOC (China); CEA, LAM (France); IFSI, Univ."
" Padua (Italy); IAC (Spain); Stockholm Observatory (Sweden); Imperial College London, RAL, UCL-MSSL, UKATC, Univ."," Padua (Italy); IAC (Spain); Stockholm Observatory (Sweden); Imperial College London, RAL, UCL-MSSL, UKATC, Univ."
" Sussex (UK); Caltech, JPL, NHSC, Univ."," Sussex (UK); Caltech, JPL, NHSC, Univ."
 Colorado (USA)., Colorado (USA).
" This development has been supported by national funding agencies: CSA (Canada); NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain); SNSB (Sweden); STFC (UK); and NASA (USA)."," This development has been supported by national funding agencies: CSA (Canada); NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain); SNSB (Sweden); STFC (UK); and NASA (USA)."
 The data presented in this paper will be released through the Herschel Database in Marseille HeDaM (hedam.oamp.fr/HerMES), The data presented in this paper will be released through the Herschel Database in Marseille HeDaM $/$ HerMES)
calculated the scattered liebt using the same results of lvdvodvuamic ealeulatious as Naravananetal.(2006).,calculated the scattered light using the same results of hydrodynamic calculations as \citet{Narayanan06}.
 They have shown that the vertically iutegrated surface density and the scattered light have no correlation. since strong turbulence induced by a gravitational instability produces blobs iu the upper laver of the disk. aud the scattered light is stronely affected by them.," They have shown that the vertically integrated surface density and the scattered light have no correlation, since strong turbulence induced by a gravitational instability produces blobs in the upper layer of the disk, and the scattered light is strongly affected by them."
 Iu this paper. we consider a laminar disk where the vertical hivdrostatie balance is achieved.," In this paper, we consider a laminar disk where the vertical hydrostatic balance is achieved."
 Therefore. our model is not useful in a stronely turbulent disk where the Toomre’s Q-paraincter (defined in the later section) becomes smaller than ~1.52. when nou-axisviunetric eravitational imstabilitv causes the disk to be in a turbulent state (e.g. Mejíaetal.(2005):Boley (2006))).," Therefore, our model is not useful in a strongly turbulent disk where the Toomre's $Q$ -parameter (defined in the later section) becomes smaller than $\sim1.5-2$, when non-axisymmetric gravitational instability causes the disk to be in a turbulent state (e.g., \citet{Mejia05,Boley06}) )."
 The paper is coustructed as follows., The paper is constructed as follows.
 Tn Section 2.. we describe the disk model aud the methods of calculations.," In Section \ref{sec:model}, we describe the disk model and the methods of calculations."
 Iu Section 3.. we show the results of caleulatious.," In Section \ref{sec:obsimage}, we show the results of calculations."
 Iu this section. we also describe the features that appear in the NIR scattered Πο tage and are unique to scl-eravitating disks.," In this section, we also describe the features that appear in the NIR scattered light image and are unique to self-gravitating disks."
 Iu Section [. we discuss the observational implications.," In Section \ref{sec:diskmass}, we discuss the observational implications."
 Section 5 is for sununary and further discussions., Section \ref{sec:discuss} is for summary and further discussions.
 Iu this section. we describe the basic setup of our model.," In this section, we describe the basic setup of our model."
 We cousider a model where the outer disk has a gapped structure., We consider a model where the outer disk has a gapped structure.
 We consider a exlindrical coordinate svsteni (roo.:) with the origin at the central star.," We consider a cylindrical coordinate system $(r,\phi,z)$ with the origin at the central star."
 As a simple model that features the effects of sclteravity. we ASSL an axisvnumuetrie disk with a gap: where foohance IX a paraincter that coutrols the overall Surface deusity of the disk aud αμ). is a fuuction that determines the shape of the gap.," As a simple model that features the effects of self-gravity, we assume an axisymmetric disk with a gap: where $f_{\rm enhance}$ is a parameter that controls the overall surface density of the disk and $f_{\rm gap}(r)$ is a function that determines the shape of the gap."
 We vary the disk mass bv chaneime άμμος from 1 to 8.15. but in the subsequent section. we expecially consider the case of Foubauccl and 6.6 as illustrative examples.," We vary the disk mass by changing $f_{\rm enhance}$ from $1$ to $8.15$, but in the subsequent section, we especially consider the case of $f_{\rm enhance}=1$ and $6.6$ as illustrative examples."
 The case of fubanee=6.6 is rather extreme. since the typical Q- value in this case is close to 2. where turbulence may take place.," The case of $f_{\rm enhance}=6.6$ is rather extreme, since the typical $Q$ -value in this case is close to $2$, where turbulence may take place."
 However. the effects of selt-exavitv ou the vertical structure is pronunent in this extreme case.," However, the effects of self-gravity on the vertical structure is prominent in this extreme case."
" We assuue a Gaussian eap profile. with à=0.95. ra,=I00AU. and Aya,=30AU."," We assume a Gaussian gap profile, with $\alpha=0.95$, $r_{\rm gap}=100\mathrm{AU}$, and $\Delta_{\rm gap}=30\mathrm{AU}$."
" The parauneters o. τραμ. aud Aga, paraucterize the depth. location aud width of the eap. respectively,"," The parameters $\alpha$, $r_{\rm gap}$, and $\Delta_{\rm gap}$ parameterize the depth, location and width of the gap, respectively."
 Iu this paper. iu order to make the problem simple and tractable. we make a simple assuniption for the temperature of the disk.," In this paper, in order to make the problem simple and tractable, we make a simple assumption for the temperature of the disk."
 We assume αᾱ vertically isothermal disk where the temperature T(r) is giveu by Although we do not question ourselves about the origin of such surface density and temperature structure in detail. we conunent here that it is reported that a massive planet emedded d a disk causes a deep gap iu the surface density profile. while the temperature does nof vary as siguificautly as the surface density (D'Angeloetal.2003).," We assume a vertically isothermal disk where the temperature $T(r)$ is given by Although we do not question ourselves about the origin of such surface density and temperature structure in detail, we comment here that it is reported that a massive planet emedded in a disk causes a deep gap in the surface density profile, while the temperature does not vary as significantly as the surface density \citep{DAngelo03}."
. By varving the overall surface deusitv bv foshance While keeping the temperature profile fixed. we vary the disk Toommre’s Q-paramcter. Q=Oc/xGYX.," By varying the overall surface density by $f_{\rm enhance}$ while keeping the temperature profile fixed, we vary the disk Toomre's $Q$ -parameter, $Q = \Omega c/\pi G \Sigma$."
 The Q-piriuueter is also varied in the racial direction due to the gap., The $Q$ -parameter is also varied in the radial direction due to the gap.
 Figure Al shows one example of the disk model., Figure \ref{fig:diskmodel} shows one example of the disk model.
 Iu this figure. we show the profile of surface density aud Q-value with Fonhance=1.," In this figure, we show the profile of surface density and $Q$ -value with $f_{\rm enhance}=1$."
 From the surface density aud the temperature profile. we obtain the vertical structure of the disk.," From the surface density and the temperature profile, we obtain the vertical structure of the disk."
 We assume that the livcrostatic equilibrimim is reached and that the disk is geometrically thin., We assume that the hydrostatic equilibrium is reached and that the disk is geometrically thin.
 We solve where p=τρ is the pressure and Or)=(GM.pyle is the Ieplerian angular frequency., We solve where $p=c^2 \rho$ is the pressure and $\Omega(r)=(GM_{\ast}/r^3)^{1/2}$ is the Keplerian angular frequency.
 The sound speed ο) is obtained from the temperature profile given by equation bv Iu equation (1).. cq denotes the eravitational poteutial of the disk. which is determined by Poisson equation We solve equations and iteratively to fud a consistent solution.," The sound speed $c(r)$ is obtained from the temperature profile given by equation by In equation , $\psi_{\rm d}$ denotes the gravitational potential of the disk, which is determined by Poisson equation We solve equations and iteratively to find a consistent solution."
 If the variation of the surface density is slow in the radial direction. the local approach bv DPaczvuski(1978). vields a reasonable model.," If the variation of the surface density is slow in the radial direction, the local approach by \citet{Pac78}
 yields a reasonable model."
 However. in the model preseuted in this paper. such variation iav be rapid due to the presence of the gap. so it is necessary to solve full Poisson equation.," However, in the model presented in this paper, such variation may be rapid due to the presence of the gap, so it is necessary to solve full Poisson equation."
 Once we obtain the deusitv structure of the disk. we solve the equatious of radiative transfer to obtain the observed image.," Once we obtain the density structure of the disk, we solve the equations of radiative transfer to obtain the observed image."
 We consider the thermal emission of the disk and the Πο from the ceutral star scattered by the disk surface., We consider the thermal emission of the disk and the light from the central star scattered by the disk surface.
 The former is iuportaut in the observations in πα waveleneth aud the latter is important in the NIR imaging observations., The former is important in the observations in sub-mm wavelength and the latter is important in the NIR imaging observations.
 For simplicity. we assume that the dust aud gas are welbliuixed aud the optical properties of the dust particles do uot vary in the disk.," For simplicity, we assume that the dust and gas are well-mixed and the optical properties of the dust particles do not vary in the disk."
 We also assume that the disk is face-on., We also assume that the disk is face-on.
 We follow the approach by D'Alessioetal.(1999). iu calculating the thermal emission aud the scattered light., We follow the approach by \citet{DAlessio99} in calculating the thermal emission and the scattered light.
 The thermal emission may be caleulated by integrating where Z is the coordinate along the line of sight. & is the absorption cross section por unit gas mass. BCL) is the Planck function. aud 7 is the optical depth alone the line of sight. which is calculated by where y is the total extinction Gucliding absorption aud scattering) cross section per unit eas mass.," The thermal emission may be calculated by integrating where $Z$ is the coordinate along the line of sight, $\kappa$ is the absorption cross section per unit gas mass, $B(T)$ is the Planck function, and $\tau$ is the optical depth along the line of sight, which is calculated by where $\chi$ is the total extinction (including absorption and scattering) cross section per unit gas mass."
 Iu calculating the scattered light. we assmme a single aud isotropic scattering.," In calculating the scattered light, we assume a single and isotropic scattering."
 Then. thescattered light cau be calculated by iutegratiug," Then, thescattered light can be calculated by integrating"
to the N-rav source are clearly visible in each baud.,to the X-ray source are clearly visible in each band.
 The nebula appears exteuded along the East-West direction., The nebula appears extended along the East-West direction.
" The ellipse drawn ou the figure has a major diameter of 2.07 aud a minor diameter of 1.1"" corresponding to a physical size of 50 pe by 35 pe at a distance of 5.1 Alpe (deVaucouleurs1979).", The ellipse drawn on the figure has a major diameter of $2.0\arcsec$ and a minor diameter of $1.4\arcsec$ corresponding to a physical size of 50 pc by 35 pc at a distance of 5.1 Mpc \citep{deVaucouleurs79}.
". We performed aperture photometry on the N-rav source counterparts. extracting counts from a circle with a radius of 0.15"","," We performed aperture photometry on the X-ray source counterparts, extracting counts from a circle with a radius of $0.15\arcsec$."
 We use the WFPC2 FSIIN image. rather than the ACS/WTEC one. since the source may )o variable and the three WEPC2 images were obtained over a span of less than 5 hours.," We use the WFPC2 F814W image, rather than the ACS/WFC one, since the source may be variable and the three WFPC2 images were obtained over a span of less than 5 hours."
 We subtracted the jebular flux. scaled. by area. from the point source flux.," We subtracted the nebular flux, scaled by area, from the point source flux."
" If. instead. we use a backeround region exterual to the jebula. rretain the uebular contribution within the 115"" extraction circle. the FUV counts iucrease by aud the BVI couuts ierease by ~30%."," If, instead, we use a background region external to the nebula, retain the nebular contribution within the $0.15\arcsec$ extraction circle, the FUV counts increase by and the BVI counts increase by $\sim 30\%$."
 We relied ou the package iu IRAF for the aperture correction or the ACS data., We relied on the package in IRAF for the aperture correction for the ACS data.
" For the WEPC2. we determined heaperature correction (from a radius of 0.15"" to 1.07) using the star USNO D1.0 1501-0283181 also ocated on the planetary camera chip."," For the WFPC2, we determined theaperature correction (from a radius of $0.15\arcsec$ to $1.0\arcsec$ ) using the star USNO B1.0 1501-0283184 also located on the planetary camera chip."
 The same star was also used to check our photometry in the D aud I vans., The same star was also used to check our photometry in the B and I bands.
 Nebular-backerouud subtracted count rates for he point source for cach continua filter are given in Table L., Nebular-background subtracted count rates for the point source for each continuum filter are given in Table \ref{rates}.
" We estimate a error on the source count rates, primarily duc to uncertainties iu the backerouud subtraction (due to the presence of the nebula)."," We estimate a error on the source count rates, primarily due to uncertainties in the background subtraction (due to the presence of the nebula)."
 The backeround-subtracted count rate in the FLIOLP filter for the nebula. excluding the poiut source. is 0.78 cfs. For the point source. we find equivalent iiagnuitudes of B=22.73 and V=22.70.," The background-subtracted count rate in the F140LP filter for the nebula, excluding the point source, is 0.78 c/s. For the point source, we find equivalent magnitudes of $B=22.73$ and $V=22.70$."
 The V maguitucde is iu good agreement with V.—22.61 ax quoted by Blair.Fesen.&Schlegel (2001).. while our Danaguitude is somewhat brighter than their value of B=23.10.," The V magnitude is in good agreement with $V=22.64$ as quoted by \citet{Blair01}, while our B-magnitude is somewhat brighter than their value of $B=23.10$."
 We imodeled the iutrimsic source flux as a power-law. FyxΑ and used the command in the svuphot package to relate the measured count rates to intrinsic source fluxes.," We modeled the intrinsic source flux as a power-law, $F_{\lambda} \propto \lambda^{-2}$ and used the command in the synphot package to relate the measured count rates to intrinsic source fluxes."
 We note that the iferrec fluxes change by less than if we chanee the power-law spectral iudex to FyXA|., We note that the inferred fluxes change by less than if we change the power-law spectral index to $F_{\lambda} \propto \lambda^{-4}$.
 The filter pivot wavelcnetlis are 1533 for FLLOLP. 1557 for PFI50W. 5113 for F555W. and 7996 for FallW. The Galactic interstellar extinction along he line of sight to AIFLG is Ay=LE nag (Schlegeletal. 1998).," The filter pivot wavelengths are 1533 for F140LP, 4557 for F450W, 5443 for F555W, and 7996 for F814W. The Galactic interstellar extinction along the line of sight to MF16 is $A_V = 1.14$ mag \citep{Schlegel98}."
. Reported values of the additional. intrinsic extinction. inostlv found by measuring decrements οσοι Balmer lines iu the nebular spectrüuni vary roni 0.2 to 0.66 maenitudes (Blair.Fesen.&Schlegel2001:Roberts&Colbert2003:Abolmasovetal. 2008).," Reported values of the additional, intrinsic extinction, mostly found by measuring decrements between Balmer lines in the nebular spectrum, vary from 0.2 to 0.66 magnitudes \citep{Blair01,Roberts03,Abolmasov08}."
. Following Abolmasovotal.(2008). we adopt au extinction of Ay=1.51 mae aud the extinction. curve of Cardelli.Clayton.&Mathis(1989).," Following \citet{Abolmasov08}, we adopt an extinction of $A_V = 1.54$ mag and the extinction curve of \citet{Cardelli89}."
.. The point source fluxes are given in Table E. for ety=1.5 Linag., The point source fluxes are given in Table \ref{rates} for $A_V = 1.54$ mag.
 The lower extinction value discussed by Aboluasoyetal.(2008).. Ay=LOM mas. decreases the intrinsic flux bv im the FUV. in D. in V. aud in L The lower extinction may be valid for the point source. if source prefereutiallv ionizes the nearby regions of the nebula.," The lower extinction value discussed by \citet{Abolmasov08}, $A_V = 1.34$ mag, decreases the intrinsic flux by in the FUV, in B, in V, and in I. The lower extinction may be valid for the point source, if source preferentially ionizes the nearby regions of the nebula."
" EYopting au extinction a the high end of the reported values. ον=Ls mae (Dhuür.Feseu.&Schlegel2001:Roberts&Colbert 2003).. increases the tutrinsic flux by in the FUV. iu D. in V. and in 1. FUV imaging of ME 16 reveals both a poiut-like source coincident with the X-rav source and a surroundiug nebula with an augular extent of 2.0"" by 1.1""."," Adopting an extinction at the high end of the reported values, $A_V = 1.8$ mag \citep{Blair01,Roberts03}, increases the intrinsic flux by in the FUV, in B, in V, and in I. FUV imaging of MF 16 reveals both a point-like source coincident with the X-ray source and a surrounding nebula with an angular extent of $2.0\arcsec$ by $1.4\arcsec$."
 The point source has a flux of 5&LOeresbem2A+ and the nebula has a flux of 1.6«10.herestem7A1H. both quoted at 1533À.. correspondiug to the pivot waveleneth of the FILOLP filter and assunüug an extinction of Ay=151.," The point source has a flux of $5 \times 10^{-16} \rm \, erg \, s^{-1} \, cm^{-2} \, \AA^{-1}$ and the nebula has a flux of $1.6 \times 10^{-15} \rm \, erg \, s^{-1} \, cm^{-2} \, \AA^{-1}$, both quoted at 1533, corresponding to the pivot wavelength of the F140LP filter and assuming an extinction of $A_V = 1.54$."
" The monochromatic Iuuinositv estimate is ALY=LO«10eyes.+, suggesting the presence of an ultrahuuineus UV source (ULUV)."," The monochromatic luminosity estimate is $\lambda L_{\lambda} = 1.0 \times 10^{40} \rm \, erg \, s^{-1}$, suggesting the presence of an ultraluminous UV source (ULUV)."
 The measured ftux is close to that predicted by Abolinasovetal.(2008) near 1000A., The measured flux is close to that predicted by \citet{Abolmasov08} near 1000.
. Figure 2. shows the spectral cucrey distribution of the point-like source at the ceuter of AIF 16 using the data from the FUV (FLLOLP). D (Finsow). V (F555W). aud I (FstWW) bauds.," Figure \ref{mspec} shows the spectral energy distribution of the point-like source at the center of MF 16 using the data from the FUV (F140LP), B (F450W), V (F555W), and I (F814W) bands."
 The dereddeued fixes are from Table 1.., The dereddened fluxes are from Table \ref{rates}. .
 We note that the BVI data are quasi-inmltauecous. obtained within 5 hours. while the FUV data were obtained 7 vears later.," We note that the BVI data are quasi-simultaneous, obtained within 5 hours, while the FUV data were obtained 7 years later."
 Thus. the nonu- of the observations should be of couceru in fitting the data.," Thus, the non-simultaneity of the observations should be of concern in fitting the data."
 ILowever. FEridiikssonuetal.(2008) ," However, \citet{Fridriksson08} "
2000.,.
. By using two-dimensional radiation hvdrocdynamic simulations. Ohsugaetal.(2005) demonstrated for. the fist time that quasi-steady super-Edcineton disc accretion is possible. because both the radiation [field ancl the mass accretion [ow around a DII are non-spherical and the large. number of photons generated. in the clisc is swallowed by the BIL without being racdiated away due to photon-trapping (see Ohsuga&Alineshige 2007)).," By using two-dimensional radiation hydrodynamic simulations, \citet{Oh05} demonstrated for the fist time that quasi-steady super-Eddington disc accretion is possible, because both the radiation field and the mass accretion flow around a BH are non-spherical and the large number of photons generated in the disc is swallowed by the BH without being radiated away due to photon-trapping (see \citealt{OM07}) )."
 This conclusion has recently been confirmed by two-dimensional radiation magnetohycrodynamic simulations (Ohsugaetal. 2009)., This conclusion has recently been confirmed by two-dimensional radiation magnetohydrodynamic simulations \citep{Oh09}.
. Llowever. it is still unclear whether super-Exddington accretion can significantly contribute to the formation of SMDIIS. because the radiation pressure and/or disc wind associated with super-Edcineton aceretion might play an important role in the eas accretion process in host. galaxies (c.g. Silk&Rees1998:Fabian1999:Ilxing2003:Ohsuga 2007)).," However, it is still unclear whether super-Eddington accretion can significantly contribute to the formation of SMBHs, because the radiation pressure and/or disc wind associated with super-Eddington accretion might play an important role in the gas accretion process in host galaxies (e.g., \citealt{SR98,Fa99,Ki03,Oh07}) )."
 In the local Universe. we have discovered. a class of AGNs with a higher Edcington ratio among the ACN population. Le. Narrow Line Sevfert 1 galaxies (NLSIs).," In the local Universe, we have discovered a class of AGNs with a higher Eddington ratio among the AGN population, i.e., Narrow Line Seyfert 1 galaxies (NLS1s)."
 They have characteristic properties such as narrow Balmer ines (e.g.. Osterbrock&Posee1985:Posse 2000)). strong soft. N-rav excess (e.g.. Pounds.Done.&OsborneBolleretal.1996. 2003)) and rapid. variability (c.g.. Otani.al. 20043).," They have characteristic properties such as narrow Balmer lines (e.g., \citealt{OP85,Po00}) ), strong soft X-ray excess (e.g., \citealt{Po95,Bo96,Bo03}) ) and rapid variability (e.g., \citealt{Ot96,Le99,Bo02,Ga04}) )."
 These properties indicate that they have a small Mack hole mass ancl high Eddington ratio (e.g. Brandt&Boller1998:Llavashicla2000:Mineshigeοἱal. 2000)).," These properties indicate that they have a small black hole mass and high Eddington ratio (e.g., \citealt{BB98,Ha00,Mi00}) )."
 The observed. bolometric luminosity of NLS1s saturates at a few imes the Eddington luminositw (e.g. Collin&Ixawaguchiraetal. 2009)). which is consistent with predictions of 10 super-Leleington accretion disc (e.g.. \Watarai.Fukue. 2009)).," The observed bolometric luminosity of NLS1s saturates at a few times the Eddington luminosity (e.g., \citealt{CK04,Co06,Mu08,La09}) ), which is consistent with predictions of the super-Eddington accretion disc (e.g., \citealt{Wa00,Oh05,Oh09}) )."
 In acdition. 1e star formation rate of NLSI hosts is higher than that of BLSIs (Sanietal. 20000).," In addition, the star formation rate of NLS1 hosts is higher than that of BLS1s \citealt{Sa10}) )."
 Thus. NLSIs may be the early hase of rapid. BLL growth via super-Edcdington accretion (c... Mathur2000:Kawaguchietal. 2004)).," Thus, NLS1s may be the early phase of rapid BH growth via super-Eddington accretion (e.g., \citealt{Ma00,Ka04}) )."
 In the high-z universe. we may speculate that super-Edcineton AGNs are a more dominant population of AGNs because the lifetime of SALBL growth is closer to the cosmic age if the formation oE[SMDIIS is mainly a gas accretion process.," In the $z$ universe, we may speculate that super-Eddington AGNs are a more dominant population of AGNs because the lifetime of SMBH growth is closer to the cosmic age if the formation of SMBHs is mainly a gas accretion process."
 To confirm this. it is important to determine whether the fraction of super-Ecclington ACGNs increases with redshift.," To confirm this, it is important to determine whether the fraction of super-Eddington AGNs increases with redshift."
 Some observations to date have indicated that the average Exldington ratio increases slightly with redshift (o... Ixollmeieretal.2006:Shenetal.2008:Willott 2010)).," Some observations to date have indicated that the average Eddington ratio increases slightly with redshift (e.g., \citealt{Ko06,Sh08,Wi10}) )."
 However. we have not observed super-Ecddington XGNs among high-z quasars (e.g.. Mebure&Dunlop2004:Shenetal. 2008)) and ACGNs in high-z massive galaxies (0.8... Yamadaetal. 20093).," However, we have not observed super-Eddington AGNs among $z$ quasars (e.g., \citealt{MD04,Sh08}) ) and AGNs in $z$ massive galaxies (e.g., \citealt{Ya09}) )."
 Thus. it ds essential to [find candidate super-Edcineton ACGNs before starting detailed observations.," Thus, it is essential to find candidate super-Eddington AGNs before starting detailed observations."
 According to unification models of XGNs. the accretion disc is surrounded: by a custy structure (ο. Antonucci1993:Elitzur 2008)).," According to unification models of AGNs, the accretion disc is surrounded by a dusty structure (e.g., \citealt{An93,El08}) )."
 A significant. fraction of the emitted ultraviolet. ancl optical radiation of the accretion disc is absorbed by the dust and is re-emittecl at LR wavelengths., A significant fraction of the emitted ultraviolet and optical radiation of the accretion disc is absorbed by the dust and is re-emitted at IR wavelengths.
 In particular. the hot dust. ds. directly heated. by. the central engine and. produces near-LR (NLR) emission (e.g... lüeke1978:Llaasctal. 2003)).," In particular, the hot dust is directly heated by the central engine and produces near-IR (NIR) emission (e.g., \citealt{Ri78,Ha03}) )."
 Since it has been reported that the polar angle. dependence of the radiation Hux for super-IExidington accretion flow deviates from that for sub-Iddington Low (e.g. Pukue2000:Ohsugaetal.2005:Wataraictal.2005/— hereafter. N05). the reprocessed. LR emission may be useful for exploring candidates of super-Ecelington aceretion Dow in AGNs.," Since it has been reported that the polar angle dependence of the radiation flux for super-Eddington accretion flow deviates from that for sub-Eddington flow (e.g., \citealt{Fu00,Oh05,Wa05} hereafter W05), the reprocessed IR emission may be useful for exploring candidates of super-Eddington accretion flow in AGNs."
 In this paper. we investigate the properties of Ht (as well as NI) emission from the inner edge of the dusty torus. emploving radiative lux from the dise aceretion flows by W05. in which the radiation Iuxes of sub- and super-IEcdcdington accretion cliscs are given as functions of the polar angle and the mass accretion rate.," In this paper, we investigate the properties of IR (as well as NIR) emission from the inner edge of the dusty torus, employing radiative flux from the disc accretion flows by W05, in which the radiation fluxes of sub- and super-Eddington accretion discs are given as functions of the polar angle and the mass accretion rate."
 We brielly summarize the relation between the radiation [Dux from the aceretion disc and the mass accretion rate into an SAIBIL., We briefly summarize the relation between the radiation flux from the accretion disc and the mass accretion rate into an SMBH.
 For a small mass accretion rate. Len mPhpH—Alou(heapfe)=1 10. where Mpgg and Lea are the mass accretion rate into an ΛΙΠΗΣΕ and the Ecelington luminositv. respectively. if the accretion disc is geometrically thin. it is known as a standard disc (Shakura&SyunvaevLOT3).," For a small mass accretion rate, i.e., $\dot{m}_{\rm BH}\equiv
\dot{M}_{\rm BH}/(L_{\rm Edd}/c^2)=1$ $10$, where $\dot{M}_{\rm BH}$ and $L_{\rm Edd}$ are the mass accretion rate into an SMBH and the Eddington luminosity, respectively, if the accretion disc is geometrically thin, it is known as a standard disc \citep{SS73}."
. As the mass accretion rateincreases (ipa:LO). the disc becomes geometrically thick via strong radiation pressure. e. a slim dise. as first introciuced by Abramowiezctal. (1988).," As the mass accretion rateincreases $\dot{m}_{\rm BH}\geq 10$ ), the disc becomes geometrically thick via strong radiation pressure, i.e., a slim disc, as first introduced by \citet{Ab88}."
. The maximum thickness of the disc is about45° because the ratio of the scale. height and. the disc. size is approximately unity (e.g. Ixato.Fukue.& Mineshige 2008..," The maximum thickness of the disc is about$45^{\circ}$ because the ratio of the scale height and the disc size is approximately unity (e.g., \citealt{Ka08}."
 W05 examined the observed spectra for various mass accretion rates and included the οσοι of disc. geometry., W05 examined the observed spectra for various mass accretion rates and included the effect of disc geometry.
 They found that the radiation Εαν is proportional to cos8 for the sub-IEXddington regime. mig=1 10. where 8 is the polar angle (see Fig.," They found that the radiation flux is proportional to $\cos\theta$ for the sub-Eddington regime, $\dot{m}_{\rm BH}=1$ $10$, where $\theta$ is the polar angle (see Fig."
 1)., 1).
 The factor cos@ represents the change of the effective area via a change of 8. iie. the projection effect.," The factor $\cos{\theta}$ represents the change of the effective area via a change of $\theta$, i.e., the projection effect."
 For a super-Eddington accretion disc with ig107. although the radiation [lux is also proportional to cos& at 0« 45. it decreases more sienilicantly as @ increases in regions with 9z 45°.Vhe bolometric luminosity of an aceretion disc. is expressed as a function of mpg (Watarai.Pukuc. as follows: Dased on the results of W05.. we can describe the radiation Ilux from the accretion disc. £(8). as follows: For 6«45 For 8&45," For a super-Eddington accretion disc with $\dot{m}_{\rm BH} \geq 10^2$, although the radiation flux is also proportional to $\cos\theta$ at $\theta < 45^{\circ}$ , it decreases more significantly as $\theta$ increases in regions with $\theta \geq 45^{\circ}$ .The bolometric luminosity of an accretion disc is expressed as a function of $\dot{m}_{\rm BH}$ \citep{Wa00} as follows: Based on the results of W05, we can describe the radiation flux from the accretion disc, $F(\theta)$ , as follows: For $\theta < 45^{\circ}$ , For $\theta \geq 45^{\circ}$ ,"
of radius is given by μερα)=No The L/P profile. sometimes refered to as a Meste clistribution ((Alestel 1963). is also. motivated. by observations.,"of radius is given by (R)=N_0 The $1/R$ profile, sometimes refered to as a Mestel distribution \nocite{mest:63}( 1963), is also motivated by observations."
 Radio observations 1981) have shown that the surface density of LIL gas is proportiona to the projected surface density of the total mass. which for a perfectly Dat rotation curve would imply a 1/4 distribution.," Radio observations \nocite{bosma:81}( 1981) have shown that the surface density of HI gas is proportional to the projected surface density of the total mass, which for a perfectly flat rotation curve would imply a $1/R$ distribution."
 We parameterize the Moestel disc in terms of the truncation radius Z2; and the column density at tha radius Ny=maf1... (£2) ," We parameterize the Mestel disc in terms of the truncation radius $R_t$ and the column density at that radius $N_t
\equiv m_{\rm gas}/(2\pi \mu m_H R_t^2)$: (R)=N_t."
1n the limit of infinitely thin discs. we can calculate the cross section for these distributions analytically and use them to check our numeric code.," In the limit of infinitely thin discs, we can calculate the cross section for these distributions analytically and use them to check our numeric code."
 For an exponential disc. the inclination averaged cross section 1s ((Bartelmann 1996).," For an exponential disc, the inclination averaged cross section is \nocite{bl:96}( 1996)."
" The variable 5,,= is introduced. because a column density of NV when ANcNy can only be reached if the disc is inclined such that cos8as."," The variable $\gamma_m=\min[{{N_0}\over{N}},1]$ is introduced because a column density of $N$ when $N > N_0$ can only be reached if the disc is inclined such that $\cos{\theta} >
\gamma_m$."
 For Mestel dises the corresponding expression is for NoxNy., For Mestel discs the corresponding expression is for $N > N_t$.
 The fiducial SPP 7standard SAAIS” provide us with a list of ealaxies contained within a halo of à given mass or circular velocity at a given redshift., The fiducial SPF “standard SAMs” provide us with a list of galaxies contained within a halo of a given mass or circular velocity at a given redshift.
 For cach of these galaxies. we are also provided with the internal circular velocity. racial distance from the halo centre. stellar exponential scale length. and the cold. gas. stellar. and metal content of its disc.," For each of these galaxies, we are also provided with the internal circular velocity, radial distance from the halo centre, stellar exponential scale length, and the cold gas, stellar, and metal content of its disc."
 We distribute the galaxies randomly on circular orbits (we discuss the importance of this simplification in section 5)) and assign them random inclinations., We distribute the galaxies randomly on circular orbits (we discuss the importance of this simplification in section \ref{depend}) ) and assign them random inclinations.
" We create twenty realizations of a grid. of halos with circular.. velocities2. between 50. kms"" and 500 kmsL", We create twenty realizations of a grid of halos with circular velocities between 50 $\kms$ and 500 $\kms$.
 These correspond to dillerent. Monte Carlo realizations of the halos’ merging histories., These correspond to different Monte Carlo realizations of the halos' merging histories.
" We then choose a model for the racial distribution of the gas and calculate the surface density distribution. constrained bv the total gas mass as determined by the SXMs,"," We then choose a model for the radial distribution of the gas and calculate the surface density distribution, constrained by the total gas mass as determined by the SAMs."
 We create twenty randoni realizations of the satellite orbits and. inclinations in each of the four hundred halos and calculate the column density along each line of sight., We create twenty random realizations of the satellite orbits and inclinations in each of the four hundred halos and calculate the column density along each line of sight.
 The number of lines of sight passed through each halo is determined by the cross-section weighted probability of intersecting a halo of that mass., The number of lines of sight passed through each halo is determined by the cross-section weighted probability of intersecting a halo of that mass.
 Phe total number of lines of sight is chosen to produce about ten thousand DLAS., The total number of lines of sight is chosen to produce about ten thousand DLAS.
 Each line of sight that passes through a total column density exceeding 2«10em7 is then saved (along with all the properties of the halo that it is found in) and analyzed using the methods of PWO7., Each line of sight that passes through a total column density exceeding $2 \times 10^{20} \cm2$ is then saved (along with all the properties of the halo that it is found in) and analyzed using the methods of PW97.
 To create. synthesized. spectra we must include substructure in the gas discs. which we do by assuming ju the gas is distributed in small clouds within the disc.," To create synthesized spectra we must include substructure in the gas discs, which we do by assuming that the gas is distributed in small clouds within the disc."
 The necessary parameters are: σης the internal velocity ispersion of each cloud: AY. the number of clouds: and e... jeir isotropic random. motions.," The necessary parameters are: $\sigma_{int}$, the internal velocity dispersion of each cloud; $N_c$, the number of clouds; and $\sigma_{cc}$, their isotropic random motions."
" Following PW9T. we take Tn=+3kms and JN.=5: both values were derived rom Voigt profile fits to the observations with IN,=5 being 1o mininium acceptable number of individual components."," Following PW97, we take $\sigma_{int}=4.3 \kms$ and $N_c=5$; both values were derived from Voigt profile fits to the observations with $N_c=5$ being the minimum acceptable number of individual components."
 Increasing the cloud number. NV to as high as 60 does not improve the goodness of fit (DW97) for a dise model like we are considering here because our model discs are relatively hin.," Increasing the cloud number, $N_c$ to as high as 60 does not improve the goodness of fit (PW97) for a disc model like we are considering here because our model discs are relatively thin."
" Also we take ao,=IOkms|. since we assume that the eas disces are cold."," Also we take $\sigma_{cc}=10 \kms$, since we assume that the gas discs are cold."
 These internal velocities are in addition o the circular velocity of the disc and the motions between discs., These internal velocities are in addition to the circular velocity of the disc and the motions between discs.
 For every line of sight the positions of the clouds are chosen by taking the continuous density distribution to be a wobability distribution: i.e. the likelihood of a cloud being at a position in space is proportional to the σας density at hat point., For every line of sight the positions of the clouds are chosen by taking the continuous density distribution to be a probability distribution; i.e. the likelihood of a cloud being at a position in space is proportional to the gas density at that point.
 Synthetic metal-line profiles are produced taking into account the varving metallicity of the eas in the multiple discs along the sightline as given by the SAAISs., Synthetic metal-line profiles are produced taking into account the varying metallicity of the gas in the multiple discs along the sightline as given by the SAMs.
 The spectrum is smoothed to the resolution of the ΤΗ10 spectrograph ((Voet 1992). noise is added and then the four statistics of PAVOT are applied.," The spectrum is smoothed to the resolution of the HIRES spectrograph \nocite{vogt:92}( 1992), noise is added and then the four statistics of PW97 are applied."
 Finally. a Ixolmagornoy.SmirnolT (INS) test is performed to ascertain the probability that the data of (1998) and (2000) could be a random subset of the model.," Finally, a Kolmagornov–Smirnoff (KS) test is performed to ascertain the probability that the data of \nocite{pw:98}{ (1998) and \nocite{wp:00}{ (2000) could be a random subset of the model."
 1t should be noted that while we try to include all of the relevant. physics in the modeling. there are a number of simplifications.," It should be noted that while we try to include all of the relevant physics in the modeling, there are a number of simplifications."
 The kinematics of sub-halos within the host halos assumes that the sub-halos are on circular orbits ancl utilizes an approximate formula for the elects of vnamical friction., The kinematics of sub-halos within the host halos assumes that the sub-halos are on circular orbits and utilizes an approximate formula for the effects of dynamical friction.
 We assume that the gas disces have a simple radial profile and are axisvmnmietric., We assume that the gas discs have a simple radial profile and are axisymmetric.
 Also we assume that the distribution of the gas does not depend on galaxy environment or. Hubble type., Also we assume that the distribution of the gas does not depend on galaxy environment or Hubble type.
 We expect that gas. clises should be distorted by the presence of other galaxies in the same halo or by previous merger events ((cfDonald ," We expect that gas discs should be distorted by the presence of other galaxies in the same halo or by previous merger events \nocite{mm:99,kola:99}( (cf."
1999: 1999) vet we ignore these ellects., 1999; 1999) yet we ignore these effects.
 In the spirit of SAAIs we hope that these assumptions will capture the essential properties of the resulting DLAS to first. order. and we investigate the sensitivity. of our results to some of these assumptions.," In the spirit of SAMs we hope that these assumptions will capture the essential properties of the resulting DLAS to first order, and we investigate the sensitivity of our results to some of these assumptions."
 In section 6.. we note the good. agreement of some of the features of our moclel with the results of recent hydrodynamical simulations. and in the future we hope to refine our. modelling by further comparisons with simulations.," In section \ref{compare}, we note the good agreement of some of the features of our model with the results of recent hydrodynamical simulations, and in the future we hope to refine our modelling by further comparisons with simulations."
DDO photoelectric photometry of 5 red giants was obtained by Clari&(1985)..,DDO photoelectric photometry of 5 red giants was obtained by \citet{claria}.
 From 4 menmbers of these red giants. he conclude that this cluster has a metallicity [Fe/1I]—0.1. E(B-W)=0.0G40.02 and a distance of 530 pe.," From 4 members of these red giants, he conclude that this cluster has a metallicity [Fe/H]=0.1, $E(B-V) = 0.06 \pm 0.02$ and a distance of 530 pc."
 Mermilliod(L981) derived an age of 0.30 Gvr based on a svnthetic composite color-magnitude diagram., \citet{merm} derived an age of 0.30 Gyr based on a synthetic composite color-magnitude diagram.
 Wuetal.(2002) determined absolute proper motions and membership probabilities for 501 stars in the field of MAS., \citet{wu} determined absolute proper motions and membership probabilities for 501 stars in the field of M48.
 More recently. Rideretal.(2004) obtained a new result of MÁS taken in the igriz SDSS filter svstem.," More recently, \citet{rid03} obtained a new result of M48 taken in the $u^{'}g^{'}r^{'}i^{'}z^{'}$ SDSS filter system."
 Thev find that a distance of 100 pe. an age of 0.40 Gyr and a metallicity of |Fe/IH]—0.0 can fit their data best.," They find that a distance of 700 pc, an age of 0.40 Gyr and a metallicity of [Fe/H]=0.0 can fit their data best."
 In this paper we present a new photometric result of MAS taken with BATC Multi-Color survey Photometric svstem., In this paper we present a new photometric result of M48 taken with BATC Multi-Color Survey Photometric system.
 The BATC filler svstem consists of 15 filters of band-widths 150 350A that cover the wavelength range 33004 — 10000A.. which avoids strong and variable sky emission lines (Fanetal.1996).," The BATC filter system consists of 15 filters of band-widths 150 – 350Å that cover the wavelength range 3300Å – 10000Å,, which avoids strong and variable sky emission lines \citep{fan}."
.. As the first object in BATC survey. the old open cluster M67 has been studied based on color-magnitude diagram (CMD)(Fan 1996).," As the first object in BATC survey, the old open cluster M67 has been studied based on color-magnitude diagram \citep{fan}."
. Using the DATC filter svstem. Chenetal.(2000). studied the globular cluster NGC 288 by comparing SEDs of bright stars with Ixurucz models.," Using the BATC filter system, \citet{ch01} studied the globular cluster NGC 288 by comparing SEDs of bright stars with Kurucz models."
 The estimated effective temperatures and average value of [Fe/I] for these stars are consistent with spectroscopic determinations., The estimated effective temperatures and average value of [Fe/H] for these stars are consistent with spectroscopic determinations.
 Dased on the BATC survey observations. (he main aim of this stacy is to determine simultaneously the fundamental parameters of M4AS. such as age. distance. metallicitv and reddening. bv comparing observational SEDs of cluster stars with theoretical stellar evolutionary models.," Based on the BATC survey observations, the main aim of this study is to determine simultaneously the fundamental parameters of M48, such as age, distance, metallicity and reddening, by comparing observational SEDs of cluster stars with theoretical stellar evolutionary models."
 The observations and reduction of the MÁS data are described in Sec., The observations and reduction of the M48 data are described in Sec.
 2., 2.
 In Sec., In Sec.
 3. we derive fundamental parameters of MAS.," 3, we derive fundamental parameters of M48."
 Conclusions and summary are presented in Sec., Conclusions and summary are presented in Sec.
 4., 4.
"(2005),, and source 29 in the X-ray study of the cluster by Preibisch&Zinnecker",", and source 29 in the X-ray study of the cluster by \cite{2002AJ....123.1613P}."
 It has a bolometric temperature of KK and (2002)..luminosity of 1.6Lo.," It has a bolometric temperature of K and luminosity of $1.6\,L_\odot$."
 Walawenderetal.(2006) also detect several H» shocks on either side of the star., \citet{2006AJ....132..467W} also detect several $H_{2}$ shocks on either side of the star.
 We detect both strong continuum and line emission toward IRAS 03410+3152., We detect both strong continuum and line emission toward IRAS 03410+3152.
 The continuum is a point source at the resolution of our data (Figure 2a[Online- but the CO 2-1 emission shows prominent only))red-blue outflow lobes with a moderate opening angle (Figure 3)) characteristic of Class I sources (Arce&Sar-gent 2006)., The continuum is a point source at the resolution of our data (Figure \ref{fig:mosaic1}) ) but the CO 2–1 emission shows prominent red-blue outflow lobes with a moderate opening angle (Figure \ref{fig:outflow}) ) characteristic of Class I sources \citep{2006ApJ...646.1070A}.
. IC348 abuts an active star forming region in the Perseus cloud., IC348 abuts an active star forming region in the Perseus cloud.
 IRAS 03410+3152 belongs to this younger population but the SMA field contained a second source (Lada source ID 468) that is a bona-fide Class II YSO., IRAS 03410+3152 belongs to this younger population but the SMA field contained a second source (Lada source ID 468) that is a bona-fide Class II YSO.
 Our detection of this object in the same field provides a graphic illustration of the tremendous difference in the millimeter flux from a envelope-dominated to disk-dominated YSO., Our detection of this object in the same field provides a graphic illustration of the tremendous difference in the millimeter flux from a envelope-dominated to disk-dominated YSO.
 The rest of this paper concerns the nature of the latter objects., The rest of this paper concerns the nature of the latter objects.
 Nine of the remaining 84 YSOs in the 22 SMA fields were detected with 1.3mmm fluxes ranging from 2 to mmJy., Nine of the remaining 84 YSOs in the 22 SMA fields were detected with mm fluxes ranging from 2 to mJy.
 The 3c limits on the fluxes of the 75 non-detections were typically less than 2mmJy., The $3\sigma$ limits on the fluxes of the 75 non-detections were typically less than mJy.
 We stacked the emission of these non-detections to constrain their average properties., We stacked the emission of these non-detections to constrain their average properties.
 There was no clear detection in the stacked map with a 380 limit to the average flux of mmJy., There was no clear detection in the stacked map with a $3\sigma$ limit to the average flux of mJy.
" As the millimeter emission is optically thin, it provides a direct measure of the dust mass."," As the millimeter emission is optically thin, it provides a direct measure of the dust mass."
" We convert fluxes to total disks masses through the simple formula, where the dust opacity, Ko30Guz=0.023cm?g7, follows the commonly used prescription from Beckwithetal.(1990) and implicitly assumes an ISM gas-to-dust ratio of 100."," We convert fluxes to total disks masses through the simple formula, where the dust opacity, $\kappa_{230{\rm GHz}}=0.023\,{\rm cm}^2\,{\rm g}^{-1}$, follows the commonly used prescription from \cite{1990AJ.....99..924B}
 and implicitly assumes an ISM gas-to-dust ratio of 100."
 We use a distance d=320 ppc based on Herbig(1998).., We use a distance $d=320$ pc based on \cite{1998ApJ...497..736H}.
" Although the dust in the disks lies over a range of radii and therefore temperatures, the Planck function, B,, is calculated at a a single temperature 20KK based on the modeling by Andrews&Williams(2005) who showed this provides an excellent fit to detailed modeling of the infrared-millimeter spectral energy distribution (SED)."," Although the dust in the disks lies over a range of radii and therefore temperatures, the Planck function, $B_\nu$, is calculated at a a single temperature K based on the modeling by \cite{2005ApJ...631.1134A}
 who showed this provides an excellent fit to detailed modeling of the infrared-millimeter spectral energy distribution (SED)."
" 'The vernacular of disk mass, as derived in this way, is standard in the literature but is really a shorthand for a procedure that has much more validity in the interstellar, as opposed to circumstellar, medium."," The vernacular of disk mass, as derived in this way, is standard in the literature but is really a shorthand for a procedure that has much more validity in the interstellar, as opposed to circumstellar, medium."
 It is important to keep in mind that we are really only estimating the mass of micron to millimeter sized particles and then extrapolating the gas content by multiplying by two orders of magnitude!, It is important to keep in mind that we are really only estimating the mass of micron to millimeter sized particles and then extrapolating the gas content by multiplying by two orders of magnitude!
 The inferred disk masses vary between 2 and 6Mjup.," The inferred disk masses vary between 2 and $6\,M_{\rm Jup}$."
 'Their properties are tabulated in Table 2.., Their properties are tabulated in Table \ref{tab:det}. .
 T'he survey is complete for disk masses Mg>2.3Mjup.," The survey is complete for disk masses $M_{\rm d}>2.3\,M_{\rm Jup}$."
" The limit on the stacking of the non-detections implies that the average mass of the non-detections, (Ma)«0.27Mjup."," The limit on the stacking of the non-detections implies that the average mass of the non-detections, $\langle M_{\rm d}\rangle < 0.27\,M_{\rm Jup}$."
" The nine detections, although only a small fraction of the entire sample, provide the first disk mass measurements in 1C348."," The nine detections, although only a small fraction of the entire sample, provide the first disk mass measurements in IC348."
 They provide important constraints on the disk evolution at 2-3MMyr with commensurate implications for understanding planet formation., They provide important constraints on the disk evolution at Myr with commensurate implications for understanding planet formation.
 We place these measurements in context with disk masses in younger star forming regions below after first comparing the properties of the detected and detected sources., We place these measurements in context with disk masses in younger star forming regions below after first comparing the properties of the detected and non-detected sources.
 Figure 4 plotsthe SED of the nine detected sources., Figure \ref{fig:sed} plotsthe SED of the nine detected sources.
 The mid-infrared fluxes in IRAC bands 1-4 (3.6—8.0 um) and MIPS band 1 (24 jum) are from Ladaetal. (2006)..," The mid-infrared fluxes in IRAC bands 1–4 $3.6-8.0\,\mu$ m) and MIPS band 1 $24\,\mu$ m) are from \cite{2006AJ....131.1574L}."
" Optical and near-infrared flux densities in the Rc, Ic, J, H, and K,-bands were obtained by cross-referencing the source positions in Ladaetal.(2006) with tables published by Ciezaetal.(2007) and Luhmanetal. (2003).."," Optical and near-infrared flux densities in the $R_{\rm C}$ , $I_{\rm C}$, $J$, $H$, and $K_{\rm s}$ -bands were obtained by cross-referencing the source positions in \cite{2006AJ....131.1574L}
 with tables published by \cite{2007ApJ...667..308C}
 and \cite{2003ApJ...593.1093L}. ."
" We estimate the extinction, Ay=4.76E(R—I)IEc)o], as in Ciezaetal. where 4.Τθ](Ποthe-- ἴσολονε--intrinsic (Rastellar colors, (Rc—Ic)o, are from (2007)Kenyon&Hartmann(1995) for the spectral types as determined by Luhmanetal.(2003).."," We estimate the extinction, $A_{\rm V} = 4.76E(R-I)
 = 4.76[(R_{\rm C}-I_{\rm C})_{\rm obs} - (R_{\rm C}-I_{\rm C})_{0}]$, as in \cite{2007ApJ...667..308C}
 where the intrinsic stellar colors, $(R_{\rm C}-I_{\rm C})_{0}$, are from \cite{1995ApJS..101..117K}
 for the spectral types as determined by \cite{2003ApJ...593.1093L}."
" The infrared SEDs are similar, but marginally lower in a handful of cases, than the median SED of classical T- stars (CTTS) out to 24 um, which is over-plotted in each panel."," The infrared SEDs are similar, but marginally lower in a handful of cases, than the median SED of classical T-Tauri stars (CTTS) out to $24\,\mu$ m, which is over-plotted in each panel."
 Many of the millimeter non-detections have quite similar infrared SEDs as the millimeter detections inFigure 4.., Many of the millimeter non-detections have quite similar infrared SEDs as the millimeter detections inFigure \ref{fig:sed}. .
" However, there is a statistical difference in the accretion properties of the two samples."," However, there is a statistical difference in the accretion properties of the two samples."
" Figure 5 plots the cumulative distribution ofthe Ha equivalent widths, W)(Ha), for 60 ofthe 84 disks in our survey from Luhmanetal. (2003).."," Figure \ref{fig:EW} plots the cumulative distribution ofthe $\alpha$ equivalent widths, $W_\lambda({\rm H}\alpha)$ , for 60 ofthe 84 disks in our survey from \cite{2003ApJ...593.1093L}. ."
 All nine disks detected, All nine disks detected
 y-ray y-ray (AGILE) 4-rays ambiguous y-raycompared γ-ταγς 4- (Muslimov pair-starved(Frackowiak2005; 2009). mainly by photon-photon (7-7) pair production., $\gamma$ $\gamma$ $\gamma$ $\gamma$ $\gamma$ $\gamma$ mainly by photon-photon $\gamma$ $\gamma$ ) pair production.
" The produced pairs polarize owing to the magnetic-field-aligned electric field, Ej."," The produced pairs polarize owing to the magnetic-field-aligned electric field, $\Ell$."
" If the rotation and magnetic axes reside in the same hemisphere, a positive Ej is exerted to accelerate e*'s (or e~’s) outward (or inwards), increasing the real charge density outwards."," If the rotation and magnetic axes reside in the same hemisphere, a positive $\Ell$ is exerted to accelerate $e^+$'s (or $e^-$ 's) outward (or inwards), increasing the real charge density outwards."
" The inward-migrating, relativistic e~’s radiate curvature -rays, some of which (nearly head-on) collide with the X-rays emitted from the NS surface to materialize as pairs."," The inward-migrating, relativistic $e^-$ 's radiate curvature $\gamma$ -rays, some of which (nearly head-on) collide with the X-rays emitted from the NS surface to materialize as pairs."
" In a pulsar magnetosphere, there is a surface called the ‘null surface’, on which the Goldreich-Julian charge density (Goldreich Julian 1969) 2—B- Ω/(2πο) changes sign, where denotes thep"," In a pulsar magnetosphere, there is a surface called the null surface', on which the Goldreich-Julian charge density (Goldreich Julian 1969) $\rhoGJ \equiv -\mbox{\boldmath$ $}\cdot \mbox{\boldmath$ $}/(2\pi c)$ changes sign, where $\mbox{\boldmath$ $}$ denotes the local magnetic field,$\mbox{\boldmath$ $}$ the NS angular-velocity vector, and $c$ the speed of light."
a; local , There is another characteristic surface called the light cylinder' beyond which the co-rotational velocity exceeds $c$.
magnetic," The distance of the light cylinder from the rotation axis is called the light-cylinder radius', $\rlc \equiv c/\Omega$, where $\Omega \equiv \vert \mbox{\boldmath$ $} \vert$ denotes the NS rotational angular frequency."
 field," On each magnetic azimuthal angle $\varphi_\ast$ (measured around the magnetic axis counter-clockwise), there is a magnetic field line that crosses the light cylinder tangentially; they are called the the last-open field lines'."
", "," An OG is essentially located between the null surface and the light cylinder in the higher colatitudes (i.e., in the closer region to the magnetic axis) than the last-open field lines."
"'The characteristic energy of curvature emission is given by (e.g., Rybicki Lightman 1979) where y refers to the Lorentz factor of the electron or positron, h the Planck constant divided by 27, and Qc the curvature radius of particle’s three-dimensional (3-D) motion."," The characteristic energy of curvature emission is given by (e.g., Rybicki Lightman 1979) where $\gamma$ refers to the Lorentz factor of the electron or positron, $\hbar$ the Planck constant divided by $2\pi$ , and $\varrho_{\rm c}$ the curvature radius of particle's three-dimensional (3-D) motion."
" In the gap, the e+’s achieve electrostatic force balance, eE,lactor.=2e?4*/(392), and saturate at the terminal Lorentz where e designates the charge on the positron, and Ej the electricfield component projected along themagneticfieldline."," In the gap, the $e^\pm$ 's achieve electrostatic force balance, $ e\Ell= 2e^2\gamma^4 / (3\varrho_{\rm c}^2)$ , and saturate at the terminal Lorentz factor, where $e$ designates the charge on the positron, and $\Ell$ the electricfield component projected along themagneticfieldline."
"Combiningequations (1)) and (2)), we obtain Since e*’s are saturated at y, the spectral density of their curvature emission declines sharply as (E/E.)\/? abovethe cutoffenergy, E> E.","Combiningequations \ref{eq:Ec})) and \ref{eq:term}) ), we obtain Since $e^\pm$ 's are saturated at $\gamma$ the spectral density of their curvature emission declines sharply as $(E/E_{\rm c})^{1/2} \exp(-E/E_{\rm c})$ abovethe cutoffenergy, $E \gg E_{\rm c}$ ."
" Therefore, exp(—E/E.)for the same Ej, equation (3))"," Therefore, for the same $\Ell$ , equation \ref{eq:Ec2}) )"
appears that the predicted number may still be too high (Stringeretal.2009) even with this generous allocation of feedback energy.),appears that the predicted number may still be too high \cite{Stringer09} even with this generous allocation of feedback energy.)
" In conjunction with the inclusion of AGN heating, this assumption leads to a system with a much greater hot gas component: As soon as the supply of gas to the disk can no longer be replenished by cooling, all the gas in the system quickly ends up in the hot component."," In conjunction with the inclusion of AGN heating, this assumption leads to a system with a much greater hot gas component: As soon as the supply of gas to the disk can no longer be replenished by cooling, all the gas in the system quickly ends up in the hot component."
" This effect is additionally enforced by the low star formation efficiency (see below), which prevents disk mass being locked up into stars."," This effect is additionally enforced by the low star formation efficiency (see below), which prevents disk mass being locked up into stars."
 The star formation assumptions inGASOLINE and the etal.(2000) model have already been contrasted in refStarFormation.., The star formation assumptions in and the \scite{Cole00} model have already been contrasted in \\ref{StarFormation}.
" The assumed efficiency in the Boweret model is to be so low (e,= 0.0029) that, even after allowing for the structure and gas fraction, the final conversion rate is about a factor of 10lower than in the simulation."," The assumed efficiency in the \scite{Bower06} model is to be so low $\epsilon_\star=0.0029$ ) that, even after allowing for the structure and gas fraction, the final conversion rate is about a factor of 10 than in the simulation."
" (M,/Mcaia& 0 to 0.2/Gyr as opposed to the rate of 0.5 to 3/Gyr seen in Fig. ??)).", $\dot{M}_\star/M_{\rm cold}\approx$ 0 to 0.2/Gyr as opposed to the rate of 0.5 to 3/Gyr seen in Fig. \ref{SFtimescale}) ).
 This stands out in Fig ??.., This stands out in Fig \ref{Bower}.
 The large mass of cold gas which is supported by the disk in the semi-analytic realisation (except when cooling is shut off by AGN) cannot be sustained in the simulation due to much more effective conversion to stars., The large mass of cold gas which is supported by the disk in the semi-analytic realisation (except when cooling is shut off by AGN) cannot be sustained in the simulation due to much more effective conversion to stars.
" The merger history of a simulated, Milky Way-like galaxy has been used to recreate the same system using the semi-analytic model."," The merger history of a simulated, Milky Way-like galaxy has been used to recreate the same system using the semi-analytic model."
" The incoming trajectories of satellite halos were investigated and found to be consistent with the usual, assumed distributions."," The incoming trajectories of satellite halos were investigated and found to be consistent with the usual, assumed distributions."
 Information on the different modes of accretion onto the main halo is shown to be inherently contained in the merger tree (and would be similarly contained in a statistically generated merger tree which may be used in the study of larger volumes)., Information on the different modes of accretion onto the main halo is shown to be inherently contained in the merger tree (and would be similarly contained in a statistically generated merger tree which may be used in the study of larger volumes).
" 'The subsequent fate of accreted gas is not modelled equivalently by the two methods, but the analytic frameworkdoes account for both shocked and unshocked gas."," The subsequent fate of accreted gas is not modelled equivalently by the two methods, but the analytic framework account for both shocked and unshocked gas."
" Furthermore, the estimated formation time of the hot halo is in excellent agreement with the simulation."," Furthermore, the estimated formation time of the hot halo is in excellent agreement with the simulation."
 'The single-parameter analytic distributions which are assumed in are found to be a satisfactory description of the radial distribution of gas and stars inside the simulated halo., The single-parameter analytic distributions which are assumed in are found to be a satisfactory description of the radial distribution of gas and stars inside the simulated halo.
 It was also shown that the choice of hot halo scalelength is of minor importance in the calculation of cooled gas mass., It was also shown that the choice of hot halo scalelength is of minor importance in the calculation of cooled gas mass.
" However, the enforcement of spherical symmetry in these analytic distributions means that any subsequent cold accretion along filaments (persisting in spite of shocks elsewhere) is neglected."," However, the enforcement of spherical symmetry in these analytic distributions means that any subsequent cold accretion along filaments (persisting in spite of shocks elsewhere) is neglected."
" Thus, thetotal quantity of unshocked gas onto the central galaxy is underestimated, as expected by Brooksetal.(2009)."," Thus, the quantity of unshocked gas onto the central galaxy is underestimated, as expected by \scite{Brooks09}."
". Some gas is consequently delayed in its arrival at the central galaxy, though it is found to proceed at later times."," Some gas is consequently delayed in its arrival at the central galaxy, though it is found to proceed at later times."
 The disk structure is predicted analytically by conserving the angular momentum of the cooled gas., The disk structure is predicted analytically by conserving the angular momentum of the cooled gas.
 Resulting scalelengths are within a factor of two of the simulation results and the circular velocities even closer., Resulting scalelengths are within a factor of two of the simulation results and the circular velocities even closer.
" 'This is quite satisfactory for the case of a single halo, but more work is needed to establish whether the two techniques will consistently agree to this level for an arbitrary set of initial conditions (and whether any disagreement is biased)."," This is quite satisfactory for the case of a single halo, but more work is needed to establish whether the two techniques will consistently agree to this level for an arbitrary set of initial conditions (and whether any disagreement is biased)."
" With regard to the subsequent evolution of the disk, the code was altered to apply the same fundamental assumptions as the simulation by considering theirstructure."," With regard to the subsequent evolution of the disk, the code was altered to apply the same fundamental assumptions as the simulation by considering their."
 Subsequent agreement was excellent given this immense reduction in complexity from many thousands of SPH particles to a few one-dimensional equations., Subsequent agreement was excellent given this immense reduction in complexity from many thousands of SPH particles to a few one-dimensional equations.
" So, by assuming the same physics and using the same initial conditions, the two techniques predict a final system which is recognisably the same, despite fundamental differences in the way that these assumptions are applied and the evolution followed."," So, by assuming the same physics and using the same initial conditions, the two techniques predict a final system which is recognisably the same, despite fundamental differences in the way that these assumptions are applied and the evolution followed."
 This suggests that equations which attempt to characterise the emergent behaviour of, This suggests that equations which attempt to characterise the emergent behaviour of
"transition width caused bv the energy shift AZ, ds (Daxetal.1959).. Stark iuisine becomes dominant over the ladder transition when RkR,. where R,, (Gu nuits of Dolir radius day) ls given by. Or R, can be larger than the intermolecular distauce iu a solid or a liquid.","transition width caused by the energy shift $\Delta
E_{Stark} = \langle n, n-2 | e\vec{E}\cdot\vec{r} | n, n-1 \rangle$ is \citep{day59}, Stark mixing becomes dominant over the ladder transition when $R \le R_n$, where $R_n$ (in units of Bohr radius $a_0$ ) is given by, or $R_n$ can be larger than the intermolecular distance in a solid or a liquid."
 This leads to suppression of the ladder N-rayv transitions of interest., This leads to suppression of the ladder X-ray transitions of interest.
 Dowever. molecules iu a eas follow the Boltzimannu distribution and preseut different iupact parameters to the exotic atom.," However, molecules in a gas follow the Boltzmann distribution and present different impact parameters to the exotic atom."
 For this collisional process. the rate of Stark mixing is given by. Nds the muaber deusity of atoms.," For this collisional process, the rate of Stark mixing is given by, $N_{a}$ is the number density of atoms."
 We represent the velocity of atoms e in the gas by their thermal velocity bath<P—\BRENUM where s ds the atomic weight of the eas.," We represent the velocity of atoms $v$ in the gas by their thermal velocity $v_{th}=\sqrt{\case{3kT}{Am_p}}$, where $A$ is the atomic weight of the gas."
" ""Therefore. where p and T are the density [gem5| aud temperature [I] of the gas."," Therefore, where $\rho$ and $T$ are the density $[\mbox{gcm}^{-3}]$ and temperature [K] of the gas."
 When an exotic atom is highly ionized. charge trausfer from other neutral atoms cau refill shells in the exotic atomn.," When an exotic atom is highly ionized, charge transfer from other neutral atoms can refill shells in the exotic atom."
 Charee transfer may cause a continual cvele of Auecr ionization at rates hieh οποιο] to suppress the radiative transitions., Charge transfer may cause a continual cycle of Auger ionization at rates high enough to suppress the radiative transitions.
" The cross-section for charge trausfer from adjacent atoms is in the rango of a,=10i10Deus (Rvutukuetal.1980).", The cross-section for charge transfer from adjacent atoms is in the range of $\sigma_r=10^{-14}-10^{-15} \mbox{cm}^2$ \citep{ryufuku80}.
". When p, is the uuuber density. of clectrous and 6,αι Is the cross section of electron refilling iu units of 10. Henz. the rate of electron refilling is given by. Aletals are not appropriate target materials because their fast Fermi velocities (~10 cui/s) lead to rapid refilling rates."," When $n_e$ is the number density of electrons and $\sigma_{r,-14}$ is the cross section of electron refilling in units of $10^{-14} \mbox{cm}^2$ , the rate of electron refilling is given by, Metals are not appropriate target materials because their fast Fermi velocities $\sim 10^8$ cm/s) lead to rapid refilling rates."
" The stroug nuclear decay rate of the 2/7 state iu protoniun is Pop=35 meV (BRoifeuróther&Ileurpt 1989).. correspouding to Tap=5.3«101? 1,"," The strong nuclear decay rate of the $2P$ state in protonium is $\Gamma_{2P}=35$ meV \citep{reifenrother89}, corresponding to $\Gamma_{2P}=5.3\times10^{13}$ $^{-1}$."
 Nuclear aunihilation for the 2P state can be as fast as its radiative deexcitation., Nuclear annihilation for the $2P$ state can be as fast as its radiative deexcitation.
 The 2)>1 ladder trausition is also excluded from observable line candidates since its transition energv is much higher than the energv baud of a moderately thick XN-raw detector., The $2\rightarrow1$ ladder transition is also excluded from observable line candidates since its transition energy is much higher than the energy band of a moderately thick X-ray detector.
 The strong interaction in D-states is negligible (Reitenrother&Πιο1989)., The strong interaction in D-states is negligible \citep{reifenrother89}.
. The initial capture angular momentum fy is not well nucderstood presently., The initial capture angular momentum $l_0$ is not well understood presently.
 However. we can assimnae a statistical distribution for /y. i.c. a probability for the capture of an antiparticle iuto states /j is given by 2th," However, we can assume a statistical distribution for $l_0$, i.e. a probability for the capture of an antiparticle into states $l_0$ is given by $\case{2l_0+1}{{n_0}^2}$."
 An autiprotou captured iuto fy=O5is not relevant. since it decays iuto a low à» circular state which does not give several lÓladder transition plotons observable by a thin N-raw detector.," An antiproton captured into $l_0=0-5$ is not relevant, since it decays into a low $n$ circular state which does not give several ladder transition photons observable by a thin X-ray detector."
 However. the probability of capture mto such a low Jy state is less thanLOY%.," However, the probability of capture into such a low $l_0$ state is less than."
. Coulomb. deexcitation. a process by which transition cucrey is trausfered to kinetic energv of the exotic atomi and a nearby atom. reduces the vield of ladder transition plotous (Aschenaucretal.1995).," Coulomb deexcitation, a process by which transition energy is transfered to kinetic energy of the exotic atom and a nearby atom, reduces the yield of ladder transition photons \citep{aschenauer95}."
. However. the significance of the Coulomb deexcitation is still not well understood aud is likely a simall effect.," However, the significance of the Coulomb deexcitation is still not well understood and is likely a small effect."
 Tnexperiments it has been shown that the vield of ladder ransitions is depeudent on the gas pressure., In experiments it has been shown that the yield of ladder transitions is dependent on the gas pressure.
 The vield of adder trausitious (729T) was measured to be ~ or relatively low pressure (~10? atin) gases forming antiprotonic atoms (Bacheretal.1988)..On the other waned. Lyiuan lines were measured at high gas pressures (5 170 atin) for muonic atoms (Jacot-Caullarmiedetal.L9ss).," The yield of ladder transitions $n\ge7$ ) was measured to be $\sim$ for relatively low pressure $\sim 10^{-2}$ atm) gases forming antiprotonic atoms \citep{bacher88}.On the other hand, Lyman lines were measured at high gas pressures $\le$ 170 atm) for muonic atoms \citep{jg88}."
". Ilereafter. we set the vield of the overall ladder rausitions Y,,= (n, isthe Ist circular state of he ladder transitions of interest) aud Y,= (n<my. lower à ladder transitions)."," Hereafter, we set the yield of the overall ladder transitions $Y_{n_1}=$ $n_1$ isthe 1st circular state of the ladder transitions of interest) and $Y_n=$ $n<n_1$ , lower $n$ ladder transitions)."
 We will investigate the xessure dependence of the vield of lacder transitions at antiproton beaming facilities., We will investigate the pressure dependence of the yield of ladder transitions at antiproton beaming facilities.
 Atomic calculation sets an upper limit ou the acceptable eas deusitv required to eenerate the characteristic X-ravs froin the ladder trausitious., Atomic calculation sets an upper limit on the acceptable gas density required to generate the characteristic X-rays from the ladder transitions.
 At hieh density. Stark. musing aud electron refillime cau suppress the radiative deexcitatious.," At high density, Stark mixing and electron refilling can suppress the radiative deexcitations."
 However. Stark mixing does not set iu when electrons are preseut in the exotic atom aud so the electron refilhug cau suppress the Stark mixiug as well.," However, Stark mixing does not set in when electrons are present in the exotic atom and so the electron refilling can suppress the Stark mixing as well."
 Therefore. the optimal situation is for the ladder transition rate to be faster than Stark mixing rate by the time electrons no longer refill the exotic atom (more precisely. refill tle K-shell).," Therefore, the optimal situation is for the ladder transition rate to be faster than Stark mixing rate by the time electrons no longer refill the exotic atom (more precisely, refill the K-shell)."
 Refilling of the L-shell cau still be faster than ladder transitions., Refilling of the L-shell can still be faster than ladder transitions.
 However. ata<T. the fluorescence trausitions of L-shell electrons to the I-shell become slower thau the ladder transition in autiprotonic atoms.," However, at $n\le7$, the fluorescence transitions of L-shell electrons to the K-shell become slower than the ladder transition in antiprotonic atoms."
 This condition is explicitly describect as. Therefore. We plot the niacin density for several gases im figure 12..," This condition is explicitly described as, Therefore, We plot the maximum density for several gases in figure \ref{max_density}."
 Ibvdrogen aud helium are ruled out due to their fast Stark mining rate at the eas density desirable for stoppiug antiparticles., Hydrogen and helium are ruled out due to their fast Stark mixing rate at the gas density desirable for stopping antiparticles.
 Ou the other hand. Kr and Ne are not optimal because of high photo-absorption of ladder N-ravs observable with N-rav detectors.," On the other hand, Kr and Xe are not optimal because of high photo-absorption of ladder X-rays observable with X-ray detectors."
 Therefore. we limit our choice of the eas to No. Oo. Ne and Ar.," Therefore, we limit our choice of the gas to $_2$, $_2$, Ne and Ar."
 Cas selection will be discussed further in the following sections., Gas selection will be discussed further in the following sections.
 We assiunie a cubic detectorfor a preliminary calculation of detection efficiency., We assume a cubic detectorfor a preliminary calculation of detection efficiency.
 À detailed description of the detector desigu will be found in rofdesieu.., A detailed description of the detector design will be found in \\ref{design}.
 Hereafter. we denote the kinetic energy of an incident particle as E y. its incident direction as =(0.0) and an incident poiut on the detector surface S as," Hereafter, we denote the kinetic energy of an incident particle as $E_{\bar{X}}$ , its incident direction as $\hat{n}=(\theta, \phi)$ and an incident point on the detector surface $S$ as"
by AMiralda-Escudé/(1999).,by \cite{m99}.
. We follow his result here., We follow his result here.
" Considering a variable X-ray source at angular diameter distance D.. and an intervening dust laver (e.g. a galaxy or a eloud) at redshift z; and with angular diameter distance D. the time delay /; of photons received al an angle 0 [rom the source scattered by the dust laver is given by where D, is the angular diameter distance from the dust laver (ο the source."," Considering a variable X-ray source at angular diameter distance $D_s$, and an intervening dust layer (e.g. a galaxy or a cloud) at redshift $z_d$ and with angular diameter distance $D_d$ , the time delay $t_d$ of photons received at an angle $\theta$ from the source scattered by the dust layer is given by where $D_{ds}$ is the angular diameter distance from the dust layer to the source."
 See Figure 1 for the geometry of the scattering., See Figure 1 for the geometry of the scattering.
 For simplicity. (he variable source is assumed (o emit most of its energv in a narrow energv band (e.g. 210 keV) with a given fInence Sy). in a time much shorter than /y.," For simplicity, the variable source is assumed to emit most of its energy in a narrow energy band (e.g. 2–10 keV) with a given fluence $S_0$, in a time much shorter than $t_d$."
 Thus we ean approximate the source emission as a pulse of monochromatic light with wavelength A., Thus we can approximate the source emission as a pulse of monochromatic light with wavelength $\lambda$.
" If an experiential size distribution of grains in the intervening dust laver is assumecl. 1.e.. T(a)xa""?. where 7(«) is the scattering optical depth of dust grains with radius a. the total flux of the X-ray echo due to the small-angle scattering can be estimated by where a,©1.5A4D,4,/(2xD.,0) is approximately satisfied. Vv, is the mean gas column density over the laver. and Z, is the total metallicity in dust grains."," If an experiential size distribution of grains in the intervening dust layer is assumed, i.e., $\tau(a)\propto
a^{-0.5}$, where $\tau(a)$ is the scattering optical depth of dust grains with radius $a$, the total flux of the X-ray echo due to the small-angle scattering can be estimated by where $a_p\simeq1.5\lambda D_{ds}/(2\pi D_s \theta)$ is approximately satisfied, $\bar{N_b}$ is the mean gas column density over the layer, and $Z_t$ is the total metallicity in dust grains."
 Here the differential cross section is (treated as an approximation of the Mie solution. which is given in the form of the “Ravlei¢h-Gans” approximation in § 3.," Here the differential cross section is treated as an approximation of the Mie solution, which is given in the form of the “Rayleigh-Gans"" approximation in $\S$ 3."
 Using the relation between { and @ (eq. [, Using the relation between $t_d$ and $\theta$ (eq. [
1]) in equation (2) vields where the echo flux is a function of the time delay /;.,1]) in equation (2) yields where the echo flux is a function of the time delay $t_d$.
 We should point ont that (1) (his equation is a good estimate of the flux of X-ray echoes. from which we can get a preliminary idea of how much the flux will be at a given /; and whether it will be detectable al the tme of interest.," We should point out that (1) this equation is a good estimate of the flux of X-ray echoes, from which we can get a preliminary idea of how much the flux will be at a given $t_d$ and whether it will be detectable at the time of interest."
 We discuss this issue later in this section. (, We discuss this issue later in this section. (
2) Considering the twpical scattering angle implied bv the value of αρ. equation (3) is only valid before or around L40LOMSOAP327υύοD.)/0.1jn)7. where ¢ is the speed of light. (,"2) Considering the typical scattering angle implied by the value of $a_p$, equation (3) is only valid before or around $t_d\sim
10^{10}\,{\rm s}(9\lambda^2/32\pi^2
c)(1+z_d)(D_{ds}D_d/D_s)(a_p/0.1\,\mu{\rm m})^{-2}$, where $c$ is the speed of light. ("
3) This equation implies (hat a pulse of light propagating through a dusty region will produce a long-term delaved emission. which is like an vallerelow’ of (hispulse (see Fig.,"3) This equation implies that a pulse of light propagating through a dusty region will produce a long-term delayed emission, which is like an “afterglow"" of thispulse (see Fig."
 1)., 1).
 Here we call (his a dust-scatteringdriven alterglow., Here we call this a dust-scattering–driven afterglow.
 We work out its detailed light curve in § 3., We work out its detailed light curve in $\S$ 3.
in the calculation of gravitational force to decrease relaxation [rom close encounters of particles ancl to make the simulation numerically feasible.,in the calculation of gravitational force to decrease relaxation from close encounters of particles and to make the simulation numerically feasible.
 We use spline kernel gravitational softening (?).., We use spline kernel gravitational softening \citep{HernquistANDKatz89}.
 The force on each particle is integrated using a timestep M=7Vt. where the timestep criterion 7=0.115 and « is the acceleration.," The force on each particle is integrated using a timestep $\Delta t_{grav} = \eta \sqrt{\frac{\epsilon}{a}}$, where the timestep criterion $\eta = 0.175$ and $a$ is the acceleration."
 For gas particles. the timestep also satisfies the Courant condition (2)..," For gas particles, the timestep also satisfies the Courant condition \citep{Wadsley04}."
 Each gas particle has an SPII smoothing radius. h. over which gas properties are samplecl.," Each gas particle has an SPH smoothing radius, $h$, over which gas properties are sampled."
 The value for h is set for each gas particle such that above a minima value of /xe. the same number of particles are contained in each smoothing kernel: when hi is equal to the minimun value. more particles may be contained within the kernel.," The value for $h$ is set for each gas particle such that above a minimum value of $h \propto \epsilon$, the same number of particles are contained in each smoothing kernel; when $h$ is equal to the minimum value, more particles may be contained within the kernel."
 In our simulations >0.1e and 32 particles ave contained in each kernel when h>0.1e., In our simulations $h \geq 0.1 \epsilon$ and 32 particles are contained in each kernel when $h > 0.1\epsilon$.
 We implement gas cooling assuming ionization equilibrium ancl an ideal gas of primordial composition., We implement gas cooling assuming ionization equilibrium and an ideal gas of primordial composition.
 The cooling rate and the ion abundances are solved for using collisional ionization rates (?).. radiative recombination (??).. bremsstrahlung. and line cooling (?)..," The cooling rate and the ion abundances are solved for using collisional ionization rates \citep{Abel97}, radiative recombination \citep{Black81, VernerANDFerland96}, bremsstrahlung, and line cooling \citep{Cen92}."
 As in SOG. the ellects of the UV background radiation. metal line cooling and molecular hvdrogen cooling are not included.," As in S06, the effects of the UV background radiation, metal line cooling and molecular hydrogen cooling are not included."
 We modeled SF using the sub-grid recipe outlined in 506., We modeled SF using the sub-grid recipe outlined in S06.
 The formation of each star particle (representing a simple stellar population of a set metallicity) is a statistical process based on the properties of the local gas., The formation of each star particle (representing a simple stellar population of a set metallicity) is a statistical process based on the properties of the local gas.
 In (liis recipe. gas particles must meet the following set of criteria to be capable of forming stars: the gas density must be greater than a minimum densitv. μην. the temperature must beless than a maxinum temperature. 000Ix. and the gas must be in a convergine flow (7)..," In this recipe, gas particles must meet the following set of criteria to be capable of forming stars: the gas density must be greater than a minimum density, $n_{min}$, the temperature must beless than a maximum temperature, $T < T_{max} = 15,000$ K, and the gas must be in a converging flow \citep{Katz92}."
" Whether a potentially star forming eas particle spawns a star particle in a (ime step. Af. is determined stochastically with the folowing probability p: in which m4; 1s the mass of the gas particle. Πω, 1s Che initial mass of the potential star particle. &* is a constant star forming efficiency [actor (the likelihood a valid eas particle will make a star particle in a formation tme). and the formation (ime. /;,,,. is equal to the dynamical time (SOG)."," Whether a potentially star forming gas particle spawns a star particle in a time step, $\Delta t$, is determined stochastically with the following probability $p$: in which $m_{gas}$ is the mass of the gas particle, $m_{star}$ is the initial mass of the potential star particle, $c^*$ is a constant star forming efficiency factor (the likelihood a valid gas particle will make a star particle in a formation time), and the formation time, $t_{form}$, is equal to the dynamical time (S06)."
" We chose ma4, to be of the initial gas particle mass as a compromise between (he resolution of the galaxys stellar component ancl computational expense.", We chose $m_{star}$ to be of the initial gas particle mass as a compromise between the resolution of the galaxy's stellar component and computational expense.
" This set of minimum star forming criteria and stochastic SF parameters leaves (wo (unable parameters: 2,5, and c*.", This set of minimum star forming criteria and stochastic SF parameters leaves two tunable parameters: $n_{min}$ and $c^*$.
" The tests in ? determined best values of 5,=0.1em* and οἳ=0.05 and we use a? initial mass function to our star particles. as thev clic."," The tests in \citet{Governato07} determined best values of $n_{min} = 0.1 \mathrm{cm}^{-3}$ and $c^* = 0.05$ and we use a \citet{Kroupa93} initial mass function to our star particles, as they did."
 Feedback from star particles represents the energy injected into the interstellar medium (ISM) by Type Ia and Type II supernovae., Feedback from star particles represents the energy injected into the interstellar medium (ISM) by Type Ia and Type II supernovae.
 To calculate the effect of the released. energy. we use (he blastwave recipe described in SOG based on work of ?..," To calculate the effect of the released energy, we use the blastwave recipe described in S06 based on work of \citet{McKee77}. ."
 Energy. [rom a given star, Energy from a given star
"gas morphology, we focus on CC clusters.","gas morphology, we focus on CC clusters."
 We started analyzing a small sample of eight objects extracted from the Chandra dataset investigated in Bonamente et al. (, We started analyzing a small sample of eight objects extracted from the Chandra dataset investigated in Bonamente et al. (
2006).,2006).
" The study of a central region, commonly known as the core region, has been challenged by high-resolution numerical simulations (Navarroetal.1995,Henningetal. 2009))."," The study of a central region, commonly known as the core region, has been challenged by high-resolution numerical simulations \cite{Navarro95, Borgani04, Kay04, Nagai07, Henning09}) )."
" These works lead to an agreement on whether there is à cooling core in the very central denser gas region (r« 0.11599) of some clusters, as well as a slower decline in the temperature at large radii (r> 0.2rgoo)."," These works lead to an agreement on whether there is a cooling core in the very central denser gas region $r<0.1r_{500}$ ) of some clusters, as well as a slower decline in the temperature at large radii $r>0.2r_{500}$ )."
 As usual we refer to rsoo as the radius of the cluster that defines a volume with mean density 500 times the critical density pci; at cluster redshift., As usual we refer to $r_{500}$ as the radius of the cluster that defines a volume with mean density $500$ times the critical density $\rho_{crit}$ at cluster redshift.
 The choice of T5090 is motivated by simulation results from Evrard et al. (, The choice of $r_{500}$ is motivated by simulation results from Evrard et al. (
1996) showing the gas within this radius relaxed and in hydrostatic equilibrium.,1996) showing the gas within this radius relaxed and in hydrostatic equilibrium.
" Moreover, many observational studies in X-rays have shown that the X-ray surface brightness, hence the underlying density, cannot be represented correctly by a profile."," Moreover, many observational studies in X-rays have shown that the X-ray surface brightness, hence the underlying density, cannot be represented correctly by a profile."
 A second component should be added or a peaked central part introduced in order to properly fit the observation., A second component should be added or a peaked central part introduced in order to properly fit the observation.
 The observed deprojected density profiles are peaked for CC systems and flatter for morphologically disturbed clusters., The observed deprojected density profiles are peaked for CC systems and flatter for morphologically disturbed clusters.
 The πε cluster profile can be described by a double beta--model profile where the parameters’ data have been taken from Bonamente et al. (, The $n_e$ cluster profile can be described by a double -model profile where the parameters' data have been taken from Bonamente et al. (
"2006) and adapted to this work to have a symmetric standard deviation (D’Agostini, 2003).","2006) and adapted to this work to have a symmetric standard deviation (D'Agostini, 2003)."
" This distribution is a generalization of a double -model profile of the electron number density, developed by La Roque et al. ("," This distribution is a generalization of a double -model profile of the electron number density, developed by La Roque et al. ("
"2005), but instead using the same 3 parameter for both the central region and the outskirts, as in Bonamente et al. (","2005), but instead using the same $\beta$ parameter for both the central region and the outskirts, as in Bonamente et al. ("
2006).,2006).
" The ray=04/D4 and 12=0.3/D4 are the core radii of the inner and outer distributions, and f is a parameter defined between 0 and 1 that represents how the core region dominates the outer region."," The $r_{c1}=\theta_{c1}/D_A$ and $r_{c2}=\theta_{c2}/D_A$ are the core radii of the inner and outer distributions, and $f$ is a parameter defined between 0 and 1 that represents how the core region dominates the outer region."
" These parameters, together with n-o, are taken from Bonamente et al. ("," These parameters, together with $n_{e0}$, are taken from Bonamente et al. ("
2006) and summarized in Table for the selected clusters.,2006) and summarized in Table \ref{table_neCC} for the selected clusters.
" We describe the temperature profile as where A,» are determined by fixing the position at which the two power laws, described by the mi and mo parameters, intersect each other, as explained in Appendix 7?"," We describe the temperature profile as where $A_{1,2}$ are determined by fixing the position at which the two power laws, described by the $m_1$ and $m_2$ parameters, intersect each other, as explained in Appendix \ref{sec_appendix}. ."
Probe (AWALAP)satellite’.,Probe (WMAP).
. The expected CMB polarization meastcients from will allow to push both CZ aud D) polarization results well bevoid the preseut knowledge aud considerably coustraiu the tensor (D modes) to scalar ratio paraiueter r. if not to obtaira detection on it.," The expected CMB polarization measurements from will allow to push both $E$ and $B$ ) polarization results well beyond the present knowledge and considerably constrain the tensor $B$ modes) to scalar ratio parameter $ r $, if not to obtain a detection on it."
 Iu this respect. the way of extracting aud plivysically iuterpretiug cosmological paraincters (once the CMD data cleaned οι the different astroplivsical foregrounds) will be important.," In this respect, the way of extracting and physically interpreting cosmological parameters (once the CMB data cleaned from the different astrophysical foregrounds) will be important."
 Iu other words. thePlanck actu:d performance for the crucial primordial parameter kc will depend ou the adopted physical modeling aud on the qiality of data analysis aud interpretation.," In other words, the actual performance for the crucial primordial parameter $ r $ will depend on the adopted physical modeling and on the quality of data analysis and interpretation."
 It is then maportaut aud timely to make forecasts for thePlanch determivation of r aud other cosmological parameters taking muto accotut the theoretical progress in the field and WALIAP results., It is then important and timely to make forecasts for the determination of $ r $ and other cosmological parameters taking into account the theoretical progress in the field and WMAP results.
" The Standard Model of the Uuverse (or ""concordauce nodel) provides the current realistic ¢context to analyze the CAIB aud other cosinological'astroplivsical datzv", The Standard Model of the Universe (or “concordance” model) provides the current realistic context to analyze the CMB and other cosmological/astrophysical data.
 Iuflatk1n (quasi-exponcutia accelerated expaion) of the early Universe is a part of this model aud one inportaut goal of CAIB experiments is probing the physics of it., Inflation (quasi-exponential accelerated expansion) of the early Universe is a part of this model and one important goal of CMB experiments is probing the physics of it.
" Tidlation solves the shortcomines of the decelerate expanding costnology (rorizon problem. flaness, entropv of the Universe). aud explains the observed CMD auisotropies providing the mechausin for the generation οscalar and tensor perturbations seeding the large scale structures (LSS) and priiioxlial (still uudetected) eravitational waves (B mode polarization)."," Inflation solves the shortcomings of the decelerated expanding cosmology (horizon problem, flatness, entropy of the Universe), and explains the observed CMB anisotropies providing the mechanism for the generation of scalar and tensor perturbations seeding the large scale structures (LSS) and primordial (still undetected) gravitational waves $B$ mode polarization)."
 The current CMD | LSS data suport the standard inflationary predictious of a nearly spatially flat Universe with adiabatic aud nearly scale iuvariaw initial deusitv ])orturbaions., The current CMB $+$ LSS data support the standard inflationary predictions of a nearly spatially flat Universe with adiabatic and nearly scale invariant initial density perturbations.
 These data are vaidating the sinele field. slow-roll inflationary scenario (INomatsuetal.2009)., These data are validating the single field slow-roll inflationary scenario \citep{WMAP5}.
. Single field slow-roll models provide aLrappealing. simple aud fairly ecucric description o: iuflation (Dodelsou2003:Bovanovekyetal.," Single field slow-roll models provide an appealing, simple and fairly generic description of inflation \citep{libros,reviu}."
2009).. The inflationary scenario is 1uplemented using a scalar field. tlie with a poteitial Wey) selfuiedsteutv coupled to the space-tiue iuetric.," The inflationary scenario is implemented using a scalar field, the with a potential $ V(\varphi) $, self-consistently coupled to the space-time metric."
 In the effective theory based ou the Cainzbure-Landau (C-L) approach to inflation (Bovanovekyctal.2009).. the poteutial is a polwuouiial in the field sarting bv a constant term.," In the effective theory based on the Ginzburg-Landau (G-L) approach to inflation \citep{reviu}, the potential is a polynomial in the field starting by a constant term."
 Linear terius cau always be clininmated by a constant shift of the inflaton field., Linear terms can always be eliminated by a constant shift of the inflaton field.
 The mass (quadratic) term can have a positivο Op d ieative sign associated to unbrokeu sviunetry (chaotic inflation) or to broken svuunetrv (new inflation). respectively.," The mass (quadratic) term can have a positive or a negative sign associated to unbroken symmetry (chaotic inflation) or to broken symmetry (new inflation), respectively."
 The fourth degree doublewell iufaton potential οἼνος an excelleut üt of the present CAIB | LSS data (Bovanovskyetal.2009)., The fourth degree double–well inflaton potential gives an excellent fit of the present CMB $+$ LSS data \citep{reviu}.
.. A cubic term docs not improve the fit and cau be omitted (Destrietal.2008a)., A cubic term does not improve the fit and can be omitted \citep{mcmc1}.
. Adding higher order terim with additional parameters does not improve siguificautly the fits (Destrietal.2009)., Adding higher order terms with additional parameters does not improve significantly the fits \citep{high}.
.. T1ο CEL framework is uot. just a class of plivsicallvy we] motivated inflaton potentials. ;unong them the double aud singe well potentials.," The G-L framework is not just a class of physically well motivated inflaton potentials, among them the double and single well potentials."
 This approach provides the effective theory for iuflation. with powerful exin iu the plivsical insight aud analysis of the data.," This approach provides the effective theory for inflation, with powerful gain in the physical insight and analysis of the data."
 The preseut set of data with the effective theory of inflation favor the double well potential (Dovanovskyetal.2009:Destrict2008a).," The present set of data with the effective theory of inflation favor the double well potential \citep{reviu,mcmc1}."
. Analwzing the preseut data without the relation between r aud iy docs rot allow to discriminate amoung differeut classes of models for the inflaton potential iu the considered framework., Analyzing the present data without the relation between $r$ and $n_s$ does not allow to discriminate among different classes of models for the inflaton potential in the considered framework.
 Altrough the G-L effective theory approach to inflation is quite general. it predicts precise order of magnitude estimate dor pus r aud the ramming of the spectral iudex div.fdluk here AN~60 is the number of efolds since the cosimologically reevant modes exit t1C rorizon till inflation cuds.," Although the G-L effective theory approach to inflation is quite general, it predicts precise order of magnitude estimates for $ n_s $, $ r $ and the running of the spectral index $ dn_s/d \ln k $ \citep{reviu}
 here $ N \sim 60 $ is the number of efolds since the cosmologically relevant modes exit the horizon till inflation ends."
" The WMAP values for a, aud the upper bounds for r and digfdlak agree with these estimates.", The WMAP values for $ n_s $ and the upper bounds for $ r $ and $ dn_s/d \ln k $ agree with these estimates.
 Since iu this framework the estimated ruuniug. dii/dluk~3.101 ls very snuall. iu thIs paper we will conceitvate on iἐς aud or.," Since in this framework the estimated running, $ dn_s/d \ln k \sim 3 \times 10^{-4} $, is very small, in this paper we will concentrate on $ n_s $ and $ r $."
" Tn this work. we ev""dluate the accuracy in the recovery of the cosnologica pirinueters expected from thePlanck data."," In this work, we evaluate the accuracy in the recovery of the cosmological parameters expected from the data."
 First. we do this forecast without iucludiug the svstematic effects of iustiuiieutal audfx astroplivsical origin or their coupling. affectine tjePlanck measurements. aud theà by includiie the systematic effecs.," First, we do this forecast without including the systematic effects of instrumental and/or astrophysical origin or their coupling, affecting the measurements, and then by including the systematic effects."
 Iu this study we exploit thePlanch scusitivity atid resolution at its three favorite» cosmoloegical chaunes. Le. at tie frequencies of 70. 100. and 115 Gz.," In this study we exploit the sensitivity and resolution at its three favorite cosmological channels, i.e. at the frequencies of 70, 100, and 143 GHz."
 Table 1 repors thePlanck performance at these frequencies. based«πι PlanckColaboration(2006) and. for the LFI chaunel at 70 €Uz. ax updated in \laudolesieal.(2010)... Bersanelictal. (2010).. Saudictal. (2010).," Table \ref{table:sens} reports the performance at these frequencies, based on \citet{PlanckBlueBook} and, for the LFI channel at 70 GHz, as updated in \citet{prelauM}, \citet{prelauB}, \citet{sandri_etal_2010}."
.. These seusitivities do not 1uclude the degradation in accuracy that could come froii variot SSCmrces of systematic effects; of both instrumental aidfor astrophysical origin. or their coupling.," These sensitivities do not include the degradation in accuracy that could come from various sources of systematic effects, of both instrumental and/or astrophysical origin, or their coupling."
 Iu Sect., In Sect.
" we discuss tj6 current »iblished estimates for the residuals of svsteimatic effects aud foreerowks affecting thePlanch CAB incasurements: stravlight. main beam asviuuetry, leakage. time constants. elitches. aud fexeerounds."," \ref{toym} we discuss the current published estimates for the residuals of systematic effects and foregrounds affecting the CMB measurements: straylight, main beam asymmetry, leakage, time constants, glitches, and foregrounds."
 In general. we do not τιse dm this work a precise (stil uot completely available) description of the consiclercκα systematic effects. mt only suitadle represeitatious of them. as described in Sect. L.," In general, we do not use in this work a precise (still not completely available) description of the considered systematic effects, but only suitable representations of them, as described in Sect. \ref{toym}."
 This is done in a parametric a»proach. ideutifvius tje corresponding levelsat which the coutro of he systematic effects is necessary not to spoil thePlanck data scientific accuracy.," This is done in a parametric approach, identifying the corresponding levelsat which the control of the systematic effects is necessary not to spoil the data scientific accuracy."
 We techinically iamplement this rescaling with a multiplicative constant on the residuals of the svstematic effects ou the CMD ultiles CY., We technically implement this rescaling with a multiplicative constant on the residuals of the systematic effects on the CMB multipoles $ C_\ell $.
 Obviously. the real analvsis ofPlanck data will have to properv consider all possible svstematic effects of οical. thermal. ane instrumental (vacdiometzric aud bolometric) origm. wi haneven better accuracy than those acliievec in past projects.," Obviously, the real analysis of data will have to properly consider all possible systematic effects of optical, thermal, and instrumental (radiometric and bolometric) origin, with an even better accuracy than those achieved in past projects."
 Iu parallel. a significantly inproved separation of CMD from astrophysical components will be uceded. a task iu principle possible for thauks to its wide frequeney covοσο.," In parallel, a significantly improved separation of CMB from astrophysical components will be needed, a task in principle possible for thanks to its wide frequency coverage."
 Iistrunental svstematics on CAIB teusors-to-scalar have been studied by C2010)..," Instrumental systematics on CMB tensors-to-scalar have been studied by \citet{whu,shimon,yadav}. ."
"to tight, circular orbits.","to tight, circular orbits."
" Thus, even marginal dynamical captures of former terrestrial planets can produce stable exomoons."," Thus, even marginal dynamical captures of former terrestrial planets can produce stable exomoons."
" In order to understand how exomoon orbits may evolve after being captured, we created a numerical Kozai Cycle and Tidal Friction (KCTF) model."," In order to understand how exomoon orbits may evolve after being captured, we created a numerical Kozai Cycle and Tidal Friction (KCTF) model."
 Kozai cycles in this context are the secular oscillations in eccentricity and inclination of the exomoon’s orbit caused by stellar torques (Kozai1962)., Kozai cycles in this context are the secular oscillations in eccentricity and inclination of the exomoon's orbit caused by stellar torques \citep{Kozai1962}.
". In isolation, these oscillations preserve the semimajor axis of the orbit and the quantity Hy=cos(I)x V1—e?, where I is the inclination of the exomoon’s orbit relative to the planet’s stellarcentric orbit, and e is the eccentricity of the exomoon’s orbit."," In isolation, these oscillations preserve the semimajor axis of the orbit and the quantity $H_k=cos(I)\times\sqrt{1-e^2}$ , where $I$ is the inclination of the exomoon's orbit relative to the planet's stellarcentric orbit, and $e$ is the eccentricity of the exomoon's orbit."
" In the scenarios simulated, these cycles can initially be as fast as just a few years (see Figure 1)), causing very rapid evolution of the exomoon’s orbit."," In the scenarios simulated, these cycles can initially be as fast as just a few years (see Figure \ref{fig:tale}) ), causing very rapid evolution of the exomoon's orbit."
" Since all the initial orbits for this study were highly elliptical, only very low initial inclinations (within about 15° of coplanar) produced values of Hy, sufficiently low that Kozai cycles did not occur."," Since all the initial orbits for this study were highly elliptical, only very low initial inclinations (within about $15^{\circ}$ of coplanar) produced values of $H_k$ sufficiently low that Kozai cycles did not occur."
" Because the Kozai oscillations from the star attempt to preserve H;, a drop in the inclination of the orbit can cause the eccentricity to become very high."," Because the Kozai oscillations from the star attempt to preserve $H_k$, a drop in the inclination of the orbit can cause the eccentricity to become very high."
" As the eccentricity of the orbit increases, the periapse of the exomoon's orbit becomes much closer to the planet."," As the eccentricity of the orbit increases, the periapse of the exomoon's orbit becomes much closer to the planet."
" Tidal forces become much stronger at these close distances (going as a? in this model), and so a close periapse due to eccentricity causes the orbit to shrink and thus decay further."," Tidal forces become much stronger at these close distances (going as $a^{-8}$ in this model), and so a close periapse due to eccentricity causes the orbit to shrink and thus decay further."
" This sets up a positive feedback loop, which progressively shrinks and circularizes the orbit."," This sets up a positive feedback loop, which progressively shrinks and circularizes the orbit."
" The stair-step semimajor axis decay in Figure 1 happens because the obit is initially only decaying at high eccentricities; once the apoapse is close enough to also experience tidal forces, the oscillations stop."," The stair-step semimajor axis decay in Figure \ref{fig:tale} happens because the obit is initially only decaying at high eccentricities; once the apoapse is close enough to also experience tidal forces, the oscillations stop."
 The result is to cause semimajor axis decay and circularization much faster than if the star were not exciting eccentricity., The result is to cause semimajor axis decay and circularization much faster than if the star were not exciting eccentricity.
" In addition, since the high eccentricity part of the Kozai cycles are at low inclinations, orbits are preferentially frozen near the plane of the stellarcentric orbit."," In addition, since the high eccentricity part of the Kozai cycles are at low inclinations, orbits are preferentially frozen near the plane of the stellarcentric orbit."
" To simulate this process, we used a KCTF model based on that of Eggleton&Kiseleva-Eggleton(2001),, which allows the direct integration of the exomoon orbit’s specific angular momentum vector h and Laplace-Runge-Lenz eccentricity vector e, as well as the spin vectors of the planet and satellite."," To simulate this process, we used a KCTF model based on that of \citet{Eggleton2001}, which allows the direct integration of the exomoon orbit's specific angular momentum vector $\mathbf{h}$ and Laplace-Runge-Lenz eccentricity vector $\mathbf{e}$, as well as the spin vectors of the planet and satellite."
 The tidal properties of the giant planets were based on those presented in Fabrycky&Tremaine(2007) for a Jupiter-mass planet., The tidal properties of the giant planets were based on those presented in \citet{Fabrycky2007} for a Jupiter-mass planet.
" The tidal properties of the exomoons used the formulation of Ragozzine(2009),, along with his addition of solid-body quadrapole gravity."," The tidal properties of the exomoons used the formulation of \citet{Ragozzine2009}, along with his addition of solid-body quadrapole gravity."
" To integrate the system of equations, we used the Burlisch-Stoer method with vector-rational interpolation 1998) and error control based on the algorithm(Sweatman given in Pressetal.(2007),, with a per-timestep tolerance of 10-10."," To integrate the system of equations, we used the Burlisch-Stoer method with vector-rational interpolation \citep{Sweatman1998} and error control based on the algorithm given in \citet{Press2007}, with a per-timestep tolerance of $10^{-10}$."
 This model does not include any dynamical effects from objects external to the exomoon-exoplanet system other than the star itself., This model does not include any dynamical effects from objects external to the exomoon-exoplanet system other than the star itself.
" In addition, we assumed that the planet did not migrate over the course of the simulation."," In addition, we assumed that the planet did not migrate over the course of the simulation."
" If it were to migrate, the effect would be to shrink the Hill radius of the planet and shorten the period of the Kozai cycles."," If it were to migrate, the effect would be to shrink the Hill radius of the planet and shorten the period of the Kozai cycles."
" The smaller Hill radii would allow more captured satellites to escape, while on the other hand, the faster Kozai cycles would decrease the decay timescale."," The smaller Hill radii would allow more captured satellites to escape, while on the other hand, the faster Kozai cycles would decrease the decay timescale."
" Thus, if this model works fast enough for a static case, it should also be applicable to a slowly migrating planet."," Thus, if this model works fast enough for a static case, it should also be applicable to a slowly migrating planet."
" 'To find in what conditions a captured exomoon may survive, we generated 18 sets of synthetic Star-Planet-Moon systems and performed KCTF simulations on them."," To find in what conditions a captured exomoon may survive, we generated 18 sets of synthetic Star-Planet-Moon systems and performed KCTF simulations on them."
" Each set contained 1000 synthetic systems, with common masses for all objects, and randomized initial exomoon orbits and spin states."," Each set contained 1000 synthetic systems, with common masses for all objects, and randomized initial exomoon orbits and spin states."
" To simulate a loose capture, the initial exomoon orbits all had apoapses beyond of the planet's Hill radius, and eccentricities greater than 0.85."," To simulate a loose capture, the initial exomoon orbits all had apoapses beyond of the planet's Hill radius, and eccentricities greater than 0.85."
" This is most consistent with a low mass ratio momentum-exchange capture etal. 2004), but generally similar to any low (Funatodelta-v, non-disruptive capture (or eccentricity excitation due to a close encounter)."," This is most consistent with a low mass ratio momentum-exchange capture \citep{Funato2004}, but generally similar to any low delta-v, non-disruptive capture (or eccentricity excitation due to a close encounter)."
" Again consistent with a capture, both the satellite's initial orbital plane and spin vector were initially pointed at random directions on the sky."," Again consistent with a capture, both the satellite's initial orbital plane and spin vector were initially pointed at random directions on the sky."
 This means an approximate equipartition between prograde and retrograde initial orbits and between prograde and retrograde initial spin states for the satellite., This means an approximate equipartition between prograde and retrograde initial orbits and between prograde and retrograde initial spin states for the satellite.
 The planets had à random obliquity less than 5? from their stellarcentric orbit., The planets had a random obliquity less than $5^{\circ}$ from their stellarcentric orbit.
 Each planet-moon system was at a stellarcentric distance such that the equilibrium temperature was equal to Earth., Each planet-moon system was at a stellarcentric distance such that the equilibrium temperature was equal to Earth.
" The stars and stellarcentric orbits used were the Sun (G2) at 1.0 AU, a main-sequence F0 (1.7 Msun) at 2.1 AU, and a main-sequence MO at 0.25 AU."," The stars and stellarcentric orbits used were the Sun (G2) at 1.0 AU, a main-sequence F0 (1.7 $M_{Sun}$ ) at 2.1 AU, and a main-sequence M0 (0.47 $M_{Sun}$ ) at 0.28 AU."
" The planets we assumed to have a (0.47mass Mgun)equal to either Jupiter or Neptune, using thetidal parameters given in Fabrycky&Tremaine (2007),, though nearly all dissipation was in the exomoon."," The planets we assumed to have a mass equal to either Jupiter or Neptune, using thetidal parameters given in \citet{Fabrycky2007}, , though nearly all dissipation was in the exomoon."
" The simulated exomoons had the mass of either Earth, Mars, or Titan (with Mars uncompressed density), using a tidal Q of 100, modulus of rigidity of 3x101?N/m? (Gladmanetal.1996), and J of 0.001."," The simulated exomoons had the mass of either Earth, Mars, or Titan (with Mars uncompressed density), using a tidal $Q$ of $100$, modulus of rigidity of $3\times10^{10} N/m^2$ \citep{Gladman1996}, and $J_2$ of $0.001$."
 The initial rotational periods of the planet and satellite were varied randomly between 0.1 and 48 hours., The initial rotational periods of the planet and satellite were varied randomly between 0.1 and 48 hours.
 The simulations were rununtil they reached either 10° years or an eccentricity below 10-?., The simulations were rununtil they reached either $10^9$ years or an eccentricity below $10^{-5}$ .
" However, the simulations were terminated early if the periapse went below the Roche limit(thus potentially causing breakup of the exomoon) or the apoapse exceeded the Hill radius"," However, the simulations were terminated early if the periapse went below the Roche limit(thus potentially causing breakup of the exomoon) or the apoapse exceeded the Hill radius"
metal-poor GCs in NGC 1399 (e.g.. Forbes et al.,"metal-poor GCs in NGC 1399 (e.g., Forbes et al."
 1997: Also see the Fig.2 in Llilker et al., 1997; Also see the Fig.2 in Hilker et al.
 2000)) anc rather shallower than that of the metal-rich ones 1.7+0.2: See Dirsch et al., 2000) and rather shallower than that of the metal-rich ones $-1.7\pm 0.2$; See Dirsch et al.
 2003 for the latest value)., 2003 for the latest value).
 One possible reason for the negative gradient in the ICGC distribution is that GC's are more likely to be stripped in the inner regions of the cluster., One possible reason for the negative gradient in the ICGC distribution is that GCs are more likely to be stripped in the inner regions of the cluster.
 Ehe derived value of the slope implies that only the blue GCs in the central region of Fornax cluster (equally. GC's of NGC 1399) were previously NGC 1404s GCs.," The derived value of the slope implies that only the blue GCs in the central region of Fornax cluster (equally, GCs of NGC 1399) were previously NGC 1404's GCs."
 Blue GCs in elliptical galaxies show more extended distributions or shallower density profiles than red GC's (e.g. Geisler et al 1996: Forbes. Broclie Lluchra 1997).," Blue GCs in elliptical galaxies show more extended distributions or shallower density profiles than red GCs (e.g., Geisler et al 1996; Forbes, Brodie Huchra 1997)."
 Our simulations demonstrate that the outer GCs are more likely (ο be stripped during the dynamical evolution than inner ones., Our simulations demonstrate that the outer GCs are more likely to be stripped during the dynamical evolution than inner ones.
 Thus. although GC's stripped from NGC. 1404 consists of only some fraction of GC's in the central region of Fornax cluster. it is reasonable to claim that the tidal stripping processes of GCs fron NGC 1404 can contribute partly to the observed power-law slope of blue GC's in the central region of Fornax cluster.," Thus, although GCs stripped from NGC 1404 consists of only some fraction of GCs in the central region of Fornax cluster, it is reasonable to claim that the tidal stripping processes of GCs from NGC 1404 can contribute partly to the observed power-law slope of blue GCs in the central region of Fornax cluster."
 Fig., Fig.
 4 furthermore demonstrates that the ICCGCOSs show wider wings of the (projected) velocity. (35) distribution in the edge-on view whereas they have the sharp distribution in the face-one one., 4 furthermore demonstrates that the ICGCs show wider wings of the (projected) velocity $V_{\rm p}$ ) distribution in the edge-on view whereas they have the sharp distribution in the face-one one.
 This result just reflects that the edege-on view ts nearly coincident with the orbital planes for most of drifting LOGCs., This result just reflects that the edge-on view is nearly coincident with the orbital planes for most of drifting ICGCs.
 A Iarger fraction of large projected. velocity ol LOGCs with respect to the eluster centre (2 400 kim 1j can be found in the edge-on distribution. in particular. in its negative velocity side.," A larger fraction of large projected velocity of ICGCs with respect to the cluster centre $>$ 400 km $^{-1}$ ) can be found in the edge-on distribution, in particular, in its negative velocity side."
" The velocity dispersion for LOCCs within i, (and. with theabsolute magnitude of projected velocity V, less than 1000. Καὶ s 1) can be estimated to νο 340 km s| for the edge-on. view and. 174 km + for ace-one one."," The velocity dispersion for ICGCs within $r_{\rm
s}$ (and with theabsolute magnitude of projected velocity $V_{\rm p}$ less than 1000 km $^{-1}$ ) can be estimated to be 340 km $^{-1}$ for the edge-on view and 174 km $^{-1}$ for face-one one."
 Some of LOGCs with [Vj] 2 400 km 1 are ocated close to the central NGC 1399 (i.e. within 5 times the elective racius of NGC 1399) so that these can be identified as NGC 1390's GCs (not as LOGCs) with an unusually large mojected velocity., Some of ICGCs with $|V_{\rm p}|$ $>$ $400$ km $^{-1}$ are located close to the central NGC 1399 (i.e. within 5 times the effective radius of NGC 1399) so that these can be identified as NGC 1399's GCs (not as ICGCs) with an unusually large projected velocity.
 The derived velocity dispersion of 340 km s.+ of LOGCs. some fraction of which are close to the centre of thecluster (Le. NGC 1399). is significantly larger than the observed outer stellar velocity dispersion (~ 200 km + for the region L-Skpe from the centre) in NCC 1399 (e.g. Bicknell ct al.," The derived velocity dispersion of 340 km $^{-1}$ of ICGCs, some fraction of which are close to the centre of thecluster (i.e., NGC 1399), is significantly larger than the observed outer stellar velocity dispersion $\sim$ 200 km $^{-1}$ for the region 1 - 8 kpc from the centre) in NGC 1399 (e.g., Bicknell et al."
 1989)., 1989).
 Based on the radial velocity measurement of 74 GCs around NGC 1399. [xissler-Patig et al. (," Based on the radial velocity measurement of 74 GCs around NGC 1399, Kissler-Patig et al. ("
1999) found. that he velocity dispersion for the whole sample is 373+35 km Land the velocity dispersion depends on radius in such a wav that it with radius from the NGC 1399 centre tween. 10 - 30 kpe (e.g.. 26392 at 2 and 408+107 kms at 5 ).,"1999) found that the velocity dispersion for the whole sample is $373 \pm 35$ km $^{-1}$ and the velocity dispersion depends on radius in such a way that it with radius from the NGC 1399 centre between 10 - 30 kpc (e.g., $263
\pm 92$ at 2 $\acute{}$ and $408 \pm 107$ km $^{-1}$ at 5 $\acute{}$ )."
 This outer increase of velocity dispersion jas also been found in the kinematical analysis of planetary nebulae (o.g.. Napolitano. Arnaboldi. and Capacetoli 2002).," This outer increase of velocity dispersion has also been found in the kinematical analysis of planetary nebulae (e.g., Napolitano, Arnaboldi, and Capaccioli 2002)."
 The observed. velocity dispersion of GC's (and PNe) is larger han the (outer) stars of NGC 1399. which is considered to of à possible evidence that GC's (PNe) are orbiting in a eravitational potential of the Fornax cluster rather than in GC 1399 itself (Curillmair et al.," The observed velocity dispersion of GCs (and PNe) is larger than the (outer) stars of NGC 1399, which is considered to be a possible evidence that GCs (PNe) are orbiting in a gravitational potential of the Fornax cluster rather than in NGC 1399 itself (Grillmair et al."
 1994: Wissler-Patig et al., 1994; Kissler-Patig et al.
 1999)., 1999).
 Our results on the larger velocity. dispersion of LOGCs stronely suggests that if they originated from tidal stripping of GCs initially in NGC 1404 and. observationallv identified as GCs in NGC 1399. these GC's can have a velocity dispersion as large as the observed. one for GC's in. the very outer part of NGC 1399.," Our results on the larger velocity dispersion of ICGCs strongly suggests that if they originated from tidal stripping of GCs initially in NGC 1404 and observationally identified as GCs in NGC 1399, these GCs can have a velocity dispersion as large as the observed one for GCs in the very outer part of NGC 1399."
 These results furthermore imply that if outer GCs in NGC 1399 are composed mostly of LOGCs stripped [rom other cluster member galaxies. GCs in NGC 1399 as a whole can show a clear dillerence in velocity dispersion. between the inner “intrinsic” CC's hat are closely associated with the NGC 1399 formation itself ancl the outer acereted/stripped ones from other cluster member galaxies.," These results furthermore imply that if outer GCs in NGC 1399 are composed mostly of ICGCs stripped from other cluster member galaxies, GCs in NGC 1399 as a whole can show a clear difference in velocity dispersion between the inner “intrinsic” GCs that are closely associated with the NGC 1399 formation itself and the outer accreted/stripped ones from other cluster member galaxies."
 Our numerical simulations have »en carried outondy for the dvnamical evolution of NGC 1404 with plausible orbits., Our numerical simulations have been carried out for the dynamical evolution of NGC 1404 with plausible orbits.
 Phe (projected) radial velocity dispersion profile of LOGCs depends on from which galaxy hese are stripped. because the projected. velocity of cach stripped GC depends on the orbit of its previous host galaxy (Le. the orbit of cach GC follows the orbit of its host).," The (projected) radial velocity dispersion profile of ICGCs depends on from which galaxy these are stripped, because the projected velocity of each stripped GC depends on the orbit of its previous host galaxy (i.e., the orbit of each GC follows the orbit of its host)."
 We thus suggest that the observed radial velocity dispersion wolile of GCs in NGC 1399 may be πα fossil record” of the ustorics of GC stripping [rom other cluster member galaxies around NGC 1399., We thus suggest that the observed radial velocity dispersion profile of GCs in NGC 1399 may be “a fossil record” of the histories of GC stripping from other cluster member galaxies around NGC 1399.
 However it may be dillieult in practice o disentangle kinematically any stripping signature from. other (formation) mechanisms that give rise to the rich GC systems of ¢D galaxies., However it may be difficult in practice to disentangle kinematically any stripping signature from other (formation) mechanisms that give rise to the rich GC systems of cD galaxies.
" Although our general results on the stripping of CC's from NGC 1404 due to the strong tidal field of the Fornax cluster do not depend on the model parameters. the final οκ and Nec do depend on (1) orbital eccentricity (65) of the host ealaxy NGC 1404 and (2) initial ratio of the scale length of GC system to the cllective radius of NGC 1404 (a,.)."," Although our general results on the stripping of GCs from NGC 1404 due to the strong tidal field of the Fornax cluster do not depend on the model parameters, the final $S_{\rm N}$ and $N_{\rm GC}$ do depend on (1) orbital eccentricity $e_{\rm p}$ ) of the host galaxy NGC 1404 and (2) initial ratio of the scale length of GC system to the effective radius of NGC 1404 $a_{gc}$ )."
 The final Sx and Noe ofGCs around NGC 1400 are estimated at T = LS9 Gyr when time evolution of Sy and Nee becomes insignificant for all models (Mocdoel 1- 6).," The final $S_{\rm N}$ and $N_{\rm GC}$ of GCs around NGC 1400 are estimated at $T$ = 1.89 Gyr when time evolution of $S_{\rm N}$ and $N_{\rm
GC}$ becomes insignificant for all models (Model 1 - 6)."
 We can only discuss the final structural ane kinematical properties of ICCGCS for the models 2 and 3. in which total number of GCs are large enough. C," We can only discuss the final structural and kinematical properties of ICGCs for the models 2 and 3, in which total number of GCs are large enough. ("
Eotal number of stripped GCs in Model 1. 4. 5. and 6 are too small for us to derive velocity dispersion of ICCGCS ab oa given radius).,"Total number of stripped GCs in Model 1, 4, 5, and 6 are too small for us to derive velocity dispersion of ICGCs at a given radius)."
 Pherefore. we only briclly describe the dependences of LOGC properties on model parameters.," Therefore, we only briefly describe the dependences of ICGC properties on model parameters."
 In Figs., In Figs.
 5 and 6. we illustrate the derived dependences on the above two parameters.," 5 and 6, we illustrate the derived dependences on the above two parameters."
" We find the following: (i) Both Sp and Noo are smaller for the models with larger cj. because in the models with lager e, (Le. with smaller pericenter distance). galaxies can pass through the inner region of the cluster. where the cluster tidal field. is sulliciently strong such that a larger number of GC's can be stripped."," We find the following: (i) Both $S_{\rm N}$ and $N_{\rm GC}$ are smaller for the models with larger $e_{\rm p}$, because in the models with larger $e_{\rm p}$ (i.e., with smaller pericenter distance), galaxies can pass through the inner region of the cluster, where the cluster tidal field is sufficiently strong such that a larger number of GCs can be stripped."
 This dependence does not depend on duc. (, This dependence does not depend on $a_{\rm gc}$. (
i) Lrrespective of ej and ess. the model with smaller Noe shows smaller. Sx.,"ii) Irrespective of $e_{\rm p}$ and $a_{\rm gc}$ , the model with smaller $N_{\rm GC}$ shows smaller $S_{\rm N}$ ."
 This means that the stripping of starsin NGC 1404. which can increase οκ if stripping of GCs does not happen. is much less elficient compared. with," This means that the stripping of starsin NGC 1404, which can increase $S_{\rm N}$ if stripping of GCs does not happen, is much less efficient compared with"
an upgrade of the code described in et al. (1997)) (,an upgrade of the code described in et al.\cite{mart97}) ) (
e.g. see Perucho et al. 2005)).,e.g. see Perucho et al. \cite{pe+05}) ).
 Simulations were performed in two dual-core processors in the Max-Planck-Institut. fürr Radioastronomie., Simulations were performed in two dual-core processors in the Max-Planck-Institut fürr Radioastronomie.
 The numerical gric of the cylindrical geometry simulations is formed by 320 cells in the radial direction and 2400 cells in the axial direction in. an. uniform region. with physical dimensions of 20x300 Aj.," The numerical grid of the cylindrical geometry simulations is formed by 320 cells in the radial direction and 2400 cells in the axial direction in an uniform region, with physical dimensions of $\times\,300$ $R_{\rm j}$."
 An expanded grid with 160 cells in the transversal direction brings the boundary from 204; to 60Rj. whereas an extended grid in the axial direction composed by 480 extra cells spans the grid axially from 300R; to 450Aj.," An expanded grid with 160 cells in the transversal direction brings the boundary from $20\,R_{\rm j}$ to $60\,R_{\rm j}$, whereas an extended grid in the axial direction composed by 480 extra cells spans the grid axially from $300\,R_{\rm j}$ to $450\,R_{\rm j}$."
 The enlargement of the grid is done to take the boundary conditions far enough from the region of interest and avoid numerical reflection of waves in the boundaries affecting our results., The enlargement of the grid is done to take the boundary conditions far enough from the region of interest and avoid numerical reflection of waves in the boundaries affecting our results.
 The numerical resolution in the uniform grid is thus of 16 cells/A; in the radial direction and 8 cells/R; in the axial direction., The numerical resolution in the uniform grid is thus of 16 $R_{\rm j}$ in the radial direction and 8 $R_{\rm j}$ in the axial direction.
 Outflow boundary conditions are used on the outer boundaries of the grid. inflow at injection. and reflection at the jet axis in the cylindrical case.," Outflow boundary conditions are used on the outer boundaries of the grid, inflow at injection, and reflection at the jet axis in the cylindrical case."
 In the simulations. all the physical variables are scaled to the units of the code. which are the jet radius. Aj. the rest-mass density of the ambient medium. and the speed of light.," In the simulations, all the physical variables are scaled to the units of the code, which are the jet radius, $R_{\rm j}$, the rest-mass density of the ambient medium, and the speed of light."
" The jet is injected in the grid at a distance of 6xI0'""cm from the compact object. and its initial radius is taken to be Rin=0X10°cm."," The jet is injected in the grid at a distance of $6\times 10^{10} \rm{cm}$ from the compact object, and its initial radius is taken to be $R_{\rm j,0}=6\times 10^{9} \rm{cm}$."
 The time unit of the code is thus equivalent to 0.2 seconds. as derived from the radius of the jet at injection and the speed of light (Aj.o/c).," The time unit of the code is thus equivalent to 0.2 seconds, as derived from the radius of the jet at injection and the speed of light $R_{\rm j,0}/c$ )."
 The grid covers the distance between 6x10!em and 2x10em. i.e. a of Ras.," The grid covers the distance between $6\times 10^{10}\,\rm{cm}$ and $2\times
10^{12}\,\rm{cm}$, i.e., a of $R_{\rm orb}$."
 The ambient medium. re. the stellar wind. is composed by a gas with thermodynamical properties derived from Sect.," The ambient medium, i.e. the stellar wind, is composed by a gas with thermodynamical properties derived from Sect."
 ??. (see Table 1))., \ref{phys} (see Table \ref{tab2}) ).
 Both the jet and the ambient medium are considered to be formed by a non relativistic gas with adiabatic exponent D25/3., Both the jet and the ambient medium are considered to be formed by a non relativistic gas with adiabatic exponent $\Gamma=5/3$.
 In the simulations the jets are injected with different properties., In the simulations the jets are injected with different properties.
 The physical parameters that characterize. each simulation are listed in Table 2.. the main difference between the three cases being the kinetic luminosity (Lj): weak jet (case D of 3x107 erg s!i mild jet (case ID of 1079 erg s7!: powerful jet (case ILL) of 3x107 erg s!.," The physical parameters that characterize each simulation are listed in Table \ref{tab1}, the main difference between the three cases being the kinetic luminosity $L_{\rm
j}$ ): weak jet (case I) of $3\times 10^{34}$ erg $^{-1}$; mild jet (case II) of $10^{36}$ erg $^{-1}$; powerful jet (case III) of $3\times 10^{37}$ erg $^{-1}$."
 The selection of the jet power Is addressed to show some illustrative cases: the lowest value. that of case L is similar to the minimum power required to power radio emission under reasonable assumptions of the radio emitter (e.g. for Cygnus X-1. see Heinz 20060): that of case I] is a bit smaller than the power estimates found for Cygnus X-1 (Gallo et al. 2005)).," The selection of the jet power is addressed to show some illustrative cases: the lowest value, that of case I, is similar to the minimum power required to power radio emission under reasonable assumptions of the radio emitter (e.g. for Cygnus X-1, see Heinz \cite{heinz06}) ); that of case II is a bit smaller than the power estimates found for Cygnus X-1 (Gallo et al. \cite{gallo05}) ),"
 and similar to the Lj inferred for LS 5039 (Paredes et al. 2006)).," and similar to the $L_{\rm j}$ inferred for LS 5039 (Paredes et al. \cite{paredes06}) ),"
 and the maximum value. that of case ΠΠ. 1s between the upper limit for the Cygnus Lj and the LZ; of SS 433 (e.g. Gallo et al. 2005::," and the maximum value, that of case III, is between the upper limit for the Cygnus X-1 $L_{\rm j}$ and the $L_{\rm j}$ of SS 433 (e.g. Gallo et al. \cite{gallo05};"
 Marshall et al. 2002))., Marshall et al. \cite{marshall02}) ).
 These jets are evolved until they have reached distances similar to the binary system size., These jets are evolved until they have reached distances similar to the binary system size.
 The velocities of the Jets were selected as those of mildly relativistic flows. with Lorentz factors y;=1.1 for case L and y;—L5 for cases IL and ILL.," The velocities of the jets were selected as those of mildly relativistic flows, with Lorentz factors $\gamma_{\rm j}=1.1$ for case I, and $\gamma_{\rm j}=1.5$ for cases II and III."
 The jet Mach numbers are similar in cases I (47.02) and IL (51.05). whereas it is much smaller in ease III (9.35).," The jet Mach numbers are similar in cases I (47.02) and II (51.05), whereas it is much smaller in case III (9.35)."
 In the latter. the increase in the jet power is due to an increase in the internal energy of the jet. keeping the same jet speed as in case IL.," In the latter, the increase in the jet power is due to an increase in the internal energy of the jet, keeping the same jet speed as in case II."
 The jet temperatures also show this fact. as the jet in case III has a temperature larger than that of case II by a factor 30.," The jet temperatures also show this fact, as the jet in case III has a temperature larger than that of case II by a factor 30."
 This allows us to see the influence of changes of velocity and internal energy on the Jet evolution., This allows us to see the influence of changes of velocity and internal energy on the jet evolution.
 In the case of the slab geometry simulations. the grid size is shortened to 200A; along the jet axis. and it is doubled in the direction transversal to the jet axis (it extends from —60Rj to 60 ΛΑ).," In the case of the slab geometry simulations, the grid size is shortened to $200\,R_{\rm j}$ along the jet axis, and it is doubled in the direction transversal to the jet axis (it extends from $-60\,R_{\rm j}$ to $60\,R_{\rm j}$ )."
 The numerical resolution is the same as that used in the cylindrical simulations., The numerical resolution is the same as that used in the cylindrical simulations.
 The jet is injected in an ambient medium that mimics the presence of a spherical wind centered in a star at a distance Ry from the compact object on the orbital plane. with the properties mentioned in Sect. ??..," The jet is injected in an ambient medium that mimics the presence of a spherical wind centered in a star at a distance $R_{\rm orb}$ from the compact object on the orbital plane, with the properties mentioned in Sect. \ref{phys}."
 In this case. the boundary coiditions are injection on the side from which the wind is assuned to come and on the base of the jet. and outflow in the outer boundaries on the opposite side of the grid and at the end of the grid in the axial direction.," In this case, the boundary conditions are injection on the side from which the wind is assumed to come and on the base of the jet, and outflow in the outer boundaries on the opposite side of the grid and at the end of the grid in the axial direction."
 This kind of simulations were done for cases LandIP... since the interaction with the wind was expected to be possibly relevant for the jet dynamics.," This kind of simulations were done for cases I and, since the interaction with the wind was expected to be possibly relevant for the jet dynamics."
 The jet in case I needs t~900s to cover the distance between the injection and the outer boundary of the numerical erid.," The jet in case I needs $t_{\rm f}\sim900\, \rm{s}$ to cover the distance between the injection and the outer boundary of the numerical grid."
" The bow shock generated by the injection of the jet in the ambient medium moves at a mean speed of V, c. The jet generates a high pressure cocoon that keeps it"," The bow shock generated by the injection of the jet in the ambient medium moves at a mean speed of $V_{\rm bs}\sim 0.06\,c$ The jet generates a high pressure cocoon that keeps it"
Due to systematic effects. cach Parkes optical. spectrum sullers from ai considerable amount of noise at. the range of observed waveleneths.,"Due to systematic effects, each Parkes optical spectrum suffers from a considerable amount of noise at the range of observed wavelengths."
 Although the noise may have some waveleneth dependence. we assume it is constant.," Although the noise may have some wavelength dependence, we assume it is constant."
" Given this assumption. we have used an optimal method. that. gives. more weight to. those wavelengths in a particular spectrum. where the residual m""(ΑΧ) (see Fig. 3))"," Given this assumption, we have used an optimal method that gives more weight to those wavelengths in a particular spectrum where the residual $\left| f_{q}(\lambda)-f_{T}(\lambda)\right|$ (see Fig. \ref{schematic}) )"
 is likely to be a maximum., is likely to be a maximum.
 In other words. the galaxy subtraction process will depend most sensitively on the spectral shape around. the 4000.A. breakregion," In other words, the galaxy subtraction process will depend most sensitively on the spectral shape around the $\,$ break region."
" Qurassumplionthallbenoiscisiraectenglbindagendenfiiltareallysimpli random fyouro""confidence thegalaxysinccolheruwise. wewo"," Our assumption that the noise is wavelength independent will greatly simplify our optimisation method, since otherwise, we would have to simultaneously optimise those observed spectral regions with the highest signal-to-noise."
 The basis of this method. simply involves convolving the integrand in Eqn., The basis of this method simply involves convolving the integrand in Eqn.
" 7 with a weighting function. CA) jt gives more weight to regions on either side of the 4000A breakwithin A,À«€Aw (see Fig. 3))."," \ref{Atwo} with a weighting function $G(\lambda)$ that gives more weight to regions on either side of the $4000$ $\,$ break within $\lambda_{min}<\lambda<\lambda_{max}$ (see Fig. \ref{schematic}) )."
 We 'hoose to define CA) purely from the galaxy spectrum fyfA) Cdashed: curve in Fig. 1))., We choose to define $G(\lambda)$ purely from the galaxy spectrum $f_{g}(\lambda)$ (dashed curve in Fig. \ref{galspec}) ).
 Phis is cefinecd as the difference (or residual) between a power-law and the galaxy spectrum within the range Aji<A«Amer. containing the 40004 break: where and The function. (A). will peak at wavelengths. on either side of the 40004 break.," This is defined as the difference (or residual) between a power-law and the galaxy spectrum within the range $\lambda_{min}<\lambda<\lambda_{max}$, containing the $4000$ $\,$ break: where and The function $G(\lambda)$ will peak at wavelengths on either side of the $4000$ $\,$ break."
Thus. byconvoleingGA) with Eqn. 7..," Thus, by convolving $G(\lambda)$ with Eqn. \ref{Atwo},"
 relatively more weight will be given to spectral data at these wavelengths. where our algorithm mostly sensitively depends.," relatively more weight will be given to spectral data at these wavelengths, where our algorithm mostly sensitively depends."
 ‘Thus. the function to minimise can now be written: where CG(A) is defined. by Eqn. 22.1..," Thus, the function to minimise can now be written: where $G(\lambda)$ is defined by Eqn. \ref{G}."
 Our aim is to determine the value of e for a particular source spectrum fFr(A) such that cf is a minimum., Our aim is to determine the value of $c$ for a particular source spectrum $f_{T}(\lambda)$ such that $A$ is a minimum.
 In Fig. 4..," In Fig. \ref{mineg},"
 we show examples of the figure of merit. function. οἱ (Eqn. 9)), we show examples of the figure of merit function $A$ (Eqn. \ref{Athree}) )
 as afunction. of the parameter ¢ for two sources whose spectra are shown in Fig. 5., as afunction of the parameter $c$ for two sources whose spectra are shown in Fig. \ref{fits1}.
 Lf the weighting factor GA) in Eqn., If the weighting factor $G(\lambda)$ in Eqn.
 9. is neglected. we find that the values of e at which cf is a minimum decrease bv up to r45'A in most sources where the break feature appears relatively weak or absent.," \ref{Athree} is neglected, we find that the values of $c$ at which $A$ is a minimum decrease by up to $5\%$ in most sources where the break feature appears relatively weak or absent."
 In sources with strong 40004 breaks. thereisnosigni ficantdifference," In sources with strong $\,$ breaks, there is no significant difference."
 Civenthevalucof elhmaainAdi scayiteodflecoLobi guahtenasviéreMWoobn Ilesmliewe velenglheanm," Given the value of $c$ that minimises $A$, the fractional galaxy contribution at some wavelength can now be computed using Eqn. \ref{galfrac}."
 In spectra where the galaxy contribution is relatively weak. an upper limit on its contribution at some appropriate level of significance would be required.," In spectra where the galaxy contribution is relatively weak, an upper limit on its contribution at some appropriate level of significance would be required."
 La order to do so. both random and systematic errors need to be investigated.," In order to do so, both random and systematic errors need to be investigated."
 The magnitude of these two types of errors are estimated: and compared in section 4.2.., The magnitude of these two types of errors are estimated and compared in section \ref{sys.vs.ran}.
 Here. we brielly outline the method. in their determination.," Here, we briefly outline the method in their determination."
" 1e ""error “at sómé Vimisatiqnaediod: level in contribution is estimated by computing the value. of ¢ corresponding to the statistical error in «d (eqn. 9)).", The random error at some confidence level in the galaxy contribution is estimated by computing the value of $c$ corresponding to the statistical error in $A$ (Eqn. \ref{Athree}) ).
 Since Eqn., Since Eqn.
 9 is actually a discrete sum overwavelength bins A;. from Ain to Ayer. the statistical error is determined. by acdcding the error for cach individual bin in quadrature. so that where m;(1) is the error in the integrand £ of Eqn.," \ref{Athree} is actually a discrete sum overwavelength bins $\lambda_{i}$, from $\lambda_{min}$ to $\lambda_{max}$, the statistical error is determined by adding the error for each individual bin in quadrature, so that where $\sigma_{i}(I)$ is the error in the integrand $I$ of Eqn."
 for bin À;., \ref{Athree} for bin $\lambda_{i}$.
" e;(4) will depend. only on uncertainties. in the measured source. Duxes. fr(À,,,,) and fr(A,;,) in Eqn. 9..", $\sigma_{i}(I)$ will depend only on uncertainties in the measured source fluxes $f_{T}(\lambda_{max})$ and $f_{T}(\lambda_{min})$ in Eqn. \ref{Athree}.
" As shall be discussed in section 3.1.. these Duxes are estimated by calculating the median continuum Lux in wavelength bins centered on Aya, and À,,;;."," As shall be discussed in section \ref{apal}, these fluxes are estimated by calculating the median continuum flux in wavelength bins centered on $\lambda_{max}$ and $\lambda_{min}$."
 We estimate the corresponding uncertainties by computing the rms deviation from the mean spectral (ux in these wavelength bins., We estimate the corresponding uncertainties by computing the rms deviation from the mean spectral flux in these wavelength bins.
 Our definition of the underlying quasar continuum f(A). in each source only requires that it be smooth and contain no breaks.," Our definition of the underlying quasar continuum $f_{q}(\lambda)$, in each source only requires that it be smooth and contain no breaks."
" This implies that the shape of f£,(ÀA) is somewhat arbitrary and thus it ds. possible that our. quantitative measures of the galaxy contribution may strongly depend on its assumption in our algorithm.", This implies that the shape of $f_{q}(\lambda)$ is somewhat arbitrary and thus it is possible that our quantitative measures of the galaxy contribution may strongly depend on its assumption in our algorithm.
 “Phe uncertainty introduced w this possible systematic elfect will be investigated., The uncertainty introduced by this possible systematic effect will be investigated.
 All our calculations assume that f(A) is à power-law (ic., All our calculations assume that $f_{q}(\lambda)$ is a power-law (ie.
 Eqn. 3)).," Eqn. \ref{pl2}) ),"
" since the continua of optical quasar spectra are well represented by f;xv"" where a~0.20.8 (Erancis 1996).", since the continua of optical quasar spectra are well represented by $f_{\nu}\propto\nu^{-\alpha}$ where $\alpha\sim0.2-0.3$ \cite{Francis1996}.
". To investigate the effects of assuming a cillerent form or f(A) however. we also apply our algorithm by assuming or simplicity that f(A) is a straight line joining À,,;, and Mince W Fig.5 3. "," To investigate the effects of assuming a different form for $f_{q}(\lambda)$ however, we also apply our algorithm by assuming for simplicity that $f_{q}(\lambda)$ is a straight line joining $\lambda_{min}$ and $\lambda_{max}$ in Fig. \ref{schematic}. ."
"This is parameterised as follows: where JUSf,(À,,;,) and fF,Cus) are delined by Eqn. 5.0", This is parameterised as follows: where $f_{q}(\lambda_{min})$ and $f_{q}(\lambda_{max})$ are defined by Eqn. \ref{fqminmax}. .
 X comparison in the fractional galaxy contributions resulting from our use of a power-law (Ίσα. 3)), A comparison in the fractional galaxy contributions resulting from our use of a power-law (Eqn. \ref{pl2}) )
 and a straight line (Eqn. 11)), and a straight line (Eqn. \ref{sl}) )
 for f(A) in our algorithm. will enable us to," for $f_{q}(\lambda)$ in our algorithm, will enable us to estimate the magnitude of this systematic effect."
, Results
To combine this with the ΟΝΕ constraint. a shaded ellipse representing the 2-parameter joiut confidence region is overplottecl. asstuuitus& that the true model deseribiug the uuiverse is given by the SM aud [7.4] = [0.057. Aloy=0 .7)].,"To combine this with the CMB constraint, a shaded ellipse representing the 2-parameter joint confidence region is overplotted, assuming that the true model describing the universe is given by the SM and $\tau, A$ ] = [0.057, $A(\sigma_8 = 0.7$ )]."
" This ellipse is narrow enough that we show a uagnuilied version of Figure 1 in Figure 2.. wit1 an additional lighter band. nested within the AS band. showing the effect of varying only. fii wlile fixine f, to its SM value."," This ellipse is narrow enough that we show a magnified version of Figure \ref{fig:taua} in Figure \ref{fig:tauamag}, with an additional lighter band, nested within the AS band, showing the effect of varying only $f_{esc}$ while fixing $f_\star$ to its SM value."
 We see that the CMB does uot really coustrain 7 or A separately at all. a near-degeueracy. that was expected from the discussion in the previous sec‘tion.," We see that the CMB does not really constrain $\tau$ or $A$ separately at all, a near-degeneracy that was expected from the discussion in the previous section."
 However. tle combination of the CMB coulideuce region aud the AS band is much more constraining: this tralslates to a l-o error of about 0.02 for ox. which is noticeably better than the «‘orrespouding valie of ~ 0.2 from the CMB ellipse alone.," However, the combination of the CMB confidence region and the AS band is much more constraining: this translates to a $\sigma$ error of about 0.02 for $\sigma_8$ , which is noticeably better than the corresponding value of $\sim$ 0.2 from the CMB ellipse alone."
 We uow extend Figure 2 to connect. thugh 7. twopriort independent parameters. fi and A. shown ii Figure 3..," We now extend Figure \ref{fig:tauamag} to connect, through $\tau$, two independent parameters, $f_{esc}$ and $A$, shown in Figure \ref{fig:fesca}."
 TIe purpose of this plot is to probe the potential of cosmology aud the astroplivsies of reionization t«) constraiu each other. given a reiouization model.," The purpose of this plot is to probe the potential of cosmology and the astrophysics of reionization to constrain each other, given a reionization model."
 The ellipse here. as in Figure L. reveals the iiherent degeneracy between 7-related quantities aud A through the long narrow ellipse.," The ellipse here, as in Figure \ref{fig:taua}, reveals the inherent degeneracy between $\tau$ -related quantities and $A$ through the long narrow ellipse."
 We now «overplot the DSF permitted baud for fe. (Q.1-0.2): clearly even this approximate range ni fi. c«siderably narrows the allowed range in 4., We now overplot the DSF permitted band for $f_{esc}$ (0.1–0.2); clearly even this approximate range in $f_{esc}$ considerably narrows the allowed range in $A$.
 We note again that our choice of the DSF range for fese Was motivated [or the reasons outlined earlier. aud that other ranges for [εις are possible: 11e thai point tliat is demonstrated by Figure 3. is thepower of using auy such band. of indepeudeuly known astroplivsies to Coustrain a cosmological parameter.," We note again that our choice of the DSF range for $f_{esc}$ was motivated for the reasons outlined earlier, and that other ranges for $f_{esc}$ are possible; the main point that is demonstrated by Figure \ref{fig:fesca} is thepower of using any such band of independently known astrophysics to constrain a cosmological parameter."
 Note, Note
doppler boosted while the couuter-jet is doppler de-hoosted to such au extent it is invisible at the current scusitivity of VLBI observations.,"doppler boosted, while the counter-jet is doppler de-boosted to such an extent it is invisible at the current sensitivity of VLBI observations."
" The jets are believed to originate in the accretion disk surrounding a central supermassive black hole. which for 5501 has been suggested to have a mass of ΠΟΣΤΟ ΕΛΙΤ, (Falomo. Kotibüneu Treves 2002)."," The jets are believed to originate in the accretion disk surrounding a central super-massive black hole, which for 501 has been suggested to have a mass of $^{8.93\pm0.21}$ $_\odot$ (Falomo, Kotilainen Treves 2002)."
 A πο of differcut iuotious have becu reported for components in the 5501 jet. raneie youn 0.2720.02 vvr1 (Cabuzdaal.1991) to | (Cüovanuninietal. 1999).," A number of different motions have been reported for components in the 501 jet, ranging from $\pm$ $^{-1}$ \citep{gab94}
 to $^{-1}$ \citep{gio99}."
. These studies have relied. ou ouly à siall uunuber of epochs. typically less than four. aud. in hindsielt. have probably uncler-cstimated the errors in locating jet components.," These studies have relied on only a small number of epochs, typically less than four, and, in hindsight, have probably under-estimated the errors in locating jet components."
 As illustrated bv Pineretal.(1999). for 1121. reliable deteriunation of component notions eenerallv requires larger. iuulti-epoch data sets.," As illustrated by \citet{pin99} for 421, reliable determination of component motions generally requires larger, multi-epoch data sets."
 Tere we study the parsec-scale jet of 5501 from twelve VEDI observations spanning 1.28 vears. at frequencies of5. 8 aud GCIz.," Here we study the parsec-scale jet of 501 from twelve VLBI observations spanning 4.28 years, at frequencies of 5, 8 and GHz."
 5501 lies at a redshift of 0.051 (Wills&Wills1971) which. for the value of Uy=65 | | adopted throughout his paper. corresponds to a distance of MMpce.," 501 lies at a redshift of 0.034 \citep{wil74} which, for the value of $_0 = 65$ $^{-1}$ $^{-1}$ adopted throughout his paper, corresponds to a distance of Mpc."
 At this distance. an angular separation of Limmas corresponds to a projected linear distance of ppc.," At this distance, an angular separation of mas corresponds to a projected linear distance of pc."
 We have compiled data from three programs for this study: the Very Loug Bascline Array (VLBA) δρα survey Wellermannuetal.(1998.seealso ww.ev.anmo.edu/2emsurvey).. the Radio Reference Frame Tage Database (Fev&Charlot1997.seealsohttp:/'rorf.usuo.navy.ual ΣΕ. and a single VLBI Space Observatory Programe (VSOP:Ilrabavashietal.1998.2000) observation.," We have compiled data from three programs for this study: the Very Long Baseline Array (VLBA) cm survey \citet[][see also
http://www.cv.nrao.edu/2cmsurvey]{kel98}, the Radio Reference Frame Image Database \citep[][see also
http://rorf.usno.navy.mil/rrfid.shtml]{fey97}, and a single VLBI Space Observatory Programme \citep[VSOP;][]{hir98,hir00} observation."
 The data used in this study are suuuuarlzed in Table 1., The data used in this study are summarized in Table 1.
 The VLBA σοι survey is being undertalkeu at multiple epochs to study the properties and evolution of over LOO active galactie uuclei., The VLBA cm survey is being undertaken at multiple epochs to study the properties and evolution of over 100 active galactic nuclei.
 The 5501 observations consisted of eight scaus at oue hour intervals of typically 5 iiuutes duration., The 501 observations consisted of eight scans at one hour intervals of typically 5 minutes duration.
 Data were recordec with a bandwidth of NMIIz using 1 bit samples and left-circulay polarization., Data were recorded with a bandwidth of MHz using 1 bit samples and left-circular polarization.
 An image derived roni the March 1997 data used in this paper was mublished by(1998)., An image derived from the March 1997 data used in this paper was published by.
. The full teu-station array was used at all epochs except the last. for which the Saint Croix telescope was unavailable.," The full ten-station array was used at all epochs except the last, for which the Saint Croix telescope was unavailable."
" The Radio Reference Frame huage. Database (RRFID) of the NNaval Observatory (USNO) isa program to regularlv image the radio sources uscd for precise astrometiy,", The Radio Reference Frame Image Database (RRFID) of the Naval Observatory (USNO) is a program to regularly image the radio sources used for precise astrometry.
 The VLBA Was used at all epochs. although the October 1995 observation was iade without the Manna Ἴνοα aud North Liberty clescopes.," The VLBA was used at all epochs, although the October 1995 observation was made without the Mauna Kea and North Liberty telescopes."
 The June 1995 observation was made with the addition of the Fairbanks 2610 (Ala). Green Bau tun (West Virginia). Iokec Park 20nuu (Jawai). Medicina 321unu (Italy). Nv Alesuud 20.1 (Norway). Ousala 1211 (Sweden). aud Westford l8nunu (Massachusetts) telescopes.," The June 1998 observation was made with the addition of the Fairbanks m (Alaska), Green Bank m (West Virginia), Kokee Park m (Hawaii), Medicina m (Italy), Ny Alesund m (Norway), Onsala m (Sweden), and Westford m (Massachusetts) telescopes."
 Typically. four scans of ~3 müuutes were made. with bandwidths of NIMIIz for the first two epochs. MMIIZ for the SCCdIz observations of the last two epochs. and GLAM2 for the 15CGGIIz observation in April 1996.," Typically, four scans of $\sim$ 3 minutes were made, with bandwidths of MHz for the first two epochs, MHz for the GHz observations of the last two epochs, and MHz for the GHz observation in April 1996."
 Right circular polarization is recorded for all RRFID observations., Right circular polarization is recorded for all RRFID observations.
 Au image from the 1995 April 12 epoch at GCGIIz was published by Fey&Charlot(1997)., An image from the 1995 April 12 epoch at GHz was published by \citet{fey97}.
 The GGIIz VSOP observation. in April 1998. was lmade over a 13 hour period with the TALCA satellite. the VLBA and the Effelsbere nuu (Conuauv) telescope.," The GHz VSOP observation, in April 1998, was made over a 13 hour period with the HALCA satellite, the VLBA and the Effelsberg m (Germany) telescope."
 Luterferometric fringes to he satellite were detected from tracking passes otaling hours., Interferometric fringes to the satellite were detected from tracking passes totaling hours.
 In the standard VSOP observing node. MIAMIz of two-bi sampled. left circular xobuization data is recorded.," In the standard VSOP observing mode, MHz of two-bit sampled, left circular polarization data is recorded."
 Although made at ie lowest frequency considered here. the long vasclines to the orbiting telescope result in the svuthesized beam-size for this observation beige re Stnallest of these data.," Although made at the lowest frequency considered here, the long baselines to the orbiting telescope result in the synthesized beam-size for this observation being the smallest of these data."
 The VSOP cdataconsidered here were combined with a GCTIIz VSOP observation one dav later to derive a spectral index map of the source (Edwardsetal. 2O00a)., The VSOP dataconsidered here were combined with a GHz VSOP observation one day later to derive a spectral index map of the source \citep{edw00a}.
. The data were fringe-fit in AIPS aud we have dmaeged all data ourselves using the Difinap package., The data were fringe-fit in AIPS and we have imaged all data ourselves using the Difmap package.
 Beas were calculated using natural weieliting liu παρ). with the exception of the VSOP. observation. for which uniform weighting Givweight=2.0 iu Difimap) is more appropriate (Ilirabawvashietal. 2000).," Beams were calculated using natural weighting $-$ 1 in Difmap), with the exception of the VSOP observation for which uniform weighting (uvweight=2,0 in Difmap) is more appropriate \citep{hir00}. ."
. Au, An
calculate the vortex mobility. which we apply to obtain the growth rate of the instability.,"calculate the vortex mobility, which we apply to obtain the growth rate of the instability."
 We conclude with a discussion of the possibility that the inner-crust superlluid becomes turbulent., We conclude with a discussion of the possibility that the inner-crust superfluid becomes turbulent.
 First we caleulate the critical velocity between the superfluid and the crust that can be sustained by vortex pinnine., First we calculate the critical velocity between the superfluid and the crust that can be sustained by vortex pinning.
 We will use these results in 86 to caleulate the vortex mobility and the growth rate of the instability., We will use these results in 6 to calculate the vortex mobility and the growth rate of the instability.
 Vortex pinning fixes the local superlluid velocity in the laboratory frame., Vortex pinning fixes the local superfluid velocity in the laboratory frame.
 As the crust spins down. a velocity dilference p between the pinned vortices and the superlluid develops.," As the crust spins down, a velocity difference $v$ between the pinned vortices and the superfluid develops."
 Ehe Magnus force per unit length of vortex is where p is the superlluidmass density. &—hm is the quantum of vorticity. and m is twice the neutron mass.," The Magnus force per unit length of vortex is where $\rho$ is the superfluidmass density, $\kappa\equiv h/m$ is the quantum of vorticity, and $m$ is twice the neutron mass."
" Let £j, be the characteristic force of the vortex-nucleus potential.", Let $F_p$ be the characteristic force of the vortex-nucleus potential.
 Above a critical velocity dilference e. the Magnus force will exceed the pinning force. ancl vortex pinning is not. possible.," Above a critical velocity difference $v_c$, the Magnus force will exceed the pinning force, and vortex pinning is not possible."
" Loa vortex could bend to intersect nuclei of average spacing e. the critical velocity dilference 0, would be given by A vortex has a laree self energy (tension) that tvpically prevents it from bending over a length scale e."," If a vortex could bend to intersect nuclei of average spacing $a$, the critical velocity difference $v_c$ would be given by A vortex has a large self energy (tension) that typically prevents it from bending over a length scale $a$."
 H£ the tension were infinite. à vortex could not pin at all. since the vortex would remain straight and the forces from nuclei that surround the vortex would cancel on average.," If the tension were infinite, a vortex could not pin at all, since the vortex would remain straight and the forces from nuclei that surround the vortex would cancel on average."
 For finite tension. the vortex can bend over a length ἐν2e. and the critical velocity ds eiven instead by eiving a critical velocity Tension lowers the critical velocity by a factor afi.," For finite tension, the vortex can bend over a length $l_p>a$, and the critical velocity is given instead by giving a critical velocity Tension lowers the critical velocity by a factor $a/l_p$."
 To caleulate ἐν. let 7.(2) bea vector in ther oy plane that gives the shape of the pinned vortex.," To calculate $l_p$, let $\rbf_v(z)$ be a vector in the $x-y$ plane that gives the shape of the pinned vortex."
" Phe energy of a static vortex in a pinning field. V(r,). in the absence of an ambient superlluid How. is where 1; is vortex tension. typically 1 MeV. +."," The energy of a static vortex in a pinning field $V(\rbf_v)$, in the absence of an ambient superfluid flow, is where $T_v$ is vortex tension, typically 1 MeV $^{-1}$."
" On average. over alength ἐν the vortex bends by an amount dr. to intersect one nucleus in a volume /,sx(89r«)7."," On average, over alength $l_p$ the vortex bends by an amount $\delta r_v$ to intersect one nucleus in a volume $l_p\pi (\delta r_v)^2$."
 The quantities ἐν and or. are therefore related by The energy of the vortex per unit length. from eq. (5)).," The quantities $l_p$ and $\delta r_v$ are therefore related by The energy of the vortex per unit length, from eq. \ref{e}) ),"
" is approximately where ££, is the interaction energy. between a vortex and a single nucleus. tvpically ~1 MeV. Contributions to the potential by nuclei that the vortex does not intersect have been ignored: these contributions will largely cancel."," is approximately where $E_p$ is the interaction energy between a vortex and a single nucleus, typically $\sim 1$ MeV. Contributions to the potential by nuclei that the vortex does not intersect have been ignored; these contributions will largely cancel."
" Minimization of Eff, with respect to ἐν. using eq. (6))."," Minimization of $E_v/l_p$ with respect to $l_p$, using eq. \ref{mfp}) ),"
" gives The vortex tension 7) is due mainly to the kinetic cncrey per unit length of vortex due to circulation abut the vortex. and takes the form (??).. where £is the radius of the vortex core and A, is the characteristic bending wavenumber. 4.—x /2f,."," gives The vortex tension $T_v$ is due mainly to the kinetic energy per unit length of vortex due to circulation abut the vortex, and takes the form \citep{thomson1880,fetter67}, where $\xi$is the radius of the vortex core and $k_v$ is the characteristic bending wavenumber, $k_v=\pi/2l_p$ ."
" For typical conditions of the inner crust. the ratio /,/e is much larger than unity."," For typical conditions of the inner crust, the ratio $l_p/a$ is much larger than unity."
 Ata density p=5.107 the lattice spacing is e250 fm and the radius of the vortex core is £210 fm., At a density $\rho=5\times 10^{13}$ the lattice spacing is $a\simeq 50$ fm and the radius of the vortex core is $\xi\simeq 10$ fm.
" For £,=1 MeV. simultaneous solution of eqs. (8))"," For $E_p=1$ MeV, simultaneous solution of eqs. \ref{lp}) )"
 and. (9)) gives ἐν2 9a., and \ref{tension}) ) gives $l_p\simeq 9a$ .
 Phe ratio νέα increases for weaker pinning., The ratio $l_p/a$ increases for weaker pinning.
" For example. for £j,=0.1 MeV. the pinning length becomes ἐν2: 32a."," For example, for $E_p=0.1$ MeV, the pinning length becomes $l_p\simeq 32a$ ."
" For E,=10 MeV. unrealistically large according to recent calculations. /,,= 2a."," For $E_p=10$ MeV, unrealistically large according to recent calculations, $l_p=2a$ ."
 Combining eq. (8)), Combining eq. \ref{lp}))
 with eq. (4)), with eq. \ref{vc0}) )
 gives the eritical velocity. modified hy vortex tension.," gives the critical velocity, modified by vortex tension,"
where is the count rate for the N points in cach light curve. with errors aud is the unweighted arithmetic mean of theje.,"where is the count rate for the N points in each light curve, with errors and is the unweighted arithmetic mean of the."
 The averageOo values of the excess variance (plus dispersion) in the soft (0.1-2 keV) and hard (2-10 keV) ⋝⋜⋯≼↧⋜⋯∖↥⋅↸∖↻∪↥⋅↑↸∖≺↧↴⋝↸∖↕∪↖↖⇁≺∐⊔∐∐↑↴∖↴∪⊔∩−⋝∙ ⋅, The average values of the excess variance (plus dispersion) in the soft (0.1-2 keV) and hard (2-10 keV) band are reported below (in units of $10^{-2}$ ).
" ⋅ ≽↽⋅≻ ∪≧∖∶≺↓∙↱⊔↼↕∶∶∩∙∶≩⊇⊇≻⊼⊺≧∶∠≩∶≺↕∙∶≩⊇∩∶∶∩∙∪≼∖∖⊓soft soft σσDN = (6.292 £ 0.502) «02,4,7g = (501 + 0.016) The excess variance m both energy bands is clearly ligher for the NLSIs.", $<\sigma_{soft}^{2}>_{N}$ = (4.571 $\pm$ 0.322) $<\sigma_{soft}^{2}>_{B}$ = (1.320 $\pm$ 0.081) $<\sigma_{hard}^{2}>_{N}$ = (6.292 $\pm$ 0.592) $<\sigma_{hard}^{2}>_{B}$ = (0.504 $\pm$ 0.016) The excess variance in both energy bands is clearly higher for the NLS1s.
and therefore a lack of energetic electrons there after a typical cooling iie.,and therefore a lack of energetic electrons there after a typical cooling time.
 But the fate of the low energv electrons. necessary or the CAIB-IC EUW production. is mainly determined w the volue fraction of the low field streneth regions.," But the fate of the low energy electrons, necessary for the CMB-IC EUV production, is mainly determined by the volume fraction of the low field strength regions."
 If se. their voluue fraction is coustaut with cluster radius (nit the field streneth of the stronely magnetized region is still varving with radius}. then these electrous are still inuuerous in the cluster center.," If e.g. their volume fraction is constant with cluster radius (but the field strength of the strongly magnetized region is still varying with radius), then these electrons are still numerous in the cluster center."
 Another possible origin of an electron population secu bv CAIB-IC is in-situ acceleration by plasina waves., Another possible origin of an electron population seen by CMB-IC is in-situ acceleration by plasma waves.
 A strong wave field can be expected if there is an ou-gonig meorecr event in the center of Coma. as several authors (Colless Dunn 1996. Biviano et al.," A strong wave field can be expected if there is an on-going merger event in the center of Coma, as several authors (Colless Dunn 1996, Biviano et al."
 1996. Vikhiiuin et al.," 1996, Vikhlinin et al."
 1997. Donnelly et al.," 1997, Donnelly et al."
 L998) report., 1998) report.
 The electron προςται produced can be approximated by (Schlickeiser 198 D)). where the cutoff depends on the ratio of acceleration time scale to svuchrotron/IC cooling time.," The electron spectrum produced can be approximated by (Schlickeiser \nocite{schlickeiser84}) ), where the cutoff depends on the ratio of acceleration time scale to synchrotron/IC cooling time."
" ον is the Alfvénn velocity aud &B° the spatial diffusion cocficicut (0<6,< 1)."," $v_{\rm A}$ is the Alfvénn velocity and $\kappa \sim
B^{-\delta_\kappa}$ the spatial diffusion coefficient $0\leq\delta_\kappa\leq 1$ )."
 The cutoff p. is a monotonically increasing function of the maeuetic field streneth. which nuplies that in he case of inhomogencous fields the in-situ accelerate electron population reaches highest energies in the regions of strongest field strength.," The cutoff $p_{\rm c}$ is a monotonically increasing function of the magnetic field strength, which implies that in the case of inhomogeneous fields the in-situ accelerated electron population reaches highest energies in the regions of strongest field strength."
 This is also supported by the nore complicated dependence of à. ou the maeuetic fiek streneth: stronger fields result im larder spectra., This is also supported by the more complicated dependence of $\ale$ on the magnetic field strength: stronger fields result in harder spectra.
 The xopertv of spatial anticorrelation between clectrou aux uaeuetic field distributions. allowing a high nuniber of radio quiet. low euergv electrons. can therefore not be achieved withiu au in-situ acceleration mocel.," The property of spatial anticorrelation between electron and magnetic field distributions, allowing a high number of radio quiet, low energy electrons, can therefore not be achieved within an in-situ acceleration model."
 Iu the case of inhomogencous fields the iu-9tu acceleration iode is nore constrained than iu the homogeneous case. due to the correlation of in-situ accelerated electrons aux magnetic fields.," In the case of inhomogeneous fields the in-situ acceleration model is more constrained than in the homogeneous case, due to the correlation of in-situ accelerated electrons and magnetic fields."
 But even for homogeneous fields the optimal parameters a.=1. x=Pat). which lead to au electron ciffereutial umber iudex of 2.5 between 150 iux 300. MeVfc. and which are still allowed by the theory of in-situ acceleration developed by Schilickeiser 1).. load to an overproduction of svuchrotron emission above 100 MIIz for a Πο] strength of 6 pC compare o the observations.," But even for homogeneous fields the optimal parameters $\ale=1$, $ p_{\rm c} =
280$, which lead to an electron differential number index of 2.5 between 150 and 300 MeV/c, and which are still allowed by the theory of in-situ acceleration developed by Schlickeiser \nocite{schlickeiser84}, lead to an overproduction of synchrotron emission above 100 MHz for a field strength of 6 $\mu$ G compared to the observations."
 The reason for this is the relative softuess of the exponential cutoff of the electron population aud the extended spectral width of the svuchrotron euissio-, The reason for this is the relative softness of the exponential cutoff of the electron population and the extended spectral width of the synchrotron emission.
 We conclude that on-goimg in-situ acceleration is very unlikely the origin of a CAIB-IC scattering electron population producing the EUV excess., We conclude that on-going in-situ acceleration is very unlikely the origin of a CMB-IC scattering electron population producing the EUV excess.
 This is astonishiug in the light of evidence for iiergiue activity in the clusters ceuter., This is astonishing in the light of evidence for merging activity in the cluster's center.
 It indicates that if merecr events are respousible for the acceleration of energetic electrons. this happeus ouly caving the early phase of the merecr. probably durus the first slock-crossing time.," It indicates that if merger events are responsible for the acceleration of energetic electrons, this happens only during the early phase of the merger, probably during the first shock-crossing time."
 The difficulty of a CAMB-IC model for the EUV excess of Coma with the discrepancy between magnetic field streugth observed by Faraday rotation. aud the estimate using the IC aud svuchrotrou emission. can be overcome in a model where clectrous cool in an inhomogeucous naenetic field.," The difficulty of a CMB-IC model for the EUV excess of Coma with the discrepancy between magnetic field strength observed by Faraday rotation, and the estimate using the IC- and synchrotron emission, can be overcome in a model where electrons cool in an inhomogeneous magnetic field."
 The narrow spatial enissiou profile ix more differ to understand within such a model. since one would expect a amore extended clectron distribution. iu urtieulur since the central electrous should cool faster han the peripheral ones.," The narrow spatial emission profile is more difficult to understand within such a model, since one would expect a more extended electron distribution, in particular since the central electrons should cool faster than the peripheral ones."
" In order to decide. if this model is realistic or not. detailed simmlations of the spatial distribution of electrou acceleration during a merger core SHALE ALC required,"," In order to decide, if this model is realistic or not, detailed simulations of the spatial distribution of electron acceleration during a merger core passage are required."
 Since the EUW aud WEN excess fit roughly into a single power-law the TEX excess might also be produced by CXMB IC scattering., Since the EUV and HEX excess fit roughly into a single power-law the HEX excess might also be produced by CMB IC scattering.
 The necessary electrous should be close to or within the cucrey range visible iu the radio and therefore both populations have to be at least similar., The necessary electrons should be close to or within the energy range visible in the radio and therefore both populations have to be at least similar.
 Their cooling time due to svuchrotrou/IC losses ds of the order of 105 vears. so that continuous acceleration or very recent injection iuto this euergv range is necessary.," Their cooling time due to synchrotron/IC losses is of the order of $10^8$ years, so that continuous acceleration or very recent injection into this energy range is necessary."
 The observed TEX cussion determines the number of electrons ini the radio energv range. if one assunes the radio iudex also to be valid or the WEN producingo Jectrous.," The observed HEX emission determines the number of electrons in the radio energy range, if one assumes the radio index also to be valid for the HEX producing electrons."
 Iu order to be in agreement with the observe svuchrotron euiüssion. the volume averaged magnetic fields Sreugth has to be 0.16. Co (Fusco-Foniiano et al.," In order to be in agreement with the observed synchrotron emission, the volume averaged magnetic fields strength has to be 0.16 $\mu$ G (Fusco-Femiano et al."
 1998. 1999). comparable to Ο.Ε iC: given by Thvang (1997) for the EUW emission. which is energetically more distaut to tie radio range.," 1998, 1999), comparable to 0.4 $\mu$ G given by Hwang (1997) for the EUV emission, which is energetically more distant to the radio range."
 The ceutral maguetie field strength is roughly a factor of 3 higher. leading to z0.5 Ci. which secnus to be too low in order to be consistent with the Faraday 1ieasureineuts. of 64 nip.," The central magnetic field strength is roughly a factor of 3 higher, leading to $\approx 0.5\,\mu$ G, which seems to be too low in order to be consistent with the Faraday measurements of $6\,\mu$ $\, \h^{1/2}$."
 Also for this 3/5 GeW electron population oue might ask if (a) a sharp step in the electron spectrum. and 2) an iuliomogeneous magnetized medium might resolve this iscrepancy.," Also for this 3–5 GeV electron population one might ask if (a) a sharp step in the electron spectrum, and (b) an inhomogeneous magnetized medium might resolve this discrepancy."
 The ICAL needs to cousist mainly out of regions with very weak fields (onlv a few 0.14) onteünmeg the CAIB IC scattering electrons. which are oeivisible iu the radio due to the weak fields.," The ICM needs to consist mainly out of regions with very weak fields (only a few $0.1\,\mu$ G) containing the CMB IC scattering electrons, which are invisible in the radio due to the weak fields."
 But highly maeguetized regious have to exist (inmavbe LOgCe or more). oei order to explain the Faraday rotation. which ποσα to contain ouly few 305 Ce electrons.," But highly magnetized regions have to exist (maybe $10\mu$ G or more), in order to explain the Faraday rotation, which need to contain only few 3–5 GeV electrons."
 If the difference in clectron content of these reeious was established by different cooling the time of injection had to be a few 0.1 Cir in order to allow the several GeV electrons iu the weak field regions still to be present. but those electrons living iu the hieh field regions to have cooled to cucreics invisible," If the difference in electron content of these regions was established by different cooling the time of injection had to be a few 0.1 Gyr in order to allow the several GeV electrons in the weak field regions still to be present, but those electrons living in the high field regions to have cooled to energies invisible"
within the low redshift range.,within the low redshift range.
 This allows us to isolate the effect of the “great wall” by comparing the low redshift range to the high redshift range., This allows us to isolate the effect of the “great wall” by comparing the low redshift range to the high redshift range.
" To determine a suitable absolute magnitude cut, we define the faint limit as the absolute magnitude where the observed luminosityMy function begins to turn over because the limiting magnitude has been reached."," To determine a suitable absolute magnitude cut, we define the faint limit $M_f$ as the absolute magnitude where the observed luminosity function begins to turn over because the limiting magnitude has been reached."
" This means that for a faint limit and limiting redshift Zmaz, the comoving number densityMy of galaxies n(My) brighter than should be the approximately the same for any limiting Myredshift z<zmaz."," This means that for a faint limit $M_f$ and limiting redshift $z_{max}$, the comoving number density of galaxies $\overline{n}(M_f)$ brighter than $M_f$ should be the approximately the same for any limiting redshift $z < z_{max}$."
 We obtain by comparing for a redshift range with a lower M;redshift subset of the (My)same range., We obtain $M_f$ by comparing $\overline{n}(M_f)$ for a redshift range with a lower redshift subset of the same range.
" For the low range, we compare the range 0.04€z<0.12 with the range 0.04€z<0.10 and for the high redshift range we compare the range 0.12€z<0.20 with the range 0.12Xz<0.18."," For the low range, we compare the range $0.04\leq z\leq 0.12$ with the range $0.04\leq z\leq 0.10$ and for the high redshift range we compare the range $0.12\leq z\leq 0.20$ with the range $0.12\leq z\leq 0.18$."
 The optimal faint limit which gives us the largest complete sample occurs whereMy the compared values of n(M;) are approximately equal., The optimal faint limit $M_f$ which gives us the largest complete sample occurs where the compared values of $\overline{n}(M_f)$ are approximately equal.
" We find that the lower redshift range is complete for M,«—19.5, M,«—20.2 and M;«—20.6 while the higher redshift range is complete for M,«—20.7, M,«—21.5 and M;«—21.9."," We find that the lower redshift range is complete for $M_g < -19.5$, $M_r < -20.2$ and $M_i < -20.6$ while the higher redshift range is complete for $M_g < -20.7$, $M_r < -21.5$ and $M_i < -21.9$."
 We plot these limits on the observed luminosity function at 0.04<z0.12 and 0.12<z0.20 in figure and summarize the subsamples we use in table , We plot these limits on the observed luminosity function at $0.04 \leq z \leq 0.12$ and $0.12 \leq z \leq 0.20$ in figure \ref{fdata-lf} and summarize the subsamples we use in table \ref{tdata-zcut}.
"Here we note that the la(g),[I]. la(r) and la(i) samples have similar spatial densities, and likewise the 1b(g), 1b(r) and and and samples also have similar 1b(i),spatial densities."," Here we note that the 1a(g), 1a(r) and 1a(i) samples have similar spatial densities, and likewise the 1b(g), 1b(r) and 1b(i), and 2b(g), 2b(r) and 2b(i) samples also have similar spatial densities."
"2b(g), 2b(r)Hence any 2b(i)large differences in clustering between the samples of different color should arise from selection effects that depend on color.", Hence any large differences in clustering between the samples of different color should arise from selection effects that depend on color.
 An analysis by |SylosLabinietαἰ.(2009) found systematic variations between different subvolumes of the SDSS catalog on scales larger than 305-7! Mpc such that these subvolumes are not statistically similar., An analysis by \citet{2009A&A...508...17S} found systematic variations between different subvolumes of the SDSS catalog on scales larger than $30 h^{-1}$ Mpc such that these subvolumes are not statistically similar.
 These variations are likely to be caused by cosmic variance that our analysis should take into account., These variations are likely to be caused by cosmic variance that our analysis should take into account.
 We use two approaches to analyze the effect of cosmic variance on our results., We use two approaches to analyze the effect of cosmic variance on our results.
 'The first is simply to consider independent subfields of the survey footprint., The first is simply to consider independent subfields of the survey footprint.
" To ensure that effects of galactic latitude, distance and lookback time are constant across all subsamples, we compute and compare the counts-in-cells distribution for cells in non-overlapping quadrants in galactic longitude."," To ensure that effects of galactic latitude, distance and lookback time are constant across all subsamples, we compute and compare the counts-in-cells distribution for cells in non-overlapping quadrants in galactic longitude."
 The size of quadrants is also comparable with the size of the SDSS “great wall” which spans ~70° , The size of quadrants is also comparable with the size of the SDSS “great wall” which spans $\sim70^\circ$ \citep{2006ChJAA...6...35D}.
"In this approach, all cells that belong to a quadrant−− are fully-contained within the selected quadrant."," In this approach, all cells that belong to a quadrant are fully-contained within the selected quadrant."
 This gives us a picture of the variations among widely separated areas of the sky., This gives us a picture of the variations among widely separated areas of the sky.
" We consider subsamples in four quadrants such that for galactic longitude J, quadrant 1 (ql) covers 0.0?<|€ 90.0°, quadrant 2 covers 90.0?<|180.0°, quadrant 3 (q3) covers 180.0?(q2)<1<270.0? and quadrant 4 (q4) covers 270.0?<1x960.05."," We consider subsamples in four quadrants such that for galactic longitude $l$, quadrant 1 (q1) covers $0.0^\circ \leq l \leq 90.0^\circ$ , quadrant 2 (q2) covers $90.0^\circ \leq l \leq 180.0^\circ$, quadrant 3 (q3) covers $180.0^\circ \leq l \leq 270.0^\circ$ and quadrant 4 (q4) covers $270.0^\circ \leq l \leq 360.0^\circ$."
" For some samples, we also examine the quadrant to quadrant variations where the quadrant boundaries boundaries have been shifted by 30°, 45? and 60? in galactic longitude to check that any variation between quadrants is not caused by our choice of quadrant boundaries."," For some samples, we also examine the quadrant to quadrant variations where the quadrant boundaries boundaries have been shifted by $30^\circ$, $45^\circ$ and $60^\circ$ in galactic longitude to check that any variation between quadrants is not caused by our choice of quadrant boundaries."
" Moreover, we have varied the size of subsamples and found that for smaller subsamples, sixths rather than quadrants, the subsample to subsample variations are smaller than for larger subsamples."," Moreover, we have varied the size of subsamples and found that for smaller subsamples, sixths rather than quadrants, the subsample to subsample variations are smaller than for larger subsamples."
" However, subsamples that are too large will not be independent, and there will be too few of them to provide an accurate estimate of the cosmic variance between disconnected regions of the sky."," However, subsamples that are too large will not be independent, and there will be too few of them to provide an accurate estimate of the cosmic variance between disconnected regions of the sky."
 'Therefore quadrants are a reasonable subsample size to use., Therefore quadrants are a reasonable subsample size to use.
" The second approach is a jackknife-style approach where we leave out cells that fall within a region of the sky, selected based on a quasi-random sequence from (1994)"," The second approach is a jackknife-style approach where we leave out cells that fall within a region of the sky, selected based on a quasi-random sequence from \citet{1994ACMTMS-20-494}"
The engine of long GRBs is thought to be produced during the collapse of a massive stellar core and thus to be launching relativistic jets into a stellar envelope.,The engine of long GRBs is thought to be produced during the collapse of a massive stellar core and thus to be launching relativistic jets into a stellar envelope.
 show that (at least in some) low-luninosity GRBs it is highly unlikely that the jets successfully. punch through the star and emerge in order to produce the observed -rav emission., \cite{Bromberg11a} show that (at least in some) low-luminosity GRBs it is highly unlikely that the jets successfully punch through the star and emerge in order to produce the observed $\g$ -ray emission.
 Iu that case what can be the source of the observed gamma-ravs?, In that case what can be the source of the observed gamma-rays?
 There are several common [features observed in all low luminosity GRBs., There are several common features observed in all low luminosity GRBs.
" All show 10—10°"" ere of gamma-rays or hard X-rays that are emitted in single pulsed light curves showing a spectra (hat become softer with lime (see Kanekoefal.2007 ancl references therein).", All show $10^{48}-10^{50}$ erg of gamma-rays or hard X-rays that are emitted in single pulsed light curves showing a spectra that become softer with time (see \citealt{Kaneko07} and references therein).
 They also show radio allerglows that indicate on similar energy in mildly relativistic ejecta (xulkarniefaf1998:Soderbergal.2004.2006).," They also show radio afterglows that indicate on similar energy in mildly relativistic ejecta \citep{Kulkarni98,Soderberg04,Soderberg06}."
. This energy. is only a small fraction of the total energv observed in the associated SNe (107?—1003 erg)., This energy is only a small fraction of the total energy observed in the associated SNe $10^{52}-10^{53}$ erg).
 The smooth light curves and the energy coupled to a mildly relativistic ejecta motivated several authors to suggest that the origin of (his emission is shock breakouts (Ixulkarnie£af1998:Tan.al... 2007).. Later.," The smooth light curves and the energy coupled to a mildly relativistic ejecta motivated several authors to suggest that the origin of this emission is shock breakouts \citep{Kulkarni98,Tan01,Campana06,Waxman07,Wang07}."
 IXatz.Dudnik&Waxman(2010) have shown that the deviation from thermal equilibrium can explain the observed s-ravs. compared to the expected X-rays when equilibrium is wrongly assumed.," Later, \cite{Katz10} have shown that the deviation from thermal equilibrium can explain the observed $\g$ -rays, compared to the expected X-rays when equilibrium is wrongly assumed."
 However. (he question whether shock breakouts are indeed the source of low-Iuminosity GRBs remained open. mostly since (here was no quantitative model that could answer whether shock breakouts can indeed explain even the most basic observables. such as energy. temperature and (nme scales. of these bursts.," However, the question whether shock breakouts are indeed the source of low-luminosity GRBs remained open, mostly since there was no quantitative model that could answer whether shock breakouts can indeed explain even the most basic observables, such as energy, temperature and time scales, of these bursts."
 Below. we use our model to answer this question.," Below, we use our model to answer this question."
 In section ?? we have shown that this three observed scales. Pro. dy and [pps over-constrain the physical parameters since they depend in the case of a quasi-spherical relativislic breakout only on (wo parameters. 550 and the breakout radius.," In section \ref{sec light curve} we have shown that this three observed scales, $E_{bo}$, $T_{bo}$ and $t_{bo}^{obs}$, over-constrain the physical parameters since they depend in the case of a quasi-spherical relativistic breakout only on two parameters, $\g_{f,0}$ and the breakout radius."
 Thus. in addition to finding the value of these (wo physical parameters. the (11ος observables must satisIv the relativistic breakout relation (equation 18)).," Thus, in addition to finding the value of these two physical parameters, the three observables must satisfy the relativistic breakout relation (equation \ref{eq tbo
test}) )."
 This can be used as a test (hat anv [lue (hat originates [rom quasi-spherical relativistic breakout must pass., This can be used as a test that any flare that originates from quasi-spherical relativistic breakout must pass.
 This result is generic as it is independent of (he pre-breakout density profile (as long as il is steeply decreasing)., This result is generic as it is independent of the pre-breakout density profile (as long as it is steeply decreasing).
 Below we apply this test to the four observed Iow-Iuminosity. GBs The four well studied low Iuminosity GRBs are divided into (wo pairs with very different properties., Below we apply this test to the four observed low-luminosity GRBs The four well studied low luminosity GRBs are divided into two pairs with very different properties.
 GRBs 060218 and 100316D are soft and very long., GRBs 060218 and 100316D are soft and very long.
 Their rather soft spectrum is best fitted at early time with a cut-off around a peak energy of ~40 keV (Ixaneko and both emitted a total energv of ~5x107? erg.," Their rather soft spectrum is best fitted at early time with a cut-off around a peak energy of $\sim 40$ keV \citep{Kaneko07,Starling11} and both emitted a total energy of $\sim
5 \times 10^{49}$ erg."
 Plugging these values into equation 18 we find that if these bursts are flares Irom quasi-spherical breakouts then they should have durations of 1500 s. which is indeed comparable to the observed durations of these two bursts.," Plugging these values into equation \ref{eq
tbo test} we find that if these bursts are flares from quasi-spherical breakouts then they should have durations of $\sim
1500$ s, which is indeed comparable to the observed durations of these two bursts."
 The breakout radius in this model is ~5xLOM em and 5j21 implying," The breakout radius in this model is $\sim 5
\times 10^{13}$ cm and $\g_{f,0} \approx 1$ implying"
a correlation between peak buuinosity and jet-weak tue.,a correlation between peak luminosity and jet-break time.
 This is. to our best knowledge. the first known direct correlation between a property of the eauuna-cray burst phase (the pulse lag or σοα]ς Iuuinositv) aud the afterglow phase (the jet-weak time).," This is, to our best knowledge, the first known direct correlation between a property of the gamma-ray burst phase (the pulse lag or peak luminosity) and the afterglow phase (the jet-break time)."
 As observed timescales (variability or spectral lags) as well as peak huminosity naturally rave a strong dependence on the Lorentz factor. 5. or anele with respect to the motion of the outflow. Ü. we propose that the variety among CRBs iis a purely kincmatic origin.," As observed timescales (variability or spectral lags) as well as peak luminosity naturally have a strong dependence on the Lorentz factor, $\gamma$, or angle with respect to the motion of the outflow, $\theta$, we propose that the variety among GRBs has a purely kinematic origin."
 The cinergence of siuple treuds between jet-break time. pulse lag and huninesity gives us clues and constraints on row such a model can be coustructed.," The emergence of simple trends between jet-break time, pulse lag and luminosity gives us clues and constraints on how such a model can be constructed."
 Frailetal.(2001) infer that the cdiffereuces among observed GRB cuereics are due to a rauge of jet opening angles., \citet{fksd01} infer that the differences among observed GRB energies are due to a range of jet opening angles.
 Within this framework we rave then shown that the fastest GRBs. with the üehest 5. have the narrowest jets (Sec. 3.1)).," Within this framework we have then shown that the fastest GRBs, with the highest $\gamma$, have the narrowest jets (Sec. \ref{burstpop}) )."
" A oxedietion of this model is that not only should here be a variation in openiug augles 06; anoug nuts. but they should be related to Lorentz actor 5 as Fj;X35B/S aud jet angles 0; have o be produced with+ a probabilitvBm P~P,ηX0LA . "," A prediction of this model is that not only should there be a variation in opening angles $\theta_j$ among bursts, but they should be related to Lorentz factor $\gamma$ as $\theta_j \propto \gamma^{-3/8}$, and jet angles $\theta_j$ have to be produced with a probability $P \sim
P_{obs}/\theta_j^2 \propto \theta_j^{-4.54}$ ."
We then have explored the possibility. that variation among CRBs. all now asstmed to be norphologically the same. is caused by variation ij observer viewing angle from the jet axis (Sec. 3.2)).," We then have explored the possibility that variation among GRBs, all now assumed to be morphologically the same, is caused by variation in observer viewing angle from the jet axis (Sec. \ref{simplejet}) ),"
 or Loreutz factor of the jet as a πιοΊο. of anele from the jet axis (Sec. 3.3))., or Lorentz factor of the jet as a function of angle from the jet axis (Sec. \ref{structuredjet}) ).
 These models qualitatively show that variation of the observed Doppler factor. D. (Eqn. 1))," These models qualitatively show that variation of the observed Doppler factor, ${\cal D}$ , (Eqn. \ref{deltaeqn}) )"
" roni a sinele jet morphology can produce the observed variations among GRD huuinosities aud nueseales: Ly,XOyeWXl/AfxTyr) D."," from a single jet morphology can produce the observed variations among GRB luminosities and timescales: $L_{pk}
\propto N_{pk} \propto 1/\Delta t \propto 1/\tau_j \propto {\cal D}$ ."
 These models also provide a natural explanation or the probability distribution of observation angles Pop. consistent with that reported by Frail (2001)., These models also provide a natural explanation for the probability distribution of observation angles $P_{obs}$ consistent with that reported by \citet{fksd01}.
. A realistic model of a sinele-burst norpholoey will have both perspective effects as iu Section 3.2. and jet-structure effects as in Section UoDe»., A realistic model of a single-burst morphology will have both perspective effects as in Section \ref{simplejet} and jet-structure effects as in Section \ref{structuredjet}.
 Set inside the collapsar progenitor model (e.g.MacFadven&Woosley1999). a cohesive picture beeins to emerge., Set inside the collapsar progenitor model \cite[e.g.][]{mcfw99} a cohesive picture begins to emerge.
 A jet driven out from the ceuter of the star will vary as a function of anele from the jet axis., A jet driven out from the center of the star will vary as a function of angle from the jet axis.
 Ejecta at the ceuter will be faster and lighter and. bearing the brunut of the as vot unknown acceleration mechanism. will be more fractured.," Ejecta at the center will be faster and lighter and, bearing the brunt of the as yet unknown acceleration mechanism, will be more fractured."
 At larger angles from the jet core. ejecta will interact with the stellar wall. and thus will be slower and will cutrain more barvons.," At larger angles from the jet core, ejecta will interact with the stellar wall, and thus will be slower and will entrain more baryons."
 Thus oue expects to see Ligh luminosities. short pulse lags aud Ligh variability as well as an carly jet break time for bursts viewed on axis. while higher viewing iuclinatious will vield lower uniünositfies. longer pulse lags. smoother bursts and later jet break times.," Thus one expects to see high luminosities, short pulse lags and high variability as well as an early jet break time for bursts viewed on axis, while higher viewing inclinations will yield lower luminosities, longer pulse lags, smoother bursts and later jet break times."
 Thus the variability-muinosity relationship (Reichartetal.2001) as well as the spectral lag-jet break time aud the spectral lae-Iunünositv (Equs.," Thus the variability-luminosity relationship \citep{rlfr+01}
 as well as the spectral lag-jet break time and the spectral lag-luminosity (Eqns."
 0. 7)) cau be raturally accomodated., \ref{t_jeqn} \ref{jpnlageqn}) ) can be naturally accomodated.
 This work was performed under the auspices of he U.S. Departinent of Encrey by University. of Califoruia Lawrence Livermore National Laboratoryunder contract W-7LO5S-ENCEbs., This work was performed under the auspices of the U.S. Department of Energy by University of California Lawrence Livermore National Laboratoryunder contract W-7405-ENG-48.
 TIG acknowledges support from the Sherman Fairchild Foundation., TJG acknowledges support from the Sherman Fairchild Foundation.
 We wish to thank Daniel Reichart aud Dale Frailfor useful comments., We wish to thank Daniel Reichart and Dale Frailfor useful comments.
(IIubeny 1958: Hubeny Lanz 1995) aud (he svnthetic spectra using the code SYNSPEC version 48 (IInbeny Lanz 1995).,(Hubeny 1988; Hubeny Lanz 1995) and the synthetic spectra using the code SYNSPEC version 48 (Hubeny Lanz 1995).
 The details of our AZ (per degree of [reedom) minimization filling procedures are discussed in detail in Sionetal.(1995) and will not be repeated here., The details of our $\chi^{2}_{\nu}$ (per degree of freedom) minimization fitting procedures are discussed in detail in \citet{sio95b} and will not be repeated here.
" To estimate physical parameters. i( is common to take (he white dwarl photospheric temperature T.;;.II surface oOgravitv logSg g. and projected rotational velocity 0,,;sin? as free parameters."," To estimate physical parameters, it is common to take the white dwarf photospheric temperature $_{eff}$, surface gravity log $g$, and projected rotational velocity $v_{rot} \sin{i}$ as free parameters."
 However. in the FUV region. gravities derived from fitting the profile of Lyman Alpha must be viewed with caution since lower n transitions of II are the least sensitive to sur[ace graviv.," However, in the FUV region, gravities derived from fitting the profile of Lyman Alpha must be viewed with caution since lower n transitions of H are the least sensitive to surface gravity."
 In this work. depending upon our confidence in the distance of a given svstem. we either adopted the prelerred estimate of the distance to each system in Table 1 by fixing it αἱ a constant value in the fitting or found the seale factor-derived distance from the best fit that was closest the preferred distance from Table 1.," In this work, depending upon our confidence in the distance of a given system, we either adopted the preferred estimate of the distance to each system in Table 1 by fixing it at a constant value in the fitting or found the scale factor-derived distance from the best fit that was closest the preferred distance from Table 1."
 Then. the WD mocel fits vield a white dwarl radius for the svstem. whieh can be converted to à mass from a mass-radius relation.," Then, the WD model fits yield a white dwarf radius for the system, which can be converted to a mass from a mass-radius relation."
 For the mass-radius relation. we used the evolutionary model grid of Matt Wood (1995) for C-O cores and hydrogen envelopes.," For the mass-radius relation, we used the evolutionary model grid of Matt Wood (1995) for C-O cores and hydrogen envelopes."
 In the absence of higher quality data and to limit the complexity of (he analvsis. we fixed (he abundances al solar.," In the absence of higher quality data and to limit the complexity of the analysis, we fixed the abundances at solar."
 Our choice to fix the abundances ab solar affects somewhat our determination of the rotational velocities. because the effects of rotational velocity. ancl abundances are intertwined.," Our choice to fix the abundances at solar affects somewhat our determination of the rotational velocities, because the effects of rotational velocity and abundances are intertwined."
 ILowever. increasing (he number of fillecdl parameters is not justified here. eiven the limited statistical quality of the cata (and the fact that a number of the svstems have composite spectra).," However, increasing the number of fitted parameters is not justified here, given the limited statistical quality of the data (and the fact that a number of the systems have composite spectra)."
 The models are normalized to 1 solar radius and 1 kiloparsec such Chat the distance of a source is computed from d=1000(pc)*(Regfh. )/VS. or equivalently the scale factor S=(4)2 ds the factor by which the theoretical flux. (1ntegrated over the STIS wavelength range) has to be multiplied to equal the observed (integrated) flux.," The models are normalized to 1 solar radius and 1 kiloparsec such that the distance of a source is computed from $d = 
1000(pc)*(R_{wd}/R_{\sun})/\sqrt{S}$ , or equivalently the scale factor $S 
= \left( \frac{R_{wd}}{R_{\odot}} \right)2 \left( \frac{d}{kpc} 
\right)^{-2}$, is the factor by which the theoretical flux (integrated over the STIS wavelength range) has to be multiplied to equal the observed (integrated) flux."
" Unless otherwise specilied. the erid of white dwarf models extended over the following range ol parameters: log g=1.0.1.5.8.0.8.5.9.0: T.55/1000 (IX) = 11. 12. .... 15: abundances [ixed at solar: and 0,,sini (km !) = 100. 200. 400. 600. 800."," Unless otherwise specified, the grid of white dwarf models extended over the following range of parameters: log $g = 7.0, 7.5, 8.0, 8.5, 9.0$; $_{eff}/1000$ (K) = 11, 12, ..., 75; abundances fixed at solar; and $v_{rot} \sin{i}$ (km $^{-1}$ ) = 100, 200, 400, 600, 800."
 The accretion disk models were adopted [rom the optically thick. steady state model grid of Wade&IIubeny.(1998).," The accretion disk models were adopted from the optically thick, steady state model grid of \citet{wad98}."
. For each dwarf nova. we masked out all of the obvious emission features and artifacts in theUST STIS spectrum of each object.," For each dwarf nova, we masked out all of the obvious emission features and artifacts in the STIS spectrum of each object."
 la Table 3. we indicate these masked regions where strong emissionfeatures. artifacts. or negative [Iuxes occur.," In Table 3, we indicate these masked regions where strong emissionfeatures, artifacts, or negative fluxes occur."
lower-energv part. aud hardened the slope in σαοποίος range. which was likely in à amanuer of the prediction of the svuchrotron radiation losses.,"lower-energy part, and hardened the slope in higher-energy range, which was likely in a manner of the prediction of the synchrotron radiation losses."
 The composite SED (Figure 2). from both space aud. ground-based observatious. provides further insights into the object.," The composite SED (Figure 2), from both space and ground-based observations, provides further insights into the object."
" It is clear that the SED from radio to N-rav is possibly from only one radiation component (Svuchrotron emission) aud peaks at a higher frequeucy falling in the EUWV/sott-N-ray band (47,=1.3«Lote Iz).", It is clear that the SED from radio to X-ray is possibly from only one radiation component (Synchrotron emission) and peaks at a higher frequency falling in the EUV/soft-X-ray band $\nu_p=7.3\times10^{15}$ Hz).
 The optical and ultraviolet radiation appear to be a continuation of the radio svuchrotron spectriunu: the N-rav data are likely from a conumaion emission origin as the lower euergv parts aud represcuts a lieh cuerey tail of the svuchrotron spectrum., The optical and ultraviolet radiation appear to be a continuation of the radio synchrotron spectrum; the X-ray data are likely from a common emission origin as the lower energy parts and represents a high energy tail of the synchrotron spectrum.
 Other relevant slope pariuneters from the SED are listed in Table 2., Other relevant slope parameters from the SED are listed in Table 2.
 For a comparison. we plotted in Figure 2 the EGRET scusitivity threshold as an upper limit to the GeV flux αμακος by an arrow). since the source was never detected at 2-xav.," For a comparison, we plotted in Figure 2 the EGRET sensitivity threshold as an upper limit to the GeV flux (marked by an arrow), since the source was never detected at $\gamma$ -ray."
" It is shown that the source is well dominated by a svuchrotrou process,", It is shown that the source is well dominated by a synchrotron process.
" Putting 2316-123 on the à,;, vs 6, color-color diagram. we find it is i the intermediate ranee of BL Lacs."," Putting 2316-423 on the $\alpha_{ro}$ vs $\alpha_{ox}$ color-color diagram, we find it is in the intermediate range of BL Lacs."
" As we know. ovox can more precisely 1ieasure spectral changes from optical to soft N-rav bands. however. the values of ayox for PISS 2316-125 depend ou the assumption of the N-rayv spectral indices. being 0.18/0.26 and for a,=1.0 and 1.6. respectively."," As we know, $\alpha_{XOX}$ can more precisely measure spectral changes from optical to soft X-ray bands, however, the values of $\alpha_{XOX}$ for PKS 2316-423 depend on the assumption of the X-ray spectral indices, being 0.18/0.26 and -0.42/-0.34 for $\alpha_x=1.0$ and 1.6, respectively."
 These values should locate iu the intermediate range of the àyox distribution of receut BL Lacs samples et al., These values should locate in the intermediate range of the $\alpha_{XOX}$ distribution of recent BL Lacs samples (Laurent-Muehleisen et al.
 1999)., 1999).
 The importance of the frequency at which the svuchrotron radiation peaks is that it provides a powerful diagnostics for the physical condition of the cutting reeion., The importance of the frequency at which the synchrotron radiation peaks is that it provides a powerful diagnostics for the physical condition of the emitting region.
 Recent studies showed that amoug BL Lacs the svuchrotron peak frequencies are mvorselv correlated with their luminosities (Fossati et al., Recent studies showed that among BL Lacs the synchrotron peak frequencies are inversely correlated with their luminosities (Fossati et al.
 1998)., 1998).
 We could put, We could put
taken at comparable redshifts as the SDSS SNe and therefore refrain from using this additional information.,taken at comparable redshifts as the SDSS SNe and therefore refrain from using this additional information.
 Given the homogeneity of the SDSS-LLSN dataset. we begin our analysis by revisiting the original evidence for cosmic acceleration in the expansion of the universe.," Given the homogeneity of the SDSS-II SN dataset, we begin our analysis by revisiting the original evidence for cosmic acceleration in the expansion of the universe."
 In detail. if we simply assume the universe is homogeneous and. isotropic. and ds described. by a Robertson-Walker metric with a scale factor of e(£). then the purely kinematic deceleration parameter. q. is defined by where g«0 corresponds to acceleration. (we also assume light follows null gcocdesies and is therefore redshifted in the usual way).," In detail, if we simply assume the universe is homogeneous and isotropic, and is described by a Robertson-Walker metric with a scale factor of $a(t)$, then the purely kinematic deceleration parameter, $q$, is defined by where $q<0$ corresponds to acceleration (we also assume light follows null geodesics and is therefore redshifted in the usual way)."
 We relate q to the Hubble parameter by Luminosity distance can then be calculated directly from the expansion history via where we have assumed a fat universe., We relate $q$ to the Hubble parameter by Luminosity distance can then be calculated directly from the expansion history via where we have assumed a flat universe.
 Thus. the magnitudes and redshifts of any SNe can be used to constrain q(z) without choosing a particular clark-enerev model. or even a particular theory of. gravity.," Thus, the magnitudes and redshifts of any SNe can be used to constrain $q(z)$ without choosing a particular dark-energy model, or even a particular theory of gravity."
 The assumption of Hlatness is necessary dn practice since constraints on q(z) degrade significantly when curvature is allowed. to vary., The assumption of flatness is necessary in practice since constraints on $q(z)$ degrade significantly when curvature is allowed to vary.
 A prior on curvature [rom C'MD measurements would. of course. strengthen the constraint. but such a prior is based on the validity of GR. counter to the intention of the q(z) analysis.," A prior on curvature from CMB measurements would, of course, strengthen the constraint, but such a prior is based on the validity of GR, counter to the intention of the $q(z)$ analysis."
 The simplest deceleration mocel we can fit is a constant. ασ}=qu.," The simplest deceleration model we can fit is a constant, $q(z)=q_0$."
 In this case. the luminosity distance simplifies to Fitting Eq. (6))," In this case, the luminosity distance simplifies to Fitting Eq. \ref{eq:LumDistq0}) )"
 to the SDSS-only SN data. we find a best fit of qu=0.34 and h=0.636.," to the SDSS-only SN data, we find a best fit of $q_0 = -0.34$ and $h = 0.636$."
 Alareinalizine the joint probability density function (PDE) over /. we find qo=0.34dE 0.18. or qu«0 with probability. Le. the SDSS alone finds evidence for acceleration at 20 without concerns regarding the absolute calibration of the peak brightness of SNla andthe relative calibration between SN surveys.," Marginalizing the joint probability density function (PDF) over $h$, we find $q_0=-0.34\pm0.18$ , or $q_0<0$ with probability, i.e., the SDSS alone finds evidence for acceleration at $\sigma$ without concerns regarding the absolute calibration of the peak brightness of SNIa andthe relative calibration between SN surveys."
 There is no reason to expect ασ} to be constant: indeed. underACDAL. q evolves from 1/2 during matter domination to 1 when vacuum energy dominates.," There is no reason to expect $q(z)$ to be constant; indeed, under, $q$ evolves from 1/2 during matter domination to $-1$ when vacuum energy dominates."
 However. we find that including additional parameters in q(z) (e.g. by Tavlor expanding q) degrades our constraints to the point of being uninteresting.," However, we find that including additional parameters in $q(z)$ (e.g. by Taylor expanding $q$ ) degrades our constraints to the point of being uninteresting."
 When trving to reconstruct an unknown function from noisy data. there is the concern that particular features of the reconstruction are not indicative of the true underlying function. but an artifact of the chosen parameterization.," When trying to reconstruct an unknown function from noisy data, there is the concern that particular features of the reconstruction are not indicative of the true underlying function, but an artifact of the chosen parameterization."
 This concern is magnified for a function like q(2). which is related: to the cata pels) by two derivatives.," This concern is magnified for a function like $q(z)$, which is related to the data $\mu(z)$ by two derivatives."
 Therefore. we need a more robust way to determine if the universe has accelerated à wav that does not depend on the naive assumption of a constant q(z).," Therefore, we need a more robust way to determine if the universe has accelerated — a way that does not depend on the naive assumption of a constant $q(z)$."
 A principal-component analysis (ίσος&Starkman2003) can be used to aclelress this issue in a parameter-independent way., A principal-component analysis \citep{2003PhRvL..90c1301H} can be used to address this issue in a parameter-independent way.
 Principal components are a unique set of orthogonal basis functions (e;(2)). such that which allows us to specify q(z) using the coellicients. a;.," Principal components are a unique set of orthogonal basis functions $e_i(z)$ ), such that which allows us to specify $q(z)$ using the coefficients, $\alpha_i$."
" The principal components. or ""modes"" ο(2). are explicitly constructed so. that. the coellicients. can be measured independently. of cach other. be. they have uncorrelated error bars."," The principal components, or “modes” $e_i(z)$, are explicitly constructed so that the coefficients can be measured independently of each other, i.e., they have uncorrelated error bars."
 To construct these modes. we startwith a piecewise-constant. parameterization of q(z) in bins of 0.01. and we use our data to caleulate a Fisher matrix for this parameter set and the Hubble. parameter {ο," To construct these modes, we startwith a piecewise-constant parameterization of $q(z)$ in bins of $dz=0.01$ , and we use our data to calculate a Fisher matrix for this parameter set and the Hubble parameter $H_0$."
 After thareinalizine over Ly. the modes ce;(2) are given hy the cigenvectors of the resulting matrix.," After marginalizing over $H_0$, the modes $e_i(z)$ are given by the eigenvectors of the resulting matrix."
 We are [ree to normalize these functions so that which now specifies cach function up to an overall sign convention., We are free to normalize these functions so that which now specifies each function up to an overall sign convention.
 This procedure is completely. general: in. the limit of dzÜ. we can specily any continuous function using these modes.," This procedure is completely general; in the limit of $dz\rightarrow 0$, we can specify any continuous function using these modes."
 The procedure is paramoeter-independent in the sense that we do not specify the modes a priori: they are determined primarily by the data., The procedure is parameter-independent in the sense that we do not specify the modes a priori: they are determined primarily by the data.
 We do choose q(2)=0 as the fiducial model for our Fisher matrix calculation. but the resulting modes are insensitive to this choice.," We do choose $q(z)=0$ as the fiducial model for our Fisher matrix calculation, but the resulting modes are insensitive to this choice."
 Phe six modes constrained best by the SDSS-only SN data are shown in Fig. 2., The six modes constrained best by the SDSS-only SN data are shown in Fig. \ref{fig:6modes}.
 Each mode is accompanied bv an estimate of the error bar on its coellicient. a(a;). calculated. from the eigenvalues of our Fisher matrix.," Each mode is accompanied by an estimate of the error bar on its coefficient, $\sigma(\alpha_i)$ , calculated from the eigenvalues of our Fisher matrix."
 The first coellicient. ay. has the tightest constraint. while higher mode coellicients have inercasingly [larger errors and so," The first coefficient, $\alpha_1$ , has the tightest constraint, while higher mode coefficients have increasingly larger errors and so"
"The limiting source flux, Sspc, is then determined from the modeled source counts.","The limiting source flux, $S_{\rm SDC}$, is then determined from the modeled source counts."
" If one were to think of the Fermi-LAT as an ordinary telescope with Omin=067% where 0s7o, is the half angle for a beam the contains 67% of the photons.", If one were to think of the -LAT as an ordinary telescope with $\theta_{\rm min} = \theta_{67\%}$ where $\theta_{67\%}$ is the half angle for a beam the contains $67\%$ of the photons.
 The source density criterion would correspond to 1/10 sources per beam., The source density criterion would correspond to $\sim 1/10$ sources per beam.
" For sources with Fy99€11077phcm?s7!, the probability for finding another source of similar or greater flux within the error circle is quite high; hence, many sources that, in principle, should be resolvable are, in fact,unresolved."," For sources with $F_{100} \la 1 \times 10^{-7} \, \mbox{ph cm}^{-2} \mbox{ s}^{-1}$, the probability for finding another source of similar or greater flux within the error circle is quite high; hence, many sources that, in principle, should be resolvable are, in fact,."
". At fluxes close to the sensitivity limit, this probability is so high that faint sources are and will contribute to the measured EGB of the detector in question (for the Fermi-LAT resolution at 100 MeV, the source criterion corresponds to Ngpc~1.3x1073sources/deg?; for reference, at F~2x107?phem-?s-!, our model predicts a much larger source density of ~0.3sources/deg?)."," At fluxes close to the sensitivity limit, this probability is so high that faint sources are and will contribute to the measured EGB of the detector in question (for the -LAT resolution at $100$ MeV, the source criterion corresponds to $N_{\rm SDC} \sim 1.3\times 10^{-3} \, {\rm sources}/{\rm deg}^{2}$; for reference, at $F \sim 2 \times 10^{-9} \, {\rm ph}\, {\rm cm}^{-2} \, {\rm s}^{-1}$, our model predicts a much larger source density of $\sim 0.3 \, {\rm sources}/{\rm deg}^{2}$ )."
" 'Thus, the effect of source confusion compounded with the separate effect of detector sensitivity is to flatten the faint end of an observationally derived source count distribution."," Thus, the effect of source confusion compounded with the separate effect of detector sensitivity is to flatten the faint end of an observationally derived source count distribution."
 We should emphasize that this definition of source confusion is not the same as that employed by ?.., We should emphasize that this definition of source confusion is not the same as that employed by \citet{pop10}.
" Source confusion as discussed in this section refers to the probability that for a given source density, a source has a nearest neighbor within the angular resolution of the detector with a flux greater than or equal to the flux limit."," Source confusion as discussed in this section refers to the probability that for a given source density, a source has a nearest neighbor within the angular resolution of the detector with a flux greater than or equal to the flux limit."
" In ?,, the term is applied to a source associated with a given source for which the measured flux of the detected source is greater than the flux of the real source (plus three standard deviations) by a given amount (the analysis was performed on simulated data to determine the impact on the actual "," In \citet{pop10}, the term is applied to a source associated with a given source for which the measured flux of the detected source is greater than the flux of the real source (plus three standard deviations) by a given amount (the analysis was performed on simulated data to determine the impact on the actual data)."
"In effect, the former definition identifies the limit at data).which the source density is sufficiently large so that individual sources be resolved, and only fluctuations are observed."," In effect, the former definition identifies the limit at which the source density is sufficiently large so that individual sources be resolved, and only fluctuations are observed."
 The latter definition applies to the probability that a given detected source is in fact the superposition of several sources., The latter definition applies to the probability that a given detected source is in fact the superposition of several sources.
" However, the ? criterion is based on the assumption that a source can beresolved.."," However, the \citet{pop10} criterion is based on the assumption that a source can be."
" As such, the source would have to be significantly brighter than the background, including the sources within the error circle of the detector."," As such, the source would have to be significantly brighter than the background, including the sources within the error circle of the detector."
" In the case of many sources with similar fluxes, none of the sources would be resolved."," In the case of many sources with similar fluxes, none of the sources would be resolved."
" In the case of one bright source and several fainter sources, the flux of the detected source would be dominated by the flux of the brightest source and likely would escape the ?.."," In the case of one bright source and several fainter sources, the flux of the detected source would be dominated by the flux of the brightest source and likely would escape the \citet{pop10}."
" Thus, under this criterion, it is not surprising that they conclude that there is very little source confusion."," Thus, under this criterion, it is not surprising that they conclude that there is very little source confusion."
" A common treatment of source confusion in measured source counts is a fluctuation analysis, but as yet, such an analysis has not be performed onFermi source counts."," A common treatment of source confusion in measured source counts is a fluctuation analysis, but as yet, such an analysis has not be performed on source counts."
" 'The spectral indices of the population of blazars form a distribution with a given finite spread (??),, resulting in curvature in the spectrum of the collective intensity of unresolved blazars due to the increasing relative importance at high energies of blazars with harder spectral indices (??).."," The spectral indices of the population of blazars form a distribution with a given finite spread \citep{ste96,vp07}, resulting in curvature in the spectrum of the collective intensity of unresolved blazars due to the increasing relative importance at high energies of blazars with harder spectral indices \citep{ste96,pv08}."
" Thus, the determination of the blazar SID is crucial in determining the correct spectrum of the unresolved blazar contribution."," Thus, the determination of the blazar SID is crucial in determining the correct spectrum of the unresolved blazar contribution."
" In determining the blazar SID from survey data, one must carefully account for uncertainties in measurement of the spectral indices."," In determining the blazar SID from survey data, one must carefully account for uncertainties in measurement of the spectral indices."
" In so doing, we follow the likelihood method of ?,, fitting 1FGL FSRQs toa gaussian SID."," In so doing, we follow the likelihood method of \citet{vp07}, fitting 1FGL FSRQs toa gaussian SID."
" We determined the maximum-likelihood gaussian SID parameters (mean, I'o, and spread, oo) to be [9=2.45 and oo—0.15."," We determined the maximum-likelihood gaussian SID parameters (mean, $\Gamma_0$, and spread, $\sigma_0$ ) to be $\Gamma_0 = 2.45$ and $\sigma_0 = 0.15$."
 We must also account for the effect of spectral bias of the 1FGL catalog., We must also account for the effect of spectral bias of the 1FGL catalog.
" Even in a flux-limited catalog, low-luminosity, high-redshift blazars that would be most likely to appear are those with spectral indices that are harder than most of the population."," Even in a flux-limited catalog, low-luminosity, high-redshift blazars that would be most likely to appear are those with spectral indices that are harder than most of the population."
" The 1FGL catalog is not actually flux limited (see Section ??)), rather it consists of sources above a given threshold in teststatistic!!,, which depends on source spectra and the background."," The 1FGL catalog is not actually flux limited (see Section \ref{sec:datalogNlogS}) ), rather it consists of sources above a given threshold in test, which depends on source spectra and the background."
" However, as demonstrated in ?,, applying the likelihood analysis to the sample of FSRQs with fluxes >7x10-5phcm-?s-! and galactic latitudes =10? (asper?) does not appreciably change the SID parameters (Γρ=2.45, σο=0.16)!."," However, as demonstrated in \citet{ven10a}, , applying the likelihood analysis to the sample of FSRQs with fluxes $\ga 7\times 10^{-8} \, {\rm ph \, cm^{-2} \, s^{-1}}$ and galactic latitudes $\ga 10^{\circ}$ \citep[as per][]{pop10} does not appreciably change the SID parameters $\Gamma_0 = 2.45$, $\sigma_0 = 0.16$."
". Thus, even though the 1FGL catalog is not flux limited, we can assume that the sample of 1FGL FSRQs is approximatelyflux limited."," Thus, even though the 1FGL catalog is not flux limited, we can assume that the sample of 1FGL FSRQs is approximatelyflux limited."
 We can then follow the method of ? in correcting for the sample bias inherent in a flux-limited catalog., We can then follow the method of \citet{ven09} in correcting for the sample bias inherent in a flux-limited catalog.
" In doing so, we apply a correction factor, M(o), to the SID (for ?):: where βία) is the SID corrected for measurement uncertainty in the spectral indices, and"," In doing so, we apply a correction factor,$\hat{M}(\alpha)$ , to the SID \citep[for derivation, see][]{ven09}: : where $\hat{p}(\alpha)$ is the SID corrected for measurement uncertainty in the spectral indices, and"
the stratification of elemental abundances due to diffusion processes (LeBlancetal.2010).,the stratification of elemental abundances due to diffusion processes \citep{Leblanc10}.
. llowever. a svstematic difference in one of these parameters. alfecting wCCen but not the other clusters. would be even more puzzling. and it is harder to postulate at this stage.," However, a systematic difference in one of these parameters, affecting $\omega$ Cen but not the other clusters, would be even more puzzling, and it is harder to postulate at this stage."
 It must also be noted that the mass underestimate does not completely follow the trend of the difference in gravity: while this seems to vary. with temperature as noted before. the masses ino CCen are constantly underestimated by 20.15 AL. for Tar2100000 Ix. For example. at Teg0160000-18 0000 Ix there is a clear offset in masses. but not in gravities.," It must also be noted that the mass underestimate does not completely follow the trend of the difference in gravity: while this seems to vary with temperature as noted before, the masses in $\omega$ Cen are constantly underestimated by $\sim$ 0.15 $_\sun$ for $_\mathrm{eff}\geq$ 000 K. For example, at $_\mathrm{eff}\sim$ 000 K there is a clear offset in masses, but not in gravities."
 This indicates that low gravities nav play a role. but at least another effect should be at work. causing the observed mass underestimate.," This indicates that low gravities may play a role, but at least another effect should be at work, causing the observed mass underestimate."
 All the comparison clusters have similar metallicity (|Fe/IHI]g —1.5 1996). while ο CCen shows a large spread up to |Fe/1I]2—0.6 (e.g..Sollimaetal.2005).," All the comparison clusters have similar metallicity \citep[{[Fe/H]}$ $\approx -$ , while $\omega$ Cen shows a large spread up to $-$ 0.6 \citep[e.g.,][]{Sollima05}."
".. Nevertheless, ibis unlikelv that the higher metallicity is the origin of the peculiar results in w CCen: the largest differences are found. for stars hotter than ~115500 Ix. whose surface abunclances are altered by diffusion processes."," Nevertheless, it is unlikely that the higher metallicity is the origin of the peculiar results in $\omega$ Cen: the largest differences are found for stars hotter than $\sim$ 500 K, whose surface abundances are altered by diffusion processes."
 Behr(2003). ancl Paceetal.(2006). showed that. in the presence of diffusion. the surface abundance patterns are verv sinülar in clusters of very different. initial metallicity.," \citet{Behr03} and \citet{Pace06} showed that, in the presence of diffusion, the surface abundance patterns are very similar in clusters of very different initial metallicity."
 The stars in all the clusters should. therefore show the same behavior independently of their primordial metal content. especially at 0000. Ix where diffusion reaches its maximum strength (MoniBiclinetal.2009).," The stars in all the clusters should therefore show the same behavior independently of their primordial metal content, especially at 000 K where diffusion reaches its maximum strength \citep{Moni09}."
. Moreover. etal.(2000) found no peculiarity in the measured gravity. and mass of two stus in Tuc —0.7). although using low-metallicity models.," Moreover, \citet{Moehler00} found no peculiarity in the measured gravity and mass of two stars in Tuc $-$ 0.7), although using low-metallicity models."
 We repeated (he measurements assuming different. values of the model metallicity. to test how this parameter can allect the results.," We repeated the measurements assuming different values of the model metallicity, to test how this parameter can affect the results."
 We found small differences in the stellar parameters. but the general behavior was unaltered: a higher model metallicity indeed returned sliehtly higher gravities. but hieher temperatures loo.," We found small differences in the stellar parameters, but the general behavior was unaltered: a higher model metallicity indeed returned slightly higher gravities, but higher temperatures too."
 As a consequence. the points were shifted almost parallel to the theoretical tracks in the Toy-log(gy) plane. while the masses were increased by less (han 0.05 M...," As a consequence, the points were shifted almost parallel to the theoretical tracks in the $_\mathrm{eff}$ $\log{\mathrm (g)}$ plane, while the masses were increased by less than 0.05 $_\sun$."
 The blanketing effect should be lower in metal-poor stus than in the solar-metallicity stars used to calibrate (he adopted TarDC relation. and this could cause the underestimate of the DC ancl of the mass.," The blanketing effect should be lower in metal-poor stars than in the solar-metallicity stars used to calibrate the adopted $_\mathrm{eff}$ -BC relation, and this could cause the underestimate of the BC and of the mass."
 The adopted BC should be a good approximation for the stars hotter than the Grundahl jump (Grundahlοἱal.1999).. because radiative levitation increases (heir surface abundances to super-solar values.," The adopted BC should be a good approximation for the stars hotter than the Grundahl jump \citep{Grundahl99}, because radiative levitation increases their surface abundances to super-solar values."
 As already noted. the effect should be independent of their primordial metallicity. (hus the BCdoes not explain the ollset of o CCen with respect io the other GCs.," As already noted, the effect should be independent of their primordial metallicity, thus the BCdoes not explain the offset of $\omega$ Cen with respect to the other GCs."
 At cooler temperatures. the offset could be explained by the BC ifw Cen stars were more metal-poor Chan in the other GCs. thus decreasing their BC by 20.4 mag.," At cooler temperatures, the offset could be explained by the BC if $\omega$ Cen stars were more metal-poor than in the other GCs, thus decreasing their BC by $\sim$ 0.4 mag."
 Indeed. &: CCen hosts a metal-poor sub-population at |Fe/1I]|——2 (e.g..Pancinoetal.2011).. but the BC varies by less than 0.15 mag for stars at 0000 Ix in the wholerange between solar metallicity and |Fe/I1I|2—2 (Cassisietal.1999:Alonso 1999)...," Indeed, $\omega$ Cen hosts a metal-poor sub-population at $-$ 2 \citep[e.g.,][]{Pancino11}, but the BC varies by less than 0.15 mag for stars at 000 K in the wholerange between solar metallicity and $-$ 2 \citep{Cassisi99,Alonso99}. ."
 A wrong distance modulus or reddening coukl also cause wrong mass estimates. but the required," A wrong distance modulus or reddening could also cause wrong mass estimates, but the required"
where j; is the spherical Bessel function of order € and A;=Om) is the moment of the dimensionless power spectrum or power per Ink.,where $j_\ell$ is the spherical Bessel function of order $\ell$ and $\Delta_\ell^2\equiv k^3P_\ell(k)/(2\pi^2)$ is the moment of the dimensionless power spectrum or power per $\ln k$.
 While the relative amplitudes of the Legendre moments &(s) are also given by Eq. 21..," While the relative amplitudes of the Legendre moments $\xi_{\ell}(s)$ are also given by Eq. \ref{pmoments},"
 in configuration space the moments depend differently on s (see solid lines in Fig. 45., in configuration space the moments depend differently on $s$ (see solid lines in Fig. \ref{fig:streaminglin}) ).
 For example. using the recurrence relations among the spherical Bessel functions we can write & as the average value of £ up to s minus the value at s.," For example, using the recurrence relations among the spherical Bessel functions we can write $\xi_2$ as the average value of $\xi$ up to $s$ minus the value at $s$."
 Non-linear effects. and especially tingers-of-god on small scales. can cause the €>4 moments to be non-zero.," Non-linear effects, and especially fingers-of-god on small scales, can cause the $\ell > 4$ moments to be non-zero."
 The of-god are evident in Fig., The fingers-of-god are evident in Fig.
 |. as the stretching of contours localised near R=0., \ref{fig:butterfly} as the stretching of contours localised near $R=0$.
 The presence of strong fingers-of-god also drives €» positive. whereas super-cluster infall on larger scales generally predicts &«0.," The presence of strong fingers-of-god also drives $\xi_2$ positive, whereas super-cluster infall on larger scales generally predicts $\xi_2<0$."
" Generally the monopole (©= 0) and quadrupole (€= 2) dominate in terms of signal-to-noise. and the ratio P»/P; (tor an analogous combination of £j, and £:) encodes information about B=//bapproach)."," Generally the monopole $\ell=0$ ) and quadrupole $\ell=2$ ) dominate in terms of signal-to-noise, and the ratio $P_2/P_0$ (or an analogous combination of $\xi_0$ and $\xi_2$ ) encodes information about $\beta=f/b$."
 As is so often the case. our strongest statistical constraints come from the smaller scales where we have a large number of independent samples within our survey.," As is so often the case, our strongest statistical constraints come from the smaller scales where we have a large number of independent samples within our survey."
 These are also the scales for which the simple linear theory described in the last section is the least applicable., These are also the scales for which the simple linear theory described in the last section is the least applicable.
 A standard approach to include small-scale non-linearities manifest as tingers-of-god is to use a streaming model. where linear theory is spliced together with an approximation for random motion of particles in collapsed objects1992).," A standard approach to include small-scale non-linearities manifest as fingers-of-god is to use a streaming model, where linear theory is spliced together with an approximation for random motion of particles in collapsed objects."
. This arises from ignoring the scale-dependence of the mapping between real and redshift space separations and assuming an isotropic velocity dispersion and amounts to convolving the linear theory result with a LOS smearing or multiplying the power spectrum by the Fourier transform of the small-scale velocity PDF., This arises from ignoring the scale-dependence of the mapping between real and redshift space separations and assuming an isotropic velocity dispersion and amounts to convolving the linear theory result with a LOS smearing or multiplying the power spectrum by the Fourier transform of the small-scale velocity PDF.
 modelled the small-scale velocity field as an incoherent Gaussian scatter which amounts to A more realistic distribution for the pairwise velocities is exponential1983).. because galaxies populate halos of a wide range of masses and velocity dispersions 2001).," modelled the small-scale velocity field as an incoherent Gaussian scatter which amounts to A more realistic distribution for the pairwise velocities is exponential, because galaxies populate halos of a wide range of masses and velocity dispersions ."
 While this extension. from linear theory can improve agreement with observations and N-body simulation results. this simple approach ignores the well-known fact that the velocity divergence field 8 deviates from linear theory on extremely large scales (k<0.03Mpe! see e.g. fig.," While this extension from linear theory can improve agreement with observations and $N$ -body simulation results, this simple approach ignores the well-known fact that the velocity divergence field $\theta$ deviates from linear theory on extremely large scales $k \lesssim 0.03\,h\,{\rm Mpc}^{-1}$ ); see e.g., fig."
 8 of for à recent comparison., 8 of for a recent comparison.
 therefore proposed a simple extension of linear theory for the matter power spectrum in redshift space: where the power spectra on the right hand side are evaluated in perturbation theory in space. and o7 is the linear velocity dispersion.," therefore proposed a simple extension of linear theory for the matter power spectrum in redshift space: where the power spectra on the right hand side are evaluated in perturbation theory in space, and $\sigma_v^2$ is the linear velocity dispersion."
 have extended this model for the redshift space matter power spectrum to different dark energy scenarios by fitting a relation between the three non-linear power spectra in Eq., have extended this model for the redshift space matter power spectrum to different dark energy scenarios by fitting a relation between the three non-linear power spectra in Eq.
 24. to N-body simulation outputs. but treat σε as a free parameter1999).," \ref{scoccPT} to $N$ -body simulation outputs, but treat $\sigma_v^2$ as a free parameter."
 However. the relation between the real and redshift space statistics depends on non-trivial correlations at high orders. which Eq.," However, the relation between the real and redshift space statistics depends on non-trivial correlations at high orders, which Eq."
 24. ignores: many terms from the standard. perturbation theory expression for the redshift space power spectrum have been dropped2008a)., \ref{scoccPT} ignores; many terms from the standard perturbation theory expression for the redshift space power spectrum have been dropped.
. However. it is well-known that the standard perturbation theory expression for the redshift space clustering of matter is inaccurate (e.g..(1999):: see also and. for recent comparisons to N-body simulations and observed galaxy power spectra. respectively).," However, it is well-known that the standard perturbation theory expression for the redshift space clustering of matter is inaccurate (e.g.,; see also and for recent comparisons to $N$ -body simulations and observed galaxy power spectra, respectively)."
 The importance of higher-order terms ean be seen most easily from the expression for the correlation function. Eq. I4..," The importance of higher-order terms can be seen most easily from the expression for the correlation function, Eq. \ref{eqn:allorders},"
 where we see that going beyond linear theory involves many moments of ὁ and v.., where we see that going beyond linear theory involves many moments of $\delta$ and $v_z$.
 showed that the bispectrum terms have oscillatory features which affect the BAO feature. and used them to obtain a better fit to jV- simulation data for the matter redshift space power spectrum.," showed that the bispectrum terms have oscillatory features which affect the BAO feature, and used them to obtain a better fit to $N$ -body simulation data for the matter redshift space power spectrum."
 However. still needed. to introduce a Gaussian damping term with free parameter 07; in order to fit the N-body data well.," However, still needed to introduce a Gaussian damping term with free parameter $\tilde{\sigma}_v^2$ in order to fit the $N$ -body data well."
" take a different approach.and try to reconstruct P(A) and P, from P'(k) for both matter and halos in N-body simulations. assuming a relation. like Eq. 24..."," take a different approach,and try to reconstruct $P_{\delta \delta}^{r}(k)$ and $P_{\delta \theta}^{r}(k)$ from $P^s(k)$ for both matter and halos in $N$ -body simulations, assuming a relation like Eq. \ref{scoccPT}, ,"
 butwith a more general multiplicative, butwith a more general multiplicative
he observed CAIRDs of all samples are consistent with waving been drawn from the IAIF are given in Table 2..,the observed CMRDs of all samples are consistent with having been drawn from the IMF are given in Table \ref{table:KS test}.
" With regard to the Pleiades. a Per aud Taurus (KS test xobabilities of 10.150, aud 10.Ἐν, respectively) we can reject the hvpothesis of raudoni pairing. while he higher probability (174)) between the IME aud CNIBD in Chaimacleou I does uot allow us to rule out he null hvpothesis in this case."," With regard to the Pleiades, $\alpha$ Per and Taurus (KS test probabilities of $10^{-4}$, and $10^{-11}$, respectively) we can reject the hypothesis of random pairing, while the higher probability ) between the IMF and CMRD in Chamaeleon I does not allow us to rule out the null hypothesis in this case."
 However. the number of objects (13) in Chamacleon I sample is quite low.," However, the number of objects (13) in Chamaeleon I sample is quite low."
 As we have shown in Section 3.3 the KS test can oulv distinguish extreme differences in distributions from such small samples., As we have shown in Section \ref{KS_test} the KS test can only distinguish extreme differences in distributions from such small samples.
 We should also note that in Taurus the IME is peaked toward higher masses with respect to the field IME (Lulunanetal.2009)., We should also note that in Taurus the IMF is peaked toward higher masses with respect to the field IMF \citep{Luhman2009}.
. The use of the proper iass distribution would bius the CAIRD in Taurus in closer agreement with random pairing from the IME., The use of the proper mass distribution would bring the CMRD in Taurus in closer agreement with random pairing from the IMF.
 Hence we should be cautious in interpreting this preliminary result., Hence we should be cautious in interpreting this preliminary result.
 To summarize. at the monent in the Pleiades and à Per we cau reject the possibility of the CAIRD beine drawn from the IME for orbital separation between 20 and 600 AU. whereas concerning much vounger reeions. we rule out the random pairing from the field IME oulv in Taurus in the separation range 5-5000 AU.," To summarize, at the moment in the Pleiades and $\alpha$ Per we can reject the possibility of the CMRD being drawn from the IMF for orbital separation between 20 and 600 AU, whereas concerning much younger regions, we rule out the random pairing from the field IMF only in Taurus in the separation range 5-5000 AU."
 Data from larger and different samples are needed to better constrain the result as a function of age aud environment., Data from larger and different samples are needed to better constrain the result as a function of age and environment.
 Using the same Monte Carlo method. we have tested also whether the observed CAIRD as fiction of primary lass and enviroment is consistent with other analytic forms of the CAIRD.," Using the same Monte Carlo method, we have tested also whether the observed CMRD as function of primary mass and environment is consistent with other analytic forms of the CMRD."
" First of all we have considered a linearly fat companion mass ratio distribution (see Section 3.3)) and second we have tested the companion unass (istribution dN/dAtsxAL,1 sneecsted by MIT.", First of all we have considered a linearly flat companion mass ratio distribution (see Section \ref{KS_test}) ) and second we have tested the companion mass distribution $dN/dM_{2}\propto M_{2}^{-0.4}$ suggested by MH09.
 In Table 2. we report the KS probabilities. for cach dataset., In Table \ref{table:KS test} we report the KS probabilities for each dataset.
 Coucerning the flat distribution. oulv for the sample of A and late B-type primary biuaries in Sco OD2 we can reject the hypothesis that the two distributions are consistent.," Concerning the flat distribution, only for the sample of A and late B-type primary binaries in Sco OB2 we can reject the hypothesis that the two distributions are consistent."
 The comparison iu this case is shown in the op panel of Figure L., The comparison in this case is shown in the top panel of Figure \ref{cmr_field_flat}.
 Note that the ScoODB2 dataset iS he largest sample. placing the strongest coustraiunts on YOSSI)e differeuces.," Note that the ScoOB2 dataset is the largest sample, placing the strongest constraints on possible differences."
 The AIC simulations of a flat CALF or the voung regions match well the observations (see e.g. bottom panel of Figure 1))., The MC simulations of a flat CMF for the young regions match well the observations (see e.g. bottom panel of Figure \ref{cmr_field_flat}) ).
" Regarding the CMBD xovided bv MIIOS we find a KS probability exceeding for all samples (see Table 2)) except for Taurus (2144),", Regarding the CMRD provided by MH09 we find a KS probability exceeding for all samples (see Table \ref{table:KS test}) ) except for Taurus ).
 We should keep in ind that the ISS test is not suitxd o evaluate which is the best fit distribution., We should keep in mind that the KS test is not suited to evaluate which is the best fit distribution.
 If we take as an example the results for the FAL92 sample. the difference iu the probability from to between he dat aud M09 CAIF in the context of the INS test does not have any significance.," If we take as an example the results for the FM92 sample, the difference in the probability from to between the flat and M09 CMF in the context of the KS test does not have any significance."
 Furthermore. the sample size of our datasets in the majority of cases prevents us frou discriminating between loe-normal. flat or other distiibutious (see Section 3.3)).," Furthermore, the sample size of our datasets in the majority of cases prevents us from discriminating between log-normal, flat or other distributions (see Section \ref{KS_test}) )."
 For this reason we have utilized a chi-square procedure to determine the best fit for the CAIRD for a combined sample including all primary lasses., For this reason we have utilized a chi-square procedure to determine the best fit for the CMRD for a combined sample including all primary masses.
 Motivated by the fact that the CAIRD appears to be independent of angular separation over the rauge we are considere and that the distributious are uot cistinguishable. we combined together. over the common range of mass ratios (q4—0.2-1). the samples of AL dwarts and € stars in the field aud intermediate mass stars m ScoOD2 even though the separation ranges vary across the samples.," Motivated by the fact that the CMRD appears to be independent of angular separation over the range we are considering and that the distributions are not distinguishable, we combined together, over the common range of mass ratios $q$ =0.2-1), the samples of M dwarfs and G stars in the field and intermediate mass stars in ScoOB2 even though the separation ranges vary across the samples."
 We then used the composite q distribution to find the best fit., We then used the composite $q$ distribution to find the best fit.
 According to the chi-square test for, According to the chi-square test for
the 8.2-m2 VLT tlocated on Cerro Paranal in the Atacama Desert. northern. Chile). using the Infrared Spectrometer And. Array Camera (ISAAC). a 1024. 1024 Hawaii Rockwell array with a pixel scale of 0.15 aresec.,"the 8.2-m VLT (located on Cerro Paranal in the Atacama Desert, northern Chile), using the Infrared Spectrometer And Array Camera (ISAAC), a 1024 $\times$ 1024 Hawaii Rockwell array with a pixel scale of 0.15 arcsec."
 This pixel scale provided the necessary good sampling of even the very best seeing disc. while still providing sufficient field-of-view to include the required nearby field stars for accurate on-image PSF characterisation.," This pixel scale provided the necessary good sampling of even the very best seeing disc, while still providing sufficient field-of-view to include the required nearby field stars for accurate on-image PSF characterisation."
 The reliable detection of host-galaxy light around luminous quasars is difficult at all redshifts. but especially at high redshift. due to he small angular size of any reasonable anticipated host galaxy. he faintness of any potential host galaxy relative to the scattered nuclear light in the wings of the PSF. and the fact that the optical light has been redshifted to the near-infrared.," The reliable detection of host-galaxy light around luminous quasars is difficult at all redshifts, but especially at high redshift, due to the small angular size of any reasonable anticipated host galaxy, the faintness of any potential host galaxy relative to the scattered nuclear light in the wings of the PSF, and the fact that the rest-frame optical light has been redshifted to the near-infrared."
 Predicting surface brightnesses (and hence exposure times) or high-z quasar host galaxies is complicated by unknown amounts of stellar and morphological evolution to offset against the cosmological dimming and anticipated &-corrections., Predicting surface brightnesses (and hence exposure times) for $z$ quasar host galaxies is complicated by unknown amounts of stellar and morphological evolution to offset against the cosmological dimming and anticipated $k$ -corrections.
 In an attempt ο best define our observational strategy we generated synthetic VLT/ISAAC exposures of a >=+ quasar with a nuclear/host ratio of 15. 0.6-aresec seeing. and background-limited noise.," In an attempt to best define our observational strategy we generated synthetic VLT/ISAAC exposures of a $z=4$ quasar with a nuclear/host ratio of 15, 0.6-arcsec seeing, and background-limited noise."
 The synthetic quasar had a total A s-band apparent magnitude of A4=16.5 (consistent with the anticipated A s-band magnitudes of our target SDSS quasars}. comprising a nuclear point source with Avy=16.8 and a de Vaucouleurs host galaxy with As=19.7 (consistent with existing studies of the A> relation for massive galaxies).," The synthetic quasar had a total $K_S$ -band apparent magnitude of $K_S=16.5$ (consistent with the anticipated $K_S$ -band magnitudes of our target SDSS quasars), comprising a nuclear point source with $K_S=16.8$ and a de Vaucouleurs host galaxy with $K_S=19.7$ (consistent with existing studies of the $K-z$ relation for massive galaxies)."
 For the simulated host galaxy we assumed a half-light radius of Skkpe., For the simulated host galaxy we assumed a half-light radius of kpc.
 Separate PSFs were used to construc the synthetic quasar and to carry out nuclear/host separation. with both PSFs obtained from the publicly-available ISAAC image of the GOODS-South tield. on which they were  30-aresee apar (thus accurately reflecting our observing strategy of using a nearby star to characterise the on-image PSF).," Separate PSFs were used to construct the synthetic quasar and to carry out nuclear/host separation, with both PSFs obtained from the publicly-available ISAAC image of the GOODS-South field, on which they were $\sim30$ -arcsec apart (thus accurately reflecting our observing strategy of using a nearby star to characterise the on-image PSF)."
 In order to separate hos and nuclear components and reliably determine total luminosities for such galaxies. we found that we needed to follow the surface brightness profile out to ~3 galaxy half-light radii.," In order to separate host and nuclear components and reliably determine total luminosities for such galaxies, we found that we needed to follow the surface brightness profile out to $\sim 3$ galaxy half-light radii."
 This requirec integrating down to a 30 surface brightness level of fry=25 7., This required integrating down to a $3\sigma$ surface brightness level of $\mu_K \simeq 25$ $^{-2}$.
 The VLT ISAAC Exposure Time Calculator indicated that we could achieve this level in 5-5 hours on source. and so this was set as the minimum on-source exposure time for our observations.," The VLT ISAAC Exposure Time Calculator indicated that we could achieve this level in $\simeq 5$ hours on source, and so this was set as the minimum on-source exposure time for our observations."
 For the reasons explained above. we wished to obtain emission-line free images of the most luminous quasars at zc4.," For the reasons explained above, we wished to obtain emission-line free images of the most luminous quasars at $z \simeq 4$."
 Contamination of the observed ἐν ραπ light by [ΟΠΗ and H? line emission is avoided by insisting on +3.95. and so we aimed to select quasar targets at only slightly higher redshift to ensure that the dy s-band filter was still essentially filled by the continuum light starting at Άργος&4000À.," Contamination of the observed $K$ -band light by [OIII] and $\beta$ line emission is avoided by insisting on $z > 3.95$, and so we aimed to select quasar targets at only slightly higher redshift to ensure that the $K_S$ -band filter was still essentially filled by the continuum light starting at $\lambda_{rest} \simeq 4000$."
. We confined our attention to the most luminous quasars in the SDSS quasar catalogue at these redshifts. corresponding to absolute magnitudes 1;<25.," We confined our attention to the most luminous quasars in the SDSS quasar catalogue at these redshifts, corresponding to absolute magnitudes $M_i < -28$."
" These objects are rare: the tinal SDSS quasar catalogue contains only 2200 such quasars in the redshift range 4«25 over the full 9380 ddeg"" of the SDSS quasar survey (Schneider et al.", These objects are rare; the final SDSS quasar catalogue contains only $\simeq 200$ such quasars in the redshift range $4 < z < 5$ over the full $\simeq 9380$ $^2$ of the SDSS quasar survey (Schneider et al.
 2010)., 2010).
" They thus have a comoving number density of 2 per comoving Gpe? at zc4. which corresponds to | object in the cosmological volume simulated in the ""Massive Black"" simulation recently discussed by Di Matteo et al. ("," They thus have a comoving number density of $\simeq 2$ per comoving $^3$ at $z \simeq 4$, which corresponds to $\simeq 1$ object in the cosmological volume simulated in the “Massive Black” simulation recently discussed by Di Matteo et al. ("
2011) (see Section 5).,2011) (see Section 5).
 Finally. we further limited our observational options. to quasars from within this sample which happened to have at least one bright star within a radius of 40 arcsec. in order to provide a high signal:noise ratio representation of the real form of the point-spread-function (PSF) on the final science image.," Finally, we further limited our observational options to quasars from within this sample which happened to have at least one bright star within a radius of 40 arcsec, in order to provide a high signal:noise ratio representation of the real form of the point-spread-function (PSF) on the final science image."
 It was deemed essential that the nearby stars were at least as bright as the quasar at /vs-band. but not so bright as to risk saturation or non-linearity in the background limited imaging.," It was deemed essential that the nearby stars were at least as bright as the quasar at $K_S$ -band, but not so bright as to risk saturation or non-linearity in the background limited imaging."
 This final rather stringent requirement reduced our potential target list to less than 10 quasars., This final rather stringent requirement reduced our potential target list to less than 10 quasars.
 Our insistence on long (>Shhr) integrations in the very best seeing conditions meant that we only obtained the required deep imaging for two of the potential target quasars. but all data were taken in photometric conditions and consistent high-quality (<0.4-aresec) atmospheric seeing.," Our insistence on long $> 5$ hr) integrations in the very best seeing conditions meant that we only obtained the required deep imaging for two of the potential target quasars, but all data were taken in photometric conditions and consistent high-quality $<0.4$ -arcsec) atmospheric seeing."
 Our observations of these two z+ quasars are summarised in Table I., Our observations of these two $z \simeq 4$ quasars are summarised in Table 1.
 Data reduction was performed using standardIRAP packages., Data reduction was performed using standard packages.
 Dark frames of equal integration time to the science data were taken on all nights., Dark frames of equal integration time to the science data were taken on all nights.
 Each dark-subtracted image was divided by a normalised flat-tield. derived though the scaled sigma-clipped median combination of neighbouring science frames observed within 20 minutes of each image.," Each dark-subtracted image was divided by a normalised flat-field, derived though the scaled sigma-clipped median combination of neighbouring science frames observed within 20 minutes of each image."
 Registering of the offset frames was performed using the brightest stars in the images., Registering of the offset frames was performed using the brightest stars in the images.
 A map of the bad pixels in the array was obtained and used to exclude these pixels during image combination., A map of the bad pixels in the array was obtained and used to exclude these pixels during image combination.
 Median filtering was used to reject cosmic rays and produce an initial reduced image which was then processed to create a mask of all source flux detected in the mosaiced data., Median filtering was used to reject cosmic rays and produce an initial reduced image which was then processed to create a mask of all source flux detected in the mosaiced data.
 The science frames were then reprocessed to create an improved flat-field. and then the data were re-reduced to produce the final science sub-images.," The science frames were then reprocessed to create an improved flat-field, and then the data were re-reduced to produce the final science sub-images."
" As accurate matching of both the quasar and PSF image quality is of critical importance in this work, separate mosaics were created for the quasars and for the PSF stars using offset shifts derived from each object individually."," As accurate matching of both the quasar and PSF image quality is of critical importance in this work, separate mosaics were created for the quasars and for the PSF stars using offset shifts derived from each object individually."
 Final image registration was performed using a Spline3 interpolation in the IRAF routine IMSHIFT., Final image registration was performed using a Spline3 interpolation in the IRAF routine IMSHIFT.
 Because of the expected overwhelming dominance of the nuclear light. separation of nuclear and host-galaxy light was first performed using scaled PSF subtraction.," Because of the expected overwhelming dominance of the nuclear light, separation of nuclear and host-galaxy light was first performed using scaled PSF subtraction."
 We subtracted an, We subtracted an
the file hip.jj.dat.,the file j.dat.
 No useful data are lost because of the errors. but special measures may be needed to read the corresponding records.," No useful data are lost because of the errors, but special measures may be needed to read the corresponding records."
 The interface programs described oeji Sect., The interface programs described in Sect.
 6 automatically correct these errors., \ref{sec:soft} automatically correct these errors.
 Further etails can be found at the Iuternet address given iu Sect. 6.., Further details can be found at the Internet address given in Sect. \ref{sec:soft}.
 The IHipparcos satellite was uot cesigued for laine alu did not contain anv nuaegiug device such as a CCD camnera., The Hipparcos satellite was not designed for imaging and did not contain any imaging device such as a CCD camera.
 The combination of a modulating grid aux the IDT. while well adaptee to the observation of isolated point sources. was far from ideal for the observation of more conplex resolved objects;," The combination of a modulating grid and the IDT, while well adapted to the observation of isolated point sources, was far from ideal for the observation of more complex resolved objects."
 As explained in Sect., As explained in Sect.
 2.2 the grid esseutiallv extracted two spatial frequencies (with periods 1.2071 arcsec and 0.6037 arcsoc du the direction of cach scan) of whatever intensity distribution was within the 30 arcsec sensitive spot., \ref{sec:grid} the grid essentially extracted two spatial frequencies (with periods 1.2074 arcsec and 0.6037 arcsec in the direction of each scan) of whatever intensity distribution was within the 30 arcsec sensitive spot.
 This is much more roiuiniscent of spatial interferometry than of normal optical imaging., This is much more reminiscent of spatial interferometry than of normal optical imaging.
 Iu fact. the two harmonics of the detector signal correspoud to the fringes produced by iu object iu two interferometers with baselines of 9£ nuu aud LSS nua. respectively (assuming an effective waveleugth of 550 um).," In fact, the two harmonics of the detector signal correspond to the fringes produced by an object in two interferometers with baselines of 94 mm and 188 mm, respectively (assuming an effective wavelength of 550 nm)."
 hnuage reconstruction techniques using iuterterometric observations have for a long tine been standard in the radio astronomical commmuity., Image reconstruction techniques using interferometric observations have for a long time been standard in the radio astronomical community.
 It was therefore quite natural to apply these techniques to the Wipparcos data (Lindeeren 1982: Quist ct citequist))., It was therefore quite natural to apply these techniques to the Hipparcos data (Lindegren \cite{lind}; ; Quist et \\cite{quist}) ).
 Equation (13) gives the signal for a point source located at the position r=(rc.y) relative the reference point.," Equation \ref{eq:ikphi}) ) gives the signal for a point source located at the position $\vec{r}=(x,y)$ relative the reference point."
 Now consider an extended object with the ecneral brightucss distribution B(r)., Now consider an extended object with the general brightness distribution $B(\vec{r})$.
 Sumunine up the coutributious to the detector signal frou each clement of the sky we fiud (apart fron a constant scaling factor) where S(r) ids the sensitivity profile of the instantaneous field of view., Summing up the contributions to the detector signal from each element of the sky we find (apart from a constant scaling factor) where $S(\vec{r})$ is the sensitivity profile of the instantaneous field of view.
 Introducing the Fouricr transfor. we fud that Eq. (8)), Introducing the Fourier transform we find that Eq. \ref{eq:ip}) )
 can be written Comparison with Eq. (1)), can be written Comparison with Eq. \ref{eq:ik}) )
 shows that where the asterisk denotes the complex conjugate., shows that where the asterisk denotes the complex conjugate.
" Since Af, and Mo are conventional constants (Sect. 2.3))", Since $M_1$ and $M_2$ are conventional constants (Sect. \ref{sec:phase}) )
 it is seen that a single transit defines the complex function Vin the five points 0. +f. 2f. of the spatial frequency plane.," it is seen that a single transit defines the complex function $V$ in the five points $\vec{0}$, $\pm\vec{f}$, $\pm 2\vec{f}$, of the spatial frequency plane."
 However. the conjugate sviunietry of V. means that there are ouly three independent complex visibilitics per transit.," However, the conjugate symmetry of $V$ means that there are only three independent complex visibilities per transit."
 Successive ransits of the same object are made at different spatia frequencies f=(fy.fy) aud oei the course of the mission knowledge of 1ο function —1 is built up in a munber of different points.," Successive transits of the same object are made at different spatial frequencies $\vec{f}=(f_x,f_y)$ and in the course of the mission knowledge of the function $V$ is built up in a number of different points."
 To the xteut that Str)Bir) remains coustaut over the mission. --- iav then be recoverable from V. using standi nuage recoustruction techniques.," To the extent that $S(\vec{r})B(\vec{r})$ remains constant over the mission, it may then be recoverable from $V$ using standard image reconstruction techniques."
 Iu thecontest of radio interferometry aud :)orture svuthesis. we may identify V(f) with the complex visibility function associated with the source xiehtuess distribution Br} and siugle-auteuna reception pattern Str) (Thompson ct citethom)).," In thecontext of radio interferometry and aperture synthesis, we may identify $V(\vec{f})$ with the complex visibility function associated with the source brightness distribution $B(\vec{r})$ and single-antenna reception pattern $S(\vec{r})$ (Thompson et \\cite{thom}) )."
 The visibility Tuaction Is usually expressed iu terms of coordinates αν0) which give the projection of the interferometer baseline ou the sky plane aud are expressed in wavceleusths.," The visibility function is usually expressed in terms of coordinates $(u,v)$ which give the projection of the interferometer baseline on the sky plane and are expressed in wavelengths."
 The relation to the TD spatial frequency components is simply Our reference point is equivalent to the phase reference position used in counected-clemenut radio interferometry (Thompson ct citethom)). or to the strong reference point ποιος (typically a quasar) used in phase-reterenced VLBI observations (Lestrade οἳ citelestr)).," The relation to the TD spatial frequency components is simply Our reference point is equivalent to the phase reference position used in connected-element radio interferometry (Thompson et \\cite{thom}) ), or to the strong reference point source (typically a quasar) used in phase-referenced VLBI observations (Lestrade et \\cite{lestr}) )."
 The distribution of the observations iu the «c plane is albiuportaut for the possibility to reconstruc complicated images from the measured visibilities V(u.0).," The distribution of the observations in the $uv$ plane is all-important for the possibility to reconstruct complicated images from the measured visibilities $V(u,v)$."
 Untortunately the Ilipparcos scanning law aud the use of a niodulatiug exid with just a single period seriously. Hliuüt the we coverage of the TD., Unfortunately the Hipparcos scanning law and the use of a modulating grid with just a single period seriously limit the $uv$ coverage of the TD.
 According to Eqs. (3)), According to Eqs. \ref{eq:fxfy}) )
 iux (12)) the coverage is luted to he ceutral point (a.0)=(0.0) and two concentric nues with radi zLTO830 ane 311660 waveleneths.," and \ref{eq:uv}) ) the coverage is limited to the central point $(u,v)=(0,0)$ and two concentric rings with radii $\simeq 170\,830$ and $341\,660$ wavelengths."
 Moreover. for objects in the ecliptic reeion of the sky (ecliptic latitude ||zi 457) the scanning law constraius the scan angle 0 such as to produce a eap of τσμας scans roughly in the eastwest direction.," Moreover, for objects in the ecliptic region of the sky (ecliptic latitude $|\beta| \la 45^\circ$ ) the scanning law constrains the scan angle $\theta$ such as to produce a gap of `missing' scans roughly in the east–west direction."
 At E)~[T there is instead a surplus of scaus in the eastwest direction Fie. 3))., At $|\beta| \simeq 47^\circ$ there is instead a surplus of scans in the east--west direction Fig. \ref{fig:uv}) ).
 For high-latitude objects. finally. the coverage is usually more wuiform inu.," For high-latitude objects, finally, the coverage is usually more uniform in $\theta$ ."
 Tn continuing the analogy with radio interferometry. we will discuss in the nextsection how images can be produced using the Transit Data.," In continuing the analogy with radio interferometry, we will discuss in the nextsection how images can be produced using the Transit Data."
 Utilizing the experieuce developed for aperture svutlesis iuaging. we use only," Utilizing the experience developed for aperture synthesis imaging, we use only"
aree. L.702.03 (significantly larger than the -uncertaintv range of the average P iu Table 2)).,"large, 1.70–2.03 (significantly larger than the -uncertainty range of the average $\Gamma$ in Table \ref{t:fits}) )."
 We have then investigated which plwvsical ILOCCSS can be responsible for the observed vchavior., We have then investigated which physical process can be responsible for the observed behavior.
 A natural canidate here is thermal Comptonization in a hot plasina cloud. irradiated oa variable fix of soft seed photons. as suggested w MOI.," A natural candidate here is thermal Comptonization in a hot plasma cloud irradiated by a variable flux of soft seed photons, as suggested by M91."
 If the power supplied. to the cloud aud its Thomson optical depth. 7. are constant. its clectrou teniperatire. KT. will adjust itself to the variable seed flux as to satisfy energy balance.," If the power supplied to the cloud and its Thomson optical depth, $\tau$, are constant, its electron temperature, $kT$, will adjust itself to the variable seed flux as to satisfy energy balance."
 Then. he higher the seed Hux. the lower AL. and (since a decrease of &T at a constaut 7 corresponds ο an increase ofthe N-rav sctral index. D) the softer the X-ray spectrum.," Then, the higher the seed flux, the lower $kT$, and (since a decrease of $kT$ at a constant $\tau$ corresponds to an increase of the X-ray spectral index, $\Gamma$ ) the softer the X-ray spectrum."
 Since he Comptonized spectruni joius at low cnereics to the peak of je spectrum of fie seed: photous. this will also correspond to an increase of he soft N-rav flux.," Since the Comptonized spectrum joins at low energies to the peak of the spectrum of the seed photons, this will also correspond to an increase of the soft X-ray flux."
 To test whether this process ca1 iudeed reproduce 16 observed spectral variability. we have used a Comptonization code by Cop3 (1992. 1999). with its present version described in Cierlinsski et ((1999).," To test whether this process can indeed reproduce the observed spectral variability, we have used a Comptonization code by Coppi (1992, 1999), with its present version described in Gierlińsski et (1999)."
 This model is relatively simular to the one used in $3.2 except for &T being calculated, This model is relatively similar to the one used in \ref{s:results} except for $kT$ being calculated
has much less short wavelength [ιν than do the more realise models: the short wavelength extension of the spectra of the latter arises because of changes in the opacity wilh wavelength that expose the deeper. hotter atmospheric lavers al these wavelengths.,"has much less short wavelength flux than do the more realistic models: the short wavelength extension of the spectra of the latter arises because of changes in the opacity with wavelength that expose the deeper, hotter atmospheric layers at these wavelengths."
 Addition of Fe with an abundance equal to the best eurrent upper limit (Table 1) however. dramatically reduces the soft N-rav αν; clearly any. photospheric soft X-ray emission from acid similar stars is critically dependent on (he photospheric abundances of heavier elements. as was also demonstrated earlier in the case of ELV fIuxes at longer wavelengths for slightly cooler stars (e.g.Darstowetal.1995).," Addition of Fe with an abundance equal to the best current upper limit (Table 1) however, dramatically reduces the soft X-ray flux: clearly any photospheric soft X-ray emission from and similar stars is critically dependent on the photospheric abundances of heavier elements, as was also demonstrated earlier in the case of EUV fluxes at longer wavelengths for slightly cooler stars \citep[e.g.][]{Barstow.etal:95}."
. Coronal and photospheric model spectra are compared to the observations in Figure 4.., Coronal and photospheric model spectra are compared to the observations in Figure \ref{f:compare}.
 Coronal models computed for the IHe-rich. composition listed in Table 1 are dominated by the Ile bound-free and [ree-Iree continua. and by the lines of IHe-like C near 40A.," Coronal models computed for the He-rich composition listed in Table 1 are dominated by the He bound-free and free-free continua, and by the lines of He-like C near 40."
. While it is trempting to interpret the peak in the observed Πας between 40 and 45 aas being due to these C V lines. there are no such lines visible in (he spectrum where the observed counts ave more smoothly distributed.," While it is tempting to interpret the peak in the observed flux between 40 and 45 as being due to these C V lines, there are no such lines visible in the spectrum where the observed counts are more smoothly distributed."
 The absence of C V lines does not necessarily argue against the coronal interpretation of (he spectrum. however. since one might expect heavier elements (o eravilationally settle out of such a plasma.," The absence of C V lines does not necessarily argue against the coronal interpretation of the spectrum, however, since one might expect heavier elements to gravitationally settle out of such a plasma."
 We therefore also computed models in which the metal abundances were reduced by a factor of 100., We therefore also computed models in which the metal abundances were reduced by a factor of 100.
" We conlivm the result of Flemingetal.(1993) that the best-fit coronal spectrum has a temperature in the range 2-3xLO? Ix: our best-lit model corresponds to T—2xLO? Kk. with an interstellar absorption component represented by a column of neutral hydrogen of 3.5x107"" 7."," We confirm the result of \citet{Fleming.etal:93} that the best-fit coronal spectrum has a temperature in the range $3\times 10^5$ K; our best-fit model corresponds to $T=2\times 10^5$ K, with an interstellar absorption component represented by a column of neutral hydrogen of $3.5\times 10^{20}$ $^{-2}$."
 Qualitative comparison of the best-fit coronal model and data suggest. however. that the latter have a more complex shape than can be achieved by the former.," Qualitative comparison of the best-fit coronal model and data suggest, however, that the latter have a more complex shape than can be achieved by the former."
 Photospheric models both with and without Fe were computed. as was a test mocel including 9: a small subsection of the models investigated are illustrated in the lower panel ol Figure 4..," Photospheric models both with and without Fe were computed, as was a test model including S; a small subsection of the models investigated are illustrated in the lower panel of Figure \ref{f:compare}."
 It is clear that models with only C. O and Ne greatly overpredict the observed X-ray flux by. factors of 10 or more.," It is clear that models with only C, O and Ne greatly overpredict the observed X-ray flux by factors of 10 or more."
 Based on Figure 2.. we see that the emergent soft. [lux can be reduced by addition of Fe aud the opacity this element provides at these wavelengths.," Based on Figure \ref{f:fluxes}, we see that the emergent soft X-ray flux can be reduced by addition of Fe and the opacity this element provides at these wavelengths."
" This is borne out in practice. and the photospheric spectrum that was found to best match the observations corresponds roughly to the parameters 7,;;=120000 Ix and Fe/Ile=—5.5. with a neutral hydrogen column density of 3κ105)20) em27."," This is borne out in practice, and the photospheric spectrum that was found to best match the observations corresponds roughly to the parameters $T_{eff}=120000$ K and $=-5.5$, with a neutral hydrogen column density of $3\times 10^{20}$ $^{-2}$."
 It is clear from Figure 4 that there is no obvious reason to prefer coronal models over photospheric ones., It is clear from Figure \ref{f:compare} that there is no obvious reason to prefer coronal models over photospheric ones.
 While we have not succeeded in producing a photospheric model (hat perfectly matches the observations. we have achieved a good qualitative match by adding only plausible amounts of Fe to the C. O and Ne abundance mixture.," While we have not succeeded in producing a photospheric model that perfectly matches the observations, we have achieved a good qualitative match by adding only plausible amounts of Fe to the C, O and Ne abundance mixture."
" Trace amounts of other metals such as Na. Mg. ο, Ar. Ca and Ni could also affect signilicantly the soft X-ray spectrum. and we could doubtless obtain a better match through their addition and subsequent optimisation"," Trace amounts of other metals such as Na, Mg, S, Ar, Ca and Ni could also affect significantly the soft X-ray spectrum, and we could doubtless obtain a better match through their addition and subsequent optimisation"
telescope and reported in N10 (see Figures and py.,telescope and reported in N10 (see Figures \ref{single_profiles1} and \ref{single_profiles2}) ).
" Each target received simultaneous 3x 110 s integrations with the g-filter and x 60 integrations with the i-filter using the blue and red arms of LRIS, respectively."," Each target received simultaneous $\,\times\,$ 110 s integrations with the -filter and $\,\times\,$ 60 integrations with the -filter using the blue and red arms of LRIS, respectively."
" A 20"" dither pattern was employed to generate on-sky flat-field frames.", A $''$ dither pattern was employed to generate on-sky flat-field frames.
 We performed photometric calibration using 1s g- and band observations of bright stars in each field., We performed photometric calibration using 1 s - and -band observations of bright stars in each field.
 The data were reduced using IDL routines and combined and analyzed using standard IRAF tasks; the seeing FWHM of the final science exposures is ~0.8”., The data were reduced using IDL routines and combined and analyzed using standard IRAF tasks; the seeing FWHM of the final science exposures is $\sim$ $''$.
 To construct models of the light profiles for each lensing system we use GALFIT (Peng et al., To construct models of the light profiles for each lensing system we use GALFIT (Peng et al.
" 2002), which allows multiple profiles per object and performs a simultaneous nonlinear minimization."," 2002), which allows multiple profiles per object and performs a simultaneous nonlinear minimization."
" Prior to fitting profiles to SDP.81 and SDP.130 in the IRAC data, de Vaucouleurs profile (de Vaucouleurs 1948) models are constructed for the Keck +band imaging, to take advantage of the comparatively higher resolution of these images, and the presence of only one component per system i.e. the lens galaxy."," Prior to fitting profiles to SDP.81 and SDP.130 in the IRAC data, de Vaucouleurs profile (de Vaucouleurs 1948) models are constructed for the Keck -band imaging, to take advantage of the comparatively higher resolution of these images, and the presence of only one component per system i.e. the lens galaxy."
" To look for lensed structure in the IRAC data, the anyKeck models potentiallyare used to fit the IRAC Channels 1 and 2 data, keeping the effective radius and ellipticity fixed, and using the appropriate IRAC PSF 36"," To look for any potentially lensed structure in the IRAC data, the Keck models are used to fit the IRAC Channels 1 and 2 data, keeping the effective radius and ellipticity fixed, and using the appropriate IRAC PSF $^{36}$."
" On subtraction of the results, and in comparison to the model subtracted Keck data and SMA contours, the IRAC band residuals strongly suggest that a more complex structure, associated with the background SMG, is present for both SDP.81 and SDP.130 (see Figures and We [2))."," On subtraction of the results, and in comparison to the model subtracted Keck data and SMA contours, the IRAC band residuals strongly suggest that a more complex structure, associated with the background SMG, is present for both SDP.81 and SDP.130 (see Figures \ref{single_profiles1} and \ref{single_profiles2}) )."
"verify that these residuals are significant, and not an artifact of imperfections in the IRAC PSF, by comparing them with residuals derived for three (non-lensing) elliptical-like galaxies in the same field, after fitting them with single Sérrsic profiles."," We verify that these residuals are significant, and not an artifact of imperfections in the IRAC PSF, by comparing them with residuals derived for three (non-lensing) elliptical-like galaxies in the same field, after fitting them with single Sérrsic profiles."
 For both the lensing systems and the comparison ellipticals aperture flux ratios were determined for the residual image and the corresponding un-subtracted data., For both the lensing systems and the comparison ellipticals aperture flux ratios were determined for the residual image and the corresponding un-subtracted data.
" To consider only positive structure, pixel values >20 below the local background were replaced with the median local sky value."," To consider only positive structure, pixel values $>\,2\,\sigma$ below the local background were replaced with the median local sky value."
 The SMG residuals were found to have flux ratios 3—5 greater than those for the random elliptical," The SMG residuals were found to have flux ratios $\sim\,3-5$ greater than those for the random elliptical"
 (ce.Frailetal.1991:Green1997).. 2001).. (παν1992).. |50 (c.g...Sjouwerman&," \citep[e.g.,][]{frail94,green97}. \citep[e.g.,][]{macleod97,szymczak04,niezurawska04}, \citep{gray91,gray92}."
Pillstróni2008.hereafterPaperlh.," $+50$ \citep[e.g.,][hereafter
Paper\,I]{sjouwerman08}."
 GGklauss+2008).., $+34\leq V_{\rm LSR}\leq 66$ $^{-1}$.
 Iu addition. two separate eroups of niasers are located near he Cicunuuuclear Disk (CND) that are not directly explained by the SNR/ISM interaction model.," In addition, two separate groups of masers are located near the Circumnuclear Disk (CND) that are not directly explained by the SNR/ISM interaction model."
 These wo groups of nisers have velocities that are offset roni the EEast mascrs: τω|L30kkunss | and Vpagc L30kkurss| respectively Π)., These two groups of masers have velocities that are offset from the East masers; $V_{\rm LSR}\simeq+130$ $^{-1}$ and $V_{\rm LSR}\simeq-130$ $^{-1}$ respectively I).
 In Paper we argue that these masers are unlikely to ο punrped by a shock produced bv EEast., In I we argue that these masers are unlikely to be pumped by a shock produced by East.
 Other plausible pumping scenarios inchide local shocks xoduced. by random motions of clumps or turbulence. supported by the presence of strong I] (1-0) (1) cluission (Yusef-Zacdeletal.2001).. or bv infrared (IR) Munpineg smi to conditious observed in star forming regions (Carayetal.1992).," Other plausible pumping scenarios include local shocks produced by random motions of clumps or turbulence, supported by the presence of strong $_2$ (1-0) $S(1)$ emission \citep{yusefzadeh01}, or by infrared (IR) pumping similar to conditions observed in star forming regions \citep{gray92}."
. The cuviromment in the CND is likely to be different roni that in à SNR/ISM post-shock region., The environment in the CND is likely to be different from that in a SNR/ISM post-shock region.
 Variability studies of masers can be used to probe the euviroumenut and could further shed πο on the differences in excitation mcchanisins aud conditions for CND versus SNR /ISADmasers., Variability studies of masers can be used to probe the environment and could further shed light on the differences in excitation mechanisms and conditions for CND versus SNR/ISM masers.
 Iu the CND. the pumping may inclide IR pumping routes. since local IR peaks are observed within the CND (Latvakoskietal.1999).," In the CND, the pumping may include IR pumping routes, since local IR peaks are observed within the CND \citep{latvakoski99}."
. This would be in contrast to the standard SNR iasers. which are pimped by collisions only.," This would be in contrast to the 'standard' SNR masers, which are pumped by collisions only."
 Also. differcuces in collision rate along the line of sight. due to chuupiuess of the medi. could result iu a different maser flux variability of the CND masers as compared to the SNR/ISM imasers," Also, differences in collision rate along the line of sight, due to clumpiness of the medium, could result in a different maser flux variability of the CND masers as compared to the SNR/ISM masers"
Analvzing the past orbit of the Magellanie Clouds relative to the Milky Way is an interesting problem in dvnanmücs. and it may be a useful preliminary exercise for analvses of ongoing advances in measurements of galaxy clistances and proper motions (hat will test ideas about (he evolving mass distribution around galaxies.,"Analyzing the past orbit of the Magellanic Clouds relative to the Milky Way is an interesting problem in dynamics, and it may be a useful preliminary exercise for analyses of ongoing advances in measurements of galaxy distances and proper motions that will test ideas about the evolving mass distribution around galaxies."
 Since the positions. velocities and accelerations of the Magellanie Clouds are fairly well constrained at low redshifts (Ixallivavalil 2006. Piatek. Prvor Olszewsk 2008). the challenge here is to understand how the Clouds were directed to (their present paths by interactions among newly forming galaxies ab high redshift.," Since the positions, velocities and accelerations of the Magellanic Clouds are fairly well constrained at low redshifts (Kallivayalil 2006, Piatek, Pryor Olszewsk 2008), the challenge here is to understand how the Clouds were directed to their present paths by interactions among newly forming galaxies at high redshift."
 This analvsis of the motion of the Large Magellanic Cloud (LMC) around the Milky Way (MW) focuses on the influence of galaxies in and near the Local Group., This analysis of the motion of the Large Magellanic Cloud (LMC) around the Milky Way (MW) focuses on the influence of galaxies in and near the Local Group.
 The Small Cloud is taken to be an unimportant perturbation at the hoped for accuracy of this study., The Small Cloud is taken to be an unimportant perturbation at the hoped for accuracy of this study.
 The first steps to the interpretation of the measured proper motion of the Large Magellanic Cloud considered its interaction with AIW and. M31. (Besla οἱ al., The first steps to the interpretation of the measured proper motion of the Large Magellanic Cloud considered its interaction with MW and M31 (Besla et al.
 2007: Shattow Loeb 2009: Wallivavalil 2009)., 2007; Shattow Loeb 2009; Kallivayalil 2009).
 Peebles (2009. P9) added. (he initial condition indicated by the gravitational growth of structure in the standard Dig Bane cosmology. that peculiar velocities of the protogalaxies al hieh redshift are growing.," Peebles (2009, P9) added the initial condition indicated by the gravitational growth of structure in the standard Big Bang cosmology, that peculiar velocities of the protogalaxies at high redshift are growing."
 P9 applied this condition to a illustrative mass model., P9 applied this condition to a illustrative mass model.
 Here the analvsis is extended to a more realistie model that takes, Here the analysis is extended to a more realistic model that takes
Type I X-ray bursts are well-understood as thermonuclear flashes on the surface of an accreting neutron star (see Lewin. van Paradijs. Taam 1995; Strohmayer Bildsten 2003 for reviews).,"Type I X-ray bursts are well-understood as thermonuclear flashes on the surface of an accreting neutron star (see Lewin, van Paradijs, Taam 1995; Strohmayer Bildsten 2003 for reviews)."
 The basic theory was outlined more than twenty years ago (for an overview. see Fujimoto. Hanawa. Miyajt 1981: Bildsten 1998). successfully explaining the typical burst energies (10— 1079 eres) durations (~ 10-100 s). and recurrence times (hours to days).," The basic theory was outlined more than twenty years ago (for an overview, see Fujimoto, Hanawa, Miyaji 1981; Bildsten 1998), successfully explaining the typical burst energies $10^{39}$ $10^{40}\ {\rm
ergs}$ ), durations $\sim 10$ $100\ {\rm s}$ ), and recurrence times (hours to days)."
 However. detailed comparisons of theory and observations have had mixed success (Fujimoto et al.," However, detailed comparisons of theory and observations have had mixed success (Fujimoto et al."
 1987: van Paradijs. Penninx. Lewin 1988; Bildsten 2000: Galloway et al.," 1987; van Paradijs, Penninx, Lewin 1988; Bildsten 2000; Galloway et al."
 2003)., 2003).
" The ultracompact binary 4U 1820-30 (P,=11.4min: Stella. Priedhorsky. White 1987) is à promising system for such a comparison."," The ultracompact binary 4U 1820-30 $P_{\rm orb}=11.4\ {\rm min}$; Stella, Priedhorsky, White 1987) is a promising system for such a comparison."
 It has a known distance. being located in the metal-rich globular cluster NGC 6624 (|Fe/H]z—0.4. distance 7.6x0.4 kpe): Rich. Minniti. Liebert 1993; Kuulkers et al.," It has a known distance, being located in the metal-rich globular cluster NGC 6624 $[$ $/$ $]\approx -0.4$, distance $7.6\pm
0.4\ {\rm kpc}$ ); Rich, Minniti, Liebert 1993; Kuulkers et al."
 2003)., 2003).
 It undergoes a regular z176 day accretion cycle (Priedhorsky Terrell 1984). switching between high and low states differing by a factor z3 in accretion rate (e.g. see the RXTE/ASM light curve in Figure | of Strohmayer Brown 2002).," It undergoes a regular $\approx 176$ day accretion cycle (Priedhorsky Terrell 1984), switching between high and low states differing by a factor $\approx 3$ in accretion rate (e.g. see the /ASM light curve in Figure 1 of Strohmayer Brown 2002)."
 In the low state. regular Type I bursts are seen (3-23 days around the minimum luminosity: Chou Grindlay 2001). with recurrence times z 2-4 hours (Clark et al.," In the low state, regular Type I bursts are seen $\pm 23$ days around the minimum luminosity; Chou Grindlay 2001), with recurrence times $\approx 2$ $4$ hours (Clark et al."
 1976; Clark et al., 1976; Clark et al.
 1977: Haberl et al., 1977; Haberl et al.
 1987: Cornelisse et al., 1987; Cornelisse et al.
 2003)., 2003).
" Ne bursts are seen during the rest of the cycle. implying a “non-bursting"" mode of burning (e.g. Bildsten 1995)."," No bursts are seen during the rest of the cycle, implying a ``non-bursting'' mode of burning (e.g. Bildsten 1995)."
" In the low state. 4U 1820-30 has also shown an extremely energetic ?107 erg) ""superburst"". likely due to deep ignition of a carbon layer (Strohmayer Brown 2002; see Kuulkers et al."," In the low state, 4U 1820-30 has also shown an extremely energetic $\sim 10^{42}\ {\rm erg}$ ) “superburst”, likely due to deep ignition of a carbon layer (Strohmayer Brown 2002; see Kuulkers et al."
 2002 and Strohmayer Bildsten 2003 for a summary of superburst properties)., 2002 and Strohmayer Bildsten 2003 for a summary of superburst properties).
 Different evolutionary scenarios for 4U 1820-30 have been proposed., Different evolutionary scenarios for 4U 1820-30 have been proposed.
 Those involving direct collision of a red giant and a neutron star (Verbunt 1987). or formation of a neutron sequence binary by tidal capture (Bailyn Grindlay 1987) or an exchange interaction (Rasio. Pfahl. Rappaport 2000) followed by à common envelope phase. lead to accretion of pure helium from the helium core of the red giant.," Those involving direct collision of a red giant and a neutron star (Verbunt 1987), or formation of a neutron sequence binary by tidal capture (Bailyn Grindlay 1987) or an exchange interaction (Rasio, Pfahl, Rappaport 2000) followed by a common envelope phase, lead to accretion of pure helium from the helium core of the red giant."
 A different picture is that the mass transfer starts Just after central hydrogen exhaustion (Tutukov et al., A different picture is that the mass transfer starts just after central hydrogen exhaustion (Tutukov et al.
 1987: Fedorova Ergma 1989: Podsiadlowski. Rappaport. Pfahl 2002). in which case the accreted material is again mostly helium. but contains some hydrogen (~ 5-35% by mass: Fedorova Ergma 1989; Podsiadlowski. Rappaport. Pfahl 2002).," 1987; Fedorova Ergma 1989; Podsiadlowski, Rappaport, Pfahl 2002), in which case the accreted material is again mostly helium, but contains some hydrogen $\approx 5$ $35$ by mass; Fedorova Ergma 1989; Podsiadlowski, Rappaport, Pfahl 2002)."
 The observed energetics and recurrence times of the Type I bursts from 4U 1820-30 fit well with a picture of accumulation and burning of heltum-rich. material., The observed energetics and recurrence times of the Type I bursts from 4U 1820-30 fit well with a picture of accumulation and burning of helium-rich material.
 The ratio of persistent fluence (persistent flux integrated over the burst recurrence time) to burst fluence is àz120 (Haberl et al., The ratio of persistent fluence (persistent flux integrated over the burst recurrence time) to burst fluence is $\alpha\approx 120$ (Haberl et al.
 1987)., 1987).
 For a gravitational energy GM/Rz200 MeV per nucleon. this implies à nuclear energy release ~1.6 MeV per nucleon. as expected for helium burning to tron group elements.," For a gravitational energy $GM/R\approx 200$ MeV per nucleon, this implies a nuclear energy release $\approx 1.6$ MeV per nucleon, as expected for helium burning to iron group elements."
 The mass of helium required to power the observed energetics is roughly the mass accreted in the burst recurrence time at the inferred accretion rate., The mass of helium required to power the observed energetics is roughly the mass accreted in the burst recurrence time at the inferred accretion rate.
 The burst fluence z3.5«1077ergem? (Haberl et al., The burst fluence $\approx 3.5\times 10^{-7}\ {\rm erg\ cm^{-2}}$ (Haberl et al.
 1987) implies a total burst energy 2.5«10?erg(d/7.6 kpey. or a mass of helium AM~1.6«107!e.," 1987) implies a total burst energy $2.5\times 10^{39}\ {\rm erg}\ (d/7.6\
{\rm kpc})^2$ , or a mass of helium $\Delta M\approx 1.6\times 10^{21}\
{\rm g}$."
" The persistent luminosity when bursts are seen is Lyz2.8«1077ereso!(d/7.6kper(Fy/4«ergem? s). giving an accretion rate Mz1.5«10""es!2.4&1077M...yr! for a L4M.. neutron star with radius R=10km."," The persistent luminosity when bursts are seen is $L_X\approx 2.8\times 10^{37}\ {\rm erg\ s^{-1}}\ (d/7.6\ {\rm
kpc})^2(F_X/4\times 10^{-9}\ {\rm erg\ cm^{-2}\ s^{-1}})$ , giving an accretion rate $\dot M\approx 1.5\times 10^{17}\ {\rm g\
s^{-1}}\approx 2.4\times 10^{-9}\ M_\odot\ {\rm yr^{-1}}$ for a $1.4\
M_\odot$ neutron star with radius $R=10\ {\rm km}$."
 At this rate. a mass AM is accreted in 3 hours. in excellent agreement with the observed recurrence times.," At this rate, a mass $\Delta M$ is accreted in 3 hours, in excellent agreement with the observed recurrence times."
 A previous comparison of theory with observations of 4U 1820-30 was carried out by Bildsten (1995. hereafter B95). who conducted time-dependent simulations of pure helium burning on accreting neutron stars.," A previous comparison of theory with observations of 4U 1820-30 was carried out by Bildsten (1995, hereafter B95), who conducted time-dependent simulations of pure helium burning on accreting neutron stars."
 He found good agreement with observed energetics and recurrence times. but for a somewhat hotter base temperature than expected.," He found good agreement with observed energetics and recurrence times, but for a somewhat hotter base temperature than expected."
 In this paper. Icaleulate ignition models for Type I bursts from 4U," In this paper, Icalculate ignition models for Type I bursts from 4U"
 In this paper. Icaleulate ignition models for Type I bursts from 4U.," In this paper, Icalculate ignition models for Type I bursts from 4U"
"In the last section, we have pointed out that the observed M—Z relation could indicate more star bursts or longer duration of each burst or higher SFE in more massive galaxies, if no outflow takes place.","In the last section, we have pointed out that the observed $M-Z$ relation could indicate more star bursts or longer duration of each burst or higher SFE in more massive galaxies, if no outflow takes place."
" Now we examine the possibility of the normal wind being the explanation of the M—Z relation, as suggested by many previous authors."," Now we examine the possibility of the normal wind being the explanation of the $M-Z$ relation, as suggested by many previous authors."
" In Fig. 9,,"," In Fig. \ref{Fig:wdMsZ},"
" we take the models with 7 bursts for example (i.e., Model M8n4, M9n4, M10n4 in no-wind case) but different strengthes of normal wind are introduced (A,— 0,0.2,0.5,1,3)."," we take the models with 7 bursts for example (i.e., Model M8n4, M9n4, M10n4 in no-wind case) but different strengthes of normal wind are introduced $\lambda_{w}=0,
0.2, 0.5, 1, 3$ )."
 It is evident from the upper panel of Fig., It is evident from the upper panel of Fig.
" 9 that by varying only the efficiency of a normal wind one cannot reproduce the M—Z relation, unless other parameters, such as the efficiency of SF or the number of bursts, are assumed to vary as functions of the galactic mass."," \ref{Fig:wdMsZ} that by varying only the efficiency of a normal wind one cannot reproduce the $M-Z$ relation, unless other parameters, such as the efficiency of SF or the number of bursts, are assumed to vary as functions of the galactic mass."
" When a long infall timescale is adopted (r—10 Gyr, lower panel of Fig. 9)),"," When a long infall timescale is adopted $\tau=10$ Gyr, lower panel of Fig. \ref{Fig:wdMsZ}) ),"
" less stars are formed in each model due to the slow gas accretion process, and the wind could not be induced in the high mass systems (Mingz 10'°Mo))."," less stars are formed in each model due to the slow gas accretion process, and the wind could not be induced in the high mass systems $M_{inf}\approx10^{10}$ )."
" However in the low mass galaxy (Ming 105M5)), where the wind could develop, the newly infalling gas dilutes the ISM effectively."," However in the low mass galaxy $M_{inf}=10^8$ ), where the wind could develop, the newly infalling gas dilutes the ISM effectively."
" Therefore, the general trend of the M—Z relation could be reproduced by combining the normal wind with a slow accretion process."," Therefore, the general trend of the $M-Z$ relation could be reproduced by combining the normal wind with a slow accretion process."
" In summary, the normal wind can strongly reduce the gas fraction but it cannot reduce sensibly the O/H. To explain the spread observed in O/H at the same µ we should invoke other mechanisms, such as a continuous supplement of primordial gas or a metal-enhanced wind (see Sect."," In summary, the normal wind can strongly reduce the gas fraction but it cannot reduce sensibly the O/H. To explain the spread observed in O/H at the same $\mu$ we should invoke other mechanisms, such as a continuous supplement of primordial gas or a metal-enhanced wind (see Sect."
" 4.3), both of which imply a lower mass loss rate for H and He relative to metals."," 4.3), both of which imply a lower mass loss rate for H and He relative to metals."
" The galactic wind, mainly induced by SN explosion, could blow preferentially the metal-enriched gas out of the"," The galactic wind, mainly induced by SN explosion, could blow preferentially the metal-enriched gas out of the"
LIRGs compared to those in normal galaxies rule out the possibility that the eiant regions are just ageregations of “normal” regions.,LIRGs compared to those in normal galaxies rule out the possibility that the giant regions are just aggregations of “normal” regions.
 Another possibility. would be that all regions im LIRGs were extremely vounug., Another possibility would be that all regions in LIRGs were extremely young.
" However. previously published spectroscopy of a few eiat reeions in Arp 299 and NCC 3256 has shown a rauge of equivalent widths of Πα, in other words. a range of ages."," However, previously published spectroscopy of a few giant regions in Arp 299 and NGC 3256 has shown a range of equivalent widths of $\alpha$, in other words, a range of ages."
 A more plausible explanation for the preseuce of extremely huninous regions ΡΕ is that the regious of hieli gas pressure and deusitv present oein LIRGs. ULIRGs aud interacting ealaxies provide the necessary conditions for the formation of a large umber of massive star (oniziug) clusters.," A more plausible explanation for the presence of extremely luminous regions in LIRGs is that the regions of high gas pressure and density present in LIRGs, ULIRGs and interacting galaxies provide the necessary conditions for the formation of a large number of massive star (ionizing) clusters."
 Such extreme couditious are not Likely to occur in normal isolated ealaxies., Such extreme conditions are not likely to occur in normal isolated galaxies.
 Despite the large mmubers of near-intrared SSCs and regions identified in LIRCs. there is ouly a simall fraction of coincidences between and of the total number of detected regious aud star clusters.," Despite the large numbers of near-infrared SSCs and regions identified in LIRGs, there is only a small fraction of coincidences – between and of the total number of detected regions and star clusters."
 This is sugeestive of an evolutionary sequence similar to that seen in obseured Galactic giaut regions., This is suggestive of an evolutionary sequence similar to that seen in obscured Galactic giant regions.
 For the first few million years of the cluster evolution. there will be significaut amounts of dust which may preclude the detection of the ioniziug cluster.," For the first few million years of the cluster evolution, there will be significant amounts of dust which may preclude the detection of the ionizing cluster."
 As the cluster ages. stellar winds aud superuovae will dissipate part of the gas and dust. aud the star cluster will become visible. whereas at the same tine the reeion cauission will be less Iuniinous.," As the cluster ages, stellar winds and supernovae will dissipate part of the gas and dust, and the star cluster will become visible, whereas at the same time the region emission will be less luminous."
 We have used evolutionary svuthesis models to reproduce the observed relative fractions of vouug and intermediate reeious/clusters and old) clusters in Arp 299 and NGC 3256., We have used evolutionary synthesis models to reproduce the observed relative fractions of young and intermediate regions/clusters and old clusters in Arp 299 and NGC 3256.
 Based on the observed fractions we have concluded that most likely the star formation occurs in instantaneous bursts rather than more extended periods of star formation., Based on the observed fractions we have concluded that most likely the star formation occurs in instantaneous bursts rather than more extended periods of star formation.
 Using iustautauceous star formation aud a Salpeter IAIF we have derived photometric masses of the detected star clusters of between σον104M. and 105M...," Using instantaneous star formation and a Salpeter IMF we have derived photometric masses of the detected star clusters of between $5\times 10^4\,{\rm M}_{\odot}$ and $10^6\,{\rm M}_{\odot}$."
 The fact that the peak of the f-hand Iuniuositv occurs after approximately ADIr whereas at the same time the uumiber of ionizing photons has dropped by about two orders of magnitudes from the maximum explains the limited nuunber of coincidences., The fact that the peak of the $H$ -band luminosity occurs after approximately Myr whereas at the same time the number of ionizing photons has dropped by about two orders of magnitudes from the maximum explains the limited number of coincidences.
 Within the prescut detection Bnits in Arp 299 and NGC 3256. we can only detect both region cussion aud a star cluster for the most massive clusters (7109 AL.) duriug the first MMyr.," Within the present detection limits in Arp 299 and NGC 3256, we can only detect both region emission and a star cluster for the most massive clusters $\simeq 10^6\,{\rm M}_{\odot}$ ) during the first Myr."
 The near-intrared clusters with no detected region cluission will be older than approximately MM., The near-infrared clusters with no detected region emission will be older than approximately Myr.
 The regions with uo detected cluster counterpart are most likely vounger than MMyr. and have intermediate mass (~ος10!10 NL.) ioniziug clusters.," The regions with no detected cluster counterpart are most likely younger than Myr, and have intermediate mass $\simeq 5\times10^4-10^5\,{\rm M}_\odot$ ) ionizing clusters."
 If. as observed in obscured Calactie regions. there are significant αλλος of extinction during the first million wears of tie evolution of clusters aud reeious. then we max I© luissine the vouugest star forming reeious. and hence t16 observed fractions of regious aud coiucideuces will I)o lower liits.," If, as observed in obscured Galactic regions, there are significant amounts of extinction during the first million years of the evolution of clusters and regions, then we may be missing the youngest star forming regions, and hence the observed fractions of regions and coincidences will be lower limits."
 An estimate of the age distribution of theobserved clusters can be interred from the relative nuubers of regions. aud uearinfrared star clusters: the higher the fraction of near-intrared clusters. the older the ages of the detected star clusters will be.," An estimate of the age distribution of the clusters can be inferred from the relative numbers of regions, and near-infrared star clusters: the higher the fraction of near-infrared clusters, the older the ages of the detected star clusters will be."
 We find that the ages of the detected star clusters in Arp 299 aud NCC 3256 rauge up to 20)10 MM., We find that the ages of the detected star clusters in Arp 299 and NGC 3256 range up to $20-40$ Myr.
 Older clusters possibly created iu this or previous episodes of star formation are Likely to exist iu these systems but cannot be identified with the preseut detection threshold., Older clusters possibly created in this or previous episodes of star formation are likely to exist in these systems but cannot be identified with the present detection threshold.
 Another possibility to explain the apparent vouth of the clusters in Arp 299 and NGC 3256 would be destruction of clusters., Another possibility to explain the apparent youth of the clusters in Arp 299 and NGC 3256 would be destruction of clusters.
 Iu that case. if the clusters have been created at a coustaut rate for the last ~LOOM vy. then roughly of the clusters are destroved during that time to account for the observed fraction of clusters in these two systems.," In that case, if the clusters have been created at a constant rate for the last $\simeq 100\,$ Myr, then roughly of the clusters are destroyed during that time to account for the observed fraction of clusters in these two systems."
 The data preseuted in this paper does not allow us to distinguish between these two possibilities., The data presented in this paper does not allow us to distinguish between these two possibilities.
 Frou the preseut observations aud modelling it is clear that a huge population of the vouugest clusters (that is the ionizing clusters of the regious with ages of «hb GMyvr) will not be detected from near-infrared continuuni imaging alone. as only some 816% of these regions in our sample of LIRCGs (both isolated aud interacting/mereiueg systems) appear to have near-infrared. cluster counterparts.," From the present observations and modelling it is clear that a large population of the youngest clusters (that is the ionizing clusters of the regions with ages of $<5-6\,$ Myr) will not be detected from near-infrared continuum imaging alone, as only some $8-16\%$ of these regions in our sample of LIRGs (both isolated and interacting/merging systems) appear to have near-infrared cluster counterparts."
 This suggests that studies of the voung star clusters in galaxies performed usiug only near-infrared continu nuaegnmeg may be biased against the vouugest star forming regious. which may account for up to 10€ of the detected population of star chisters aud regions iu LIRGs.," This suggests that studies of the young star clusters in galaxies performed using only near-infrared continuum imaging may be biased against the youngest star forming regions, which may account for up to $40\%$ of the detected population of star clusters and regions in LIRGs."
 It is a pleasure to thank Valentin Ivanov aud Joliu Black for enliehteniug discussions on the paper., It is a pleasure to thank Valentin Ivanov and John Black for enlightening discussions on the paper.
 We are also erateful to an anonvinous referee for helpful comets. which resulted i au improved paper.," We are also grateful to an anonymous referee for helpful comments, which resulted in an improved paper."
 This work has been partially supported by the National Acronautics and Space Administration eraut NAC 5-3012 through the University of Arizona and Contract 960785 through the Jet Propulsion Laboratory., This work has been partially supported by the National Aeronautics and Space Administration grant NAG 5-3042 through the University of Arizona and Contract 960785 through the Jet Propulsion Laboratory.
 This research bas made use of the NASA/TPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Achuinistration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
Falgarone 1996: Abergel et al.,Falgarone 1996; Abergel et al.
 1996: Elmeereen 1997: Heiles 1997: Ileithausen οἱ al., 1996; Elmegreen 1997; Heiles 1997; Heithausen et al.
 1998: Falearone οἱ al., 1998; Falgarone et al.
 1993: Chappell Scalo 2001; Welly Fitzpatrick 2001: Faison Goss 2001: Andrews. Mever. Lauroesch 2001).," 1998; Chappell Scalo 2001; Welty Fitzpatrick 2001; Faison Goss 2001; Andrews, Meyer, Lauroesch 2001)."
 Other galaxies (Stanimirovié et al., Other galaxies (Stanimirović et al.
 1999 for the SAIC: Westplahl et al., 1999 for the SMC; Westpfahl et al.
 1999 [or the M81 group) show the same phenomena in I 1. The best observed reflection nebulae are clumpyv (e.g. Selleren. Werner. Dinerstein 1992: Martini. Selleren. DePov 1999: IXnauth et al.," 1999 for the M81 group) show the same phenomena in H I. The best observed reflection nebulae are clumpy (e.g. Sellgren, Werner, Dinerstein 1992; Martini, Sellgren, DePoy 1999; Knauth et al."
 2001)., 2001).
 We have examined the images in the SIMDAD database of the reflection nebulae with which we are acquainted (about a dozen): all appear to be elumpy., We have examined the images in the SIMBAD database of the reflection nebulae with which we are acquainted (about a dozen); all appear to be clumpy.
 The observed thermal pressure in the ISM as judged from C I (Jenkins Tripp 2001) varies by over an order of magnitude within a sample of stars against which the column densities of C I lines can be determined., The observed thermal pressure in the ISM as judged from C I (Jenkins Tripp 2001) varies by over an order of magnitude within a sample of stars against which the column densities of C I lines can be determined.
 The hierarchical structure of the ISM is a natural consequence of turbulence. which is scale-free.," The hierarchical structure of the ISM is a natural consequence of turbulence, which is scale-free."
 Turbulent models of the ISAT generate hierarchical censitv structure (e.g.. Norman Ferrera 1996).," Turbulent models of the ISM generate hierarchical density structure (e.g., Norman Ferrera 1996)."
 Doissé (1990) determined (he transport of radiation through a (wo-phased cbunipy medium will isotropic scattering., Boissé (1990) determined the transport of radiation through a two-phased clumpy medium with isotropic scattering.
 Witt Gordon (1996) extended the calculations to anisotropic scaltering in a centrally illuminated sphere., Witt Gordon (1996) extended the calculations to anisotropic scattering in a centrally illuminated sphere.
 They investigated the effects of changing various parameters such as (he contrast between (he density contrast between (the phases or the sizes of the clumps relative to the radius of the sphere., They investigated the effects of changing various parameters such as the contrast between the density contrast between the phases or the sizes of the clumps relative to the radius of the sphere.
 Many of the qualitative effects we [ind [from varying parameters for hierarchical models are discussed by them as well., Many of the qualitative effects we find from varying parameters for hierarchical models are discussed by them as well.
 As the volume fraction of the dense phase was increased. (heir models varied from widely separated dense clumps up to continuous structures will small holes.," As the volume fraction of the dense phase was increased, their models varied from widely separated dense clumps up to continuous structures with small holes."
 Since all cells have either of two densities. we will reler to this tvpe of clumping as “two-phase models.," Since all cells have either of two densities, we will refer to this type of clumping as “two-phase models”."
 By contrast. our hierarchical models have an almost continuous distribution of densities.," By contrast, our hierarchical models have an almost continuous distribution of densities."
 Witt Gordon (2000) have investigated. radiative transfer in galaxies by means of a specific Gvo-phased. ISM. one being 100 times as dense as the other and occupying of the volume.," Witt Gordon (2000) have investigated radiative transfer in galaxies by means of a specific two-phased ISM, one being 100 times as dense as the other and occupying of the volume."
 In order to compare wwith hierarchical. we adopt this recipe. with 16 cells along the radius of the sphere (Witt Gordon used 15).," In order to compare with hierarchical, we adopt this recipe, with 16 cells along the radius of the sphere (Witt Gordon used 15)."
 All of our assumed (e.g) = (0.6.00.6). which are values tvpical for optical wavelengths suggested bv observations (see plots in Witt Gordon 2000).," All of our assumed $a,\,g)$ = 0.6), which are values typical for optical wavelengths suggested by observations (see plots in Witt Gordon 2000)."
 The fact that the ISM ishierarchically clumped is important for the propagation of scatlerecl light. since such a structure has relatively open spaces through which the radiation can move rather Ireelv.," The fact that the ISM is clumped is important for the propagation of scattered light, since such a structure has relatively open spaces through which the radiation can move rather freely."
 The main thrust of (his paper is not to trv to put limits on the grain properties of real reflection nebulae. but to show the among hierarchical models as (μον are viewed [rom various angles.," The main thrust of this paper is not to try to put limits on the grain properties of real reflection nebulae, but to show the among hierarchical models as they are viewed from various angles."
 We will not be verv concerned. with the averaged, We will not be very concerned with the averaged
CRS 1915|105. discovered by WATCII iustuuieut ou board in 1992 (7) al located in our ealaxy at an estimated distance of 943 kpe (2). is a low-mass X-ray binary containiug a spiuuiung accreting black hole (€?) of mass about LlL+| AL. and a IN-AI HII eiaut star of mass (0.8+0.5 ML. as the donor (ανν ,"GRS 1915+105, discovered by WATCH instrument on board in 1992 \citep{Castro92} and located in our galaxy at an estimated distance of $9\pm3$ kpc \citep{Chapuis04}, is a low-mass X-ray binary containing a spinning accreting black hole \citep{Zhang97} of mass about $14\pm4$ $_{\odot}$ and a K-M III giant star of mass $0.8\pm0.5$ $_{\odot}$ as the donor \citep{Harlaftis04, Greiner01a}."
The orbital separation aud period of this binary are. respectively. abou los—1 R. aud 33.5 davs (2)..," The orbital separation and period of this binary are, respectively, about $108\pm4$ $_{\odot}$ and 33.5 days \citep{Greiner01b}."
" Serving as à famous niücroquasar, GRS 1915|105 produces superluminal radio jet.. ον"," Serving as a famous microquasar, GRS 1915+105 produces superluminal radio jets \citep{Mirabel94, Fender99}."
 It shows various A-aav elt curves and complex timing phenomena., It shows various X-ray light curves and complex timing phenomena.
 Based on the appearance of light curves aud color-color diagrams. the behaviors of CRS 1915|105 can be classified into 12 classes.," Based on the appearance of light curves and color-color diagrams, the behaviors of GRS 1915+105 can be classified into 12 classes."
 The variability of the source can be further reduced to transitions between three basic states (A. D. aud. C) (2)..," The variability of the source can be further reduced to transitions between three basic states (A, B, and C) \citep{Belloni00}."
 Of these 12. classes. class X (state) is most commonly observed (2)..," Of these 12 classes, class $\chi$ (state) is most commonly observed \citep{Belloni00}."
 It shows characteristics exclusively of state C. the state which is steady in the X-rays and lies in a rather hard part of the color-color diagram.," It shows characteristics exclusively of state C, the state which is steady in the X-rays and lies in a rather hard part of the color-color diagram."
 It is the state when the low-frequency (~0.510 Iz) QPOs (LFQPOs) are most frequently observed (e.g...2).. providing au idea site for y.udyiug LFQPOs.," It is the state when the low-frequency $\sim0.5-10$ Hz) QPOs (LFQPOs) are most frequently observed \citep[e.g.,][]{Muno99}, providing an idea site for studying LFQPOs."
 Mich effort has been made for exploring the origius of the LFQPOs of GRS 1915)105., Much effort has been made for exploring the origins of the LFQPOs of GRS 1915+105.
 It was found that 16 QPO frequency was positively correlated with the 1uxes of ftthe individuallividual componentst audl spectraltral totatotal Hux (e.g.222?27)..," It was found that the QPO frequency was positively correlated with the fluxes of the individual components and spectral total flux \citep[e.g.,][]{Chen97, Markwardt99, Muno99, Trudolyubov99, Reig00, Tomsick01}."
 2 confirmed that OPO frequency was tightly correlated with the source flux. aud. rowever. they also found that for some observations tle OPO frequency was not correlated with the flux.," \citet{Muno01} confirmed that QPO frequency was tightly correlated with the source flux, and, however, they also found that for some observations the QPO frequency was not correlated with the flux."
 It was ound that the QPO amplitude was inversely correlated with the source £ux or QPO frequency (e.g..227).," It was found that the QPO amplitude was inversely correlated with the source flux or QPO frequency \citep[e.g.,][]{Muno99, Reig00, Trudolyubov99}."
 ? and ? reported that as ΟΡΟ frequency increased. he temperature of the mner accretion disk iucreased and the radius of the imucr accretion disk decreased.," \citet{Muno99} and \citet{Rodriguez02a} reported that as QPO frequency increased, the temperature of the inner accretion disk increased and the radius of the inner accretion disk decreased."
efficiencies. {he missing AGN detections below 10 keV. have two main implications which are lightly linked with each other. concerning the covering [actor of the X-ray absorber and the coupling of the gas ancl dust components in the outskirts of the compact active source.,"efficiencies, the missing AGN detections below 10 keV have two main implications which are tightly linked with each other, concerning the covering factor of the X-ray absorber and the coupling of the gas and dust components in the outskirts of the compact active source."
 Also. we discuss our results in the context of the AGN/SB connection. searching for possible relations between the excess of AGN obseuration and the strenght of the simultaneous burst of star formation.," Also, we discuss our results in the context of the AGN/SB connection, searching for possible relations between the excess of AGN obscuration and the strenght of the simultaneous burst of star formation."
 As pointed out earlier. the upper limits that can be placed on the possible ACN reflected emission implv a complete covering of the nuclear region.," As pointed out earlier, the upper limits that can be placed on the possible AGN reflected emission imply a complete covering of the nuclear region."
 In the X-ray obseured. AGN that are classified in (hie optical as (wpe 2s. the direct optical to soft X-ray radiation blocked bv (he ]aimed axisvinmetric custy absorber can be scattered into the line of sight after the interactions wilh material located above (he equatorial plane.," In the X-ray obscured AGN that are classified in the optical as type 2s, the direct optical to soft X-ray radiation blocked by the claimed axisymmetric dusty absorber can be scattered into the line of sight after the interactions with material located above the equatorial plane."
 This process strongly. affects (he V.iape of the emerging spectrum in (he N-ravs. and (he usual elliciency in terms of the μαeceived [Iux is of the order of a few per cent. depending on geometrical effects.," This process strongly affects the shape of the emerging spectrum in the X-rays, and the usual efficiency in terms of the received flux is of the order of a few per cent, depending on geometrical effects."
 Except [ου 1e two AGN detections. in all the other sources of our sample we assess a reflection efficiency lower than ~2xLO7. assuming an intrinsic 2LO keV [lux consistent with the results of the mic-IR analvsis.," Except for the two AGN detections, in all the other sources of our sample we assess a reflection efficiency lower than $\sim 2 \times 10^{-3}$, assuming an intrinsic 2–10 keV flux consistent with the results of the mid-IR analysis."
 Thus. regardless of the large and manifold uncertainties involved in {he single cases. the X-ray follow-up observations of buried ACN presented or reviewed here eenerallv hint at a cocoon-like geometry of the absorber. which is not surprising in extreme and chaotic svstems such as ULIBGs.," Thus, regardless of the large and manifold uncertainties involved in the single cases, the X-ray follow-up observations of buried AGN presented or reviewed here generally hint at a cocoon-like geometry of the absorber, which is not surprising in extreme and chaotic systems such as ULIRGs."
 Indeed. the detection of so prominent CO 5gas absorption in objects like νο 00182— 7112 (Spoon et al.," Indeed, the detection of so prominent CO gas absorption in objects like IRAS $-$ 7112 (Spoon et al."
 2004) and IRAS 00397— 1312 could be related with full covering rather than toroidal obscuration. since (his feature is not usually found among tvpe 2 dSevlerts (e.g. Lutz et al.," 2004) and IRAS $-$ 1312 could be related with full covering rather than toroidal obscuration, since this feature is not usually found among type 2 Seyferts (e.g. Lutz et al."
 2004b)., 2004b).
 As a consequence. the most interesting aspect (o investigate concerns the location of the absorber aud (the properties of its gas and dust content.," As a consequence, the most interesting aspect to investigate concerns the location of the absorber and the properties of its gas and dust content."
 By adopting a typical Galactic extinction curve (e.g. Draine 2003: Nishivama 2008. 2009) and the standard Galactic dust-to-gas ratio (Bohlin. Savage Drake 1973). one can easily obtain the relation between the optical. depth at 6. sam and the A-rayr column density.− that is− Ny)-8 Ll amx10772 >5.," By adopting a typical Galactic extinction curve (e.g. Draine 2003; Nishiyama 2008, 2009) and the standard Galactic dust-to-gas ratio (Bohlin, Savage Drake 1978), one can easily obtain the relation between the optical depth at 6 $\mu$ m and the X-ray column density, that is $N_\rmn{H} \simeq 8$ –11 $\tau_6 \times 10^{22}$ $^{-2}$."
 Then. even questioning (he accuracy of our measure. anv reasonable value of the 7; would not be consistent wilh a Compton-thick environment with dust and gas components," Then, even questioning the accuracy of our measure, any reasonable value of the $\tau_6$ would not be consistent with a Compton-thick environment with dust and gas components"
population properties (such as age and metallicity).,population properties (such as age and metallicity).
 This also allows us to track the physical reasons for the vvariations. which in turn should be reproduced. by. galaxy formation moclels.," This also allows us to track the physical reasons for the variations, which in turn should be reproduced by galaxy formation models."
 Dillerenth from what happens for colour. gradients which are mostly driven from metallicity gradients. from the right column of Fig.," Differently from what happens for colour gradients which are mostly driven from metallicity gradients, from the right column of Fig."
 2. it is evident that the rrelation is driven mainly byage. in both LRGs and EPCs.," \ref{fig:fig2} it is evident that the relation is driven mainly by, in both LTGs and ETGs."
 The trends of wwith stellar mass and velocity. dispersion are shown in he top panels of Fie. 3.., The trends of with stellar mass and velocity dispersion are shown in the top panels of Fig. \ref{fig:fig3}.
 As with the trend of colour and metallicity gracients with mass (see comments in 77 also T|10). we find that L'ECis have a decreasing rend with mass and show OforlogM./M.<9 and dncontrast. (Oallargermasses," As with the trend of colour and metallicity gradients with mass (see comments in \ref{sec:intro} also T+10), we find that LTGs have a decreasing trend with mass and show $>0$ for $\log \mst/\Msun \lsim 9$ and $<0$ at larger masses."
"ETOshaveposilieeNy, at low masses. which decreases down to a minimuntr of Vy,~Ql at logM,fA.~10.3 and then invert he trend. with the most massive EPCs having gradients around. zero or slight positive."," In contrast, ETGs have positive $\gML \sim 0.25$ at low masses, which decreases down to a minimum of $\gML \sim -0.1$ at $\log \mst/\Msun \sim 10.3$ and then invert the trend, with the most massive ETGs having gradients around zero or slight positive."
" When plotted as a function of velocity dispersion. the rrun with συ is generally Hatter. with ΕΝ having no trend and EPCs showing a two-fold trend with a minimum around loge,=2.1 km/s and a mild increase toward the low and high συ sides. always staving consistent with Ve.O."," When plotted as a function of velocity dispersion, the run with $\sigc$ is generally flatter, with LTGs having no trend and ETGs showing a two-fold trend with a minimum around $\log \sigc=2.1$ km/s and a mild increase toward the low and high $\sigc$ sides, always staying consistent with $\lsim 0$."
 Since galaxies are morphologically selected. by using criteria on concentration and Sérsic index. the steeper eradients found. for LPGs when compared. with LLCs indicate a sequence in terms of such structural parameters.," Since galaxies are morphologically selected by using criteria on concentration and $\rm \acute{e}$ rsic index, the steeper gradients found for LTGs when compared with ETGs indicate a sequence in terms of such structural parameters."
 In fact. at fixed mass. galaxies with lower C orn have steeper colour and eeracdients.," In fact, at fixed mass, galaxies with lower $C$ or $n$ have steeper colour and gradients."
 As observed. for the stellar population gradients. (see T10). the central age is responsible for much of the scatter of the trends versus mass ancl velocity. dispersion. which is illastratecd by the median distribution for centrally old. 6CGvr) and voung (« 6GCwvr) galaxies in Fig. 3..," As observed for the stellar population gradients (see T+10), the central age is responsible for much of the scatter of the trends versus mass and velocity dispersion, which is illustrated by the median distribution for centrally old $> 6 \, \rm Gyr$ ) and young $< 6 \, \rm Gyr$ ) galaxies in Fig. \ref{fig:fig3}."
" In particular. voung L'TGs have almost null eracicnts. while the oldest ones have very steep values of Vx,~ 0.310. 0.4."," In particular, young LTGs have almost null gradients, while the oldest ones have very steep values of $\gML \sim -0.3$ to $-0.4$ ."
 EPCs show a similar but olfset. elfect., ETGs show a similar but offset effect.
" Older ΙΤς have slightly negative values of ~0.1. while the vounger ones have steep positive Ve,0.5 at very low masses. decreasing to 0.1 at larger κ..."," Older ETGs have slightly negative values of $\sim -0.1$, while the younger ones have steep positive $\gML \sim 0.5$ at very low masses, decreasing to $\sim 0.1$ at larger masses."
 As a comparison. we also show in Fig.," As a comparison, we also show in Fig."
 3. (top panels) the results when galaxy colours are fitted to à synthetic spectra model which has no age gradient (fixed age=10 vr).," \ref{fig:fig3} (top panels) the results when galaxy colours are fitted to a synthetic spectral model which has no age gradient (fixed $\rm age = 10 \, \rm Gyr$ )."
" In this case. L'ICGis have wwhich are null up to logA,fAl.~ 9.5. ancl decrease a larger masses. while gracients in EEGs show a milder trenc when compared with the reference fit."," In this case, LTGs have which are null up to $\log
\mst/\Msun \sim 9.5$ , and decrease at larger masses, while gradients in ETGs show a milder trend when compared with the reference fit."
 When plotted. versus both. LLCs and. EPCs have similar ~—0.1 for al velocity clispersions., When plotted versus both LTGs and ETGs have similar $\sim -0.1$ for all velocity dispersions.
 Note that these results are consisten with what was found for the older galaxiesin the reference fit., Note that these results are consistent with what was found for the older galaxiesin the reference fit.
We also compare our results with e-band Al/Ls derives,We also compare our results with g-band s derived
"where G and m,, are the gravitational constant and proton mass and pp=0.6.",where G and $m_p$ are the gravitational constant and proton mass and $\mu=0.6$.
 The temperature profile seen in Fig 5. is modelled by a polytropic model with ©? and core radius fixed to the values obtained from the surface brightness fit to the data giving 21.09., The temperature profile seen in Fig \ref{tprofile} is modelled by a polytropic model with $\beta$ and core radius fixed to the values obtained from the surface brightness fit to the data giving $\gamma\simeq 1.09$.
 It is then straightforward to derive an analytical description for the total mass using the expressions for density under the polytropic assumption., It is then straightforward to derive an analytical description for the total mass using the expressions for density under the polytropic assumption.
 However the error in the mass should be obtained indirectly., However the error in the mass should be obtained indirectly.
 We use a Monte-method which is described in Khosroshahi.Jones&Pon-man(2004) and Neumann&Bohringer(1995). for our error estimation., We use a Monte-Carlomethod which is described in \citet{kjp04} and \citet{nb95} for our error estimation.
 In brief. the formula for the mass profile consists of two parts. the variation in gas density (In ndn) and in he temperature (207) and dln1άμα ry.," In brief, the formula for the mass profile consists of two parts, the variation in gas density $d\ln n/d\ln r$ ) and in the temperature $T(r)$ and $d\ln T/d\ln
r$ )."
 For the latter we have generated 1000 physical temperature profiles with a temperature envelope detined by the observed errors in the temperature and inear interpolation., For the latter we have generated 1000 physical temperature profiles with a temperature envelope defined by the observed errors in the temperature and linear interpolation.
 A physical temperature profile is one which guarantees a monotonically increasing mass with radius., A physical temperature profile is one which guarantees a monotonically increasing mass with radius.
 In this simulation the density profile parameters were fixed at the values rom the «?-model fit., In this simulation the density profile parameters were fixed at the values from the $\beta$ -model fit.
 In order to estimate the contribution of he density variation (dInnz ry to the mass error. the surface brightness profile at the location of each data point was fitted with dInn/cdlnr as the free parameter instead of «7.," In order to estimate the contribution of the density variation $d\ln n/d\ln r$ ) to the mass error, the surface brightness profile at the location of each data point was fitted with $d\ln n/d\ln r$ as the free parameter instead of $\beta$."
 The error was estimated at the confidence level using the error matrix orovided by the fitting program., The error was estimated at the confidence level using the error matrix provided by the fitting program.
 We checked the results by monitoring the variation., We checked the results by monitoring the variation.
 The total error in the mass at each data point was then derived by adding the two uncertainties quadratically., The total error in the mass at each data point was then derived by adding the two uncertainties quadratically.
 Motivated by the results of numerous cosmological N-body simulations and to enable us to make a direct comparison with mass profiles from other studies. we attempt to fit à NFW protile (Navarroetal.1995) to the total gravitational mass density.," Motivated by the results of numerous cosmological N-body simulations and to enable us to make a direct comparison with mass profiles from other studies, we attempt to fit a NFW profile \citep{nfw95}
 to the total gravitational mass density."
 Integrated the mass profile for a spherical mass distribution leads to. ΠΕ. where ο)ης) Is the density at r..," Integrated the mass profile for a spherical mass distribution leads to, _m(r_s)r_s^2[r_s where $\rho_m(r_s)$ is the density at $r_s$."
 The mass concentration parameter can then be defined as σου=rs., The mass concentration parameter can then be defined as $c_{200}=r_{200}/r_s$.
 By definition. rogo Is the radius within which the mean gravitational mass density is 200 times the critical density. p.(2).," By definition, $r_{200}$ is the radius within which the mean gravitational mass density is 200 times the critical density, $\rho_c(z)$."
" For our assumed cosmology (0,,,20.3. A=0.7). the critical density is p.(z) )."," For our assumed cosmology $\Omega_m$ =0.3, $\Lambda$ =0.7), the critical density is $\rho_c(z)=\frac{3H_0^2}{8\pi G}(1+z)^2(1+z\Omega_m)$ ."
 We tit the NFW mass profile (equation 5)) to the five observed data points with errors obtained from MC simulation., We fit the NFW mass profile (equation \ref{nfw}) ) to the five observed data points with errors obtained from MC simulation.
 The best fit profile has rs.=97 kpe., The best fit profile has $r_s=97$ kpc.
 reoo is then calculated by extrapolating the profile to 200p..., $r_{200}$ is then calculated by extrapolating the profile to $200\rho_c$.
 This gives ους=1.22+0.06 Mpc.," This gives $r_{200}=
1.22\pm0.06$ Mpc."
 The extrapolated total mass. using the best fit NFW profile. at requ is Abou=3141.0).1011 M...," The extrapolated total mass, using the best fit NFW profile, at $r_{200}$ is $M_{200}= 3.1(\pm1.0)\times 10^{14}$ $_{\odot}$ ."
 To obtain the concentration of the dark matter we take into account the contribution of the stellar mass of the central dominant galaxy. for which we have the R-band luminosity and assume a stellar M/L;;~ 5 M./L. from the recent fundamental plane studies (Cappellarietal.2005).," To obtain the concentration of the dark matter we take into account the contribution of the stellar mass of the central dominant galaxy, for which we have the R-band luminosity and assume a stellar $_R\sim$ 5 $_\odot$ $_\odot$ from the recent fundamental plane studies \citep{capp05}."
".. This gives r,=LOS kpe and therefore a concentration parameter c»).=11.2c4.5. which is high compared to the values of coo; for poor clusters studied by Pratt&Arnaud(2005)."," This gives $r_s=108$ kpc and therefore a concentration parameter $c_{200}= 11.2\pm4.5$, which is high compared to the values of $c_{200}$ for poor clusters studied by \citet{pratt05}."
. Implications of this high mass concentration is discussed in section 35., Implications of this high mass concentration is discussed in section 5.
 The integrated gas mass can be derived by integrating the gas density., The integrated gas mass can be derived by integrating the gas density.
 Fig 8 shows the gas mass profile., Fig \ref{massprof} shows the gas mass profile.
 The totalgas mass within raga Is ~d10*M--., The totalgas mass within $r_{200}$ is $\sim 4\times 10^{13}$$\odot$.
 With the extrapolated total gravitational ↘⊓↰∙⋅ same radius 3.13107 Me. the gas fraction is calculated to be ~ 0.13.," With the extrapolated total gravitational mass within the same radius $\sim 3.1\times 10^{14}$ $\odot$ , the gas fraction is calculated to be $\sim 0.13$ ."
from the 5000 Gaussian simulations described above.,from the 5000 Gaussian simulations described above.
" Since the £C statistic of a map is a maximumvalue (4/07=mar(iC, )). the confidence regions are one-sided."," Since the $HC$ statistic of a map is a maximumvalue $HC=max_i(HC_{n,i})$ ), the confidence regions are one-sided."
 As one can see. there is no detection of deviations from Caussianity in the Jianchi corrected WALADP data (factor of 1.0 in the ligure).," As one can see, there is no detection of deviations from Gaussianity in the Bianchi corrected WMAP data (factor of $1.0$ in the figure)."
 This corroborates the results obtained by Jalle et al., This corroborates the results obtained by Jaffe et al.
 2005b., 2005b.
 Subtraction of the Bianchi template corrects most of the anomalies that have been observe in the first vear WALAP data., Subtraction of the Bianchi template corrects most of the anomalies that have been observed in the first year WMAP data.
 The maps of €! values for the WALAP and the Bianchi corrected \WALAD data at a wavelet. scale of 5 degrees are »esented in Figures 2 and 3., The maps of $HC$ values for the WMAP and the Bianchi corrected WMAP data at a wavelet scale of 5 degrees are presented in Figures 2 and 3.
 As one can immediately see the detected non-CGaussianity in the WNLADP data is dominated o» the values of a ring of pixels at the spot. centered. at Pr200° b~5T (note that the rine structure is likely to ne caused by the convolution with the wavelet).," As one can immediately see the detected non-Gaussianity in the WMAP data is dominated by the values of a ring of pixels at the spot centered at $l\sim 209^{\circ}$, $b\sim -57^{\circ}$ (note that the ring structure is likely to be caused by the convolution with the wavelet)."
 Subtraction of the Bianchi template drops the values of the {1Ο statistic at that spot. below the 68% οἱ. (, Subtraction of the Bianchi template drops the values of the $HC$ statistic at that spot below the $68\%$ c.l. (
see Figure 1).,see Figure 1).
 However. he spot still appears as one of the most prominent regions.," However, the spot still appears as one of the most prominent regions."
 Even if the relevance of this feature is debatable we decided ο see whether cüllerent. normalizations of the subtracted emplate can inlluence its amplitude., Even if the relevance of this feature is debatable we decided to see whether different normalizations of the subtracted template can influence its amplitude.
 X factor of 1.2 reduces the amplitude of the spot making it comparable with the surrounding values as can be seen in Figure 4., A factor of $1.2$ reduces the amplitude of the spot making it comparable with the surrounding values as can be seen in Figure 4.
 We have also studied whether this normalization factor alfects the other anomalies observed in the WALAD first vear data., We have also studied whether this normalization factor affects the other anomalies observed in the WMAP first year data.
 The etct on the alignment between the quadrupole ancl the octopole and the planarity as defined in Oliveira-Costa et al., The effect on the alignment between the quadrupole and the octopole and the planarity as defined in Oliveira-Costa et al.
 2004 is oesented in Table 1., 2004 is presented in Table 1.
 These large scale anomalies are not very much alfected by a change in the normalization factor of the Sianchi template., These large scale anomalies are not very much affected by a change in the normalization factor of the Bianchi template.
 The ratio of power between. two jemispheres for a certain range of/ was calculated. as in Eviksen et al., The ratio of power between two hemispheres for a certain range of $l$ was calculated as in Eriksen et al.
 2004a., 2004a.
 The probability of having a maximum »ower asvmmetry ratio larger than that of a given map (CLher the WALAD or the Bianchi corrected WALAP maps) is ga10wn in Table 2., The probability of having a maximum power asymmetry ratio larger than that of a given map (either the WMAP or the Bianchi corrected WMAP maps) is shown in Table 2.
 This probability is slightly alfected by he normalization factor., This probability is slightly affected by the normalization factor.
 There is no much of a cillerence tween the results lor a factor of 1: and a factor of 1.2., There is no much of a difference between the results for a factor of 1 and a factor of $1.2$.
 However. above this. an increase in the normalization factor starts worsening the results.," However, above this, an increase in the normalization factor starts worsening the results."
 This agrees with the trend observed for the fC statistic as can be seen in Figure 1., This agrees with the trend observed for the $HC$ statistic as can be seen in Figure 1.
V533 Lor. where the signal is no longer visible) highly stable mocdulations at 63.63. 71.07 and 351.34 s respectively (see. e.g.. the review in Warner 1995).,"V533 Her, where the signal is no longer visible) highly stable modulations at 63.63, 71.07 and 351.34 s respectively (see, e.g., the review in Warner 1995)."
 If. the 258 s. periodicity proves to be persistent and coherent then W3732 Set will join this group of rapid rotators., If the 258 s periodicity proves to be persistent and coherent then V373 Sct will join this group of rapid rotators.
" The most important result from our survey is that we have detected. eclipses in three old novae: BY Cir with Pas = 6.176 h. DD Cir with {δν = 2.339 h (placing it squarely in the ""orbital period gap""). and CP Cru with {δν = 22.7 h. which is the third longest period known for a classical nova."," The most important result from our survey is that we have detected eclipses in three old novae: BY Cir with $P_{orb}$ = 6.76 h, DD Cir with $P_{orb}$ = 2.339 h (placing it squarely in the “orbital period gap”), and CP Cru with $P_{orb}$ = 22.7 h, which is the third longest period known for a classical nova."
 We also find a fourth orbital period. from the presence of a moculation at 3.686 h in the old nova V992 Sco.," We also find a fourth orbital period, from the presence of a modulation at 3.686 h in the old nova V992 Sco."
 In addition. DD Cir has a ~ 670 s low amplitude photometric modulation and the old nova V373 Set has à 258.3 s modulation: these are signatures of magnetic primaries with these as their spin period. or orbital side bands.," In addition, DD Cir has a $\sim$ 670 s low amplitude photometric modulation and the old nova V373 Sct has a 258.3 s modulation; these are signatures of magnetic primaries with these as their spin period, or orbital side bands."
 There is some evidence that DD Cir has a 0.3 mag rellection effect that is partly obscured at phase 0.5. producing a secondary. eclipse caused. by. the optically thick accretion disc.," There is some evidence that DD Cir has a 0.3 mag reflection effect that is partly obscured at phase 0.5, producing a secondary eclipse caused by the optically thick accretion disc."
 The oleh novae V8S42. Cen and V655 CrA are highly active photometrically but. show no clear periodic modulation. indicating that they are. low inclination systenis.," The old novae V842 Cen and V655 CrA are highly active photometrically but show no clear periodic modulation, indicating that they are low inclination systems."
" We have also measured. 2,5 1.509 h for the SU UMa star PY Cry from an orbital modulation seen curing quiescence.", We have also measured $P_{orb}$ = 1.509 h for the SU UMa star TV Crv from an orbital modulation seen during quiescence.
 Only the superhump. period. observed: during superoutburst. was previously known for this star.," Only the superhump period, observed during superoutburst, was previously known for this star."
The distributions of cluster early-type and late-type galaxies (ETGs and LTGs hereafter) have long been known to be different (?:: see also ? for a review).,The distributions of cluster early-type and late-type galaxies (ETGs and LTGs hereafter) have long been known to be different \citealt{Dressler80}; see also \citealt{Biviano00} for a review).
" Most striking. and hence the first to have been discovered. ts the difference in the distributions. ETGs living in higher density regions than LTGs. the so-called ""morphology-density relation” (MDR hereafter)."," Most striking, and hence the first to have been discovered, is the difference in the distributions, ETGs living in higher density regions than LTGs, the so-called “morphology-density relation” (MDR hereafter)."
 In relaxed clusters density is an almost monotonic decreasing function of clustercentrie radius. hence the relation is often described as a morphology-radius relation (e.g.?)..," In relaxed clusters density is an almost monotonic decreasing function of clustercentric radius, hence the relation is often described as a morphology-radius relation \citep[e.g.][]{WGJ93}."
 The MDR has been found to exist in rich clusters up to redshift z~| (??).. but the global fractions of ETGs and LTGs evolve with z.," The MDR has been found to exist in rich clusters up to redshift $z \sim 1$ \citep{Postman+05,Smith+05}, but the global fractions of ETGs and LTGs evolve with $z$."
 The fraction of ETGs in clusters decreases with z quite rapidly cp to z~0.5 (22?2).. then seems to flatten out to z~1 (?)..," The fraction of ETGs in clusters decreases with $z$ quite rapidly up to $z
\sim 0.5$ \citep{Dressler+97,Fasano+00,Postman+05,Smith+05}, then seems to flatten out to $z \sim 1$ \citep{Desai+07}."
 Another aspect of the segregation. of different. galaxy populations in. clusters is the difference in the velocity Cistributions of ETGs and LTGs., Another aspect of the segregation of different galaxy populations in clusters is the difference in the velocity distributions of ETGs and LTGs.
 In clusters. the velocity Cuistribution of LTGs is broader than that of ETGs. te. LTGs have a larger line-of-sight (los hereafter) velocity dispersion hereafter) than ETGs (????)..," In clusters, the velocity distribution of LTGs is broader than that of ETGs, i.e. LTGs have a larger line-of-sight (los hereafter) velocity dispersion hereafter) than ETGs \citep{Tammann72,MD77,Sodre+89,Biviano+92}."
" The difference is not only n theσως. but also in the oy),,.--profiles. ""hotter"" and steeper (at least in the inner regions) for the LTGs (???).."," The difference is not only in the, but also in the -profiles, 'hotter' and steeper (at least in the inner regions) for the LTGs \citep{Carlberg+97-equil,Biviano+97,ABM98}."
 While the z-evolutionof the MDR has been extensively studied and deseribed. less is known about the z-evolution of the morphological segregation in velocity space. because of the difficulty of obtaining large spectroscopic samples of cluster galaxies at high-z.," While the $z$ -evolutionof the MDR has been extensively studied and described, less is known about the $z$ -evolution of the morphological segregation in velocity space, because of the difficulty of obtaining large spectroscopic samples of cluster galaxies at $z$."
" As remarked by ?.. the oy,,--profiles of both ΕΤΟΣ and LTGs in z0.07 clusters (fromENACS.theESONearbyAbellClusterSurvey.??) are remarkably similar to the profiles of red. and respectively. blue galaxies in z~0.3 clusters (fromCNOC.theCanadianNetworkforObservationalCosmologysurvey. ?).. but little is known about the evolution at still higher z."," As remarked by \citet{Biviano06}, the -profiles of both ETGs and LTGs in $z \sim 0.07$ clusters \citep[from ENACS, the ESO Nearby Abell
Cluster Survey,][]{Katgert+96,Biviano+97} are remarkably similar to the profiles of red, and respectively, blue galaxies in $z \sim 0.3$ clusters \citep[from CNOC, the Canadian Network for Observational
Cosmology survey,][]{Carlberg+97-equil}, but little is known about the evolution at still higher $z$."
 Instead of looking separately at the spatial and. velocity segregation of cluster galaxies. it is possible to use the joint information coming from their 2-dimensional projected phasespace distribution in the los velocity vs. clustercentric radius diagram.," Instead of looking separately at the spatial and velocity segregation of cluster galaxies, it is possible to use the joint information coming from their 2-dimensional projected phase-space distribution in the los velocity vs. clustercentric radius diagram."
 The projected phase-space distribution of a given class of cluster galaxies is the observable that enters the Jeans equation for the equilibrium of a galaxy system??).. hence separating different cluster galaxy populations on the basis of their projected phase-space distributions ts a way of identifying different. independent tracers of the cluster gravitational potential.," The projected phase-space distribution of a given class of cluster galaxies is the observable that enters the Jeans equation for the equilibrium of a galaxy system, hence separating different cluster galaxy populations on the basis of their projected phase-space distributions is a way of identifying different, independent tracers of the cluster gravitational potential."
 ? have shown that cluster ETGs and LTGs have significantly different projected phase-space distributions: ? have then used cluster ETGs to determine the average mass profile of massive clusters., \citet{Biviano+02} have shown that cluster ETGs and LTGs have significantly different projected phase-space distributions; \citet{KBM04} have then used cluster ETGs to determine the average mass profile of massive clusters.
 In order to do that. they first had to constrain the velocity anisotropy profile of ETGs.," In order to do that, they first had to constrain the velocity anisotropy profile of ETGs."
 They did so by comparing the velocity distribution of ETGs with distribution function models from ?.. and found that ETGs move on nearly isotropic orbits.," They did so by comparing the velocity distribution of ETGs with distribution function models from \citet{vanderMarel+00}, and found that ETGs move on nearly isotropic orbits."
 On the other hand. LTGs were found to move on slightly radially anisotropic orbits. with an increasing radial anisotropy at larger clustercentric radit (?)..," On the other hand, LTGs were found to move on slightly radially anisotropic orbits, with an increasing radial anisotropy at larger clustercentric radii \citep{BK04}."
 At higher redshift. in the CNOC clusters. red galaxies were found to move along nearly isotropic orbits (?).. similarly to ETGs in nearby ENACS clusters. and blue galaxies were found to be in equilibrium within the cluster potential. despite having a different projected phase-space distribution from that of red galaxies (?)..," At higher redshift, in the CNOC clusters, red galaxies were found to move along nearly isotropic orbits \citep{vanderMarel+00}, similarly to ETGs in nearby ENACS clusters, and blue galaxies were found to be in equilibrium within the cluster potential, despite having a different projected phase-space distribution from that of red galaxies \citep{Carlberg+97-equil}."
 Although a solution for the velocity anisotropy of the blue CNOC cluster galaxies has not been derived. the similarity of their projected phase-space distribution to that of LTGs in ENACS suggests that they similarly move on slightly radial orbits (22). ," Although a solution for the velocity anisotropy of the blue CNOC cluster galaxies has not been derived, the similarity of their projected phase-space distribution to that of LTGs in ENACS suggests that they similarly move on slightly radial orbits \citep{Biviano06,Biviano08}. ."
Hence. while there is significant evolution in the relative fractions of early-type (red) and late-type (blue) cluster galaxies. their orbital anisotropies do not seem to evolve over the same 0—0.3 redshift range.," Hence, while there is significant evolution in the relative fractions of early-type (red) and late-type (blue) cluster galaxies, their orbital anisotropies do not seem to evolve over the same 0–0.3 redshift range."
 ? do claim significant orbital evolution for the cluster galaxy population. on the basis of the analyses of three low-z and two z~0.2—0.3 clusters.," \citet{Benatov+06} do claim significant orbital evolution for the cluster galaxy population, on the basis of the analyses of three $z$ and two $z \sim 0.2-0.3$ clusters."
 Specifically. they find galaxies in the higher-z clusters to have more radially anisotropic orbits than galaxies in the lower-z clusters.," Specifically, they find galaxies in the $z$ clusters to have more radially anisotropic orbits than galaxies in the $z$ clusters."
" Since they consider all cluster galaxies together in their analysis. and since LTGs are known to be characterized by radial orbital anisotropy. 2""s result could be explained by the increasing fraction of LTGs with =. without the need for any evolution of the orbits of either ETGs or LTGs."," Since they consider all cluster galaxies together in their analysis, and since LTGs are known to be characterized by radial orbital anisotropy, \citet{Benatov+06}' 's result could be explained by the increasing fraction of LTGs with $z$, without the need for any evolution of the orbits of either ETGs or LTGs."
 The lack of a significant evolution in the orbital anisotropy of. separately. early-type (red) and late-type (blue) cluster galaxies from z~0.05 to z~ 0.3. coupled with the significant evolution in the relative fractions of these two populations over the same redshift range is intriguing.," The lack of a significant evolution in the orbital anisotropy of, separately, early-type (red) and late-type (blue) cluster galaxies from $z \sim 0.05$ to $z \sim 0.3$ , coupled with the significant evolution in the relative fractions of these two populations over the same redshift range is intriguing."
 It suggests that the change of class fromblue. LTG. to red. ETG. goes together," It suggests that the change of class fromblue, LTG, to red, ETG, goes together"
essentially bypasses this complication.,essentially bypasses this complication.
 The issue could also be circumvented by binnine the visibilities rradiallv in (α.0) distance). to increase the signal-to-noise. ancl then performing a A? fit to the binned values.," The issue could also be circumvented by binning the visibilities radially in $(u,v)$ distance), to increase the signal-to-noise, and then performing a $\chi^2$ fit to the binned values."
 We avoid this solution because the visibility profiles for the various models are falling steeply al the baselines of interest. ancl binning (he data inevitably introduces a bias in the fit due to contracting the range of baseline vectors to a sinele length at the center of the bin.," We avoid this solution because the visibility profiles for the various models are falling steeply at the baselines of interest, and binning the data inevitably introduces a bias in the fit due to contracting the range of baseline vectors to a single length at the center of the bin."
 We use binning only as a graphical tool in order to demonstrate the fit quality of a partüeular mocel., We use binning only as a graphical tool in order to demonstrate the fit quality of a particular model.
 The PBI dataset covers baselines [rom 12 to 80 kA. with an overall rms noise of 0.9 mJv.," The PdBI dataset covers baselines from 12 to 80 $\lambda$, with an overall rms noise of 0.9 mJy."
 The visibility data are consistent with the noise., The visibility data are consistent with the noise.
 The long baselines covered by the PdBI data therefore provide an important constraint on (he maximum compact component of the [lux from the L6942 core. F<2.7 mJy (3 60).," The long baselines covered by the PdBI data therefore provide an important constraint on the maximum compact component of the flux from the L694–2 core, $F_c < 2.7$ mJy (3 $\sigma$ )."
" For optically thin dust. this limiting flux corresponds to an implied (disk) mass limit of AJX5x10.1 M.(60 N/T) for an opacity Hpaaun—0.02 en? 1H,"," For optically thin dust, this limiting flux corresponds to an implied (disk) mass limit of $M \lesssim 5 \times 10^{-4}$ $_{\odot} \, (60$ ${}/T_{\mathrm{disk}})$ for an opacity $\kappa_{1.3~\mathrm{mm}} = 0.02$ $^2$ $^{-1}$."
 This limit is roughly an order of magnitude lower than the AL. range of disk masses observed around T-Tauri stars by Beekwith et ((1990). and further demonstrates the “starless” nature of the L6942 core.," This limit is roughly an order of magnitude lower than the 0.002--0.3 $_{\odot}$ range of disk masses observed around T-Tauri stars by Beckwith et (1990), and further demonstrates the “starless” nature of the L694–2 core."
 The [αν limit allows an estimate of the maximum bolometric haminosity for any compact component embedded in the LGQ42 core., The flux limit allows an estimate of the maximum bolometric luminosity for any compact component embedded in the L694–2 core.
" For a simple estimate. we assume a dust opacity spectral index of unity. aud model the compact component as a eravbocly of the form: where D,(I4)) denotes the Planck function at [requencey v for a mean dust temperature Causa). Tr is the dust optical depth. and O5 the solid angle subtended by the source BBeckwith et 11990)."," For a simple estimate, we assume a dust opacity spectral index of unity, and model the compact component as a graybody of the form: where $B_{\nu}(\langle T_{{\rm dust}} \rangle)$ denotes the Planck function at frequency $\nu$ for a mean dust temperature $\langle T_{{\rm dust}} \rangle$, $\tau_{\nu}$ is the dust optical depth, and $\Omega_{S}$ the solid angle subtended by the source Beckwith et 1990)."
 Since the envelope is entirely optically thin at 1.3 min. the gravbodxy must have a flux that is <2.7 mJy al this wavelength.," Since the envelope is entirely optically thin at 1.3 mm, the graybody must have a flux that is $\leq 2.7$ mJy at this wavelength."
 For a given mean dust temperature. (his constraint fixes the mass of the compact component.," For a given mean dust temperature, this constraint fixes the mass of the compact component."
 The only remaining parameter is (he radius A of the component. which essentially identilies the wavelength at which the enission becomes optically thick.," The only remaining parameter is the radius $R$ of the component, which essentially identifies the wavelength at which the emission becomes optically thick."
 For a compact component (to remain undetected in our PdDBlI observations implies that its bolometric luminosity is:, For a compact component to remain undetected in our PdBI observations implies that its bolometric luminosity is:
 (Sperecle£al2007) (Reissetαἱ.1998.2001. (Dradacctal2008).. 2005).. af2003).. c£αἱ.2006.2008).," \citep{WMAP} \citep{Reiss, Per1999} \citep{BulletCluster} \citep{OtherCluster}. \citep{Kamin07, Ber2005}. \citep{PQ1977,Torres00}. \citep{Jeans1,Jeans3,Jeans2}. \citep{K1985,edwm, bala2006, bala2008}."
. like boson stars., like boson stars.
 Structure formation on scales sinualler than the spreading of an individual boson (the Compton wavelength of one particle) is forbidden bv quautiuu mechanics (Ine£al2000:Sahui&WangLee&Lim 2008).," Structure formation on scales smaller than the spreading of an individual boson (the Compton wavelength of one particle) is forbidden by quantum mechanics \citep{Hu,Varun,Lee08}."
". Talos formed from ultralight scalars witli Compton wavelength of galactic scales thus do not lead to over-abuudance of dwarf galaxies uulike cold dark matter siauulations with heavier bosons (Navarroοἱal,1996:Saluccietaf 2003).."," Halos formed from ultralight scalars with Compton wavelength of galactic scales thus do not lead to over-abundance of dwarf galaxies unlike cold dark matter simulations with heavier bosons \citep{overabundance1, Matos01,miguel02,overabundance2}."
 Scalar field halos have been fit to rotation curves of spiral galaxies (Schuuck&Liddle1997:Cowimanbevefαἱ.2001:Doebliner&Tarko 2007).," Scalar field halos have been fit to rotation curves of spiral galaxies \citep{Franz1,GL,Matos00,arby1,Boehmer07}."
. By using the mass as a free parsaineter to fit rotation curves Arbevctel.(2001) lave obtained in2(0.1.1.6)«10.22 eV for nuon-nuteractius ultra-light bosous., By using the mass as a free parameter to fit rotation curves \cite{arby1} have obtained $m = (0.4-1.6) \times 10^{-23}$ eV for non-interacting ultra-light bosons.
 Iu a followup paper. Arbevcfαἱ.(2002). performed a non-thermal analysis of ultralight bosonic halos.," In a followup paper, \cite{arby} performed a non-thermal analysis of ultralight bosonic halos."
 Ureua-Lopez(2009) pointed out that scalars field particles can Bose-coudense at finite temperatures resturecting previous work on relativistic Bosc-Einstcin condensation by Parker&Zhaue(1991.1993)..," \cite{urena-lopez} pointed out that scalars field particles can Bose-condense at finite temperatures resurrecting previous work on relativistic Bose-Einstein condensation by \cite{ParkerZhang1, ParkerZhang2}."
 A condensate is considered relativistic when the temperature of the condensate is significantly larger thui the mass of one boson., A condensate is considered relativistic when the temperature of the condensate is significantly larger than the mass of one boson.
 Parker&Zhaug(1993) discuss inflationary expansion driven by a relativistic Dosc-Einstein condensate that then evolves iuto a radiation doninated uuiverse., \cite{ParkerZhang2} discuss inflationary expansion driven by a relativistic Bose-Einstein condensate that then evolves into a radiation dominated universe.
 Tn this paper we perform a thermal analysis of the postandationary cosmological behavior of scalar field dark iuatter formec from ultra-light bosous., In this paper we perform a thermal analysis of the post-inflationary cosmological behavior of scalar field dark matter formed from ultra-light bosons.
 We use the quantum field theory formalisin of Parker&Fulling(1971). and extend the analysis that [hi(1982) used for a description of finite temperature effects in the carly universe., We use the quantum field theory formalism of \cite{PF} and extend the analysis that \cite{Hu1980s} used for a description of finite temperature effects in the early universe.
 The bosons are described by a complex scalar field to provide a conserved charge., The bosons are described by a complex scalar field to provide a conserved charge.
 We asstune the scalar particles decouple after inflation in the carly universe. after which the field has a simple," We assume the scalar particles decouple after inflation in the early universe, after which the field has a simple"
Soon alter its ciscovery (Zwitterοἱal.2004)... six low-resolution optical spectra of SN 2004et were taken at the David Dunlap Observatory. Canada. with the Casscerain spectrograph mounted. on the 74. telescope.,"Soon after its discovery \citep{discov04et}, six low-resolution optical spectra of SN 2004et were taken at the David Dunlap Observatory, Canada, with the Cassegrain spectrograph mounted on the 74"" telescope."
 Phe spectra covered the 4000 - SOOO rregime with a resolution of SOO at 6000 A(seethedetailsonthedatareductioninVinkóetal. 2006).., The spectra covered the 4000 - 8000 regime with a resolution of $\sim 800$ at 6000 \citep[see the details on the data reduction in][]{vinko04dj}.
 Due to the fixed North-South slit direction. the spectra could not be taken at the parallactie angle. thus. the slope of the continuum in the blue is allected by differential refraction.," Due to the fixed North-South slit direction, the spectra could not be taken at the parallactic angle, thus, the slope of the continuum in the blue is affected by differential refraction."
 Table AL contains the journal of these observations. and the spectra are plotted in Fig. Al.," Table \ref{04ettable} contains the journal of these observations, and the spectra are plotted in Fig. \ref{04etsp}."
" In Table Bl owe present the velocities obtained with (0,1. sce 822)). along with those measured from the absorption minima of and A5169 (6g« and ep). as well as those obtained with cross-correlation method. using the observed spectra of SN 1999em as templates (603) and the CAIPGIEN mocels (6)."," In Table \ref{vel} we present the velocities obtained with $v_{model}$, see \ref{sec_synow}) ), along with those measured from the absorption minima of $\beta$ and $\lambda5169$ $v_{H\beta}$ and $v_{Fe}$ ), as well as those obtained with cross-correlation method using the observed spectra of SN 1999em as templates $v_{cc\#1}$ ) and the CMFGEN models $v_{cc\#2}$ )."
" In Table Cl owe present the parameters of the hest- SYNOW mocels. including: 7,-, ol the main atomsτοις. the power-law exponent of the optical depth function (n). the photospheric temperature (75,) and photospheric velocity (0,5,) (sce 877. for details)."," In Table \ref{synowtable} we present the parameters of the best-fitting models, including: $\tau_{ref}$ of the main atoms/ions, the power-law exponent of the optical depth function $n$ ), the photospheric temperature $T_{bb}$ ) and photospheric velocity $v_{phot}$ ) (see \ref{sec_synow} for details)."
The preceding analysis neglects the effects of acceleration due to orbital motion.,The preceding analysis neglects the effects of acceleration due to orbital motion.
 The differential Doppler shift induced over the course of the observation can result in the smearing of power across several bins of the fluctuation spectrum., The differential Doppler shift induced over the course of the observation can result in the smearing of power across several bins of the fluctuation spectrum.
 Since the shift moves frequency components a by multiplicative factor (at/c). the loss of sensitivity is greatest at higher harmonies.," Since the shift moves frequency components a by multiplicative factor $at/c$ ), the loss of sensitivity is greatest at higher harmonics."
 The acceleration spacing of the FA search is ~0.3 m s. implying a worst-case differential shift factor of 0.15 m 57«859 s/cc4.30«1077 for a pulsar with an acceleration mid- between two trial values.," The acceleration spacing of the FA search is $\sim$ 0.3 m $^{-2}$, implying a worst-case differential shift factor of $0.15$ m $^{-2}\times 859$ $/c \simeq
4.30\times 10^{-7}$ for a pulsar with an acceleration mid-way between two trial values."
 Since each bin of the fluctuation spectrum represents 1/859 s ~1.16«10? Hz. the 2500-Hz fundamental of a putative 0.4-ms sub-millisecond pulsar would experience at most 1 bin of acceleration smearing. with higher harmonies experiencing proportionately more smearing but representing a smaller fraction of the total power of the pulsar.," Since each bin of the fluctuation spectrum represents $1/859$ s $\simeq 1.16\times 10^{-3}$ Hz, the 2500-Hz fundamental of a putative $0.4$ -ms sub-millisecond pulsar would experience at most $\sim$ 1 bin of acceleration smearing, with higher harmonics experiencing proportionately more smearing but representing a smaller fraction of the total power of the pulsar."
 To examine the true sensitivity loss we processed real data from an FA search. to which we had added simulated signals from non-accelerated pulsars of various pulse widths and periods.," To examine the true sensitivity loss we processed real data from an FA search, to which we had added simulated signals from non-accelerated pulsars of various pulse widths and periods."
 We used a trial acceleration spacing of 0.05 m s in the range |a|<6 m s to examine the loss of signal to noise ratio when trial accelerations do not match the true acceleration of the pulse (in this case zero)., We used a trial acceleration spacing of 0.05 m $^{-2}$ in the range $|a|<6$ m $^{-2}$ to examine the loss of signal to noise ratio when trial accelerations do not match the true acceleration of the pulse (in this case zero).
 The results are shown in Figure 5.., The results are shown in Figure \ref{fig:accsn}.
 Even in the worst case. a 0.4 ms signal with a pulse width of 5 per cent FWHM (that is. at the extreme lower end of observed MSP pulse widths). some 80 per cent of sensitivity was retained at the outer edges of the range of acceleration offsets experienced m the FA search.," Even in the worst case, a 0.4 ms signal with a pulse width of 5 per cent FWHM (that is, at the extreme lower end of observed MSP pulse widths), some 80 per cent of sensitivity was retained at the outer edges of the range of acceleration offsets experienced in the FA search."
 It is clear that the impact of acceleration on sensitivity of the FA search was small for P20.4 ms.," It is clear that the impact of acceleration on sensitivity of the FA search was small for $P \gtrsim
0.4$ ms."
 In addition. due to scintillation it is expected that à real pulsar would experience less signal-to-noise loss than is indicated by these results.," In addition, due to scintillation it is expected that a real pulsar would experience less signal-to-noise loss than is indicated by these results."
 For the SA search a more modest acceleration spacing of 0.5 m s was chosen due to the fact that the long sample interval employed itself severely limits sensitivity to sub-millisecond pulsars., For the SA search a more modest acceleration spacing of 0.5 m $^{-2}$ was chosen due to the fact that the long sample interval employed itself severely limits sensitivity to sub-millisecond pulsars.
 In this case. any pulsar or harmonic with a period greater than 0.6 ms would experience less than one bin of acceleration-induced spectral smearing.," In this case, any pulsar or harmonic with a period greater than 0.6 ms would experience less than one bin of acceleration-induced spectral smearing."
 Of course. all of the above only applies to pulsars with accelerations that remain essentially constant throughout the observation and lie within the range of |a|<30 m s searched.," Of course, all of the above only applies to pulsars with accelerations that remain essentially constant throughout the observation and lie within the range of $|a|<30$ m $^{-2}$ searched."
 The former requirement is likely to be satisfied for the vast majority of systems: the loss of signal-to-noise ratio under a constant acceleration approximation is likely to be significant only for systems with orbital periods less than a few hours (see Johnston&Kulkarni 1991))., The former requirement is likely to be satisfied for the vast majority of systems; the loss of signal-to-noise ratio under a constant acceleration approximation is likely to be significant only for systems with orbital periods less than a few hours (see \citealt{jk91}) ).
 Nevertheless. such exotic and interesting systems do exist and it is important that general surveys maintain sensitivityto them. either through repeated observation (to exploit favorable orbital phases of nearly constant acceleration) or by the use of more involved techniques (e.g. Ransom 2000)).," Nevertheless, such exotic and interesting systems do exist and it is important that general surveys maintain sensitivityto them, either through repeated observation (to exploit favorable orbital phases of nearly constant acceleration) or by the use of more involved techniques (e.g. \citealt{ran00}) )."
 Of those systems exhibiting constant acceleration. sensitivity in this work may still have been compromised if. the acceleration exceeded 430m s.," Of those systems exhibiting constant acceleration, sensitivity in this work may still have been compromised if the acceleration exceeded $\pm 30$ m $^{-2}$."
 This range is typical of previous pulsar searches and is likely to encompass the majority of pulsar systems. however there are several known exceptions.," This range is typical of previous pulsar searches and is likely to encompass the majority of pulsar systems, however there are several known exceptions."
 The eclipsing binary of Terzan 5 (Lyneetal.1990) experiences line-of-sight accelerations greater than this for approximately 30 per cent of its orbit. with a maximum value of 33.2 m $7.," The eclipsing binary of Terzan 5 \citep{lmd+90}
 experiences line-of-sight accelerations greater than this for approximately 30 per cent of its orbit, with a maximum value of 33.2 m $^{-2}$."
 The eccentric double neutron star systems B2127+11C (in MIS). BI534«12. BI913+16 also experience strong accelerations. exceeding 30 m s7 in the line of sight for 20—50 per cent of the orbit.," The eccentric double neutron star systems B2127+11C (in M15), B1534+12, B1913+16 also experience strong accelerations, exceeding 30 m $^{-2}$ in the line of sight for 20–50 per cent of the orbit."
 Such systems have relatively long pulse periods (P230 ms) and so would be detected in the present work at accelerations up to around 40 m s7.," Such systems have relatively long pulse periods $P\gtrsim
30$ ms) and so would be detected in the present work at accelerations up to around 40 m $^{-2}$."
 However the acceleration in these systems exceeds even this value for a significant proportion of the time. reaching a peak in excess of 100 m s for the highly eccentric systems B2127+11C and B1913+16.," However the acceleration in these systems exceeds even this value for a significant proportion of the time, reaching a peak in excess of 100 m $^{-2}$ for the highly eccentric systems B2127+11C and B1913+16."
 Eccentric systems with much shorter orbital periods are expected to evolve from pulsars typical of the presently known double neutron star population via the loss of orbital energy in the form of gravitational radiation., Eccentric systems with much shorter orbital periods are expected to evolve from pulsars typical of the presently known double neutron star population via the loss of orbital energy in the form of gravitational radiation.
 The detectability of such pulsars in this and previous globular cluster searches would be severely affected by acceleration smearing., The detectability of such pulsars in this and previous globular cluster searches would be severely affected by acceleration smearing.
 The primary aim of this work was to detect or place limits on the existence of [sub-]millisecond pulsars (which are not expected to have neutron-star companions). and the computational load associated with high resolution baseband processing limited the feasibility of searching a broad acceleration range.," The primary aim of this work was to detect or place limits on the existence of [sub-]millisecond pulsars (which are not expected to have neutron-star companions), and the computational load associated with high resolution baseband processing limited the feasibility of searching a broad acceleration range."
" However. we note that future surveys with good basic sensitivity would be well-served by searching a range of at least £100 m s. perhaps at a reduced sample rate of ~1 ms and with correspondingly fewer trial dispersion measures and accelerations,"," However, we note that future surveys with good basic sensitivity would be well-served by searching a range of at least $\pm 100$ m $^{-2}$, perhaps at a reduced sample rate of $\sim$ 1 ms and with correspondingly fewer trial dispersion measures and accelerations."
 It is clear from inspection of Figure 1. and the preceding analysis that our FA search. unlike most searches in the past. had a relatively flat sensitivity function for all periods greater than ~0.4 ms.," It is clear from inspection of Figure \ref{fig:sensitivityF} and the preceding analysis that our FA search, unlike most searches in the past, had a relatively flat sensitivity function for all periods greater than $\sim0.4$ ms."
 The FD search also provided similar characteristics for nearby (DM <50 pe em™) unaccelerated pulsars., The FD search also provided similar characteristics for nearby (DM $\la 50$ pc $^{-3}$ ) unaccelerated pulsars.
 We are therefore in a stable position to analyze the period distribution of the detected population relatively free of concerns regarding selection effects., We are therefore in a stable position to analyze the period distribution of the detected population relatively free of concerns regarding selection effects.
 Unfortunately the system used was only sensitive enough to detect a few pulsars. making any assertions somewhat perilous due to small number statistics.," Unfortunately the system used was only sensitive enough to detect a few pulsars, making any assertions somewhat perilous due to small number statistics."
 However it is clear even from this sample that the majority of recycled pulsars do have pulse periods of a few milliseconds or more., However it is clear even from this sample that the majority of recycled pulsars do have pulse periods of a few milliseconds or more.
 All six pulsars detected lay in the two octaves from 3-12 ms. whilst no pulsars were detected in the preceding three octave interval. 0.375—3 ms over which we had comparable sensitivity.," All six pulsars detected lay in the two octaves from 3–12 ms, whilst no pulsars were detected in the preceding three octave interval, 0.375–3 ms over which we had comparable sensitivity."
 No new pulsars were discovered in NGC 6544. Liller 1. or Terzan 5 despite the presence of steep-spectrum emission as reported by Fruchter&Goss(2000).," No new pulsars were discovered in NGC 6544, Liller 1, or Terzan 5 despite the presence of steep-spectrum emission as reported by \citet{fg00}."
. [tis probable that several pulsars in each cluster are responsible for the emission. with each individual pulsar having a flux density below our detection limit.," It is probable that several pulsars in each cluster are responsible for the emission, with each individual pulsar having a flux density below our detection limit."
 An exception might be the “N° source of Terzan 5 (Fruchter&Goss 2000).. which was unresolved at a resolution of 2/99 (c.f.," An exception might be the `N' source of Terzan 5 \citep{fg00}, , which was unresolved at a resolution of 9 (c.f."
" the cluster core radius of11"":; Harris 1996)).", the cluster core radius of; \citealt{har96}) ).
 From the published spectral index and 20 em flux. we infer à 50 cm flux density of 9 mJy.," From the published spectral index and 20 cm flux, we infer a 50 cm flux density of 9 mJy."
 The sky temperature in this region is ~300 K at 70 em (Haslametal.1982). implying a temperature of ~100 K at 50 em (Lawsonetal.1987).," The sky temperature in this region is $\sim 300$ K at 70 cm \citep{hssw82}, implying a temperature of $\sim 100$ K at 50 cm \citep{lmo+87}."
. The sensitivity limits are thus 160/70—2.3 times greater than indicated in Figure 1.. however even so. a very broad profile and/or very short spin period would be required for the pulsar to have remained undetected at 9 mJy.," The sensitivity limits are thus $160/70 \simeq 2.3$ times greater than indicated in Figure \ref{fig:sensitivityF}, however even so, a very broad profile and/or very short spin period would be required for the pulsar to have remained undetected at 9 mJy."
 It is possible that the source Is a pulsar and was undetected due to scatter broadening., It is possible that the source is a pulsar and was undetected due to scatter broadening.
 Up to à millisecond of scattering could be induced by the intervening interstellar medium without having hampered the detection of the presently known pulsars. with pulse periodsof 11.6 and 8.4 ms (Lyneetal... 2000)..," Up to a millisecond of scattering could be induced by the intervening interstellar medium without having hampered the detection of the presently known pulsars, with pulse periodsof 11.6 and 8.4 ms \citep{lmbm00}. ."
 Such scattering would be catastrophic for the detection of very fast pulsars and could, Such scattering would be catastrophic for the detection of very fast pulsars and could
scales for different values of D.,scales for different values of $\beta$.
 For example. the passive scalar contribution to the density fluctuations in the solar wind should be suppressed when D>|. because the damping rate of the passive fluctuations (slow waves and entropy modes) is larger than the cascade rate ($2)).," For example, the passive scalar contribution to the density fluctuations in the solar wind should be suppressed when $\beta \gtrsim 1$, because the damping rate of the passive fluctuations (slow waves and entropy modes) is larger than the cascade rate \ref{sec:pred}) )."
 Thus. measurements of the density fluctuations m solar wind epochs with large P are likely to be particularly instructive: more detailed theoretical. caleulations of the suppression. of the passive scalar contribution at D=I would aid in interpreting such measurements.," Thus, measurements of the density fluctuations in solar wind epochs with large $\beta$ are likely to be particularly instructive; more detailed theoretical calculations of the suppression of the passive scalar contribution at $\beta \gtrsim
1$ would aid in interpreting such measurements."
 In addition. future progress towards understanding the inertial-range power spectra of the density fluctuations. Sunward Alfvénn waves. and anti-Sunward Alfvénn waves in imbalanced (or cross-helical) turbulence could lead to significantly tighter constraints on the heating rate contributed by KAW turbulence. particularly near the Sun where the energy in anti-Sunward Alfvénn waves greatly exceeds the energy of Sunward Alfvénn waves.," In addition, future progress towards understanding the inertial-range power spectra of the density fluctuations, Sunward Alfvénn waves, and anti-Sunward Alfvénn waves in imbalanced (or cross-helical) turbulence could lead to significantly tighter constraints on the heating rate contributed by KAW turbulence, particularly near the Sun where the energy in anti-Sunward Alfvénn waves greatly exceeds the energy of Sunward Alfvénn waves."
 Finally. the fact that the upper limit on the KAW heating rate at | AU derived from density fluctuations is nearly equal to the measured heating rate in the solar wind (33.2)) strongly motivates a more detailed analysis of the small-scale density fluctuations and their implications.," Finally, the fact that the upper limit on the KAW heating rate at 1 AU derived from density fluctuations is nearly equal to the measured heating rate in the solar wind \ref{sec:1AU}) ) strongly motivates a more detailed analysis of the small-scale density fluctuations and their implications."
 This work was supported in part by the the Center for Integrated Computation and Analysis of Reconnection and Turbulence (CICART) under DOE under grant number DE-FG02-07-ER46372. by the NSF/DOE Partnership in Basic Plasma Science and Engineering under grant number AST-0613622. and by NASA under grant numbers NNXO7AP65G and NNHOGZDAOO01N-SHPO06-0071 at the University of New Hampshire.," This work was supported in part by the the Center for Integrated Computation and Analysis of Reconnection and Turbulence (CICART) under DOE under grant number DE-FG02-07-ER46372, by the NSF/DOE Partnership in Basic Plasma Science and Engineering under grant number AST-0613622, and by NASA under grant numbers NNX07AP65G and NNH06ZDA001N-SHP06-0071 at the University of New Hampshire."
 E. Quataert was supported in part by NSF-DOE Grant PHY-0812811 and by NSF Grant ATM-0752503., E. Quataert was supported in part by NSF-DOE Grant PHY-0812811 and by NSF Grant ATM-0752503.
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(e.g. Balelryetal. 2006)),"(e.g., \citealt{baldry06}) )."
 ὃν splitting the luminosity function into central and satellite galaxies. we see tha this is driven by the rapidly increasing preponderance of blue central galaxies at faint magnitudes.," By splitting the luminosity function into central and satellite galaxies, we see that this is driven by the rapidly increasing preponderance of blue central galaxies at faint magnitudes."
" In contrast. the Luminosity function of central red. galaxies drops by abou three orders of magnitude between ""Ad, of 20 and. 17."," In contrast, the luminosity function of central red galaxies drops by about three orders of magnitude between $^{0.1}M_r$ of $-20$ and $-17$."
" ""his ellect is driven by the AGN feedback in the mocel. anc is independent of the stripping moclel."," This effect is driven by the AGN feedback in the model, and is independent of the stripping model."
 The properties of satellites in the new mocel are much more dependent on the environmental physics include in the model., The properties of satellites in the new model are much more dependent on the environmental physics included in the model.
 In. the new model. the blue and. red. Lis reach similar values at the faint end. (consistent with the results in Figure 3. showing that at the faint end. there are roughly similar numbers of τος ancl blue. satellites)," In the new model, the blue and red LFs reach similar values at the faint end (consistent with the results in Figure \ref{fig:fblue}
 showing that at the faint end there are roughly similar numbers of red and blue satellites)."
 The dramatic increase in the fraction of faint blue satellites can be seen by comparing the solid and dashed. blue lines., The dramatic increase in the fraction of faint blue satellites can be seen by comparing the solid and dashed blue lines.
 This is an alternative wav of presenting the information in Fig. 4..," This is an alternative way of presenting the information in Fig. \ref{fig:fblue_lim},"
 and underscores the importance of studying the environmental dependence of galaxy properties in order to obtain a complete picture of galaxy formation and evolution., and underscores the importance of studying the environmental dependence of galaxy properties in order to obtain a complete picture of galaxy formation and evolution.
 Until now. most. current semi-analvtic models of galaxy formation have adopted a crude modelling of the ram pressure stripping of the hot gaseous halos of satellite galaxies.," Until now, most current semi-analytic models of galaxy formation have adopted a crude modelling of the ram pressure stripping of the hot gaseous halos of satellite galaxies."
 In particular. they typically assume. complete and instantaneous stripping of the hot eas halo when the galaxy first falls in. without regard for the ealaxy’s mass. orbit. or structural properties.," In particular, they typically assume complete and instantaneous stripping of the hot gas halo when the galaxy first falls in, without regard for the galaxy's mass, orbit, or structural properties."
 Εις is at odds with results of hydrodynamic simulations. ancl also with recent X-ray observations of galaxies in massive clusters.," This is at odds with results of hydrodynamic simulations, and also with recent X-ray observations of galaxies in massive clusters."
 Lo the oesent study. we have improved the treatment of stripping » implementing the model of AleCarthyetal.(2008) (which has been shown to match simulations of ram oressure stripping to high accuracy) into the semi-analvtic model for galaxy formation.," In the present study, we have improved the treatment of stripping by implementing the model of \citet{mccarthy08} (which has been shown to match simulations of ram pressure stripping to high accuracy) into the semi-analytic model for galaxy formation."
 Although the initial stripping event does not require us to add additional xwameters to the model. subsequent stripping of the gas re-reated from the disk. requires additional parameterization.," Although the initial stripping event does not require us to add additional parameters to the model, subsequent stripping of the gas re-heated from the disk requires additional parameterization."
 We parameterize this process by assuming that most of the ejecta are retained in the galaxy after the first. pericentre )nssaSo., We parameterize this process by assuming that most of the ejecta are retained in the galaxy after the first pericentre passage.
 We find that the newly implemented: treatment of stripping leads to a significant improvement in the ability of the model to match the colours of satellite galaxies in 16 SDSS., We find that the newly implemented treatment of stripping leads to a significant improvement in the ability of the model to match the colours of satellite galaxies in the SDSS.
 The new mocel is also able to account for the environmental dependence of the colours of satellite galaxies., The new model is also able to account for the environmental dependence of the colours of satellite galaxies.
 )ur results suggest that for the majority of satellite galaxies. iis environmental process can be as important in mocifving 10 galaxy colours as intrinsic processes. such as AGN or Lgupernoyvac feedback. which operate within the satellites remselves.," Our results suggest that for the majority of satellite galaxies, this environmental process can be as important in modifying the galaxy colours as intrinsic processes, such as AGN or supernovae feedback, which operate within the satellites themselves."
 This finding is in broad agreement with previous studies that found that internal processes that quench star formation do not seem capable of explaining the full range of colour and morphology data (e.g. Weinmannetal.2006a:Daldryetal.2006).," This finding is in broad agreement with previous studies that found that internal processes that quench star formation do not seem capable of explaining the full range of colour and morphology data (e.g., \citealt{weinmann06a,baldry06}) )."
 Although we have only focused on the stripping of hot gascous halos in the present study. other environmental processes may be relevant as well.," Although we have only focused on the stripping of hot gaseous halos in the present study, other environmental processes may be relevant as well."
 We now brielly. review some of these and conclude that none of them appear to be as important as the ram pressure stripping of the hot haloes and the feedback: already incorporated. into our presen model., We now briefly review some of these and conclude that none of them appear to be as important as the ram pressure stripping of the hot haloes and the feedback already incorporated into our present model.
 In cases where the ram pressure stripping of the ho halo is complete. some stripping of the cold. gaseous. clisks may also occur (Gunn&Gott1972:Abadetal.1999:Quilisotal.2000).," In cases where the ram pressure stripping of the hot halo is complete, some stripping of the cold gaseous disks may also occur \citep{gunn72,abadi99,quilis00}."
.. However. Okamoto&Nagashima(2003) anc Lanzonietal.(2005) have explored ram pressure stripping of disks in seni-analvtic models ancl have concluded. tha the effect on the colours and star formation rates of satellite ealaxies are minimal.," However, \citet{okamoto03} and \citet{lanzoni05} have explored ram pressure stripping of disks in semi-analytic models and have concluded that the effect on the colours and star formation rates of satellite galaxies are minimal."
 VPhis most likely stems from the fac that disk stripping is only expected. to be relevant for a minority of satellite galaxies whose orbits take them into the very centres ol massive sysenis (e.g. Müssen&DeLucia 2008)).," This most likely stems from the fact that disk stripping is only expected to be relevant for a minority of satellite galaxies whose orbits take them into the very centres of massive systems (e.g., \citealt{bruggen08}) )."
 Thermal evaporation of the hot gaseous haloes (and/or isks) could. also be relevant (Cowie&Songaila1977). rut Observational studies of bubbles and. cold fronts in. N-rav groups ane clusters have placed strong constraints on 1e elliciency. of conduction. concluding that it is strongly suppressed (Markeviteh&Vikhlinin2007:McNamaratulsen 2007).," Thermal evaporation of the hot gaseous haloes (and/or disks) could also be relevant \citep{cowie77}, but observational studies of bubbles and cold fronts in X-ray groups and clusters have placed strong constraints on the efficiency of conduction, concluding that it is strongly suppressed \citep{markevitch07,mcnamara07}."
. “Purbulent stripping via the generation of Ixelvin-Helmholtz anc Bavleigh-Tavlor instabilities at the interface between the hot gaseous halo of the satellite anc he parent svstem is possible. but the timescale for this vpe of stripping is generally quite long (sce AleCarthyο ).," Turbulent stripping via the generation of Kelvin-Helmholtz and Rayleigh-Taylor instabilities at the interface between the hot gaseous halo of the satellite and the parent system is possible, but the timescale for this type of stripping is generally quite long (see \citealt{mccarthy08}) )."
 Other possibly relevant processes. include viscous stripping (unfortunately at present the viscosity of the ho gas in groups and clusters is poorly constrained: MeNamara&Nulsen 2007)). tical effects on the gas as the result of the interaction with the gravitational potential of the paren ido (Dyrd&Valtonen1990:Merritt.1983). or with the other satellites (Le. mergers (Alihos1995) or harassmen (Mooreetal. 1996))). and shock heating of the satellite's ho eas às it [alls in at transonic velocities.," Other possibly relevant processes include viscous stripping (unfortunately at present the viscosity of the hot gas in groups and clusters is poorly constrained; \citealt{mcnamara07}) ), tidal effects on the gas as the result of the interaction with the gravitational potential of the parent halo \citep{byrd90,merritt83} or with the other satellites (i.e., mergers \citep{mihos95} or harassment \citep{moore96}) ), and shock heating of the satellite's hot gas as it falls in at transonic velocities."
 However. MeCarthiyetal.(2008). have argued that ram. pressure stripping of he hot gaseous haloes is always more efficient than tida stripping by the parent halo's potential or shock heating in cases where the satellite mass is less than about ofthe went halo's mass.," However, \citet{mccarthy08} have argued that ram pressure stripping of the hot gaseous haloes is always more efficient than tidal stripping by the parent halo's potential or shock heating in cases where the satellite mass is less than about of the parent halo's mass."
 In this paper we have concentrated only on comparisons with cata at low redshifts ες« 0.1)., In this paper we have concentrated only on comparisons with data at low redshifts $z<0.1$ ).
 In a future study. we intend to test the model at other redshifts and compare the redshift evolution of galaxy colours with data from current deep survevs.," In a future study, we intend to test the model at other redshifts and compare the redshift evolution of galaxy colours with data from current deep surveys."
 Another important application of the mocel is the study of the clustering of galaxies as a function of different physical quantities. such as colour.," Another important application of the model is the study of the clustering of galaxies as a function of different physical quantities, such as colour."
 In addition to the large scale dependence driven by the relative importance of AGN feedback and ram pressure stripping as a function of halo mass. our model also predicts a radial dependence within a halo driven by the variation in the strength of stripping with the galaxy orbit.," In addition to the large scale dependence driven by the relative importance of AGN feedback and ram pressure stripping as a function of halo mass, our model also predicts a radial dependence within a halo driven by the variation in the strength of stripping with the galaxy orbit."
 The colour dependence of small-scale clustering will constrain the model ancl perhaps, The colour dependence of small-scale clustering will constrain the model and perhaps
Eq(9)) isa first order non-linear ordinary ditfereutial equation which coutaius a characteristic parameter ¢ but satisfies a two-boundary coudition: =MOT) This kind of equation cau be easily solved. by the μπαΊσα. method described in (Pressetal1992).. chapter 17.,"\ref{yofxDetermination}) ) is a first order non-linear ordinary differential equation which contains a characteristic parameter $\zeta$ but satisfies a two-boundary condition: =1 This kind of equation can be easily solved by the numerical method described in \citep[]{PressNR}, chapter 17."
 We display ournunuerical results in FIC.2.., We display ournumerical results in \ref{zetaFig}.
 Which is also depicted in the fleure is the results of Ct.νε) When dark cuerev is assuned to move svuchronously with ordinary matters both ou IInbble scale and galaxy cluster scales.," Which is also depicted in the figure is the results of $\zeta(w,\Omega_{mb,ta})$ when dark energy is assumed to move synchronously with ordinary matters both on Hubble scale and galaxy cluster scales."
 The actual fact should lio between this two extreme cases;, The actual fact should lie between this two extreme cases.
" But the result of Wane&Steinhardt1998) does not lie between this two extreme cases, so if cannot be thought asau approximation of the actual facts. Eq(9))"," But the result of \citep[]{WangSteinhardt1} does not lie between this two extreme cases, so it cannot be thought asan approximation of the actual facts. \ref{yofxDetermination}) )"
 can also be solved analytically., can also be solved analytically.
 Note SO substituting eqtc11)) iuto (9)) aud equating the two sides of the resulting equation order by order u.c. we can get ij and a similar expression for 3. ... etc.," Note so 1_- Let ] substituting \ref{yofxAnsaltz}) ) into \ref{yofxDetermination}) ) and equating the two sides of the resulting equation order by order in $x$ , we can get ] and a similar expression for $\beta$ , ... etc."
 Tn this paper we only need to know o., In this paper we only need to know $\alpha$.
 Some people may argue that if progressional solution of eq(9)) is uot of the form as we wrote 1- eqt11)) or does not exist at all. then our ansaltz —j) will be iuvalid.," Some people may argue that if progressional solution of \ref{yofxDetermination}) ) is not of the form as we wrote in \ref{yofxAnsaltz}) ) or does not exist at all, then our ansaltz \ref{yofxAnsaltz}) ) will be invalid."
 We would like to poit out that if such thiugs occur. then when ο) is substitute oeito eq(9)) we can not ect a selt-conusisteut equation with the two sides equated order by order m ο.," We would like to point out that if such things occur, then when \ref{yofxAnsaltz}) ) is substituted into \ref{yofxDetermination}) ) we can not get a self-consistent equation with the two sides equated order by order in $x$."
 Tu the ideal model. if there is au over-deuse region iu a flat backeround universe. then at very carly times. this region will expand as the background universe expands: but because this region's over-deuse. its expanding rate will decrease and stop doing so at some iniddle times: then it starts to shrink because of sclberavitating. the final fate of this over-dense reeion is a singular point.," In the ideal model, if there is an over-dense region in a flat background universe, then at very early times, this region will expand as the background universe expands; but because this region's over-dense, its expanding rate will decrease and stop doing so at some middle times; then it starts to shrink because of self-gravitating, the final fate of this over-dense region is a singular point."
 But in practice. when this region shrinks to some degree. the pressures originate from the raucom moving of particles iuside the over-dense region will balance the selteravitation aud the system will cuter the virtalization period.," But in practice, when this region shrinks to some degree, the pressures originate from the random moving of particles inside the over-dense region will balance the self-gravitation and the system will enter the virialization period."
 Iu theoretical studies. it is usually assuued that the virialization point is coincident with the collapse point of the ideal model on the time axis.," In theoretical studies, it is usually assumed that the virialization point is coincident with the collapse point of the ideal model on the time axis."
" According to Press-Sheeter theory. if an over-deuse reeion is to be virializated at some time e... its deusitv-contrast «hould be no less than ὃν0,09.de)."," According to Press-Shceter theory, if an over-dense region is to be virializated at some time $a_c$, its density-contrast should be no less than $\delta_c(w,\Omega_{m0},a_c)$."
" where ο aud py are the matter densities of cluster and backeround respectively, while Dya) isthegrowth function of liuear perturbation theory (Docdelsou 2003)."," where $\rho_{mc}$ and $\rho_{mb}$ are the matter densities of cluster and background respectively, while $D_1(a)$ isthegrowth function of linear perturbation theory \citep[]{Dodelson}, , ."
. (18) To reduce the uwmuerical computation burdens. we will use the fitting forumlacs provided by," To reduce the numerical computation burdens, we will use the fitting formulaes provided by"
brightness barred galaxies (Debattista&Sellwood2000).. it is important to verify that the full interaction with a live dark matter (DM) halo does not lead to a drastically different evolution for the structural and. kinematic parameters.,"brightness barred galaxies \citep{ds00}, it is important to verify that the full interaction with a live dark matter (DM) halo does not lead to a drastically different evolution for the structural and kinematic parameters."
 The ability of a bar to lose angular momentum to a DM halo helps make it stronger (Debattista&Sellwood2000)., The ability of a bar to lose angular momentum to a DM halo helps make it stronger \citep{ds00}.
. We therefore have run a pair of matched live and rigid halo simulations., We therefore have run a pair of matched live and rigid halo simulations.
 The live halo simulations were run on PIKXDGRAV (Stadel(2001): details of the live halo simulations will be presented elsewhere)., The live halo simulations were run on PKDGRAV \citet{s01}; details of the live halo simulations will be presented elsewhere).
 We checked that these results are nol sensitive to resolution bv running the live halo model with larger NV and smaller e., We checked that these results are not sensitive to resolution by running the live halo model with larger $N$ and smaller $\epsilon$.
 The two simulations evolved differentlv: the bar in the live halo case was slightly stronger and slows down because of dvnamical friction with the halo., The two simulations evolved differently: the bar in the live halo case was slightly stronger and slows down because of dynamical friction with the halo.
" Nonetheless. the mass density and V/o,, profiles in the rieid- ancl live-halo simulations remain «quite similar. as we illustrate in Fig."," Nonetheless, the mass density and $V/\sigma_\phi$ profiles in the rigid- and live-halo simulations remain quite similar, as we illustrate in Fig."
" 1. (we use σι, in this comparison because it is representative of what we might expect in high inclination observations).", \ref{fig:fig1} (we use $\sigma_\phi$ in this comparison because it is representative of what we might expect in high inclination observations).
 The larger central density in the live halo simulation of Fig., The larger central density in the live halo simulation of Fig.
 1. is a result of the fact that its bar can shed additional angular momentum., \ref{fig:fig1} is a result of the fact that its bar can shed additional angular momentum.
 Otherwise. (his comparison suggests that (he gross features seen in the live halo simulation are well-reproduced by (he rigid halo simulation. eivine confidence that the rigid halo simulations survey [rom which we draw the main conclusions of this paper is adecquate to describe the evolution of massive disks.," Otherwise, this comparison suggests that the gross features seen in the live halo simulation are well-reproduced by the rigid halo simulation, giving confidence that the rigid halo simulations survey from which we draw the main conclusions of this paper is adequate to describe the evolution of massive disks."
 Moreover. bv using rigid halos. we assure (hat we obtain a minimal level of bar secular evolution. which helps to disentangle the effects of different evolutionary. processes.," Moreover, by using rigid halos, we assure that we obtain a minimal level of bar secular evolution, which helps to disentangle the effects of different evolutionary processes."
too much.,too much.
 Lots of stars form between 3—7 Gyr and basically within the scatter of the data., Lots of stars form between $3-7$ Gyr and basically within the scatter of the data.
" In model mw12 (n= 11), we assume shorter bursts (d~— 0.3)."," In model mw12 $n=11$ ), we assume shorter bursts $d\sim0.2-0.3$ )."
" Stars formed between 3—6 Gyr have metallicity consistent with the observed PNe, but in this model we still have lots of stars formed at a late age."," Stars formed between $3-6$ Gyr have metallicity consistent with the observed PNe, but in this model we still have lots of stars formed at a late age."
 The results of this model are consistent with observations., The results of this model are consistent with observations.
" Compared with the former two SF regimes, the bursting one reproduces more closely the scatter in the abundance ratios, especially for the N/O ratio when 10+log(O/H)>8.2, as shown in Fig 14.."," Compared with the former two SF regimes, the bursting one reproduces more closely the scatter in the abundance ratios, especially for the N/O ratio when $>$ 8.2, as shown in Fig \ref{Fig:XOburstSF}."
" Therefore, model mw12 can also be considered to be among the best ones."," Therefore, model mw12 can also be considered to be among the best ones."
" All these models have SFRs (~0.15—0.2 yr) consistent with that derived from Ha observations, but slightly higher than the SFR obtained from FIR and 1.4GHz measurements ~0.05—0.07 (see Table 1))."," All these models have SFRs $\sim0.15-0.2$ ) consistent with that derived from $\alpha$ observations, but slightly higher than the SFR obtained from FIR and 1.4GHz measurements $\sim 0.05-0.07$ (see Table \ref{Tab:obs}) )."
" If we were to accept the SFR value derived from Ho, then the bursting SF history might also be acceptable for IC10."," If we were to accept the SFR value derived from $\alpha$, then the bursting SF history might also be acceptable for IC10."
" On the other hand, the SFR values obtained from FIR and 1.4 GHz measurements imply that a gasping SF history is more likely."," On the other hand, the SFR values obtained from FIR and 1.4 GHz measurements imply that a gasping SF history is more likely."
 We have studied the chemical evolution of the dwarf irregular galaxy IC10 for which detailed abundance data are now available (MG09)., We have studied the chemical evolution of the dwarf irregular galaxy IC10 for which detailed abundance data are now available (MG09).
 We have adopted a chemical evolution model for dwarf irregulars developed by Yin et al. (2010))., We have adopted a chemical evolution model for dwarf irregulars developed by Yin et al. \cite{yin10}) ).
 This model includes detailed metallicity-dependent stellar yields for both massive and low and intermediate mass stars., This model includes detailed metallicity-dependent stellar yields for both massive and low and intermediate mass stars.
 The model also includes feedback from SNe and stellar winds and follows the development of galactic winds., The model also includes feedback from SNe and stellar winds and follows the development of galactic winds.
" We have explored three cases: i) no winds (no feedback); ii) normal winds, namely all the gas is lost at the same rate; and iii) metal-enhanced winds where the metals are lost preferentially relative to H and He."," We have explored three cases: ) no winds (no feedback); ) normal winds, namely all the gas is lost at the same rate; and ) metal-enhanced winds where the metals are lost preferentially relative to H and He."
 We also explored several regimes of star formation: a) bursting mode with short and long bursts (a regime we called “gasping SF’); and b) a continuous but low SFR., We also explored several regimes of star formation: a) bursting mode with short and long bursts (a regime we called “gasping SF”); and b) a continuous but low SFR.
" We computed the evolutions of the abundances of several elements (He, N, O, S)."," We computed the evolutions of the abundances of several elements (He, N, O, S)."
" The parameters of our model were: 1) the number of bursts; 2) the duration of bursts; 3) the efficiency of star formation; and 4) the efficiency of the galactic wind rate, assumed to be proportional to the amount of gas present at the time of the wind, which can vary from element to element in the case of the metal-enhanced wind."," The parameters of our model were: 1) the number of bursts; 2) the duration of bursts; 3) the efficiency of star formation; and 4) the efficiency of the galactic wind rate, assumed to be proportional to the amount of gas present at the time of the wind, which can vary from element to element in the case of the metal-enhanced wind."
" The observational constraints were represented by the abundances and abundance ratios of the above-mentioned elements, by the present-time gas mass, the SFR, and the estimated total mass of IC10."," The observational constraints were represented by the abundances and abundance ratios of the above-mentioned elements, by the present-time gas mass, the SFR, and the estimated total mass of IC10."
" In spite of the large number of parameters, the number of observational constraints reproduced by our model makes us confident in concluding the following:"," In spite of the large number of parameters, the number of observational constraints reproduced by our model makes us confident in concluding the following:"
a description of the spectroscopic mapping technique and explanation of those algorithms. included. in the. package that were developed to analyse spectroscopic maps.,a description of the spectroscopic mapping technique and explanation of those algorithms included in the package that were developed to analyse spectroscopic maps.
 Some benchmarks of this new package are also given within Section 3., Some benchmarks of this new package are also given within Section 3.
 In Section 4 the empirical results for one of the long-slits observed. in this work are compared. with those obtained from the spectroscopic maps by means of the package., In Section 4 the empirical results for one of the long-slits observed in this work are compared with those obtained from the spectroscopic maps by means of the package.
 Section 5 is dedicated to our. spatially resolved results. for NGC 40. while the summary aud conclusions are in Section 6.," Section 5 is dedicated to our spatially resolved results for NGC 40, while the summary and conclusions are in Section 6."
 The observational data used in. this work includes both an llo image and optical lone-slit) spectra., The observational data used in this work includes both an $\alpha$ image and optical long-slit spectra.
 Phe observations were obtained on the night. of 2005 October Os’ a0 the 2.56 m Nordic Optical Telescope. (NOT) located at the Observatorio del Roque de los Muchachos (European Northern Observatory. La Palma. Spain).," The observations were obtained on the night of 2005 October $^{th}$, at the 2.56 m Nordic Optical Telescope (NOT), located at the Observatorio del Roque de los Muchachos (European Northern Observatory, La Palma, Spain)."
 bor both the image ancl spectra. the Andalucia Faint Object perti and Camera CALEOSC?) which has a pixel scaleof * was used.," For both the image and spectra, the Andalucia Faint Object Spectrograph and Camera (ALFOSC), which has a pixel scale of $^-$$^1$, was used."
 The n image has an exposure Dune ofmi30 s and the secing during the exposure was, The $\alpha$ image has an exposure time of 30 s and the seeing during the exposure was.
 theIn paper this iniage is used only for comparison with llo spectroscopic map that is constructed using the spectra. as described in detail below.," In this paper this image is used only for comparison with the $\alpha$ spectroscopic map that is constructed using the spectra, as described in detail below."
 The reduction of the Ho. image included only bias subtraction. in addition to Ilat-Bield.(dome‘this ancl sky) correction.," The reduction of the $\alpha$ image included only bias subtraction, in addition to flat-field (dome and sky) correction."
 No photometric calibration of image was performed., No photometric calibration of this image was performed.
 Lt is worth noting. however. that the resolution of the Ho. image is better than the resolution achieved. with the maps.," It is worth noting, however, that the resolution of the $\alpha$ image is better than the resolution achieved with the maps."
 For spectroscopy. the long-slit width used. was," For spectroscopy, the long-slit width used was."
 The ALEOSC was used with the grism number 7. which has 600 rules 1 and a spectral coverage from 3850 to 6850A.," The ALFOSC was used with the grism number 7, which has 600 rules $^-$$^1$ and a spectral coverage from 3850 to 6850."
. Phe reciprocal dispersion. of the binnecl pixel is pixel, The reciprocal dispersion of the binned pixel is 3.0 $^-$$^1$.
 The slit was located on 16 cilferent. parallel positions across the nebula., The slit was located on 16 different parallel positions across the nebula.
 The distance between them was fixed⋅ and set atδ, The distance between them was fixed and set at.
"ν, Forqe 15 of those positionss the exposure times were ο 300 s. For slit D only one 300 s exposure was taken.", For 15 of those positions the exposure times were 3 $\times$ 300 s. For slit B only one 300 s exposure was taken.
 Figure 1. shows our NOT Lla image. anc also the relative slit positions across the nebula.," Figure \ref{imhalpha} shows our NOT $\alpha$ image, and also the relative slit positions across the nebula."
 Note. in addition. that the slits do not cover the entire object.," Note, in addition, that the slits do not cover the entire object."
 Slit lengths in ligure 1. are reduced for visualization purposes with their respective labels used throughout the text., Slit lengths in Figure \ref{imhalpha} are reduced for visualization purposes with their respective labels used throughout the text.
 The data reduction was performed. using the standard proceduresfor long-slit spectroscopy of the Image Reduction Analvsis. Facility. (18AT)., The data reduction was performed using the standard proceduresfor long-slit spectroscopy of the Image Reduction Analysis Facility ).
 Jas frames. Lat-fielel. helium-neon wavelength calibrations and standard star (6191D2DB) exposures were obtained during the observation runs. ancl used in the reduction/calibration procedure.," Bias frames, flat-field, helium-neon wavelength calibrations and standard star (G191B2B) exposures were obtained during the observation runs, and used in the reduction/calibration procedure."
 Three spectra per slit. position were taken to improve the signal-to-noise ratio (S/N) and eliminate cosmic ravs., Three spectra per slit position were taken to improve the signal-to-noise ratio (S/N) and eliminate cosmic rays.
 The only non-standard procedure. used was binning on the dispersion axis., The only non-standard procedure used was binning on the dispersion axis.
 This was necessary to improve the wavelength (A) calibration due to the fact that the helium-neon [amp was resolved in the are frames., This was necessary to improve the wavelength $\lambda$ ) calibration due to the fact that the helium-neon lamp was resolved in the arc frames.
 The final calibrated. spectrum for slit C. which passes through the central star. is shown in Figure 2..," The final calibrated spectrum for slit G, which passes through the central star, is shown in Figure \ref{slitspectra}."
 Fluxes here are representative of the emission integrated alone the slit. which we will later call “WNT (whole nebula) when comparing long-slit measurements and results obtained from the spectroscopic maps. in Sections 3.2 and 4.," Fluxes here are representative of the emission integrated along the slit, which we will later call “WN” (whole nebula) when comparing long-slit measurements and results obtained from the spectroscopic maps, in Sections 3.2 and 4."
 ]5mission-line mapping is a technique in which the spatial xofiles from a set of parallel slits. for a given emission-line. are combined to create an emission-line map of the nebula (Monteiroetal.2004).," Emission-line mapping is a technique in which the spatial profiles from a set of parallel slits, for a given emission-line, are combined to create an emission-line map of the nebula \citep{b10}."
.. We applied this technique using the 16 parallel long-slit spectra described above. to obtain the emission-line maps of NGC 40.," We applied this technique using the 16 parallel long-slit spectra described above, to obtain the emission-line maps of NGC 40."
 These maps correspond to à ield of view of ((see Figure 3))., These maps correspond to a field of view of $\times$ (see Figure \ref{hacontour}) ).
 This method has been previously applied to Hubble 5. GC 6369. Mz IL and NGC 2022 (Ricectal. 2004: 2004: Monteiroetal.2005. and," This method has been previously applied to Hubble 5, NGC 6369, Mz 1 and NGC 2022 \citealt{b13}; ; \citealt{b10}; ; \citealt{b14} and"
The paper is organized as follows.,The paper is organized as follows.
 82 discusses the adopted metallicity scale. while 83 and 84 present the equation sets for both thedisk-Ghe andscenarios. respectively.," 2 discusses the adopted metallicity scale, while 3 and 4 present the equation sets for both the and, respectively."
 85 describes the computational routine which provides photometric estimates ol metallicity. reddening ancl clistance. while in 86 we test the described method on two template GC's in (he Galactic bulee and in the Large Magellanic Cloud. namely NGC 6539 and NGC L978 respectively. in order to demonstrate its reliability.," 5 describes the computational routine which provides photometric estimates of metallicity, reddening and distance, while in 6 we test the described method on two template GCs in the Galactic bulge and in the Large Magellanic Cloud, namely NGC 6539 and NGC 1978 respectively, in order to demonstrate its reliability."
 As widely discussed in Ferraroetal.(1999.2000) a correct parameterization of the RGB characteristics as fanction of the metal content of the population does require (he knowledge of the so-called.global metallicity. which takes into account the iron as well as the a element abundances.," As widely discussed in \citet{F99,F00} a correct parameterization of the RGB characteristics as function of the metal content of the population does require the knowledge of the so-called metallicity, which takes into account the iron as well as the $\alpha$ –element abundances."
 This is an important point since the location of the RGB strongly depends on the lowionization potential |Fe2-Mg4-Si/1l] abundance mixture (Straniero&Chielfi1991:Salaris&Cassisi1996) rather than on [Fe/H] abundance alone.," This is an important point since the location of the RGB strongly depends on the low–ionization potential [Fe+Mg+Si/H] abundance mixture \citep{SC91,SC96}
 rather than on [Fe/H] abundance alone."
 In fact. lowionization potential elements are the main contributors (o tree electrons which generate the I ion. (he major component responsible for the conünuum opacity in the RGB (temperature range (3000.6000Ix.Renzini1971).," In fact, low–ionization potential elements are the main contributors to free electrons which generate the $\rm H^-$ ion, the major component responsible for the continuum opacity in the RGB temperature range \citep[3000--6000~K,][]{R77}."
 In halo/disk field stars the average [o /Fe] abundance ratio shows a general enhancement of 0.30.5 dex with respect to the Solar value up to [Fe/I]z 1 (seee.g.Doesgaardetal.etal.2000:Carretta.Gratton&Sneden2000.andreferencestherein) and a linear decreasing trend towards Solar [o /Fe] with further increasing metallicity.," In halo/disk field stars the average $\alpha$ /Fe] abundance ratio shows a general enhancement of 0.3–0.5 dex with respect to the Solar value up to $\approx$ –1 \citep[see e.g.][ and references therein]{BOE99,gra00,car00}
 and a linear decreasing trend towards Solar $\alpha$ /Fe] with further increasing metallicity."
 A fa /Fe| enhancement is also found in the metal poor halo GCs (seee.g.Gratton.Snecden&Carretta2004:Snedenetal.2004.anclreferences therein)..," A $\alpha$ /Fe] enhancement is also found in the metal poor halo GCs \citep[see e.g.][ and references therein]{gra04,sne04}."
 However. until a few vears ago. onlv a few measurements were available in the high metallicity regime. to properly define the [a /Fe| ivend in GC's (seee.g.τα1994:Carney1996).," However, until a few years ago, only a few measurements were available in the high metallicity regime, to properly define the $\alpha$ /Fe] trend in GCs \citep[see e.g.][]{kra94,car96}."
. The actual position of the knee (i.e. the metallicity at which fa/Fe] begins to decrease) depends on the tvpe Ia SN timescales and it is also a function of the star formation rate. while the amount of a enhancement depends on the initial mass function of the progenitors of the (vpe I] SNe (seeMeWilliam1997).," The actual position of the knee (i.e. the metallicity at which $\alpha$ /Fe] begins to decrease) depends on the type Ia SN timescales and it is also a function of the star formation rate, while the amount of $\alpha $ enhancement depends on the initial mass function of the progenitors of the type II SNe \citep[see][]{mw97}."
. Measurements of metal rich. field and cluster giants. towards the Galactic bulge are a recent. accomplishment of high resolution optical and II. spectroscopy., Measurements of metal rich field and cluster giants towards the Galactic bulge are a recent accomplishment of high resolution optical and IR spectroscopy.
 Bulge stars are indeed ideal targets to study the behavior of the abundance patterns in the high metallicity domain. but. foreground extinction is so great as to largely preclude optical studies of any kind. particularly at high spectral resolution.," Bulge stars are indeed ideal targets to study the behavior of the abundance patterns in the high metallicity domain, but foreground extinction is so great as to largely preclude optical studies of any kind, particularly at high spectral resolution."
 The most accurate abundance determinations obtained so lar and based on hieh resolution optical spectroscopy refer (ο a sample of in the Daade's window, The most accurate abundance determinations obtained so far and based on high resolution optical spectroscopy refer to a sample of K--giants in the Baade's window
Fig.,Fig.
 3 shows the results of ligi curve fitting for image C in he region of LIAL with the £I-model., \ref{Fi3} shows the results of light curve fitting for image C in the region of HAE with the $F^1$ -model.
 Here. 44 «ala points were included corresponding to the 1289.1442] epoch intervaI (light circles).," Here, 44 data points were included corresponding to the [1289,1442] epoch interval (light circles)."
 To choose a fitting interval in a vicinity of the ight curve maximum. we used the [act that the values of fitting parameters are reasonably stable with respect to a reduction of the interval.," To choose a fitting interval in a vicinity of the light curve maximum, we used the fact that the values of fitting parameters are reasonably stable with respect to a reduction of the interval."
" This allowed us to find a preliminary estimate ofthe ""background level” 21.", This allowed us to find a preliminary estimate of the “background level” $A$.
 After that. all the points above st rave been involved in the final tre:iument.," After that, all the points above $A$ have been involved in the final treatment."
 The models £4 and £7 fit the data equally well., The models $F^1$ and $F^3$ fit the data equally well.
 Lot us disetiss the model £7 ancl the role of the second correction., Let us discuss the model $F^2$ and the role of the second correction.
" As ('ompared with the other mocels. he mocdel £7 Icas το à SOLnewhat lower value of X7. but the confidence interval for fer is our times wider than that in the case of the model i)£"". and the source size appears to be 20 per cent more than that within (ther moclels."," As compared with the other models, the model $F^2$ leads to a somewhat lower value of $\chi ^2$, but the confidence interval for $t_C \mbox{ }$ is four times wider than that in the case of the model $F^0$, and the source size appears to be 20 per cent more than that within other models."
 Here. the typical feaures of the seconc correction mentioned in the previous section become apparent.," Here, the typical features of the second correction mentioned in the previous section become apparent."
 The true (letermination of 2 from formula (28)) could allow us to estimate the ratio of the curvature and source radii: wl=PPB. B," The true determination of $D$ from formula \ref{eq21}) ) could allow us to estimate the ratio of the curvature and source radii: $\kappa L = - {2D} \mathord{\left/ {\vphantom {{2D}
B}} \right. \kern-\nulldelimiterspace} B$."
ut the obtained. value of D corresponds to the curyaUre racius Z2.=l/r that is about twice less than the source “size” L., But the obtained value of $D$ corresponds to the curvature radius $R_{c}\equiv 1/\kappa$ that is about twice less than the source “size” $L$.
 So. the condition of smallness of the correction is violated. anc the model £7 is not acceptable.," So, the condition of smallness of the correction is violated, and the model $F^2$ is not acceptable."
 On the other hand. neglecting he second correction in (28)) (Le. the transition to t1C MOel F1 corresponds to 2.zzL.," On the other hand, neglecting the second correction in \ref{eq21}) ) (i.e., the transition to the model $F^1$ ) corresponds to $R_{c}>>L$."
 Our caleulations have shown hat the assignment 2.7L and the introduction of the appropriate additional term into the model £7 have little influence on the estimates of cA.D. and €*.," Our calculations have shown that the assignment $R_{c}>L$ and the introduction of the appropriate additional term into the model $F^1$ have little influence on the estimates of $A,B$, and $C$."
 To sum up. the model £7 fits satisfactorily the data on 1AE in question within the present accuracy. though we cannot rule out competing models and/or competing effects.," To sum up, the model $F^1$ fits satisfactorily the data on HAE in question within the present accuracy, though we cannot rule out competing models and/or competing effects."
 We also analysed LIAL in the light curve of image Q390207aadO305A that happened in L999 (Woziiaketal.2000)., We also analysed HAE in the light curve of image Q2237+0305A that happened in 1999 \citep{wozniak_00}.
. However. the treatment of this ΗΛΙΟ does not allowed us to find statistically significant estimates of the corrections because of lacking the data corresponding to the caustic inner region.," However, the treatment of this HAE does not allowed us to find statistically significant estimates of the corrections because of lacking the data corresponding to the caustic inner region."
1e Norma cluster 66312) has an even Larger velocity ollset.,the Norma cluster 6312) has an even larger velocity offset.
 36hhringer et al. (, Böhhringer et al. (
1996) identified an N-ray. subgroup close the centre of the Norma cluster (see Fig. 6)).,1996) identified an X-ray subgroup close the centre of the Norma cluster (see Fig. \ref{optxray}) ).
 This subgroup (dubbed “Norma minor’) is fairly massive: Tamura et al. (, This subgroup (dubbed `Norma minor') is fairly massive; Tamura et al. (
1998) estimate that the mass of this subgroup could acd up to ~50% to the total mass of the cluster.,1998) estimate that the mass of this subgroup could add up to $\sim$ to the total mass of the cluster.
 A comparison of the S43 MlIZ radio continuum emission olPISS BBIGIO-GOS (Jones MeXdam 1992) with the N-ray contours of this central subgroup (shown in Fig. 11)), A comparison of the 843 MHz radio continuum emission of B1610-608 (Jones McAdam 1992) with the X-ray contours of this central subgroup (shown in Fig. \ref{central}) )
 shows that the racio lobes of 66269 are closely aligned with the X-ray subgroup., shows that the radio lobes of 6269 are closely aligned with the X-ray subgroup.
 The large observed peculiar velocity of the cD galaxy in the Norma cluster is most likely caused by this ongoing merger., The large observed peculiar velocity of the cD galaxy in the Norma cluster is most likely caused by this ongoing merger.
 Aased on the compactness of the X-ray subcluster. sohbhringer et al (," Based on the compactness of the X-ray subcluster, Böhhringer et al. ("
1996) argued that the merger has not progressed very [ar vet. and that most of the main component of the cluster is still undisturbed by the collision.,"1996) argued that the merger has not progressed very far yet, and that most of the main component of the cluster is still undisturbed by the collision."
 This is consistent with simulations of cluster mergers (Pinkney et al., This is consistent with simulations of cluster mergers (Pinkney et al.
 1996). which show that large peculiar velocities can be reproduced. in. the event of large-scale mereers at the time of core-crossing.," 1996), which show that large peculiar velocities can be reproduced in the event of large-scale mergers at the time of core-crossing."
 I£ this merger takes place close to the plane of the sky. it would also explain the non-results of the statistical tests.," If this merger takes place close to the plane of the sky, it would also explain the non-results of the statistical tests."
 These are least sensitive to mergers occuring perpendicular to the line-of-sight., These are least sensitive to mergers occuring perpendicular to the line-of-sight.
 Therefore. the X-ray morphology — in combination with the large peculiar motion of the central cD galaxy strongly suggests à recent or commencing merger at. the core of the cluster.," Therefore, the X-ray morphology – in combination with the large peculiar motion of the central cD galaxy -- strongly suggests a recent or commencing merger at the core of the cluster."
 For the determination of the dynamical mass of the Norma cluster. we have used both the virial theorem (Ayr) and he projected mass estimator (Adpare ). see equations 21 and 22 of Pinkney et al. (," For the determination of the dynamical mass of the Norma cluster, we have used both the virial theorem $M_{\rm VT}$ ) and the projected mass estimator $M_{\rm PME}$ ), see equations 21 and 22 of Pinkney et al. ("
1996).,1996).
 Phe use of the biweight velocity centroid and scale (Beers et al., The use of the bi-weight velocity centroid and scale (Beers et al.
 1990) in the virial theorem (instead of the velocity mean and standard ceviation) Leacks o a more robust mass estimate (Αν)., 1990) in the virial theorem (instead of the velocity mean and standard deviation) leads to a more robust mass estimate $M_{\rm RVT}$ ).
 The latter is more robust against the effects of contamination by the inclusion of possible non-members in the analysis., The latter is more robust against the effects of contamination by the inclusion of possible non-members in the analysis.
" The ojected mass estimator (Bircl 1995). on the other Πα. is sensitive to the presence of (spatiallv-separated) subclusters due to its proportionality to the projected distance between ealaxy 7 and the cluster centroid (A2 ,,;) (see equation 22 in Pinkney et al."," The projected mass estimator (Bird 1995), on the other hand, is sensitive to the presence of (spatially-separated) subclusters due to its proportionality to the projected distance between galaxy $i$ and the cluster centroid $R_{{\perp},i}$ ) (see equation 22 in Pinkney et al."
 1996)., 1996).
 Phe presence of a spatially-scparatec subcluster (e.g. in à premerger configuration) would. result in à svstematic olfset with respect to the cluster centroid: us leads to larger values of Ayo; and thus to a significantly larger mass estimate.," The presence of a spatially-separated subcluster (e.g. in a premerger configuration) would result in a systematic offset with respect to the cluster centroid; this leads to larger values of $R_{{\perp},i}$ and thus to a significantly larger mass estimate."
 For a full discussion of the appropriate use of these dynamical mass estimators we refer to Pinkney et al. (, For a full discussion of the appropriate use of these dynamical mass estimators we refer to Pinkney et al. (
1996) and Bird (1995).,1996) and Bird (1995).
 The three dynamical mass estimates (Adr. Mgyy and Alpe) determined within the three radial limits (using 16 Combined samples of No = 129. 239 and 296 galaxies. respectively) are given in Table 4..," The three dynamical mass estimates $M_{\rm VT}$ , $M_{\rm RVT}$ and $M_{\rm PME}$ ) determined within the three radial limits (using the combined samples of $N$ = 129, 239 and 296 galaxies, respectively) are given in Table \ref{massnorma}."
 On average. Mgr is ~ larger than Ayr.," On average, $M_{\rm RVT}$ is $\sim 5$ larger than $M_{\rm VT}$."
 The projected mass estimate. however. is generally about larger than Myr ancl indicates the presence of a spatially-cistinet subcluster (projected on the plane of the skv) presumably in the carly stages of merging (Pinkney et al.," The projected mass estimate, however, is generally about larger than $M_{\rm VT}$ and indicates the presence of a spatially-distinct subcluster (projected on the plane of the sky) presumably in the early stages of merging (Pinkney et al."
 1996)., 1996).
 This is consistent with our previous indications of subclustering. particularly in the form. of Norma minor (the X-rav subgroup).," This is consistent with our previous indications of subclustering, particularly in the form of Norma minor (the X-ray subgroup)."
 3ohhringer ct al. (, Böhhringer et al. (
1996) and Tamura et al. (,1996) and Tamura et al. (
1998) both eive an estimate of the gravitational mass of the Norma cluster based. on. ROSATLT and ASCA X-ray observations. respectively.,"1998) both give an estimate of the gravitational mass of the Norma cluster based on ROSAT and ASCA X-ray observations, respectively."
 La ‘Table 4.. we list the values (converted. from hog o h.l) of the mass within a specific radius as derive from X-ray observations by Bóhhringer οἱ al. (," In Table \ref{massnorma}, we list the values (converted from $h_{50}^{-1}$ to $h_{73}^{-1}$ ) of the mass within a specific radius as derived from X-ray observations by Böhhringer et al. ("
1996) anc ‘Tamura et al. (,1996) and Tamura et al. (
1998).,1998).
 Both virial mass estimates (Ady-p zux ΑΓ) ave consistent with the mass determined from the X-rav Luminosity of the cluster., Both virial mass estimates $M_{\rm VT}$ and $M_{\rm RVT}$ ) are consistent with the mass determined from the X-ray luminosity of the cluster.
 In the presence of substantia subclustering. as suggested here for the Norma cluster. al dynamical mass estimators could still overestimate the true mass of the cluster. depending on the projection angle of the clustersubcluster merger axis with respect to the line of sight. (Pinkney et al.," In the presence of substantial subclustering, as suggested here for the Norma cluster, all dynamical mass estimators could still overestimate the true mass of the cluster, depending on the projection angle of the cluster–subcluster merger axis with respect to the line of sight (Pinkney et al."
 L996)., 1996).
 This effect is smallest. [or Nyy and Alpyr lor a merger occuring perpencicular to the linc-of-sight., This effect is smallest for $M_{\rm VT}$ and $M_{\rm RVT}$ for a merger occuring perpendicular to the line-of-sight.
 We can therefore safely conclude that the mass of the Norma cluster within the Abell radius corresponds to 1lLIlores AL., We can therefore safely conclude that the mass of the Norma cluster within the Abell radius corresponds to $1-1.1 \times 10^{15} h_{73}^{-1}$ $_{\odot}$.
 This confirms the status of the Norma cluster as the most massive cluster in the Great Attractor., This confirms the status of the Norma cluster as the most massive cluster in the Great Attractor.
 A number of galaxies in the Norma cluster show clirect or indirect evidence of interaction with the intracluster medium (1CMD., A number of galaxies in the Norma cluster show direct or indirect evidence of interaction with the intracluster medium (ICM).
 Here. we will explore these galaxies in some cetail in the light of the preceding discussion.," Here, we will explore these galaxies in some detail in the light of the preceding discussion."
 Recent and observations of Why 6176 (= ESO 137-001) revealed the presence ofa ~70 kpc long X-ray tail pointing away from the cluster centre (Sun et al., Recent and observations of WKK 6176 (= ESO 137-001) revealed the presence ofa $\sim$ 70 kpc long X-ray tail pointing away from the cluster centre (Sun et al.
 2006). suggesting this galaxy is undergoing a significant," 2006), suggesting this galaxy is undergoing a significant"
In the first plot comparing SEIS vs. geas mass fraction. (upper left). we see that high SET ealaxies have low gas fractions. with the converse being true or low SEL galaxies.,"In the first plot comparing SFE vs. gas mass fraction (upper left), we see that high SFE galaxies have low gas fractions, with the converse being true for low SFE galaxies."
 Star-forming and green valley galaxies display a range of gas fractions from to4., Star-forming and green valley galaxies display a range of gas fractions from to.
. We plot SEE vs. (Capper right). and find that the highest SEIS galaxies do not nave the highest specific star formation rates. but. instead »eak at slightlv lower values.," We plot SFE vs. (upper right), and find that the highest SFE galaxies do not have the highest specific star formation rates, but instead peak at slightly lower values."
 This suggests that a high SEE is not driven by a high SER., This suggests that a high SFE is not driven by a high SFR.
 Not surprisingly. the gas raction vs. rrelation (lower right) shows that star-forming and green vallev galaxies reveal correlated gas. fraction. ancl specific star formation rates with the extreme SET galaxies Lying olf of this relation.," Not surprisingly, the gas fraction vs. relation (lower right) shows that star-forming and green valley galaxies reveal correlated gas fraction and specific star formation rates with the extreme SFE galaxies lying off of this relation."
 IL of these plots. taken together. suggest that itis low content. as opposed to an excess in star formation rate. that is responsible for the majority of high SEE galaxies.," All of these plots, taken together, suggest that it is low content, as opposed to an excess in star formation rate, that is responsible for the majority of high SFE galaxies."
 Figures S--10 can be used to isolate galaxies that show gas excess or deficiency when compared to their current star formation rates. useful for determining what causes high or low SEI in galaxies.," Figures \ref{Fig:ssfr_sfe_mass}- \ref{Fig:ssfr_gfrac_sfe} can be used to isolate galaxies that show gas excess or deficiency when compared to their current star formation rates, useful for determining what causes high or low SFE in galaxies."
 Figure Ll attempts to combine this diagnostic capability in one diagram. by plotting the olfset of a galaxv's star formation rate and mass relative to the average value for a given stellar mass.," Figure \ref{Fig:delta_ssfr_gfrac} attempts to combine this diagnostic capability in one diagram, by plotting the offset of a galaxy's star formation rate and mass relative to the average value for a given stellar mass."
" We use data from this paper and Paper LE and. perform linear regression fits to the mean SER and Mg; vs For the nmiuniasses we use the mean values for which the masses of non-detections have been set to the upper limit. (Paper L. ""Table 4. column 1)."," We use data from this paper and Paper I and perform linear regression fits to the mean SFR and $_{\hi}$ vs For the masses we use the mean values for which the masses of non-detections have been set to the upper limit (Paper I, Table 4, column 1)."
 We find that and and. define In this delta gas-fraction vs. celta specific star formation rate plot four quadrants are celineated., We find that and and define In this delta gas-fraction vs. delta specific star formation rate plot four quadrants are delineated.
" The top right quadrant contains those galaxies with high and/Ad,.. in other words galaxies that are eas-rich ancl actively star-forming."," The top right quadrant contains those galaxies with high and, in other words galaxies that are gas-rich and actively star-forming."
 The majority of the CASS detections fall here along a line where eas fraction excess equals the cxeess which corresponds approximately to a line of constant SEL., The majority of the GASS detections fall here along a line where gas fraction excess equals the excess which corresponds approximately to a line of constant SFE.
 This quadrant is also likely to include eas-rich mergers. starbursts and high surface brightness galaxies with extended: UV-clisks Clhilkeretal.2007).," This quadrant is also likely to include gas-rich mergers, starbursts and high surface brightness galaxies with extended UV-disks \citep{Thilker2007}."
.. Phe upper left. quadrant: contains eas rich galaxies with lower than average star formation rates. including the low SEE galaxies highlighted above.," The upper left quadrant contains gas rich galaxies with lower than average star formation rates, including the low SFE galaxies highlighted above."
 Low surface brightness galaxies anc galaxies that have recently accreted eas might be found here., Low surface brightness galaxies and galaxies that have recently accreted gas might be found here.
 The majority of galaxies identified as low SPE are forming stars below average rates. Consistent with their being passive galaxies with large reservoirs of gas.," The majority of galaxies identified as low SFE are forming stars below average rates, consistent with their being passive galaxies with large reservoirs of gas."
 These galaxies clo not appear to have extremely high gas masses and average star formation rates (whieh may be more typical of galaxies with lower stellar mass)., These galaxies do not appear to have extremely high gas masses and average star formation rates (which may be more typical of galaxies with lower stellar mass).
 GASS3505. an unusually gas-rich galaxy cliscussecl in Paper L belongs to this class of object and is indicated with a red star in the figure.," GASS3505, an unusually gas-rich galaxy discussed in Paper I, belongs to this class of object and is indicated with a red star in the figure."
 The lower left quadrant contains gas-poor galaxies with below average SER., The lower left quadrant contains gas-poor galaxies with below average SFR.
 Passive galaxies are found in this «quadrant., Passive galaxies are found in this quadrant.
 Finally the lower right. quadrant. contains high SEE galaxies that are relatively gas poor but with above average SER., Finally the lower right quadrant contains high SFE galaxies that are relatively gas poor but with above average SFR.
 On this plot it becomes apparent that these galaxies are distributed across the quadrant ancl appear to have either. high star formation rates when compared. to their (typical) gas mass. or they have lower than average gas mass but typical star formation rates.," On this plot it becomes apparent that these galaxies are distributed across the quadrant and appear to have either high star formation rates when compared to their (typical) gas mass, or they have lower than average gas mass but typical star formation rates."
 Ehe most gas deficient in this category have possibly very recently experienced. a process that disrupted. gas Low and/or removed gas. such as starvation or stripping.," The most gas deficient in this category have possibly very recently experienced a process that disrupted gas flow and/or removed gas, such as starvation or stripping."
 The latter scenario is likely to produce eas deficiencies even lower than we have observed., The latter scenario is likely to produce gas deficiencies even lower than we have observed.
 Later stages of such galaxies might also be found in the third quadrant below or near the line of constant SEE., Later stages of such galaxies might also be found in the third quadrant below or near the line of constant SFE.
 CASS 7050. identified in Paper | as an unusual gas-poor galaxy. appears to belong to this category and is indicated with a erecn star in the figure.," GASS 7050, identified in Paper I as an unusual gas-poor galaxy, appears to belong to this category and is indicated with a green star in the figure."
 Lastly. we return to the question of whether or not the scatter around our mean scaling relations suggests a diversity. in the processes that trigeer and quench star formation in the CLASS sample.," Lastly, we return to the question of whether or not the scatter around our mean scaling relations suggests a diversity in the processes that trigger and quench star formation in the GASS sample."
 Without question the CLASS sample does not show a tight distribution in SEIs. and it is tempting to reconsider internal quenching mechanisms for at least some fraction of svstems.," Without question the GASS sample does not show a tight distribution in SFEs, and it is tempting to reconsider internal quenching mechanisms for at least some fraction of systems."
 Alternately. our results may rule out scenarios where the inflow of material is occurring as a steady Blow (or drizzle) onto the galaxy.," Alternately, our results may rule out scenarios where the inflow of material is occurring as a steady flow (or drizzle) onto the galaxy."
 Instead. the scatter may be an indication that accretion of lis episodic. with infalline eas arriving in larger discrete chunks.," Instead, the scatter may be an indication that accretion of is episodic, with infalling gas arriving in larger discrete chunks."
 It may suggest. that large scale. processes. play a role in regulating the growth and evolution of gas and star formation in galaxies., It may suggest that large scale processes play a role in regulating the growth and evolution of gas and star formation in galaxies.
 Future work is being planned to investigate the large- and. small-scale environment of CLASS. galaxies. and in particular the outliers. which should. reveal whether enviromental processes are driving this evolution.," Future work is being planned to investigate the large- and small-scale environment of GASS galaxies, and in particular the outliers, which should reveal whether enviromental processes are driving this evolution."
 Some galaxies with high star formation rates may be undergoing a merger or interaction that is driving up the star formation rate in these galaxies and it will also be interesting to investigate the connection between signs of. interaction. merging or other disturbance and location on this diagram.," Some galaxies with high star formation rates may be undergoing a merger or interaction that is driving up the star formation rate in these galaxies and it will also be interesting to investigate the connection between signs of interaction, merging or other disturbance and location on this diagram."
" We use measurements of the ccontent. stellar mass and star formation rates in ~190 massive galaxies with M,1027AL... obtained. from the Galex Arecibo SDSS survey described in Paper LE (Catinellaetal.2010) to explore the elobal scaling relations associated with the ratio. which we call the based star formation elliciency (SEI)."," We use measurements of the content, stellar mass and star formation rates in $\sim 190$ massive galaxies with $M_\star >10^{10}$, obtained from the Galex Arecibo SDSS survey described in Paper I \citep{Catinella2010} to explore the global scaling relations associated with the ratio, which we call the based star formation efficiency (SFE)."
 We find that:, We find that:
agree well out to a radius of ~30 kpc.,agree well out to a radius of $\sim$ 30 kpc.
 The stack color values are plotted in the observed frame and are matched to the ? rest frame colors by adding a constant., The stack color values are plotted in the observed frame and are matched to the \cite{peletier_ccd_1990} rest frame colors by adding a constant.
 We note that outside of about 20 kpc the local sample is composed of only a few galaxies where sufficient depth was achievable., We note that outside of about 20 kpc the local sample is composed of only a few galaxies where sufficient depth was achievable.
" With the advantage of our deep stacks, however, we are able to study the colors of LRGs out to more than eight effective radii."," With the advantage of our deep stacks, however, we are able to study the colors of LRGs out to more than eight effective radii."
" This is the first time that the colors of massive galaxies are observed at such large radius, providing a new insight on the stars that are found in the outskirts of massive galaxies."," This is the first time that the colors of massive galaxies are observed at such large radius, providing a new insight on the stars that are found in the outskirts of massive galaxies."
" Indeed, the color profiles of LRGs change trend beyond the limit of nearby galaxy observations."," Indeed, the color profiles of LRGs change trend beyond the limit of nearby galaxy observations."
" Despite initially getting bluer in the inner 40 kpc of both plotted colors, the profiles quickly flatten out to a relatively constant value."," Despite initially getting bluer in the inner 40 kpc of both plotted colors, the profiles quickly flatten out to a relatively constant value."
" 'The outskirts of LRGs are then roughly 0.15 dex and 0.2 dex bluer than their centers in the r—4 and g—z colors, respectively."," The outskirts of LRGs are then roughly 0.15 dex and 0.2 dex bluer than their centers in the $r-i$ and $g-z$ colors, respectively."
 The color profile is different from that observed by Z05 at r>20 kpc as Z05 find that BCGs become very red at large radii., The color profile is different from that observed by Z05 at $>$ 20 kpc as Z05 find that BCGs become very red at large radii.
" As we show in the appendix, the Z05 color profile may be severely affected by PSF effects at all radii, including the stack outer parts."," As we show in the appendix, the Z05 color profile may be severely affected by PSF effects at all radii, including the stack outer parts."
" Therefore, the different radial color profiles do not necessarily imply that LRGs and BCGs are fundamentally different objects."," Therefore, the different radial color profiles do not necessarily imply that LRGs and BCGs are fundamentally different objects."
 We note that the g-r color gradient found by Z05 has a similar slope as the g-z profile presented in figure 8.., We note that the g-r color gradient found by Z05 has a similar slope as the g-z profile presented in figure \ref{fig:pelet}.
 The g-r profile of Z05 may suffer less from PSF effects than their r-i profile., The g-r profile of Z05 may suffer less from PSF effects than their r-i profile.
" The first and foremost result that arises from this study is that faint, gravitationally bound stellar light can be traced in massive elliptical galaxies out to a radius of 100 kpc."," The first and foremost result that arises from this study is that faint, gravitationally bound stellar light can be traced in massive elliptical galaxies out to a radius of 100 kpc."
 By stacking a large number of faint galaxy images we detect light at such distance from the centers of massive galaxies with good confidence., By stacking a large number of faint galaxy images we detect light at such distance from the centers of massive galaxies with good confidence.
 In figure 10 we show that the total accumulated light at 20«r/kpc«100 is negligible and accounts for roughly of the overall flux in the stack., In figure \ref{fig:lfrac} we show that the total accumulated light at $< r/\rm kpc <$ 100 is non-negligible and accounts for roughly of the overall flux in the stack.
" In fact, more than of the stack light can be detected at very large radii outside of r =100 kpc, or more than 8 effective radii."," In fact, more than of the stack light can be detected at very large radii outside of $r=$ 100 kpc, or more than 8 effective radii."
" This is especially interesting in light of recent studies that find compact massive galaxies at z~2, exhibiting effective radii of only 1 kpc (e.g.,???).."," This is especially interesting in light of recent studies that find compact massive galaxies at $z\sim$ 2, exhibiting effective radii of only 1 kpc \citep[e.g.,][]{daddi_passively_2005,
 trujillo_size_2006, van_dokkum_confirmation_2008}."
" The size growth of these objects is evidently rapid, expanding the physical scale of the galaxy by a factor of at least 5 in less than 10 Gyr."," The size growth of these objects is evidently rapid, expanding the physical scale of the galaxy by a factor of at least 5 in less than 10 Gyr."
" Unfortunately, the physical growth mechanism cannot be directly observed in the LRG stacks as any signal from individual galaxies is smoothed and averaged over the entire sample."," Unfortunately, the physical growth mechanism cannot be directly observed in the LRG stacks as any signal from individual galaxies is smoothed and averaged over the entire sample."
" Nevertheless, the lack of a clear change in the stack light profile slopes out to 100 kpc suggests that the observed light in the outskirts of LRGs is physically associated to the galaxies and their inner parts."," Nevertheless, the lack of a clear change in the stack light profile slopes out to 100 kpc suggests that the observed light in the outskirts of LRGs is physically associated to the galaxies and their inner parts."
 Further evidence for this comes from the relatively radially-independent ellipticity profiles which vary only slightly out to 100 kpc., Further evidence for this comes from the relatively radially-independent ellipticity profiles which vary only slightly out to 100 kpc.
" Any other light sources, such as background contamination or residual PSF scattering would be uncorrelated with the position angle of the LRG as measured in individual SDSS images, resulting in a circular light distribution."," Any other light sources, such as background contamination or residual PSF scattering would be uncorrelated with the position angle of the LRG as measured in individual SDSS images, resulting in a circular light distribution."
" Outside of roughly 100 kpc the light profiles of the g,r,i and z-band stacks depart from the simple Sérrsic model profile and exhibit excess light (figure 6))."," Outside of roughly 100 kpc the light profiles of the g,r,i and z-band stacks depart from the simple Sérrsic model profile and exhibit excess light (figure \ref{fig:allprofs}) )."
 This departure from a simple model is observed here for the first time in LRGs and it shows that stars at the extreme outskirts of massive galaxies follow a different gravitational potential than stars in the inner parts., This departure from a simple model is observed here for the first time in LRGs and it shows that stars at the extreme outskirts of massive galaxies follow a different gravitational potential than stars in the inner parts.
 It is known that the potential at these radii is dominated by the properties, It is known that the potential at these radii is dominated by the properties
anv corresponding svstem should be visible in nearby eroups and isn't there 2004).,any corresponding system should be visible in nearby groups and isn't there \citep{PB04}.
. The implications of our survey limit £o bevond simple numbers of galaxies., The implications of our survey limit go beyond simple numbers of galaxies.
 identilv a possible incompleteness in the census of Milkv. Wavy satellites bevond about. 110. κρο," \citet{WG04}
 identify a possible incompleteness in the census of Milky Way satellites beyond about 110 kpc."
 Within that distance (he faintest known galaxies can be lound through identifving their individual stars., Within that distance the faintest known galaxies can be found through identifying their individual stars.
 Devond. if the limit for detection of diffuse objects is 24-25 magnitudes per square arc second (their source for this figure aud (he applicable waveband are nol given). some unknown fraction of satellites would not be seen.," Beyond, if the limit for detection of diffuse objects is 24-25 magnitudes per square arc second (their source for this figure and the applicable waveband are not given), some unknown fraction of satellites would not be seen."
 If this incompleteness exists. the radial distribution of. Milkv. Wav satellites could be the same as that of M31 satellites (and simulations).," If this incompleteness exists, the radial distribution of Milky Way satellites could be the same as that of M31 satellites (and simulations)."
 Now. however. with a limit of somewhere bevond 25. we can sav that (his possible incompleteness is not there.," Now, however, with a limit of somewhere beyond 25, we can say that this possible incompleteness is not there."
 We were sensitive enough to detect all but a handful of known Local Group cdwarls. and those are more or less evenly divided between AI31 and the Milky Way.," We were sensitive enough to detect all but a handful of known Local Group dwarfs, and those are more or less evenly divided between M31 and the Milky Way."
 The radial distribution of Milkv. Way satellites different. [rom that of M31 satellites and simulations., The radial distribution of Milky Way satellites different from that of M31 satellites and simulations.
 As this paper was in the relereeing stage several new objects even smaller and fainter than Andromeda IX were reported (Zuckerοἱal.200Ga:Delokuroval. 200G6b.¢).," As this paper was in the refereeing stage several new objects even smaller and fainter than Andromeda IX were reported \citep{Z06a, B06, Z06b, Z06c}."
. Il is too early vel to draw any conclusions about their implications for the missing satellite problem., It is too early yet to draw any conclusions about their implications for the missing satellite problem.
 If they are indeed (he first of many (sav. LOO or more) galaxiesand (heir cark-matter masses are of (he same order as Chose of much brighter satellites. the problem is solved (or al least. transformed into a star-ormation puzzle).," If they are indeed the first of many (say, 100 or more) galaxies their dark-matter masses are of the same order as those of much brighter satellites, the problem is solved (or at least transformed into a star-formation puzzle)."
 If not. the problem persists.," If not, the problem persists."
 This work has benelited greatly [vom a wide variety of sources., This work has benefited greatly from a wide variety of sources.
 Partial landing support was received [rom the Intitute of Astronomy of the University of Cambridge. the Physics Department of the U. S. Naval Academy. the European Southern Observatory. and two American Astronomical Society Small Research Grants. (which last. finds originated. with NASA).," Partial funding support was received from the Intitute of Astronomy of the University of Cambridge, the Physics Department of the U. S. Naval Academy, the European Southern Observatory, and two American Astronomical Society Small Research Grants (which last funds originated with NASA)."
 G. Ik. T. IL. asknowledges financial support [rom the Chilean FONDECYT erant 1990442., G. K. T. H. asknowledges financial support from the Chilean FONDECYT grant 1990442.
" Extensive use has been made of the NASA Extragalactie Database (NED). which is operated by (he Jet Propulsion Laboratory, California Institute of Technology. wider contract with the National Aeronanties and Space Administration: and of the SIMDAD database. operated at Centre de Donnéees Astronomiques de Strasbourg. France."," Extensive use has been made of the NASA Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration; and of the SIMBAD database, operated at Centre de Donnéees Astronomiques de Strasbourg, France."
 The results are based on observations made with the Isaac Newton Telescope. operated on the island of La Palma bv the Isaac Newton Group in the Spanish Observatorio del loque de los Aluchachos of the Instituto de Astroffssica de Canarias: and al Cerro Tololo Inter-American Observatory (CTIO) and Ixitt Peak National Observatory (INPNO).," The results are based on observations made with the Isaac Newton Telescope, operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de sica de Canarias; and at Cerro Tololo Inter-American Observatory (CTIO) and Kitt Peak National Observatory (KPNO)."
 Both CTIO and IKRPNO are operated by the Association of Universities for Research in Astronomy. Ine. (AURA). under a cooperative agreement with the National Science Foundation. as the National Optical," Both CTIO and KPNO are operated by the Association of Universities for Research in Astronomy, Inc. (AURA), under a cooperative agreement with the National Science Foundation, as the National Optical"
ofile also indicate Hat inner density. profiles (Bacmann et al.,profile – also indicate flat inner density profiles (Bacmann et al.
 2000)., 2000).
 Additionally. mic-infrared observations suggest hat bevond ~ 10.000 AU the density. profile may steepen on24 CAbergel et al.," Additionally, mid-infrared observations suggest that beyond $\sim$ 10,000 AU the density profile may steepen to $\eta \ga 4\,$ (Abergel et al."
 1996. Bacmann ct al.," 1996, Bacmann et al."
 2000)., 2000).
 This steepening of the outer envelope appears to oceur at smaller racii in relatively close-packed protoclusters like p Ophiuchi han in regions of distributed star formation such as Taurus: viz., This steepening of the outer envelope appears to occur at smaller radii in relatively close-packed protoclusters like $\rho$ Ophiuchi than in regions of distributed star formation such as Taurus; viz.
 at  3.000 AU in p Ophiuchi. and at ~ 15.000 AU in Taurus (Motte André 2001).," at $\sim$ 3,000 AU in $\rho$ Ophiuchi, and at $\sim$ 15,000 AU in Taurus (Motte André 2001)."
 The asvnunetric self-absorbed. molecular-line profiles observed. in some prestellar cores. (Tafalla ct al., The asymmetric self-absorbed molecular-line profiles observed in some prestellar cores (Tafalla et al.
 1998. Williams et al.," 1998, Williams et al."
 1999. Lec. Myers Tafalla 1999. Gregersen Evans 2000) suggest that they are indeed already collapsing.," 1999, Lee, Myers Tafalla 1999, Gregersen Evans 2000) suggest that they are indeed already collapsing."
 In particular. the detailed analyses of LISA by Fafalla et al. (," In particular, the detailed analyses of L1544 by Tafalla et al. ("
1998). and Williams et al. (,1998) and Williams et al. (
1999) imply that the inner parts of the core are relatively stationary. and an approximately uniform velocity field has been established in the outer lavers.,"1999) imply that the inner parts of the core are relatively stationary, and an approximately uniform velocity field has been established in the outer layers."
 This is very reminiscent ofthe velocity fields which are set up by inward-propagating compression waves in similarity. solutions for contracting isothermal spheres (Whitworth Summers 1985)., This is very reminiscent of the velocity fields which are set up by inward-propagating compression waves in similarity solutions for contracting isothermal spheres (Whitworth Summers 1985).
 It is this apparent similarity which we explore in the present paper., It is this apparent similarity which we explore in the present paper.
 The prestellar phase terminates as soon as a star-like object forms at the centre of the dense core. and the core hen becomes a Class 0 protostar.," The prestellar phase terminates as soon as a star-like object forms at the centre of the dense core, and the core then becomes a Class 0 protostar."
 Conceptually. the Class ) protostellar phase terminates once the extended envelope contains less than half the total mass of the original core. ic. more than hall the mass is in the central star-like object(s) jus attendant dise(s) (Anedré.. Ward-Phomipson Barsony 1993. 2000): the Class L phase then begins (Lada Wilking 1984. Lada 1987. André Montmerle 1994).," Conceptually, the Class 0 protostellar phase terminates once the extended envelope contains less than half the total mass of the original core, i.e. more than half the mass is in the central star-like object(s) plus attendant disc(s) (André,, Ward-Thompson Barsony 1993, 2000); the Class I phase then begins (Lada Wilking 1984, Lada 1987, André Montmerle 1994)."
 The Class 1 otostellar phase terminates when most of the envelope has »oen accreted or dissipated. revealing a classical T Tauri star (C'PTS) accreting from a residual circumstellar disc. i.e. à Class Hl object.," The Class I protostellar phase terminates when most of the envelope has been accreted or dissipated, revealing a classical T Tauri star (CTTS) accreting from a residual circumstellar disc, i.e. a Class II object."
 Once the inner disc has been dissipated. the accretion rate is greatly. reduced and the source becomes a weak-lined E Tauri star (WETS) or Class HE object.," Once the inner disc has been dissipated, the accretion rate is greatly reduced and the source becomes a weak-lined T Tauri star (WTTS) or Class III object."
 On the basis of statistical arguments (ie. source numbers and a presumed constant star formation rate) it is inferred. (Beichmann et al., On the basis of statistical arguments (i.e. source numbers and a presumed constant star formation rate) it is inferred (Beichmann et al.
" 1986. André... Ward-Thompson Barsony 2000) that the prestellar phase lasts. tvpically. 10"" to LOS years."," 1986, André,, Ward-Thompson Barsony 2000) that the prestellar phase lasts, typically, $10^6$ to $10^7$ years."
 In close-packed protoclusters like. p Ophiuchi. the Class 0. phase appears to last a lew times 101 vears. and ids. characterized. by powerful] collimatect. outllows indicative of rapid accretion. &LO7M.vear ( Bontempset al.," In close-packed protoclusters like $\rho$ Ophiuchi, the Class 0 phase appears to last a few times $10^4$ years, and is characterized by powerful collimated outflows indicative of rapid accretion, $\ga 10^{-5} \, M_\odot \, \mbox{year}^{-1}$ (Bontemps et al."
 1996)., 1996).
" En distributed star formation regions like “Taurus. the duration. of the Class 0. phase appears to be longer. 107"" vears. and the accretion rates. lower. zo2010""AL.vear (Motte André POOL)."," In distributed star formation regions like Taurus, the duration of the Class 0 phase appears to be longer, $\sim 10^5$ years, and the accretion rates lower, $\ga 2 \times 
10^{-6} \, M_\odot \, \mbox{year}^{-1}$ (Motte André 2001)."
 The Class ] phase appears to last ~2( vears (e.g. Greene et al., The Class I phase appears to last $\sim 2 \times 10^5$ years (e.g. Greene et al.
" 1994. Ixenvon Lartmann 1995) and is characterized bv slower accretion. E10""AZ.vear and weaker loss collimated outflows (Bontemps ct al."," 1994, Kenyon Hartmann 1995) and is characterized by slower accretion, $\la 10^{-6} \, M_\odot \, \mbox{year}^{-1}$, and weaker less collimated outflows (Bontemps et al."
 1996)., 1996).
 Together. the Class HE ancl Class HE phases appear to last <10' vears. but seemingly the transition from Class HE (CTS) to Class LLL (NETS) is very short and can occur at any time: there are both very voung W'EPTSs (close to the birthline. on the right of the Hertzsprung-Itussell Diagram) and very ok οος (approaching the Main Sequence on the Left. of the Lertzsprung-Russell Diagram) (e.g. Stabler Walter 1993).," Together, the Class II and Class III phases appear to last $\ga 10^7$ years, but seemingly the transition from Class II (CTTS) to Class III (WTTS) is very short and can occur at any time; there are both very young WTTSs (close to the birthline, on the right of the Hertzsprung-Russell Diagram) and very old CTTSs (approaching the Main Sequence on the left of the Hertzsprung-Russell Diagram) (e.g. Stahler Walter 1993)."
 These observationally inferred features of prestellar a wotostellar evolution combine to form a reasonably coheren xeture., These observationally inferred features of prestellar and protostellar evolution combine to form a reasonably coherent picture.
 However. especially in star-forming clusters. this xceture is dillicult to reconcile with the standard. theory of Shu. Adams Lizano (1987) based on the inside-ou collapse of a singular isothermal sphere.," However, especially in star-forming clusters, this picture is difficult to reconcile with the standard theory of Shu, Adams Lizano (1987) based on the inside-out collapse of a singular isothermal sphere."
 Ehe strength of he standard theory is that it makes specific quantitative owedietions. but some of these predictions are clillicul o reconcile with observation. (," The strength of the standard theory is that it makes specific quantitative predictions, but some of these predictions are difficult to reconcile with observation. ("
i) ln the standard: moclel. oestellar cores should. be stronely centrally condensed. n=odínp]ídinr|~2 and voung protostars shoul o»* somewhat less centrally condensed. 5g~3/2.,"i) In the standard model, prestellar cores should be strongly centrally condensed, $\eta \equiv - 
d \ell n [\rho] / d \ell n [r] \sim 2$ and young protostars should be somewhat less centrally condensed, $\eta \sim 
3/2\,$."
 By contrast. observations sugeest that prestellar cores have rather flat central density profiles. 4X1. and Class 0 ootostars have steeper ones. (," By contrast, observations suggest that prestellar cores have rather flat central density profiles, $\eta \la 1$, and Class 0 protostars have steeper ones. ("
ii) The standard. theory owediets that prestcllar cores are static. and that. inwarel motions only develop during the protostellar phase and are initially confined. to the central regions.,"ii) The standard theory predicts that prestellar cores are static, and that inward motions only develop during the protostellar phase and are initially confined to the central regions."
 In. contrast. the observations imply that inward motions already exist during 1e prestellar phase. and that initially they are more rapid in the outer regions. (," In contrast, the observations imply that inward motions already exist during the prestellar phase, and that initially they are more rapid in the outer regions. ("
iii) The standard theory predicts that the accretion rate is roughly constant. and hence the Class O and. Class [lifetimes should be comparable.,"iii) The standard theory predicts that the accretion rate is roughly constant, and hence the Class 0 and Class I lifetimes should be comparable."
 Observations detect many more Class { sources than Class 0 ones although this statistical result may be compromised by the difficulty of measuring precisely when an object has accreted half the total mass of its initial core. (, Observations detect many more Class I sources than Class 0 ones – although this statistical result may be compromised by the difficulty of measuring precisely when an object has accreted half the total mass of its initial core. (
iv) Actelitionally. the stancard theory assumes initial conditions which are unlikely to arise in nature. because they are both singular and unstable. (,"iv) Additionally, the standard theory assumes initial conditions which are unlikely to arise in nature, because they are both singular and unstable. ("
v) Also the standard theory has a strong inbuilt’ pre-disposition to the formation of single stars in stark contrast with the high proportion of binaries ancl higher multiples observed in voung star-formation regions.,v) Also the standard theory has a strong inbuilt pre-disposition to the formation of single stars – in stark contrast with the high proportion of binaries and higher multiples observed in young star-formation regions.
" Therefore. although the standard theory. may provide a eood. zeroth-order description of protostellar collapse in sparse. quiescent star-formation regions like ""Taurus (cf."," Therefore, although the standard theory may provide a good, zeroth-order description of protostellar collapse in sparse, quiescent star-formation regions like Taurus (cf."
 Alotte André 2001). it appears that more cvnanmical models are required to understand close-packed regions like p Ophiuchi.," Motte André 2001), it appears that more dynamical models are required to understand close-packed regions like $\rho$ Ophiuchi."
 Foster Chevalier (1993). have explored. how the collapse of an isothermal core develops. if one. abandons the assumption of singularity., Foster Chevalier (1993) have explored how the collapse of an isothermal core develops if one abandons the assumption of singularity.
 Pheir simulations start from. non-singular isothermal equilibria. te. Bonnor-Ebert spheres (e.g. Bonnor 1956). and collapse is then triggered by discontinuously increasing the density.," Their simulations start from non-singular isothermal equilibria, i.e. Bonnor-Ebert spheres (e.g. Bonnor 1956), and collapse is then triggered by discontinuously increasing the density."
 The ensuing collapse results in supersonic inllow velocities in the centre of the core at the end of the prestellar phase. and in a marked decline of the accretion rate from the Class 0 to the Class 1 phase (sce Henriksen. Aadré.. Bontemps 1997 and. Whitworth Ward-Phomipson 2001 for simpler. pressure-free descriptions of this evolution)," The ensuing collapse results in supersonic inflow velocities in the centre of the core at the end of the prestellar phase, and in a marked decline of the accretion rate from the Class 0 to the Class I phase (see Henriksen, André,, Bontemps 1997 and Whitworth Ward-Thompson 2001 for simpler, pressure-free descriptions of this evolution)."
 In this paper we pursue the consequences of further., In this paper we pursue the consequences of non-singularity further.
 Our simulations also start [rom non- isothermal equilibrium cores. but. collapse is then," Our simulations also start from non-singular isothermal equilibrium cores, but collapse is then"
unweighted LAD fits and found that the results for WLSQ anc unweighted LAD fits were very similar. while the weighted LAD ft introduced. a small bias in the mean.,"unweighted LAD fits and found that the results for WLSQ and unweighted LAD fits were very similar, while the weighted LAD fit introduced a small bias in the mean."
 These results are shown in Table 4.., These results are shown in Table \ref{tbl:est}.
 We suspect that the bias occurs because any LAL) fit includes. a ‘cead-zone feature. where a range of parameter estimates give the same minimum absolute deviation.," We suspect that the bias occurs because any LAD fit includes a `dead-zone' feature, where a range of parameter estimates give the same minimum absolute deviation."
 This dead. zone is negligible when the number of estimates is large. but can be significant otherwise.," This dead zone is negligible when the number of estimates is large, but can be significant otherwise."
 Since our A estimates are dominated by a small number of AZ measurements and the results of the dillerent estimators are we chose the more standard. WLSQ it in calculating οί. similar.," Since our $A^2$ estimates are dominated by a small number of $A^2_k$ measurements and the results of the different estimators are similar, we chose the more standard WLSQ fit in calculating $\hat{A^2}$."
Although the WLSQ estimator is not maximum likelihood. it is apparently more robust in our xuticular case.," Although the WLSQ estimator is not maximum likelihood, it is apparently more robust in our particular case."
 Estimation of A is also largely independent of changes o the method of spectral analysis., Estimation of $A^2$ is also largely independent of changes to the method of spectral analysis.
 We experimented with reducing the white noise in the residuals by. smoothing each time series over a 60-dav. period before commencing he spectral analysis., We experimented with reducing the white noise in the residuals by smoothing each time series over a 60-day period before commencing the spectral analysis.
 We also tested interpolation using a constrained cubic spline of cach smoothed time series onto a ld-dav grid common to all pulsars before the spectral analysis., We also tested interpolation using a constrained cubic spline of each smoothed time series onto a 14-day grid common to all pulsars before the spectral analysis.
 “Phe results of these different approaches are given in ‘Table 5.., The results of these different approaches are given in Table \ref{tbl:specanal}.
 Since there was no statistically significant difference between the cdilferent approaches. for simplicity we elected not to smooth or interpolate the residuals.," Since there was no statistically significant difference between the different approaches, for simplicity we elected not to smooth or interpolate the residuals."
 For their detection statistic. ? caleulate the normalised cross correlation between the timing residuals of cach pulsar pair.," For their detection statistic, \citet{jhlm05} calculate the normalised cross correlation between the timing residuals of each pulsar pair."
 Phey optimise the S/N ratio using a filter designed to whiten the residuals before correlation., They optimise the S/N ratio using a filter designed to whiten the residuals before correlation.
" For a simulation of the 20 PI""EA pulsars. this approach increased the maximum achievable. detection significance for a GWD from 3a to 13m."," For a simulation of the 20 PPTA pulsars, this approach increased the maximum achievable detection significance for a GWB from $\sigma$ to $\sigma$."
 Llowever. their filter cannot be applied to real pulsar timing observations without modification.," However, their filter cannot be applied to real pulsar timing observations without modification."
 We investigated the effect of such a filter by performing simulations of the Verbiest οἱ al. (, We investigated the effect of such a filter by performing simulations of the Verbiest et al. (
2008. 2009) residuals where cach simulation included a signal from a CWB with £23.1077.,"2008, 2009) residuals where each simulation included a signal from a GWB with $A \gtrsim 3\e{-15}$."
 In the frequeney domain. the filter takes the form of a weighting factor. so we optimised this weighting factor to match the large input. GWD amplitude.," In the frequency domain, the filter takes the form of a weighting factor, so we optimised this weighting factor to match the large input GWB amplitude."
 We found that this method did not improve the S/N ratio. and we traced this undoer-performance to the problem of spectral leakage from the lowest [frequencies to the higher frequencies.," We found that this method did not improve the S/N ratio, and we traced this under-performance to the problem of spectral leakage from the lowest frequencies to the higher frequencies."
 We found that the first few cross-spectral estimates. which make the largest contribution to our detection statistic. were all more than 90 per cent correlated with the lowest spectral estimate (ic. ab frequeney f£.= dla). meaning that re-weighting cannot change the overall S/N ratio.," We found that the first few cross-spectral estimates, which make the largest contribution to our detection statistic, were all more than 90 per cent correlated with the lowest spectral estimate (i.e., at frequency $f = 1/\tol$ ), meaning that re-weighting cannot change the overall S/N ratio."
 The spectral leakage is particularly significant because of the irregular sampling ancl variable To uncertainties in these observations., The spectral leakage is particularly significant because of the irregular sampling and variable ToA uncertainties in these observations.
 We expect that an improved. spectral analysis technique (e.g..7) will eliminate the spectral leakage and enable us to take advantage of more degrees of when the CWD signal is larger than the noise.," We expect that an improved spectral analysis technique \citep[e.g.,][]{2010Coles} will eliminate the spectral leakage and enable us to take advantage of more degrees of when the GWB signal is larger than the noise."
 ‘The time series we consider in this paper have widely varying tme spans. which has not been a feature of most PTA analyses to date.," The time series we consider in this paper have widely varying time spans, which has not been a feature of most PTA analyses to date."
 As part of the pulsar parameter estimation. we fit for the pulse period and its derivative over the full duration of cach time series.," As part of the pulsar parameter estimation, we fit for the pulse period and its derivative over the full duration of each time series."
 Originally. we then computed," Originally, we then computed"
"observed quantity Ly. all other parameters are derived: 1.06xLO?<z41.3510? s. the of electrons: No =:1.0212x10%.the number densityof− electrons9.2 ney =4.1323x10!cmnumber and the thermal energy. £4 =0.7pux10erg,","observed quantity $L_1$, all other parameters are derived; $1.06 \times 10^2 < \tau_{\rm exp} < 1.35 \times 10^2$ s, the number of electrons $N_{\rm e} = 1.0^{+2.1}_{-0.5} \times 10^{45}$, the numberdensity of electrons $n_{\rm e1} = 4.1^{+9.2}_{-3.7} \times 10^{13}$ $^{-3}$ , and the thermal energy $E_1 = 0.7^{+13.7}_{-0.6}\times
10^{38}$ erg."
 Assuming that the bullet is moving at 0.26 c. the kinetic energy of the bullet is estimated to be 0.6L2x1011 erg.," Assuming that the bullet is moving at 0.26 c, the kinetic energy of the bullet is estimated to be $0.6^{+1.2}_{-0.3}\times
10^{41}$ erg."
 The assumption is Consistent wilh the observed rise time of 20 s. which is naturally explained by the time in which a bullet appears Irom a nozzle.," The assumption is consistent with the observed rise time of 20 s, which is naturally explained by the time in which a bullet appears from a nozzle."
 Eq. (5)), Eq. \ref{eq:sol}) )
 is written with these estimates as where errors are omitted., is written with these estimates as where errors are omitted.
 since a shot and (he unmocdulated component coexist. the bullets and the continuously emanaling jet may coexist.," Since a shot and the unmodulated component coexist, the bullets and the continuously emanating jet may coexist."
 In that case. (he small plasma bullets can be interpreted as bright knots in the continuous jet.," In that case, the small plasma bullets can be interpreted as bright knots in the continuous jet."
 The knots are created when the mass outflow rate or the density ol the continuous jet is temporarily increased., The knots are created when the mass outflow rate or the density of the continuous jet is temporarily increased.
 A temporal increase of temperature is not plausible. because il would result in a change of the spectrum. which has not been observed.," A temporal increase of temperature is not plausible, because it would result in a change of the spectrum, which has not been observed."
 The properties of the knots would not be much different [rom those of the small plasma bullets discussed above. and the estimates above are valid if the knots or bullets coexist with ihe continuous jet.," The properties of the knots would not be much different from those of the small plasma bullets discussed above, and the estimates above are valid if the knots or bullets coexist with the continuous jet."
 Decause (he spectral shape does not much change in the decay and because no iron line is detected in the shot component. a non-(hermal emission from expanding bullets also can account for the shot component.," Because the spectral shape does not much change in the decay and because no iron line is detected in the shot component, a non-thermal emission from expanding bullets also can account for the shot component."
 As for the steady. non-variable component. it is definitely a thin-thermal emission wilh Doppler-shifted iron lines.," As for the steady non-variable component, it is definitely a thin-thermal emission with Doppler-shifted iron lines."
 Therefore. ibis natural to interpret the shot component as a thermal emission. and an interpretation of non-thermal emission is rather eccentric.," Therefore, it is natural to interpret the shot component as a thermal emission, and an interpretation of non-thermal emission is rather eccentric."
 In the following paragraph. we show plwsical parameters of a plasma bullet assuming (hat (he bullets emit N ray via svncehrotron radiation.," In the following paragraph, we show physical parameters of a plasma bullet assuming that the bullets emit X ray via synchrotron radiation."
 A power-law distribution of svnchrotron electrons. is asstuned. where 5 is Lorentz factor of electrons. rey is the svnchrotron-electron number cdensitv. and p—2.2N is (he electron energy index derived [rom the photon index D—-(p-1)/2.," A power-law distribution of synchrotron electrons, is assumed, where $\gamma$ is Lorentz factor of electrons, $n_{\rm
e,syn}$ is the synchrotron-electron number density, and $p =
2.2^{+1.4}_{-1.0}$ is the electron energy index derived from the photon index $\Gamma = (p+1)/2$."
" The maximum Lorentz factor should be at least μας2771.4x10"" to account for the X-ray emission up to 10 keV.", The maximum Lorentz factor should be at least $\gamma_{\rm max} > 1.4 \times10^{5}$ to account for the X-ray emission up to 10 .
. Optically thinsvuchrotvon flix from such a sphere is, Optically thinsynchrotron flux from such a sphere is
the center. (log(r)—log(r)). using those stars that have al least 5 measurements in each filter and (hat have a brightness within 0.5 mag of Vacpy.,"the center, $(\log(r) - \overline{ \log(r)})$, using those stars that have at least 5 measurements in each filter and that have a brightness within 0.5 mag of $V_{RGBB}$."
 The measured slope in (D—V) color is a miniscule (0.005mc0.002) mag dex1 ellectivelv zero since the total extent of the eround-based dataset is ~1.1 dex., The measured slope in $(B-V)$ color is a miniscule $(0.005 \pm 0.002)$ mag $^{-1}$ – effectively zero since the total extent of the ground-based dataset is $\sim$ 1.1 dex.
 However. the gradient measured in (2—J) color is (0.02720.003) mag |. and is (0.02820.007) mag ! in (0)—V) color.," However, the gradient measured in $(B-I)$ color is $(0.027 \pm 0.003)$ mag $^{-1}$, and is $(0.028 \pm 0.007)$ mag $^{-1}$ in $(U-V)$ color."
 Each of these eradients is significantly smaller than that for the IIB stars. discussed in Section 5.," Each of these gradients is significantly smaller than that for the HB stars, discussed in Section 5."
 Our model results. shown in Table 3.. predict that there should be no measurable color eracdient due to temperature in the cluster at the level of the RGBB if we are only dealing with two ecqual-age. equal-metallicitv populations differing only in their initial helium abundance bv a small amount (AY~ 0.03).," Our model results, shown in Table \ref{Table:RGBBproperties}, predict that there should be no measurable color gradient due to temperature in the cluster at the level of the RGBB if we are only dealing with two equal-age, equal-metallicity populations differing only in their initial helium abundance by a small amount ${\Delta}Y \sim$ 0.03)."
 The presence of these small gradients incliciates there mar be an additional factor at play. or (hat. helium-enriched red giants might be made slightly bluer than predicted by models.," The presence of these small gradients indiciates there may be an additional factor at play, or that helium-enriched red giants might be made slightly bluer than predicted by models."
 The color variation would be consistent wilh the stars nearer the center either having a metallicity. o[Fe/I] 2 0.05 dex. or a temperature colder by OT s 17 Ιον (Alonsoetal.1999)..," The color variation would be consistent with the stars nearer the center either having a metallicity, ${\delta}$ [Fe/H] $\approx$ 0.05 dex, or a temperature colder by ${\delta}$ T $\approx$ 17 Kelvin \citep{1999A&AS..140..261A}."
 We have found ~3.60 and ~2.30 delections that the cluster RGBB stars become fainter and more numerous with increasing distance from the cluster center. and that the IIB becomes Lainter and redder at (he ον τισ and ~4.lo levels respectively.," We have found $\sim$ $\sigma$ and $\sim$ $\sigma$ detections that the cluster RGBB stars become fainter and more numerous with increasing distance from the cluster center, and that the HB becomes fainter and redder at the $\sim$ $\sigma$ and $\sim$ $\sigma$ levels respectively."
 These [our independent effects are easily explained bv stellar theory if there is a second generation of stars in 47 Tuc that is heliun-enhanced and more centrally concentrated. Briley(1997)..," These four independent effects are easily explained by stellar theory if there is a second generation of stars in 47 Tuc that is helium-enhanced and more centrally concentrated. \citet{1997AJ....114.1051B},"
 in his aualvsis of δρα indices in 283 cluster giants. Found a similar radial profile lor 47 Tuc.," in his analysis of CN-band indices in 283 cluster giants, found a similar radial profile for 47 Tuc."
 The ratio of CN-strong to CN-weak cluster stars was approximately equal ( 1.8) interior to 10. and then dropped steeply for sources separated [rom the cluster center by 10-20.," The ratio of CN-strong to CN-weak cluster stars was approximately equal $\sim$ 1.8) interior to $\arcmin$, and then dropped steeply for sources separated from the cluster center by $\arcmin$."
 If this is due to dvnamical segregation between the (wo generations. then the thickness of ihe main-sequence observed by Andersonetal.(2009) in the cluster core should not be observable in (he cluster outskirts.," If this is due to dynamical segregation between the two generations, then the thickness of the main-sequence observed by \citet{2009ApJ...697L..58A} in the cluster core should not be observable in the cluster outskirts."
 This will prove a difficult measurement to make since the surface density of stars on the sky drops steeply at these distances., This will prove a difficult measurement to make since the surface density of stars on the sky drops steeply at these distances.
 Additional ground-based data could also prove useful., Additional ground-based data could also prove useful.
 We estimate (hat there are ~30 cluster RGDD stars outside the range of our observations. (hese should be slightly fainter and with a higher equivalent width than those measured thus far.," We estimate that there are $\sim$ 30 cluster RGBB stars outside the range of our observations, these should be slightly fainter and with a higher equivalent width than those measured thus far."
 The remaining ~100 IB stars should also be fainter and redder than those within it of the cluster center., The remaining $\sim$ 100 HB stars should also be fainter and redder than those within $\arcsec$ of the cluster center.
 In spite of the evidence for an overall eraclient. it is interesting thal we do not find evidence of a gradient within the cluster center as traced by the dataset.," In spite of the evidence for an overall gradient, it is interesting that we do not find evidence of a gradient within the cluster center as traced by the dataset."
 We discuss four possible explanations., We discuss four possible explanations.
 One wav to explain tliis effect is to have the two stellar populations, One way to explain this effect is to have the two stellar populations
40.11. where 5. M. and Il represent the number of counts detected in the energy ranges 0.31.0. 1.02.0. and 2.07.0 keV. respectively.,", where S, M, and H represent the number of counts detected in the energy ranges 0.3–1.0, 1.0–2.0, and 2.0–7.0 keV, respectively."
 Gaussian errors were measured for the hackerounc-subtracted number of counts in each energy bin with the CLAO taskdinertract., Gaussian errors were measured for the background-subtracted number of counts in each energy bin with the CIAO task.
 Analvsis of the optical data initially provided a lew possible variable stars within the r2-71 error circle., Analysis of the optical data initially provided a few possible variable stars within the r2-71 error circle.
" We scrutinized each possibility through aperture photometry. completeness tests, and difference imaging."," We scrutinized each possibility through aperture photometry, completeness tests, and difference imaging."
 The results show no strong detection of optical variability ancl provide an upper-limit to the D magnitude of any hiehlv variable counterpart (to 12-11., The results show no strong detection of optical variability and provide an upper-limit to the $B$ magnitude of any highly variable counterpart to r2-71.
 The region of interest for all three of our optical observations is shown in the images in Figure L.., The region of interest for all three of our optical observations is shown in the images in Figure \ref{ims}.
 The DAOPHOT II output revealed one star in the error circle that was signilicantlv brighter in the 2004-June-14 observation. when tlie X-ray source was most likely active.," The DAOPHOT II output revealed one star in the error circle that was significantly brighter in the 2004-June-14 observation, when the X-ray source was most likely active."
 The star at .À.—00:42:32.414. decl.—41:17:03.47. was measured to have B=25.50d0.06 in the 2004-June-14 observation.," The star at R.A.=00:42:33.414, decl.=41:17:03.47, was measured to have $B=25.50\pm0.06$ in the 2004-June-14 observation."
 DAOPIIOT II fled to find this star in the 2004-Januarv- observation., DAOPHOT II failed to find this star in the 2004-January-23 observation.
" Because inspection of the 2004-Januarv-23 image reveals a source al (his location. this non-detection was likely due to the effects of the bright neighboring star 0.05"" io the southeast."," Because inspection of the 2004-January-23 image reveals a source at this location, this non-detection was likely due to the effects of the bright neighboring star $''$ to the southeast."
" Aperture photometry of the location in the 2004-Januaryv-23. observation with a 0.075"" radius aperture vields B=26.0 220.1.", Aperture photometry of the location in the 2004-January-23 observation with a $''$ radius aperture yields $B=26.0\pm0.1$ .
 DAOPIIOT II measured the star to be B=25.98+0.09 in the 2004-Aneust-15 observation., DAOPHOT II measured the star to be $B=25.98\pm0.09$ in the 2004-August-15 observation.
 Therefore according to the standard errors an increase in brightness during the 2004-June-14 at the do confidence level occurred: however. the standard errors do not take into account the uncertainty introduced by the close brighter neighbor.," Therefore according to the standard errors an increase in brightness during the 2004-June-14 at the $\sigma$ confidence level occurred; however, the standard errors do not take into account the uncertainty introduced by the close brighter neighbor."
 Anv added uncertainty due to crowding would decrease the significance ol this brightness increase., Any added uncertainty due to crowding would decrease the significance of this brightness increase.
 In addition to this suspicious counterpart candidate. there were 2 fainter stars detected bv DAOPIIOT Lin the error circle of the 2004-June-124 observation that were not detected in the other 2 observations. when the X-ray source was not active.," In addition to this suspicious counterpart candidate, there were 2 fainter stars detected by DAOPHOT II in the error circle of the 2004-June-14 observation that were not detected in the other 2 observations, when the X-ray source was not active."
 These stars had 2 magnitudes of 26.640.2 and 7250.2 in the 2004-June-14 observation., These stars had $B$ magnitudes of $\pm$ 0.2 and $\pm$ 0.2 in the 2004-June-14 observation.
 These stars may not have been detected in the other observations because crowding issues caused our photometry (o be incomplete at these faint magnitudes., These stars may not have been detected in the other observations because crowding issues caused our photometry to be incomplete at these faint magnitudes.
 We determined (he completeness of the optical data in the area of 12-71 by comparing ihe DAOPIIOT II output [rom the 2004-June-1H observation to those of the observation., We determined the completeness of the optical data in the area of r2-71 by comparing the DAOPHOT II output from the 2004-June-14 observation to those of the 2004-August-15 observation.
 Figure 3. shows (wo histograms., Figure \ref{comp} shows two histograms.
 The solid histogram shows the percentage of stars detected bv the DAOPILOT analvsis in the 2004-Jine-14 observation within 3” of the center of the errorcircle but not detected by the same analvsis in the 2004-August-15, The solid histogram shows the percentage of stars detected by the DAOPHOT analysis in the 2004-June-14 observation within $''$ of the center of the errorcircle but not detected by the same analysis in the 2004-August-15
have chaneccd sliebtly during the span of the photometric observatioi (Strassumeier ct al.,have changed slightly during the span of the photometric observations (Strassmeier et al.
 1999)., 1999).
 DDSS10 arened that the rotational period could be equal to twice the photometric pcriod. ne. 617.6 £5 davs.," DBSS10 argued that the rotational period could be equal to twice the photometric period, i.e. 617.6 $\pm$ 5 days."
 To determine independeitly the rotational period of EI Evi. we have followed the apxoach of Petit et al. (," To determine independently the rotational period of EK Eri, we have followed the approach of Petit et al. ("
2002). aud calculated a set of niaenoetIc naps. assunuidue for cach map a cdiffercut value for he rotational perio.,"2002), and calculated a set of magnetic maps, assuming for each map a different value for the rotational period."
" We impose a constant cutropy for all he images and caleulate as à goodness-ot-fit parameex the reduced 47 ""EM(xz hereafter). by comparing: the se of svuthetic Stokes V. profiles produced by ZDI to the observec time-series of Stokes V. profiles.", We impose a constant entropy for all the images and calculate as a goodness-of-fit parameter the reduced $\chi^2$ $\chi^2 _r$ hereafter) by comparing the set of synthetic Stokes $V$ profiles produced by ZDI to the observed time-series of Stokes $V$ profiles.
 We then study the variafions. of .-yp over the range of. rotation. periods. to deteΠΠο the period value producing the best magnetic mode|] (identified bv t1 lowest value of20» X7).," We then study the variations of $\chi^2 _r$ over the range of rotation periods, to determine the period value producing the best magnetic model (identified by the lowest value of $\chi^2 _r$ )."
 Tere we scaned. LOO values of the period between 200 aud 700 dave. a range which encompasses the photometric aud rotatioua periods proposed or EK Exit by DBSS1LO.," Here we scanned 400 values of the period between 200 and 700 days, a range which encompasses the photometric and rotational periods proposed for EK Eri by DBSS10."
 Our periodograms show a favoured period near 311 davs., Our periodograms show a favoured period near 311 days.
 Ilowever the ΠΜs xzD valie ds high (about 25). notwitlistaudiuο that we Πίος our observational sample to the timespa1 Novelox 2007 - March. 2010 in order to avoid the effect of the stro18o deviations observed ou 20 September 2007 and during the aitunm of 2010.," However the minimum $\chi^2 _r$ value is high (about 25), notwithstanding that we limited our observational sample to the timespan November 2007 - March 2010 in order to avoid the effect of the strong deviations observed on 20 September 2007 and during the autumn of 2010."
 We therefore itrodced a posside differentia rotation (Petit ct al., We therefore introduced a possible differential rotation (Petit et al.
 20Py, 2002).
 UItimately we could not improve the results and cdeποσα a likseligihe differenial rotation for El Evi.," Ultimately, we could not improve the results and deduced a negligible differential rotation for EK Eri."
 We aπο COsidered. tlat our a»proxinate line xofile model welt be tie source of the hie1 D xnothe current local li1e profile used il OUP Lioc elis a ο πλοίο., We also considered that our approximate line profile model might be the source of the high $\chi^2_r$: the current local line profile used in our model is a gaussian function.
 In aj attempt o improve the fit o: the line xofile. we introduced a Loreutzian fuiction and some asvmunetry. bit we still could not reduce significantly he 42.," In an attempt to improve the fit of the line profile, we introduced a Lorentzian function and some asymmetry, but we still could not reduce significantly the $\chi^2 _r$."
 Ultimatelv. we cosider that the relatively poor detailed fit to the Stokes V profiles resulting from our JOS-fit model 1 simainly die to the seasonal scatter of the naeuctic field streneth.," Ultimately, we consider that the relatively poor detailed fit to the Stokes $V$ profiles resulting from our best-fit model is mainly due to the seasonal scatter of the magnetic field strength."
 Figure Losrows the periodogram obtained for the whole data set between November 2007 aud March 2010. assnndne a itational axis inclination of EK Eri of i=QUU ancl truncating the expansion at (=3.," Figure 4 shows the periodogram obtained for the whole data set between November 2007 and March 2010, assuming a rotational axis inclination of EK Eri of $i = 60^\circ$ and truncating the expansion at $\ell=3$."
 For rotatioial periods of 311.5 d. 308.8 d. and 625 d. we obtain respectively vaues for vz of 25 (our πια). 25.3H aud 38.2.," For rotational periods of 311.5 d, 308.8 d, and 625 d, we obtain respectively values for $\chi^2 _r$ of 25 (our minimum), 25.3 and 38.2."
" T16 600 day period suggestedCoco by DDBSS10 is therefore elininated boti by our periodograim analvsis aud by the fact that their scenario to explain Poor2=2D, is uot supported by tie DB, variations (Sect.", The 600 day period suggested by DBSS10 is therefore eliminated both by our periodogram analysis and by the fact that their scenario to explain $P_{\rm rot} = 2P_{phot}$ is not supported by the $B_{\ell}$ variations (Sect.
 2.2)., 2.2).
 Since we found that phase 0.0 of the ephemeris of DBSSLO corresyouds to approxiniate‘lw the center of the magnetic spot which dominates the magnetic map (see below). we decided to use the photomevic period of 308.8 days (and TJDy = 215 3372.679) for our analysis of the maguetic topoOey.," Since we found that phase 0.0 of the ephemeris of DBSS10 corresponds to approximately the center of the magnetic spot which dominates the magnetic map (see below), we decided to use the photometric period of 308.8 days (and ${\rm HJD_0}$ = 245 3372.679) for our analysis of the magnetic topology."
" Fieure 5 prescuts the fit of our best model to the observed Stokes V LSD profiles. for P, = 308.8 d. and i=00."," Figure 5 presents the fit of our best model to the observed Stokes $V$ LSD profiles, for $P_{\rm rot}$ = 308.8 d, and $i = 60^\circ$."
 Iu the two previous subsectious. we have preseited the models aud the paraieters einiploved to infer the naenetic topology of EI Exi.," In the two previous subsections, we have presented the models and the parameters employed to infer the magnetic topology of EK Eri."
" Because ¢sin? cannot be determined preciscly, the value of the inclination / is also ouly weakly coustrained."," Because $v\sin i$ cannot be determined precisely, the value of the inclination $i$ is also only weakly constrained."
 This is our main remainiug uucertaiu parameter which we have attempted to coustrain as part ofthe ZDI reconstruction., This is our main remaining uncertain parameter which we have attempted to constrain as part of the ZDI reconstruction.
" We have therefore fit our Stokes V. data with our models and / varving frou 857 to. 103,", We have therefore fit our Stokes $V$ data with our models and $i$ varying from $^\circ$ to $^\circ$.
 For all models we find that the poloidal component colalms more than 98A of the reconstructed magnetic enerev., For all models we find that the poloidal component contains more than $\%$ of the reconstructed magnetic energy.
 When / decreases. we find the ZDI procedure to xovide us with a better magnetic model. in t1C SCISC lia the toal information content of the reconstructed naenetiec oeeonierv decreases (woe follow here the uaximuun entropy criterion. Donati 2001. explicite in Morin ot al.," When $i$ decreases, we find the ZDI procedure to provide us with a better magnetic model, in the sense that the total information content of the reconstructed magnetic geometry decreases (we follow here the maximum entropy criterion, Donati 2001, explicited in Morin et al."
 2008)., 2008).
 For high iuclinalou. /NU. the opology is dominated by a quadrupolar configuration.," For high inclination, $i > 80^\circ$, the topology is dominated by a quadrupolar configuration."
 m lis case. the two rotation poles correspond to OS]ive obhudtv while there is a rearly equatorial band. with a single enormous magnetic spot. both of negative polarity.," In this case, the two rotation poles correspond to positive polarity while there is a nearly equatorial band, with a single enormous magnetic spot, both of negative polarity."
 The magneticOo spot corresponds to the pliase of the dark. (and therefore presumably cool) photometric spot to better than 0.1 evcles., The magnetic spot corresponds to the phase of the dark (and therefore presumably cool) photometric spot to better than 0.1 cycles.
" For ;«NU"". the dipole component dominates and corresponds to about 90% of the maguctic energv at about ;/ = 607."," For $i < 80^\circ$, the dipole component dominates and corresponds to about $90\%$ of the magnetic energy at about $i$ = $^\circ$."
 Figure 6 shows the corresponding ZDI map., Figure 6 shows the corresponding ZDI map.
 This map also features a strong maguctic spot of negative polarity. corresponding to the negative pole of the magnetic dipole.," This map also features a strong magnetic spot of negative polarity, corresponding to the negative pole of the magnetic dipole."
 Again the magnetic λαΙΙΙ corresponds in phase to the photometric spot., Again the magnetic maximum corresponds in phase to the photometric spot.
 The characteristics of this model are compatible with the scenario of Aurierre et al. (, The characteristics of this model are compatible with the scenario of Aurièrre et al. (
2008) and DBSSIO. in which the maguetic/plotometric spot would correspond to the relmnant of a maguetic pole of the Ap star progenitor of EW Exi.,"2008) and DBSS10, in which the magnetic/photometric spot would correspond to the remnant of a magnetic pole of the Ap star progenitor of EK Eri."
without distinction.,without distinction.
 As the background subtraction step is no different (han a standard. ADI/LOCI speckle subtraction. we have developed a background subtraction routine (Maroisetal.2010h) based on the locally optimized combination of images algorithm (LOCI.Lalreniereetal.2007).," As the background subtraction step is no different than a standard ADI/LOCI speckle subtraction, we have developed a background subtraction routine \citep{marois10b} based on the locally optimized combination of images algorithm \citep[LOCI,][]{lafreniere07}."
. This new routine is more efficient (han the classical median subtraction as il puts different weights on the background reference images., This new routine is more efficient than the classical median subtraction as it puts different weights on the background reference images.
" Consider the background sublvaction in the pith image J, of the sequence.", Consider the background subtraction in the $p$ th image $I_p$ of the sequence.
 An annulus (called D. Fig. 1))," An annulus (called $D$, Fig. \ref{fig : f1}) )"
" is defined around (he star aud all reference background. images are selected where the star has been sufficiently dithered compared to 7,.", is defined around the star and all reference background images are selected where the star has been sufficiently dithered compared to $I_p$.
" The algorithm searches [or the linear combination of (hese images that minimizes the V (i.e. the residual noise in 2 excluding bad pixels) which can be expressed as where ÀA/ is the number of pixels in ihe image. AN ids the total number of images in (he sequence. s; is the jth pixel. J; is the /th image of the sequence. and al, the coefficients to be determined."," The algorithm searches for the linear combination of these images that minimizes the $\chi^2_p$ (i.e. the residual noise in $D$ excluding bad pixels) which can be expressed as where $M$ is the number of pixels in the image, $N$ is the total number of images in the sequence, $x_j$ is the $j$ th pixel, $I_i$ is the $i$ th image of the sequence, and $\alpha^i_p$ the coefficients to be determined."
 The inner radius of area D is chosen to avoid the bright star speckle noise (hat can bias the background. minimization algorithm. while its outer radius is chosen to avoid the dithered star images.," The inner radius of area $D$ is chosen to avoid the bright star speckle noise that can bias the background minimization algorithm, while its outer radius is chosen to avoid the dithered star images."
 In our case. (he inner and outer radii are 3 and 30A/D corresponding to 24 to 240pixels.," In our case, the inner and outer radii are 3 and $\lambda/$ D corresponding to 24 to 240pixels."
" We have then imposed al,=0 if the star is ab a similar location in J; as in I).", We have then imposed $\alpha_p^i=0$ if the star is at a similar location in $I_i$ as in $I_p$.
 The LOCI algorithm then finds the optimal a»i.best (Lafreniereetal.2007:= that minimizes the Ve," The LOCI algorithm then finds the optimal $\alpha_p^{i,best}$ \citep{lafreniere07,marois10b} that minimizes the $\chi_p^2$."
" The optimized background image D, al each pixel lor J, (not only in D) is then The generated reference background image D, is finally subtracted [rom J,.", The optimized background image $B_p$ at each pixel for $I_p$ (not only in $D$ ) is then The generated reference background image $B_p$ is finally subtracted from $I_p$.
" As ane!=()for the images where the star is registered at the same location in J; as in J,. DB,Pp does not contain any [lux from a possible companion."," As $\alpha_p^{i,best}=0$for the images where the star is registered at the same location in $I_i$ as in $I_p$, $B_p$ does not contain any flux from a possible companion."
" Moreover. if a companion exists in /,. ils impact on ihe aPpl.besl LOCI coefficients is negligible since (he companion flux is much. smaller than the thermal background. and D is large compared to the companion PSF."," Moreover, if a companion exists in $I_p$, its impact on the $\alpha_p^{i,best}$ LOCI coefficients is negligible since the companion flux is much smaller than the thermal background and $D$ is large compared to the companion PSF."
 Once the background is subtracted. we sub-pixel register each image using an iterative cross-correlation gaussian [fii to the PSF core (images are unsaturated).," Once the background is subtracted, we sub-pixel register each image using an iterative cross-correlation gaussian fit to the PSF core (images are unsaturated)."
 The speckle subtraction is performed folowing a basic ADI data reduction as described in Maroisetal. (2006)..., The speckle subtraction is performed following a basic ADI data reduction as described in \citet{marois06}. .
 For each data cube. an initial reference PSF," For each data cube, an initial reference PSF"
"values of A and B and was determined in the publication by ? to be From the velocities of stars orbiting Sgr A*, which is considered to be the centre of the Galaxy because of its small own velocity, the distance between the Sun and the GC has been determined to (?):: in agreement with previous authors (?).","values of A and B and was determined in the publication by \citet{Fuchs:2009mk} to be From the velocities of stars orbiting Sgr A*, which is considered to be the centre of the Galaxy because of its small own velocity, the distance between the Sun and the GC has been determined to \citep{Gillessen:2008qv}: in agreement with previous authors \citep{Ghez:2008ms}."
 With this Galactocentric distance one finds from Eq., With this Galactocentric distance one finds from Eq.
" 10 a rotation velocity of the Sun which is consistent with recent observations of Galactic masers in ?,, who used data from the Very Long Baseline Array (VLBA) and the Japanese VLBI Exploration of Radio Astronomy (VERA)."," \ref{oort1} a rotation velocity of the Sun which is consistent with recent observations of Galactic masers in \citet{Bovy:2009dr}, who used data from the Very Long Baseline Array (VLBA) and the Japanese VLBI Exploration of Radio Astronomy (VERA)."
 This speed determines the mass of the Galaxy inside the solar radius., This speed determines the mass of the Galaxy inside the solar radius.
 In this analysis two different rotation velocities are considered: the RC within the Galactic disc (Fig. 3)), In this analysis two different rotation velocities are considered: the RC within the Galactic disc (Fig. \ref{fig:rc_simple}) )
 and the velocity distribution for stars outside the disc with z > 4 kpc (Fig. 4))., and the velocity distribution for stars outside the disc with z $>$ 4 kpc (Fig. \ref{fig:rc_4kpc}) ).
 They are discussed separately., They are discussed separately.
 For the RC within the disc a combination of different measurements with different tracers has been summarized by ?.., For the RC within the disc a combination of different measurements with different tracers has been summarized by \citet{Sofue:2008wt}.
" The experimental data, which can be found on the author's web page, were scaled to Vo= 244 ss""! at a Galactocentric distance of 8.3 kpc."," The experimental data, which can be found on the author's web page, were scaled to $_\odot =$ 244 $^{-1}$ at a Galactocentric distance of 8.3 kpc."
" Furthermore, the rotation velocity was averaged in 17 radial bins from the GC to a radius of 22 kpc, as shown in Fig."," Furthermore, the rotation velocity was averaged in 17 radial bins from the GC to a radius of 22 kpc, as shown in Fig."
 3bb and tabulated in Table 2.., \ref{fig:rc_simple}b b and tabulated in Table \ref{tab:velocities}.
 The shape of the measured velocity distribution shows a strong increase of the rotation velocity in the inner part of the Galaxy which presumably results from the dense core of the Galaxy., The shape of the measured velocity distribution shows a strong increase of the rotation velocity in the inner part of the Galaxy which presumably results from the dense core of the Galaxy.
 For the inner Galaxy the rotation curve is dominated by the visible matter and the parametrization of section 2.1 yields a reasonable description., For the inner Galaxy the rotation curve is dominated by the visible matter and the parametrization of section \ref{subsec:para_lumi} yields a reasonable description.
" However, at the outer Galaxy the experimental data cannot be explained: all profiles predict a slow decrease of the rotation velocity in contrast to the data, which show first a decrease betwween 6 and 10 kpc and then increases again."," However, at the outer Galaxy the experimental data cannot be explained: all profiles predict a slow decrease of the rotation velocity in contrast to the data, which show first a decrease ween 6 and 10 kpc and then increases again."
" Such a peculiar change of slope cannot be explained by a smoothly decreasing DM density profile, but needs substructure, e.g. the infall of a dwarf Galaxy, as mentioned in the introduction and ?.."," Such a peculiar change of slope cannot be explained by a smoothly decreasing DM density profile, but needs substructure, e.g. the infall of a dwarf Galaxy, as mentioned in the introduction and \citet{deBoer:2005tm}."
 Such a ringlike substructure is supported by the gas flaring (?).., Such a ringlike substructure is supported by the gas flaring \citep{Kalberla:2007sr}.
" The thickness of the substructure is of the order of 1 kpc, so it should not show up for halo stars well above this height; this is indeed the"," The thickness of the substructure is of the order of 1 kpc, so it should not show up for halo stars well above this height; this is indeed the"
been proposed for hypernova SN 2003be to explain the high mass.,been proposed for hypernova SN 2003bg to explain the high mass.
 On the other hand a single star scenario cannot explain the characteristics of SN 1993J (see Sec. 2.1.2))., On the other hand a single star scenario cannot explain the characteristics of SN 1993J (see Sec. \ref{sec:1993J}) ).
 Consequently. the single star scenario cannot explain all observed type IIb supernovae.," Consequently, the single star scenario cannot explain all observed type IIb supernovae."
 The determination of the rate from single stars is sensitive to uncertainties in the stellar wind mass loss rates., The determination of the rate from single stars is sensitive to uncertainties in the stellar wind mass loss rates.
 Especially with respect to stellar winds from massive red giants our understanding ts limited., Especially with respect to stellar winds from massive red giants our understanding is limited.
 2? speculate that a superwind driven by pulsational instabilities may drive a strong mass. loss. bringing the minimum mass for type IIb supernovae down to about 20.," \cite{YoonCantiello10} speculate that a superwind driven by pulsational instabilities may drive a strong mass loss, bringing the minimum mass for type IIb supernovae down to about 20."
Μ.... Single stars with a mass above ~25 are believed to produce only faint supernovae (?).., Single stars with a mass above $\sim$ 25 are believed to produce only faint supernovae \cite[]{Fryer99}.
 Consequently. these type IIb SNe will appear different than type ΠΡ SNe formed by binary stars.," Consequently, these type IIb SNe will appear different than type IIb SNe formed by binary stars."
 Nevertheless. in the correct mass range. single stars can explode as type IIb SNe and therefore it 15 reasonable to compare the expected rate from binary and single-star progenitors.," Nevertheless, in the correct mass range, single stars can explode as type IIb SNe and therefore it is reasonable to compare the expected rate from binary and single-star progenitors."
 To be able to do this comparison we computed single stellar models with the same input physics as the binary nodels we discuss in the next section., To be able to do this comparison we computed single stellar models with the same input physics as the binary models we discuss in the next section.
 We find that the initial nass range for single stars resulting in type IIb progenitors should be within 32.5233M., We find that the initial mass range for single stars resulting in type IIb progenitors should be within 32.5–33.
.. In Table | we list the final nass and and the amount of hydrogen left in the envelope near this range of initial masses., In Table \ref{TS} we list the final mass and and the amount of hydrogen left in the envelope near this range of initial masses.
 Assuming a Kroupa initial mass function (?) we estimate that about of all single stars nore massive than 8 solar masses are within this mass range and will end their lives as a type Ib supernova., Assuming a Kroupa initial mass function \cite[]{Kroupa01} we estimate that about of all single stars more massive than 8 solar masses are within this mass range and will end their lives as a type IIb supernova.
 In this section we discuss binary progenitor models for type IIb supernova., In this section we discuss binary progenitor models for type IIb supernova.
 We compute the evolution of binary models with different initial orbital periods and different initial mass ratios., We compute the evolution of binary models with different initial orbital periods and different initial mass ratios.
 We adopt an initial primary mass of 15. in agreement with the progenitor model proposed for 1993J (?).., We adopt an initial primary mass of 15 in agreement with the progenitor model proposed for 1993J \citep{Maund04}.
 We assume that type IIb supernovae result from massive stars that undergo core collapse with an envelope which contains between 0.1 and 0.5 of hydrogen., We assume that type IIb supernovae result from massive stars that undergo core collapse with an envelope which contains between 0.1 and 0.5 of hydrogen.
 This criterion is based on observations. previous models (22?) and some test models which proved that a hydrogen mass less than 0.1 gives rise to a compact rather than an extended progenitor.," This criterion is based on observations, previous models \citep{Podslad93,Woosley94, Elmhamdi06} and some test models which proved that a hydrogen mass less than 0.1 gives rise to a compact rather than an extended progenitor."
 As an example we discuss a system with an initial orbital period of 1500 days and initial masses of 15 and 14.35 for the primary and secondary star respectively., As an example we discuss a system with an initial orbital period of 1500 days and initial masses of 15 and 14.35 for the primary and secondary star respectively.
 The initially most massive star evolves faster and experiences significant mass loss in the form of a stellar wind when it ascends the giant branch and during central helium burning., The initially most massive star evolves faster and experiences significant mass loss in the form of a stellar wind when it ascends the giant branch and during central helium burning.
 After about 13.1 Myr. when the helium mass fraction has dropped below 0.5 in the center. it fills it Roche lobe (latecaseBmassfer.?)..," After about 13.1 Myr, when the helium mass fraction has dropped below 0.5 in the center, it fills it Roche lobe \cite[late case B mass transfer, ][]{Kippenhahn67}."
 At this moment it has already lost more than 1. and has become less massive than its companion., At this moment it has already lost more than 1 and has become less massive than its companion.
 The reversal of the mass ratio before the onset of Roche-lobe overflow helps to stabilize the mass transfer., The reversal of the mass ratio before the onset of Roche-lobe overflow helps to stabilize the mass transfer.
 In Figure | we depict the mass-transfer rate as a function of the remaining envelope mass., In Figure \ref{MtrMass} we depict the mass-transfer rate as a function of the remaining envelope mass.
 We find that the mass transfer initially takes place on a timescale equal to the thermal timescale of the primary star., We find that the mass transfer initially takes place on a timescale equal to the thermal timescale of the primary star.
 The maximum mass-transfer rate during this phase is 6%107? M.yr7!., The maximum mass-transfer rate during this phase is $6\times10^{-4}~\Msun$ $^{-1}$.
 Although this phase lasts only about 0.05 Myr. about 4.5 is transferred.," Although this phase lasts only about 0.05 Myr, about 4.5 is transferred."
 After this phase the star keeps filling its Roche lobe and mass transfer continues on the nuclear timescale. at a rate comparable to the mass-loss rate in the form of a stellar wind. about 3x107 Mayr!.," After this phase the star keeps filling its Roche lobe and mass transfer continues on the nuclear timescale, at a rate comparable to the mass-loss rate in the form of a stellar wind, about $3\times10^{-6}~\Msun$ $^{-1}$."
 During this phase the primary expands on its nuclear timescale while it is burning helium in its center., During this phase the primary expands on its nuclear timescale while it is burning helium in its center.
 This phase lasts about 0.8 Myr and about 2.5 is transferred., This phase lasts about 0.8 Myr and about 2.5 is transferred.
 After central helium exhaustioi the star expands again on its thermal timescale and à second maximum in the mass-transfer rate occurs., After central helium exhaustion the star expands again on its thermal timescale and a second maximum in the mass-transfer rate occurs.
 Finally. at the onset of carbon burning the star expands again resulting in a third peak in the mass-transfer rate. see Fig. 1..," Finally, at the onset of carbon burning the star expands again resulting in a third peak in the mass-transfer rate, see Fig. \ref{MtrMass}."
 We follow the evolution of the system up to this point. when the mass of hydrogen in the envelope has decreased to 0.46M.," We follow the evolution of the system up to this point, when the mass of hydrogen in the envelope has decreased to 0.46."
.. Extrapolating the mass-loss rate we find that the amount of hydrogen in the envelope at the time of explosion will be about 0.37M., Extrapolating the mass-loss rate we find that the amount of hydrogen in the envelope at the time of explosion will be about 0.37.
... Therefore we expect that the primary star explodes as a type IIb supernova., Therefore we expect that the primary star explodes as a type IIb supernova.
 In wider systems mass transfer starts in a later phase of the evolution of the primary star., In wider systems mass transfer starts in a later phase of the evolution of the primary star.
 Because the primary stars in these systems are more evolved. there is less time available to reduce the mass of the envelope before the explosion.," Because the primary stars in these systems are more evolved, there is less time available to reduce the mass of the envelope before the explosion."
 In addition. stellar winds had more time to reduce the mass of the primary star before the onset of Roche-lobe overflow.," In addition, stellar winds had more time to reduce the mass of the primary star before the onset of Roche-lobe overflow."
 Reversal of the mass ratio stabilizes the process of mass transfer., Reversal of the mass ratio stabilizes the process of mass transfer.
 This results in à lower mass-transfer rate during the first phase of mass transfer., This results in a lower mass-transfer rate during the first phase of mass transfer.
 Vice versa we find that stars in binary systems with lower initial orbital periods remain with smaller envelope masses at the time of explosion., Vice versa we find that stars in binary systems with lower initial orbital periods remain with smaller envelope masses at the time of explosion.
 We give details of all our computed models in Tables Al--A4 in the Appendix., We give details of all our computed models in Tables \ref{TC}- \ref{TNC} in the Appendix.
ambipolar diffusion to yield Puna21 is ib the inflow vo is aligned locally towards Bua so tliat Bua/Boe42.. Even if post-shock regions are subcritical. trausieut ambipolar diffusion siguilicantly increases the imass-to-f[Iux ratio compared to the value that would hold in ideal MHD.,"ambipolar diffusion to yield $\Gamma_\mathrm{final} > 1$ is if the inflow $\mathbf{v}_0$ is aligned locally towards $\mathbf{B}_\mathrm{cloud}$ so that $B_\mathrm{cloud}/B_0 >1$ Even if post-shock regions are subcritical, transient ambipolar diffusion significantly increases the mass-to-flux ratio compared to the value that would hold in ideal MHD."
 A measure of the ünportance of this effect. is the ratio between Dgj aud Upp in the post-shock region., A measure of the importance of this effect is the ratio between $\Gamma_\mathrm{final}$ and $\Gamma_\mathrm{BE}$ in the post-shock region.
 From Equations (67)) aud (56)). is predicted.," From Equations \ref{GADmodel}) ) and \ref{GBEnum}) ), is predicted."
 The turbulent inotious iu clouds cau achieve Mt~50., The turbulent motions in clouds can achieve ${\cal M}\sim 50$.
 With 70)ου1—20. a siguilicant enhancement in tle 1uass-to-flux ratio can be expected due to trausient ambipolar (illusion.," With $\chi_{i0}\sim 1-20$, a significant enhancement in the mass-to-flux ratio can be expected due to transient ambipolar diffusion."
 The estimates of Equations (62)) aud (67)) cau be compared to the ambipolar diffusion time auc imass-to-Iux ratio as measured directly Grom time-depeudent numerical simulatious., The estimates of Equations \ref{tADmodel}) ) and \ref{GADmodel}) ) can be compared to the ambipolar diffusion time and mass-to-flux ratio as measured directly from time-dependent numerical simulations.
 Examples sliowing evolution of the measured D [or several different. parameter values are shown iu Fig. 9.., Examples showing evolution of the measured $\Gamma$ for several different parameter values are shown in Fig. \ref{spcTest}.
 To read the ambipolar diffusion time scale [rom simulations. recall that the growth rate of the uass-to-Iux ratio insicle the core cecreases at time ~/4p.," To read the ambipolar diffusion time scale from simulations, recall that the growth rate of the mass-to-flux ratio inside the core decreases at time $\sim t_\mathrm{AD}$."
 We adopt a definitiou of /4p as the time wheu the slope of the D vs. time curve drops to 25c of its inaximium value., We adopt a definition of $t_\mathrm{AD}$ as the time when the slope of the $\Gamma$ vs. time curve drops to $25\%$ of its maximum value.
" For each simulation. we neasure the mass-to-Ilux ratio X/(B,) at time /—2/xp. aud define this imass-to-Iux ratio inside he central peak (multiplied by 2z /G) as Egg"," For each simulation, we measure the mass-to-flux ratio $\Sigma/\langle B_y\rangle$ at time $t = 2t_\mathrm{AD}$, and define this mass-to-flux ratio inside the central peak (multiplied by $2\pi\sqrt{G}$ ) as $\Gamma_\mathrm{final}$."
 Table 2 shows the predicted values of /4 and Egg from Section 6.2.. as well as the simulation 'esultsfor these quautities.," Table \ref{Model} shows the predicted values of $t_\mathrm{AD}$ and $\Gamma_\mathrm{final}$ from Section \ref{sec:models}, as well as the simulation resultsfor these quantities."
 The measured ambipolar diffusion time scale is 0.3—3 Myr., The measured ambipolar diffusion time scale is $\sim 0.3-3$ Myr.
" Our moclel oreclicts the ambipolar diffusion time scale very well: the IMS value of (4p,pred—/AD.sim)//Ab.sim is 0.19. aud the range of (ap,pred—/Ab.sim)//Ap.sim Is —0.12 to 0.28."," Our model predicts the ambipolar diffusion time scale very well: the RMS value of $\left(t_\mathrm{AD,~pred} - t_\mathrm{AD,~sim}\right)/ t_\mathrm{AD,~sim}$ is $0.19$, and the range of $\left(t_\mathrm{AD,~pred} - t_\mathrm{AD,~sim}\right)/ t_\mathrm{AD,~sim}$ is $-0.42$ to $0.28$."
" The measured imass-to-fIux ""adios deviate from predicted values somewhat more. with a range of (Vinalpred—tina.sim}(Vinal.sim θε to 0.19. and RMS value 0.28."," The measured mass-to-flux ratios deviate from predicted values somewhat more, with a range of $\left(\Gamma_\mathrm{final,~pred} - \Gamma_\mathrm{final,~sim}\right)/ \Gamma_\mathrm{final,~sim}$ $-0.44$ to $0.49$, and RMS value $0.28$."
 The typical size Leone=V/Gr) of the region with enhanced uass-to-Iux ratio at time 2/4p is ~0.2—0.6 pe (Table 2.. column 6).," The typical size $L_\mathrm{core}\equiv N/\langle n\rangle$ of the region with enhanced mass-to-flux ratio at time $2 t_\mathrm{AD}$ is $\sim 0.2-0.6$ pc (Table \ref{Model}, column 6)."
" Iu most of our simulations. tle mass-to-flux ratios are higher than 0.6 (see columu 5 of Table 2)). with some cases (NIO. VI2. B02. BOL and NOL) reaching Lsj,,,> 1."," In most of our simulations, the mass-to-flux ratios are higher than $0.6$ (see column 5 of Table \ref{Model}) ), with some cases (N10, V12, B02, B04, and X01) reaching $\Gamma_{2t_\mathrm{AD}} > 1$ ."
 Recall that we assuined the effective length of the system is comparable in all directious when we define E iu the caudidate, Recall that we assumed the effective length of the system is comparable in all directions when we define $\Gamma$ in the candidate
Our first goal was to determine the optimal resolution or SPL simulations of this kind.,Our first goal was to determine the optimal resolution for SPH simulations of this kind.
 Since we want to use he results of the SPIEL simulations to study the detailed structure and evolution of the collision products. we need o have an accurate model of the distribution of the SPILL xwticles ancl their. properties.," Since we want to use the results of the SPH simulations to study the detailed structure and evolution of the collision products, we need to have an accurate model of the distribution of the SPH particles and their properties."
 On the other hand. we clo not want to waste valuable computer resources going to à ueher resolution than is necessary.," On the other hand, we do not want to waste valuable computer resources going to a higher resolution than is necessary."
 We have shown that both read-on and. olf-axis simulations with  LOO 000 particles eive essentially the same information as runs with higher resolution. and therefore we suggest that most simulations with similar requirements can confidently— use onorder 107 particles.," We have shown that both head-on and off-axis simulations with $\sim$ 100 000 particles give essentially the same information as runs with higher resolution, and therefore we suggest that most simulations with similar requirements can confidently use onorder $10^5$ particles."
 Our second goal was to use these high resolution simulations to study the outer lew percent of the collision products., Our second goal was to use these high resolution simulations to study the outer few percent of the collision products.
 In particular. we were interested in the existence of a surface convection zone. for two reasons.," In particular, we were interested in the existence of a surface convection zone, for two reasons."
 One consequence of à convection zone is mixing of elements in the convection zone., One consequence of a convection zone is mixing of elements in the convection zone.
 If the zone reaches deep enough in a star to dredge. up nuclear processed. material. for example. the surface abundances of elements like €. N and O will be different from what we would expect for primordial material.," If the zone reaches deep enough in a star to dredge up nuclear processed material, for example, the surface abundances of elements like C, N and O will be different from what we would expect for primordial material."
 In blue stragelers and other collision products. surface abuncdances could be used to determine. perhaps. the masses of the parent stars.," In blue stragglers and other collision products, surface abundances could be used to determine, perhaps, the masses of the parent stars."
 However. this can only be done if we understand. whether the abundances should. be mixed by convection or not.," However, this can only be done if we understand whether the abundances should be mixed by convection or not."
 “Lhe second consequence of a surface convection zone is angular momentum loss., The second consequence of a surface convection zone is angular momentum loss.
 Stars can Lose angular momentum through a magnetic wine (lxawaler1988)., Stars can lose angular momentum through a magnetic wind \cite{K88}.
. This process is most effective in stars with deep convection zones. since they can support stronger magnetic ficlds and since the angular momentum is drained out of the entire (unilormily rotating) convection zone.," This process is most effective in stars with deep convection zones, since they can support stronger magnetic fields and since the angular momentum is drained out of the entire (uniformly rotating) convection zone."
 We have shown that the product of the head-on stellar collision investigated. in this paper does not have à surface convection zone larger than 0.004 A7.., We have shown that the product of the head-on stellar collision investigated in this paper does not have a surface convection zone larger than 0.004 $M_{\odot}$.
 For comparison. the Sun is considered to have a fairly shallow convection zone. with a mass of 0.02 AZ. (Guentheretal.1992).," For comparison, the Sun is considered to have a fairly shallow convection zone, with a mass of 0.02 $M_{\odot}$ \cite{GDPK92}."
. X convection zone less than about Q.01 AZ. is certainly not deep enough to substantially modify the surface abundances of any element. with the possible exception of lithium. bervilium ancl boron. (which are destroyed: at. low temperatures).," A convection zone less than about 0.01 $M_{\odot}$ is certainly not deep enough to substantially modify the surface abundances of any element, with the possible exception of lithium, beryllium and boron (which are destroyed at low temperatures)."
 Angular momentum loss is also going to be very small for stars with such small convection zones., Angular momentum loss is also going to be very small for stars with such small convection zones.
 Our final goal was to determine whether the oll-axis stellar collision products had. a circumstellar disk., Our final goal was to determine whether the off-axis stellar collision products had a circumstellar disk.
 I a rotating star has a disk with a mass of about 0.01. AZ.. and it ds locked to the disk by a magnetic field. then angular momentum will be transported [rom the star to the disk by the field. (Ixonigl.1901).," If a rotating star has a disk with a mass of about 0.01 $M_{\odot}$, and it is locked to the disk by a magnetic field, then angular momentum will be transported from the star to the disk by the field \cite{K91}."
. Phe net cllect of the angular momentum transport will be to slow the stars rotation rate., The net effect of the angular momentum transport will be to slow the star's rotation rate.
" This cllect is thought to operate in pre-main-sequence stars (Sills.Pinsonneault&""Terndrup2000.andreferencestherein)", This effect is thought to operate in pre-main-sequence stars \cite[and references therein]{SPT00}.
 Unfortunately. our simulations do not show any evidence for a disk.," Unfortunately, our simulations do not show any evidence for a disk."
 Based on previous work (Benz&Lills 1987).. we expect that blue stragelers which were formecl [rom collisions with larger impact parameters will have circumstellar disks.," Based on previous work \cite{BH87}, we expect that blue stragglers which were formed from collisions with larger impact parameters will have circumstellar disks."
 However. the low impact parameter collision we mocelled in this simulation is still rotating too quickly to contract to the main sequence without losing some angular momentum.," However, the low impact parameter collision we modelled in this simulation is still rotating too quickly to contract to the main sequence without losing some angular momentum."
 Sills et al., Sills et al.
 (2001) showed that blue stragelers that are [formed by stellar collisions have an angular momentum oblem., \shortcite{SFLRW01} showed that blue stragglers that are formed by stellar collisions have an angular momentum problem.
 The collision products are formed with large total angular momentum. anc some of them are rotating near heir break-up velocity.," The collision products are formed with large total angular momentum, and some of them are rotating near their break-up velocity."
 The collision products are large objects. and they start to contract to the main sequence.," The collision products are large objects, and they start to contract to the main sequence."
 As they contract. they spin up since they have no way to ose any angular momentum.," As they contract, they spin up since they have no way to lose any angular momentum."
 They begin to reach break-up velocities and become unstable., They begin to reach break-up velocities and become unstable.
 Even by shedding material. hey cannot remove enough angular momentum to remain und. and they tear themselves apart.," Even by shedding material, they cannot remove enough angular momentum to remain bound, and they tear themselves apart."
 We are forced to one of two conclusions: either blue stragelers are not mace by stellar collisions (since head-on collisions are very rare. and all olf-axis collisions have too much angular momoentum). or collision products have some wav of losing their angular momentum.," We are forced to one of two conclusions; either blue stragglers are not made by stellar collisions (since head-on collisions are very rare, and all off-axis collisions have too much angular momentum), or collision products have some way of losing their angular momentum."
 The two stancard ways stars are thought to lose angular momentum is through a magnetic wind. or through disk-locking.," The two standard ways stars are thought to lose angular momentum is through a magnetic wind, or through disk-locking."
 We have shown that both these scenarios are not viable. since these collision products do not have surface convection zones or circumstellar disks. immediately: after the collision.," We have shown that both these scenarios are not viable, since these collision products do not have surface convection zones or circumstellar disks immediately after the collision."
 However. it is interesting to note the work of Durisen ct al. (1986)...," However, it is interesting to note the work of Durisen et al. \shortcite{DGTB86}, ,"
 who show that if a polvtrope is rotating very rapidlv. some outer material is thrown olf," who show that if a polytrope is rotating very rapidly, some outer material is thrown off"
The values of these (wo parameters were originally chosen to match the empirical Wennicutt-Schmiclt law (Ixennicutt1998a.b) using simulations of isolated. disk galaxies.,"The values of these two parameters were originally chosen to match the empirical Kennicutt-Schmidt law \citep{Kennicutt:98a,Kennicutt:98b} using simulations of isolated disk galaxies."
" Since we did not change the value of /7 in the runs with metal cooling (as well as pi, in the ""mc. runs). it is possible that our simulations may now violate the Ixennicutt law."," Since we did not change the value of $t_\star^0$ in the runs with metal cooling (as well as $\rhoth$ in the `mc' runs), it is possible that our simulations may now violate the Kennicutt law."
 However. as we show in Figure 11.. the plots of Nur vs. Nspr for the N2IGLIO and N216LI0mc runs are not so different from each other.," However, as we show in Figure \ref{fig:kenni}, , the plots of $\NHI$ vs. $\Sigma_{\rm SFR}$ for the N216L10 and N216L10mc runs are not so different from each other."
 “Phis can be understood. as follows., This can be understood as follows.
 As Figure See showed. metal cooling increases the density ancl lowers the temperature of star-forming eas.," As Figure \ref{fig:Evolv4}c c showed, metal cooling increases the density and lowers the temperature of star-forming gas."
" The fundamental scaling relationship between SER. and cold gas density does not need to change when the metal cooling is introduced: the star-forming gas particles would simply slide upward along the Kennicutt law. Mpx""Ng t. "," The fundamental scaling relationship between SFR and cold gas density does not need to change when the metal cooling is introduced; the star-forming gas particles would simply slide upward along the Kennicutt law, $\Sigma_{\rm SFR} \propto \NHI^{1.4}$ ."
"In the case of N216LI10mwv run. we varied pi, according o the gas metallicity. primarily scattering it to lower values cause higher metallicity increases the cooling rate and ολο οι."," In the case of N216L10mv run, we varied $\rhoth$ according to the gas metallicity, primarily scattering it to lower values because higher metallicity increases the cooling rate and lowers $\rhoth$ ."
" Therefore more gas particles become eligible Oo form stars and “ser is scattered: upwarcl above the Ixennicutt law for a given value of Nyy. resulting in a broader deviation from the power-law relationship of the Kennicutt aw near logVy,c20."," Therefore more gas particles become eligible to form stars and $\Sigma_{\rm SFR}$ is scattered upward above the Kennicutt law for a given value of $\NHI$, resulting in a broader deviation from the power-law relationship of the Kennicutt law near $\log \NHI \simeq 20$."
 At higher values of Na. “Msp is enhanced. but still withinthe range of the Ixennicutt aw.," At higher values of $\NHI$, $\Sigma_{\rm SFR}$ is enhanced, but still withinthe range of the Kennicutt law."
 Thisdillerence canbe alleviatedby adjusting the SE imeseale (seeSprinecl&Lernquist 2003a).. however. we," Thisdifference canbe alleviatedby adjusting the SF timescale \citep[see][]{Springel.Hernquist:03}, , however, we"
than give CC SNe.,than give CC SNe.
 ? also suggest that their SN rate estimates argue against a significant fraction of massive stars collapsing without producing a visible and detectable SN., \citet{Maoz2010} also suggest that their SN rate estimates argue against a significant fraction of massive stars collapsing without producing a visible and detectable SN.
 In other words. that the upper mass limit must be quite high.," In other words, that the upper mass limit must be quite high."
 However as shown in Fig., However as shown in Fig.
 6. the upper mass limit for CC SNe is quite unconstrianed from SN rate measurements abovemsun., \ref{sfrcom} the upper mass limit for CC SNe is quite unconstrianed from SN rate measurements above.
 The CC SN rates themselves can't. quanitatively constrain the upper mass limit within the - rrange (simply due to the steep power law nature of the IMF)., The CC SN rates themselves can't quanitatively constrain the upper mass limit within the $\sim$ range (simply due to the steep power law nature of the IMF).
 The direct detection of progenitor stars in pre-discovery images has provided identifications. mass estimates and mass limits for over 20 CC SNe (see Smartt 2009 for a review).," The direct detection of progenitor stars in pre-discovery images has provided identifications, mass estimates and mass limits for over 20 CC SNe (see Smartt 2009 for a review)."
 ? carried out a detailed and homogeneous analysis of all II-P SNe progenitor searches within 28 Mpe., \citet{Smartt2009} carried out a detailed and homogeneous analysis of all II-P SNe progenitor searches within 28 Mpc.
 A maximum-likelihood analysis gives the best fitting minimum and maximum masses for the SNe ΠΡ progenitors. 8.57! and 16.5x1.5 respectively. assuming a Salpeter IMF.," A maximum-likelihood analysis gives the best fitting minimum and maximum masses for the SNe IIP progenitors, $8.5^{+1}_{-2}$ and $16.5 \pm 1.5$ respectively, assuming a Salpeter IMF."
 The minimum mass ts consistent with our estimate within the errors., The minimum mass is consistent with our estimate within the errors.
 This lower mass limit is consistent with other studies of type II progenitor stars (2?2?22)..," This lower mass limit is consistent with other studies of type II progenitor stars \citep{2005MNRAS.364L..33M,2006ApJ...641.1060L,2008ApJ...688L..91M,2010arXiv1011.5873V,2010arXiv1011.6558F}."
 However some estimates of the hydrodynamic mass of the ejected enevlopes of ΠΡ SNe give systematically higher results. e.g. .in ? and in ?..," However some estimates of the hydrodynamic mass of the ejected enevlopes of IIP SNe give systematically higher results, e.g. in \citet{Zampieri2007} and in \citet{Utrobin2009}."
 While there is no systematic study of a large enough sample to produce an estimate of the lower mass limit. the discrepancy should be taken seriously in attempts to determine masses from both methods.," While there is no systematic study of a large enough sample to produce an estimate of the lower mass limit, the discrepancy should be taken seriously in attempts to determine masses from both methods."
 The mass dividing CC SN progenitors from WD progenitors 1s theoretically expected to lie in the mass range depending on metallicity and the degree of overshooting (?).., The mass dividing CC SN progenitors from WD progenitors is theoretically expected to lie in the mass range depending on metallicity and the degree of overshooting \citep{Siess2007}.
 Observations of most massive WD progenitor in young star clusters provide a lower limit on the value of this mass (?).., Observations of most massive WD progenitor in young star clusters provide a lower limit on the value of this mass \citep{Koester1996}.
 ? and ? studied the WD populations in the open clusters of NGC 6633. NGC 7063 and NGC 2168 and found a lower limit on the maximum mass of WD progenitors betweenmsun.," \citet{Williams2007} and \citet{Williams2009} studied the WD populations in the open clusters of NGC 6633, NGC 7063 and NGC 2168 and found a lower limit on the maximum mass of WD progenitors between."
. This result is also consistent with our estimates of the minimum mass for CC SN progenitors., This result is also consistent with our estimates of the minimum mass for CC SN progenitors.
 The ninimum mass for CC SN progenitors is an important factor 1 the study of ?.. to estimate the birth rate of the Galactic neutron star population.," The minimum mass for CC SN progenitors is an important factor in the study of \cite{Keane2008}, to estimate the birth rate of the Galactic neutron star population."
" They took an estimate of the Milky Way CC SN rate of 1.9+0.9 pper century from the measurements of the Galactic MMeV emission line from the radioactive decay of ""AI (2)..", They took an estimate of the Milky Way CC SN rate of $1.9 \pm 0.9$ per century from the measurements of the Galactic MeV emission line from the radioactive decay of $^{26}$ Al \citep{Diehl2006}. .
 Alternatively. they assumed a lower mass limit ofmsun.. and a Milky Way star-formation rate of .yyr to arrive at the same CC SN rate (1.9+0.9 pper century).," Alternatively, they assumed a lower mass limit of, and a Milky Way star-formation rate of $^{-1}$ to arrive at the same CC SN rate $1.9 \pm 0.9$ per century)."
 This appears to be significantly lower than the observed population of pulsars. rotating radio transients. X-Ray dim isolated neutron stars and magnetars. which imply a neutron star birth rate of 10.877 per century.," This appears to be significantly lower than the observed population of pulsars, rotating radio transients, X-Ray dim isolated neutron stars and magnetars, which imply a neutron star birth rate of $^{+7}_{-5}$ per century."
 Our results in this paper and those on the direct progenitor detections and WD progenitor limits. argue for lower values of mtt.," Our results in this paper and those on the direct progenitor detections and WD progenitor limits, argue for lower values of $\mlcc$."
 It appears that a value of 10-12 .for μή© is disfavoured by the combination of all these studies., It appears that a value of 10-12 for $\mlcc$ is disfavoured by the combination of all these studies.
 A value of would not be inconsistent within the three independent estimates and that would increase the Milky Way SN rate to 44xz2., A value of would not be inconsistent within the three independent estimates and that would increase the Milky Way SN rate to $4.4 \pm 2$.
 This is still below the high neutron star birth rate estimate. but just within the Io error.," This is still below the high neutron star birth rate estimate, but just within the $\sigma$ error."
 The massive star birth and death rates are tightly correlated due to their short lifetime., The massive star birth and death rates are tightly correlated due to their short lifetime.
 We can exploit the CC S rate as a diagnostic of the current SFR by assuming an IMF and à mass range of the CC SN progenitor., We can exploit the CC SN rate as a diagnostic of the current SFR by assuming an IMF and a mass range of the CC SN progenitor.
 Conversely we can obtain a significant constrain on the CC SN progenitor mass range by assuming a SER inferred through the galaxy luminosity., Conversely we can obtain a significant constrain on the CC SN progenitor mass range by assuming a SFR inferred through the galaxy luminosity.
 Only the estimate of CC SN rate in a well defined galaxy sample can provide a direct link between SN rates and different stellar populations., Only the estimate of CC SN rate in a well defined galaxy sample can provide a direct link between SN rates and different stellar populations.
 Complete and volume-limited S and galaxy samples are crucial to perform a statistically meaningful analysis and the advent of large sets of multi-wavelength observations of nearby galaxies from the 11HUGS and LVL programmes provide us. for the first time. the opportunity to compare SFRs based on CC SN rate and more established tracers in the same galaxy sample.," Complete and volume-limited SN and galaxy samples are crucial to perform a statistically meaningful analysis and the advent of large sets of multi-wavelength observations of nearby galaxies from the 11HUGS and LVL programmes provide us, for the first time, the opportunity to compare SFRs based on CC SN rate and more established tracers in the same galaxy sample."
 The data are complete enough that we can take into account the different uncertainties and biases that affect these SFR diagnostics., The data are complete enough that we can take into account the different uncertainties and biases that affect these SFR diagnostics.
 Assuming a lower mass limit cut-off of for CC SN progenitors and a Salpeter IMF for massive stars. we find that the SFR based on Car hat reproduce the observed CC SN rate while there Is a good agreement with SPR based on .in our galaxy sample.," Assuming a lower mass limit cut-off of for CC SN progenitors and a Salpeter IMF for massive stars, we find that the SFR based on can not reproduce the observed CC SN rate while there is a good agreement with SFR based on in our galaxy sample."
 The multi-wavelength data allow tto be corrected by adopting different dust extinction corrections. from either the Balmer decrement or by combining TIR and luminosity.," The multi-wavelength data allow to be corrected by adopting different dust extinction corrections, from either the Balmer decrement or by combining TIR and luminosity."
 Even with this correction. our analysis suggests that may under-estimate the total SFR in our galaxy samples. by nearly a factor of two.," Even with this correction, our analysis suggests that may under-estimate the total SFR in our galaxy samples, by nearly a factor of two."
 A future prospective of this analysis is to. study the connection between SFR tracers and CC SN rate on a galaxy-by-galaxy basis and to compare the spatial distribution of CC SNe with thatof the SFR in spiral arms., A future prospective of this analysis is to study the connection between SFR tracers and CC SN rate on a galaxy-by-galaxy basis and to compare the spatial distribution of CC SNe with thatof the SFR in spiral arms.
 Multi wavelength data also will also allow us to better constrain the dust effect and in, Multi wavelength data also will also allow us to better constrain the dust effect and in
"waveleugthns are less affected by spurious matches,",wavelengths are less affected by spurious matches.
 This males the search for counterparts much easier., This makes the search for counterparts much easier.
 We cross-correlated the DMW-IIRI catalogue with the largest catalogues available at other wavelengths. from radio to optical.," We cross-correlated the BMW-HRI catalogue with the largest catalogues available at other wavelengths, from radio to optical."
 For cross-correlation with catalogues at other wavelengths we used (unless otherwise stated) a search radius of 10 arcsec., For cross-correlation with catalogues at other wavelengths we used (unless otherwise stated) a search radius of 10 arcsec.
 This value comes from the major source of uncertaüntv i the recoustruction of DMW-IIRI source positious.which is the uncertainty iu the aspect solution of the ROSAT telescope (i.c. the boresiglt uncertainty).," This value comes from the major source of uncertainty in the reconstruction of BMW-HRI source positions,which is the uncertainty in the aspect solution of the ROSAT telescope (i.e. the boresight uncertainty)."
 FIRST Faint huages of the Radio Skv at Tweuty-cu covers 10.000. square deerees of the North Galactic Cap.," FIRST – Faint Images of the Radio Sky at Twenty-cm – covers 10,000 square degrees of the North Galactic Cap."
 The sensitivity of the survey is of — L1 imJw with an angular resolution of ~ 5 arcsec (see Decker et al. 1995))., The sensitivity of the survey is of $\sim$ 1 mJy with an angular resolution of $\sim$ 5 arcsec (see Becker et al. \cite{becker95}) ).
 A catalogue containing ~ 770.000 sources and covering  S500 square degrees las been constructed (White et al. 1997)).," A catalogue containing $\sim$ 770,000 sources and covering $\sim$ 8,500 square degrees has been constructed (White et al. \cite{white97}) )."
 The combined seusitivitv aud positional accuracy of the FIRST catalogue are unprecedented compared witli any previous wide-area radio catalogue., The combined sensitivity and positional accuracy of the FIRST catalogue are unprecedented compared with any previous wide-area radio catalogue.
 FIRST source locations lave alb ACCULACH hat matches or exceeds those of all currently available radio catalogues., FIRST source locations have an accuracy that matches or exceeds those of all currently available radio catalogues.
 The cross-correlation with the BMW-IIRI fouud. 1.019 eutries with a uusideutification probability of the order of 24 (that ix 15 nismatches).," The cross-correlation with the BMW-HRI found 1,019 entries with a misidentification probability of the order of $\%$ (that is 18 mismatches)."
 Our cross-correlation is similar to the one made using FIRST aud the catalogue of X sources WCOACAT from ROSAT PSPC observations (White. Cdommi Angelini 1991: see White et al. 1997))," Our cross-correlation is similar to the one made using FIRST and the catalogue of X--ray sources WGACAT from ROSAT PSPC observations (White, Giommi Angelini \cite{white94}; see White et al. \cite{white97}) )"
 and represeuts one of the largest lists of Nrav/radio coincidences available to date., and represents one of the largest lists of X–ray/radio coincidences available to date.
 In Fig., In Fig.
 Ll we report the distributions of the augular separation (r iun arcsec) between the radio aud the Xrav position for the 1.019 matched objects.," 11 we report the distributions of the angular separation $r$ in arcsec) between the radio and the X–ray position for the 1,019 matched objects."
 In Fig., In Fig.
 12 we plot the integrated fins densities πιοασος m niJwv versus Norav flux (fll column density) for the 1.019 cross-correlated sources.," 12 we plot the integrated flux densities measured in mJy versus X–ray flux (full column density) for the 1,019 cross-correlated sources."
 The inteerated fiux is derived by fitting au elliptical Gaussian model to all FIRST sources., The integrated flux is derived by fitting an elliptical Gaussian model to all FIRST sources.
 As inthe FIRST catalogue we have information about the source extension (the major axis. ic. the FWIIM in arcsec derived from the elliptical Gaussian model for the source) we plot in Fig.," As in the FIRST catalogue we have information about the source extension (the major axis, i.e. the FWHM in arcsec derived from the elliptical Gaussian model for the source) we plot in Fig."
 13 the radio exteusion versus, 13 the radio extension versus
well as being potentially crucial to our understuxding of the acceleration of the fast solar wind. the lower hvbrid wave instability is also of interest elsewhere in astrophvsics. and (he Sun and solar svstem offer an attractive “laboratory” for the exploration of collisionless plasma physics processes (hat are otherwise inaccessible to experiment Blackman 2003)..,"well as being potentially crucial to our understanding of the acceleration of the fast solar wind, the lower hybrid wave instability is also of interest elsewhere in astrophysics, and the Sun and solar system offer an attractive “laboratory” for the exploration of collisionless plasma physics processes that are otherwise inaccessible to experiment \citep[see e.g. the recent
debate on electron-ion equilibration in ADAFs in][]{binney03,quataert03,pariev03}."
 The work was supported by NASA Contract SI3783G and by the NRL/ONR. Solar Magnetism and the Earth's Environment 6.1 Research Option., The work was supported by NASA Contract S13783G and by the NRL/ONR Solar Magnetism and the Earth's Environment 6.1 Research Option.
 Il am grateful to Steven Crammer for enlightening discussions., I am grateful to Steven Cranmer for enlightening discussions.
vienetting iu our oobservations.,vignetting in our observations.
 However. it is a huge effect forNewton.. reducing the count rate bv at au offset ofLO”.," However, it is a large effect for, reducing the count rate by at an offset of."
 Therefore. we lave accounted for vignetting in our oobservatious by reducing the fiux-to-couuts conversion & by from the ou-axis value obtained from PIMMS. which is the iue reduction over the iuner107.," Therefore, we have accounted for vignetting in our observations by reducing the flux-to-counts conversion $\xi$ by from the on-axis value obtained from PIMMS, which is the mean reduction over the inner."
.. We estimate that our simple treatment of the vignetting mtroduces, We estimate that our simple treatment of the vignetting introduces
While the above estimate lor Typ Is no substitute for a Goerdt et al.,While the above estimate for $\tau_{df}$ is no substitute for a Goerdt et al.
 (wpe of analysis. il does allow an intercomparison among our cored mocels.," type of analysis, it does allow an intercomparison among our cored models."
 The above model with the largest core has a dvnamical friction time scale of 7.5x10? vears Which is nearly [our times that of the cluster with the smaller but higher density core., The above model with the largest core has a dynamical friction time scale of $7.5\times10^{9}$ years which is nearly four times that of the cluster with the smaller but higher density core.
 As a further check. a model with the same mass and r. but different a was made with parameters (4.25x10©.0.920.1.492. 1.0).," As a further check, a model with the same mass and $_{c}$ but different $\alpha$ was made with parameters $4.25\times10^{-6},0.920,1.49\times10^{-2},
1.0$ )."
 The reeion within which the density gradient is less than -0.1 now has a radius of only 0.05 kpe., The region within which the density gradient is less than -0.1 now has a radius of only 0.05 kpc.
 As expected. its dynamical friction time scale is shorter than the above model with the same r. but only by ~185. (," As expected, its dynamical friction time scale is shorter than the above model with the same $_{c}$ but only by $\sim18\%$. ("
Lowering a while keeping the other parameters the same lowers (he central entropy. slightly).,Lowering $\alpha$ while keeping the other parameters the same lowers the central entropy slightly).
 This simple analvsis is in accord with the conclusions of Goerdt et al., This simple analysis is in accord with the conclusions of Goerdt et al.
 and Sánnchez-Salcedo et al., and Sánnchez-Salcedo et al.
 (hat increasing the core size leads (o an increase in dvnamical [riction time scale., that increasing the core size leads to an increase in dynamical friction time scale.
" This increased lime scale approaches a Lhibble time and because our caleulated value of Q,xo is rol necessarily assumed to be a result of phase packing remains within the constraints imposed bv the sophisticated dvnamical model of Fornax by Strigar et al. (", This increased time scale approaches a Hubble time and because our calculated value of $Q_{o}$ is $not$ necessarily assumed to be a result of phase packing remains within the constraints imposed by the sophisticated dynamical model of Fornax by Strigari et al. (
2006).,2006).
 An additional check on the model comes from the recent work of Gilmore et al. (, An additional check on the model comes from the recent work of Gilmore et al. (
2007) whose analvsis of the light distribution and velocity dispersion profile of several local dwarl spheroidal galaxies shows that shallow (cored) central density. profiles will mean densities of 0.1M.pe.7 (identical to that of our model above with τς—0.385 Apc) are most consistent with the observations.,2007) whose analysis of the light distribution and velocity dispersion profile of several local dwarf spheroidal galaxies shows that shallow (cored) central density profiles with mean densities of $0.1~ M_{\odot}pc^{-3}$ (identical to that of our model above with $_{c}=0.385~ kpc$ ) are most consistent with the observations.
 A new model for dark matter halos has been proposed ancl is successfully applied to observations of objects with masses ranging from ~10? to ~LOY AL.., A new model for dark matter halos has been proposed and is successfully applied to observations of objects with masses ranging from $\sim10^{9}$ to $\sim10^{15}$ $_{\odot}$.
 Fig., Fig.
 3 is a graphical summary of the results which shows the scaling relations discussed earlier., 3 is a graphical summary of the results which shows the scaling relations discussed earlier.
 It is important to note key differences in structure occur wilh different scaling normalizations., It is important to note key differences in structure occur with different scaling normalizations.
" The [illed svinbols are structures with relatively high values of p, (~O.LAL.pe 7) while the open svinbols are structures with lower values of p, (~0.01.pe 7) and lower ρωs. Close examination ol Fig."," The filled symbols are structures with relatively high values of $\rho_{o}$ $\sim 0.1 M_{\odot}~ pc^{-3}$ ) while the open symbols are structures with lower values of $\rho_{o}$ $\sim 0.01 M_{\odot}
~ pc^{-3}$ ) and lower $\beta_{max}$ 's. Close examination of Fig."
 3 reveals a real svstematic shift between these (wo groups of objects., 3 reveals a real systematic shift between these two groups of objects.
 Objects with identical central densities would lie essentially dispersionless along; a line of the indicated slope., Objects with identical central densities would lie essentially dispersionless along a line of the indicated slope.
" Further. for (wo models with the same mass. the one with the higher Q, has the higher p, and smaller ο. (i.e. quantitatively 9 log Q,/O log p,=0.67 and 29 log Q,/O log"," Further, for two models with the same mass, the one with the higher $Q_{o}$ has the higher $\rho_{o}$ and smaller $r_{c}$ (i.e. quantitatively $\partial$ log $_{o}$ $\partial$ log $
\rho_{o}=0.67$ and $\partial$ log $_{o}$ $\partial$ log"
optimized reference PSF subtraction is a temporal filler that guarantees that no noise can have a decorrelation time longer than the time needed lor a ~1.5 EFWIIM FOV rotation.,optimized reference PSF subtraction is a temporal filter that guarantees that no noise can have a decorrelation time longer than the time needed for a $\sim 1.5$ FWHM FOV rotation.
 Fig., Fig.
 5 shows detection limits (50) in magnitude clifference as a funetion of angular separation obtained with the ADI technique for all three ADI targets., \ref{fig5} shows detection limits $\sigma$ ) in magnitude difference as a function of angular separation obtained with the ADI technique for all three ADI targets.
 To produce this ligure. data were reduced following the procedure explained in section ??..," To produce this figure, data were reduced following the procedure explained in section \ref{reduc}."
 No unsharp mask was used for the ADI reduction since multiple tests have shown that while this filter is effective lor suppressing the low-Irequency spatial noise. it does not improve candidate S/N and (he photometry is slightIv biased in the process.," No unsharp mask was used for the ADI reduction since multiple tests have shown that while this filter is effective for suppressing the low-frequency spatial noise, it does not improve candidate S/N and the photometry is slightly biased in the process."
 The detection limits are calculated using the ratio of simulated companion peak intensities over (he noise in the residual image as a function of angular separation., The detection limits are calculated using the ratio of simulated companion peak intensities over the noise in the residual image as a function of angular separation.
 The flux normalized unsaturated PSF was then used {ο simulate smeared Companion PSFs as a function of angular separation., The flux normalized unsaturated PSF was then used to simulate smeared companion PSFs as a function of angular separation.
 The noise is calculated inside aunuli of increasing diameter and width equal to 1 PSF FWIHAL, The noise is calculated inside annuli of increasing diameter and width equal to 1 PSF FWHM.
" To account for the PSF smearing effect due to FOV rotation. the image and simulated companions were convolved with elliptical Gaussian of one FWIIM in the radial direction and one EWIIM plus a smearing (erm in (he azimuthal direction. (he smearing term ranging from zero at the center to typically two EWIIM at 10"". depending on the rotation rate of the FOV."," To account for the PSF smearing effect due to FOV rotation, the image and simulated companions were convolved with elliptical Gaussian of one FWHM in the radial direction and one FWHM plus a smearing term in the azimuthal direction, the smearing term ranging from zero at the center to typically two FWHM at $^{\prime \prime}$, depending on the rotation rate of the FOV."
 Finally. detection limits were corrected [or the Altair estimated anisoplanatism following the Strehl S equation found in the Gemini web where 8 is expressed in arcsec.," Finally, detection limits were corrected for the Altair estimated anisoplanatism following the Strehl $S$ equation found in the Gemini web where $\theta$ is expressed in arcsec."
" The detection limit obtained with the ADI technique on Vega ab separations greater than ~5"" is (wo orders of magnitude deeper than the Palomar Z-band image (Metcelievetal.2003) and approximately a [actor of ten deeper (han the heck A-band image of (Macintoshetal.2003).", The detection limit obtained with the ADI technique on Vega at separations greater than $\sim$ $^{\prime \prime}$ is two orders of magnitude deeper than the Palomar $H$ -band image \citep{metchev2003} and approximately a factor of ten deeper than the Keck $K$ -band image of \citep{macintosh2003}.
. The current speckle attenuation (20-100) achieved with ADI is comparable or better to what is currently obtained (50) from one HIST orbit using the roll subtraction technique., The current speckle attenuation (20-100) achieved with ADI is comparable or better to what is currently obtained (50) from one HST orbit using the roll subtraction technique.
 Although the ADI technique is inherently optimized [or relatively large separations. the good Gemini PSF stability enables excellent. performances at arcesec separations.," Although the ADI technique is inherently optimized for relatively large separations, the good Gemini PSF stability enables excellent performances at sub-arcsec separations."
" Indeed. the ADI contrast of Am=110.1—11.9 (56) at 0.3"" obtained on IID13303 and ILD97334D equals the Am=11.0 (50) at 0.8"" obtained with the simultaneous"," Indeed, the ADI contrast of $\Delta m = 11.1 - 11.9$ $\sigma$ ) at $^{\prime \prime}$ obtained on HD18803 and HD97334B equals the $\Delta m = 11.0$ $\sigma$ ) at $^{\prime \prime}$ obtained with the simultaneous"
GAs are a family of computer models inspired on natural evolution.,GAs are a family of computer models inspired on natural evolution.
 Their basis is to assume a potential solution for a specific problem. viewed as chromosome like structure. on which are applied. genetic operators (mutation. crossover. adaptation and evolution).," Their basis is to assume a potential solution for a specific problem, viewed as chromosome like structure, on which are applied genetic operators (mutation, crossover, adaptation and evolution)."
 The use of this technique simplifies the formulation. and solution of. optimisation problems. and. parallel simultaneous procedure approach is implicit in the method. providing evaluation of the viability ofa parameter set as possible solution for complex problems (Ixoza1992.1994:Llolland1992).," The use of this technique simplifies the formulation and solution of optimisation problems, and parallel simultaneous procedure approach is implicit in the method, providing evaluation of the viability of a parameter set as possible solution for complex problems \citep{Koza92,Koza94,Holland92}."
. The implementation of a CA starts with a random. generation of sets οἱ chromosomes EINE⋅⋅⋅⋅⋅⋅↔∣⊔↴⊳ which are evaluated: ancl associated to an adaptation. probability. V. obtained by the evaluation function. which is in our case the DDNOL model.," The implementation of a GA starts with a random generation of sets of chromosomes $\left[ {S_1,S_2,\ldots,S_n } \right]$, which are evaluated and associated to an adaptation probability, $\chi^2$, obtained by the evaluation function, which is in our case the DDN01 model."
 Phe 47. values express how each individual are adapted. or how each solution are close to the best solution (Bentley&Corne2002).," The $\chi^2_i$ values express how each individual are adapted, or how each solution are close to the best solution \citep{Bentley02}."
. Then. the judgement function determines the ecnetic operator to be applied to a solution. and its values can be copy: the individual remains the same in the next generation. crossover: the individual is elected. to change a number of genes (parameters) with another individual. creating a new one. mutation: one of its genes shall be randomly changed. or termination: none of the genes may continue on next eenerations.," Then, the judgement function determines the genetic operator to be applied to a solution, and its values can be copy: the individual remains the same in the next generation, crossover: the individual is elected to change a number of genes (parameters) with another individual, creating a new one, mutation: one of its genes shall be randomly changed, or termination: none of the genes may continue on next generations."
 The chosen action is expressed bv the 9; variable. associated to each individual.," The chosen action is expressed by the $\Phi_i$ variable, associated to each individual."
 Phe further step is to evolve the current generation (4) to the next (A| 1). what is done through a procedure E that considers the solutions and the genetic operators designed. by d;.," The further step is to evolve the current generation $k$ ) to the next $k+1$ ), what is done through a procedure $\Gamma$ that considers the solutions and the genetic operators designed by $\Phi_i$."
 Formally As soon as a new generation is ready. the evaluation function is reapplied. and the algorithm repeats. the described actions until an end of loop condition is reached.," Formally As soon as a new generation is ready, the evaluation function is reapplied, and the algorithm repeats the described actions until an end of loop condition is reached."
 The end condition can be based on the number of iterations or quality (a low level for the «7 values) (Figure Al))., The end condition can be based on the number of iterations or quality (a low level for the $\chi^2_i$ values) (Figure \ref{stepsGA}) ).
" Once an acceptable mininium has been found. one computes the inverse of the Lessian matrix that has its components given by where Og(A,)/00; means the partial derivative of the SED with respect to parameter e; at AΑι and Nis the number of observed data points."," Once an acceptable minimum has been found, one computes the inverse of the Hessian matrix that has its components given by where $\partial y(\lambda_k) / \partial a_i$ means the partial derivative of the SED with respect to parameter $a_i$ at $\lambda = \lambda_k$, and $N$ is the number of observed data points."
 The components of the main diagonal of €' can be used to estimate the error bars on each. parameter by (Pressetal.1995): CGLIO2 used a Lat disc model assuming that the star and the disc are surrounced by a thin dust shell., The components of the main diagonal of $C$ can be used to estimate the error bars on each parameter by \citep{Press95}: GH02 used a flat disc model assuming that the star and the disc are surrounded by a thin dust shell.
 Phe emission from the star and from the disc are attenuated by the opacity of the envelope., The emission from the star and from the disc are attenuated by the opacity of the envelope.
 The tux is caleulated by assuming blackbocdy emission and dillerent temperature laws for the dise and the envelope., The flux is calculated by assuming blackbody emission and different temperature laws for the disc and the envelope.
 The model considers the following [ree parameters: radius A4. of the star. radius fy ane inclination ϐ of the disc. radius A. and total optical depth 7 of the envelope.," The model considers the following free parameters: radius $R_s$ of the star, radius $R_d$ and inclination $\theta$ of the disc, radius $R_e$ and total optical depth $\tau$ of the envelope."
 The stellar temperature 75 Adams&Shu(1986):: (Chiangetal.2001).. Epchteinet," The stellar temperature $T_s$ \cite{adams86}: \citep{Chiang01}. \cite{epchtein90},"
that ~10€ of the star fornation is locked im FIR and 60% in UV.,that $\sim 40\%$ of the star formation is locked in FIR and $60\%$ in UV.
 The resulting UV extinction is 0.58 mae aud PSFR0.025bhADἐνλpe?," The resulting UV extinction is 0.58 mag and $\rm \rho_{SFR}~ =~ 
0.025~h~M\odot/yr/Mpc^3$."
 Therefore WO Call COLude that the derivation of the global star formation τιite is in oagreenient with our estimate of the elobal extineion iun UV and that the same amount of star formation ra eds traced |x the global FIR and UV (not corrected for exiunction) οςτν densities., Therefore we can conclude that the derivation of the global star formation rate is in agreement with our estimate of the global extinction in UV and that the same amount of star formation rate is traced by the global FIR and UV (not corrected for extinction) luminosity densities.
 Since the IRAS sirvov. FIR bright galaxies have Όσοι the subject of numerous studies because these objects experiment an intense star formation activity.," Since the IRAS survey, FIR bright galaxies have been the subject of numerous studies because these objects experiment an intense star formation activity."
 The extreme case is that of UltraLuminous Iufrared Calaxies (ULICi) with a bolometric αποαν larger thui 1022AL-vr essentially eiudted iun FIR aud star ornation rates of several hundreds soar Hlasses por venir: hey generally are violent inergers and may represent au iuportaunt phase in the formation of large galaxies like elipticals (Mirabel Saunders 1996))., The extreme case is that of UltraLuminous Infrared Galaxies (ULIGs) with a bolometric luminosity larger than $\rm 10^{12} M\odot/yr$ essentially emitted in FIR and star formation rates of several hundreds solar masses per year: they generally are violent mergers and may represent an important phase in the formation of large galaxies like ellipticals (Mirabel Sanders \cite{sandmir}) ).
 Such objects are kiwn to he rare at low z but they mielt be far more uuucrous at high z as sugeested by the sub ΠΠ sus(vs with SCUBA (e.g. Sanders 1999))., Such objects are known to be rare at low z but they might be far more numerous at high z as suggested by the sub millimetric surveys with SCUBA (e.g. Sanders \cite{sanders}) ).
 With the launch of the ISO satellie. the seusitivity of the ISOCAAL camera vas allowed 1üxLinfrared surveys at intermediate redshift (z< 1)," With the launch of the ISO satellite, the sensitivity of the ISOCAM camera has allowed mid-infrared surveys at intermediate redshift $\rm z<1$ )."
 In particular Flores et al. (1998)), In particular Flores et al. \cite{flores}) )
 have observed one CFRS field. therefore UV (0.25411) ancl infrared data are available for these galaxies.," have observed one CFRS field, therefore UV $\mu$ m) and infrared data are available for these galaxies."
 We now compare the FIR and UW properties of these ealaxies (ULIC aud ISOCAAM/CFRS) to that of our IRAS/FOCA sample of nearby galaxies., We now compare the FIR and UV properties of these galaxies (ULIG and ISOCAM/CFRS) to that of our IRAS/FOCA sample of nearby galaxies.
 The comparison is rather straightforward since all these objects are IR selected., The comparison is rather straightforward since all these objects are IR selected.
 Treutliun e al., Trentham et al.
 (1999 ) have obtained IST observatious for three utra hnünous infrared galaxies: VII Zw031. IRAS F121]51055. IRAS F22191-18508.," \cite{trentham} ) have obtained HST observations for three ultra luminous infrared galaxies: VII Zw031, IRAS F12112+0305, IRAS F22491-1808."
 Tjese Galaxies are selected to be cool in order to avoid a uou thermal oriein for the FIR Cluission., These galaxies are selected to be cool in order to avoid a non thermal origin for the FIR emission.
 We can calcilate directly their Ley/Lv ratio using the data at 0.23 yan for the UV oenissioi," We can calculate directly their $\rm
L_{60}/L_{UV}$ ratio using the data at 0.23 $\mu$ m for the UV emission."
 The three objects are repored in figure 6 (simular to fig.2b) with empty stars for svaubols., The three objects are reported in figure 6 (similar to fig.2b) with empty stars for symbols.
 As expected for this type of objects they appear to be very uuinous at 60 gin with a high FIR to UV fx ratio., As expected for this type of objects they appear to be very luminous at 60 $\mu$ m with a high FIR to UV flux ratio.
 Such objects are not represented in our FIR seleced saniple of rearby eaaxies: this cinyhasizes how much these objects are rare iu the local Universe aud with extreue properties as often τιiderlined (c.g. Sanders Mirabel 19993)., Such objects are not represented in our FIR selected sample of nearby galaxies: this emphasizes how much these objects are rare in the local Universe and with extreme properties as often underlined (e.g. Sanders Mirabel \cite{sanders}) ).
 Since he three ULIC have also been detected atlXO dun we can estimate trein UV extinction (ve ucelect the difference in he UV wavelengths i.e. 0.23 pau versus 0.2 gan)., Since the three ULIG have also been detected at 100 $\mu$ m we can estimate their UV extinction (we neglect the difference in the UV wavelengths i.e. 0.23 $\mu$ m versus 0.2 $\mu$ m).
 We find 6.5 mag: more than 994 of the UV flux of these objects is emitted in the FIR., We find $\sim 6.5$ mag: more than $\%$ of the UV flux of these objects is emitted in the FIR.
 IIughes et al. (]908)), Hughes et al. \cite{hughes}) )
 have observed the ΠΟΤΕ field at S50 gm Wih SCUBA., have observed the HDF field at 850 $\mu$ m with SCUBA.
 5 objects detected by SCUBA in the IDF field have been tentatively associated to optical sources for whic1 photometric redshift are available but such an idetification is difficult because of the uncertainty on the 850 ja positions., 5 objects detected by SCUBA in the HDF field have been tentatively associated to optical sources for which photometric redshift are available but such an identification is difficult because of the uncertainty on the 850 $\mu$ m positions.
 Indeed. the identification of the most byieliest source (IIDFs50.1) has not beeu confirmed (Saunders. 1999) )," Indeed, the identification of the most brightest source (HDF850.1) has not been confirmed (Sanders, \cite{sanders}) )."
 Moreover the nature of these sources. starbursts or AGN. is not clear: at FIR huuinositv larger tha1 10?L. about half of the nearby ULIGs are predoninalv powered by ACNs (e.g. Sanders 1999).," Moreover the nature of these sources, starbursts or AGN, is not clear: at FIR luminosity larger than $\rm
10^{12}~L_{\odot}$ about half of the nearby ULIGs are predominantly powered by AGNs (e.g. Sanders \cite{sanders}) )."
 The 6) qna huninosity of the [| remaining galaxies is obtained roni t101 Cluission at 850 pu accounting for the redshifting aud ALL assulue spectral energv distribution chosen to be fvat of ALS2., The 60 $\mu$ m luminosity of the 4 remaining galaxies is obtained from their emission at 850 $\mu$ m accounting for the redshifting and an assumed spectral energy distribution chosen to be that of M82.
 The optical data from the TDF lea to the estimate of the UW flux at a rest frame waveleueth of 0.28 juu. Al these estimates rely on the resenmiblauce of all ULICs with M82 aud cau lead to false results (e.g. Sanders 1999) )., The optical data from the HDF lead to the estimate of the UV flux at a rest frame wavelength of 0.28 $\mu$ m. All these estimates rely on the resemblance of all ULIGs with M82 and can lead to false results (e.g. Sanders \cite{sanders}) ).
 In spite of these caveats. we have reported the L| high redshift ealaxies iu figure 6 (filled triangles).," In spite of these caveats, we have reported the 4 high redshift galaxies in figure 6 (filled triangles)."
 They appear very extreme. bee more huninous in FIR aud probably more extincted that all t1C other galaxies studied iu this paper.," They appear very extreme, being more luminous in FIR and probably more extincted that all the other galaxies studied in this paper."
 A teutative estimate of the extinction is obtained by using the Foo/Fry iustead of the FeFrv one., A tentative estimate of the extinction is obtained by using the $\rm F_{60}/F_{UV}$ instead of the $\rm F_{FIR}/F_{UV}$ one.
 We find values spanning frou 5 o ll mag., We find values spanning from 8 to 11 mag.
 As a comparison À82. which belongs to our IRAS/FOCA sample exhibit oil 5.1 nag of extinction at 0.2 jan. These high redshift ULIGs secur also to be uuch more extincted than the most hnuuinous Lyman xeak ealaxics of the TDF studied. by. Meurer et al. (19991) ," As a comparison M82, which belongs to our IRAS/FOCA sample exhibit ""only"" 5.4 mag of extinction at 0.2 $\mu$ m. These high redshift ULIGs seem also to be much more extincted than the most luminous Lyman break galaxies of the HDF studied by Meurer et al. \cite{meurheck}) )"
for which they derive an extinction. not. larger han 3.5 mae., for which they derive an extinction not larger than 3.5 mag.
 Although their UV luninosity corrected for extinction are comparable (~142D Lj these two cances of galaxies do not seeu to exhivit the same properties iu FIR aud UV as suggested by Heckiau (1999)).," Although their UV luminosity corrected for extinction are comparable $\rm 
\sim 
10^{12}~ 
L\odot$ ), these two classes of galaxies do not seem to exhibit the same properties in FIR and UV as suggested by Heckman \cite{heckman}) )."
 Iucleed. we can try to roughly locate the most luninous eaANICy. of Meurer ct al., Indeed we can try to roughly locate the most luminous galaxies of Meurer et al.
 in figure 6., in figure 6.
 The FIR luuinosity cal be estimated from ficin star formation rates and the Fou/Fuy ratio from heir extinction using the figure 1., The FIR luminosity can be estimated from their star formation rates and the $\rm F_{60}/F_{UV}$ ratio from their extinction using the figure 1.
 It gives Loy~107L- aud Foy/Fry~L2.Let for an extinction of ~:) inae., It gives $\rm L_{60}\sim 10^{12} L\odot$ and $\rm F_{60}/F_{UV}\sim 1.2-1.4$ for an extinction of $\sim 3$ mag.
 Therefore it secs that the Lyman Break CaANICS detected in the IIDF by their U-dropout do not folow the steep trend. of figure 6 fouud for FIR bright galaxies bit instead exhibit a lower increase of the extinction with the iutrinsic huninosity of the ealaxics., Therefore it seems that the Lyman Break Galaxies detected in the HDF by their U-dropout do not follow the steep trend of figure 6 found for FIR bright galaxies but instead exhibit a lower increase of the extinction with the intrinsic luminosity of the galaxies.
 Such a differeuce uav be due to the coutribution of ACN& in ULIGs., Such a difference may be due to the contribution of AGNs in ULIGs.
 Indecc the extrapolation of the mean trend found iut16 IRAS/FOCA sample (figure 3) reported as a, Indeed the extrapolation of the mean trend found in the IRAS/FOCA sample (figure 3) reported as a
the observed rotational triplets.,the observed rotational triplets.
" To this end, we have performed a simple experiment: we removed one of the observed triplets from our computations, keeping the six remainder, and searched for a solution."," To this end, we have performed a simple experiment: we removed one of the observed triplets from our computations, keeping the six remainder, and searched for a solution."
" We started with the k=22 complex, that naively could be thought to have the strongest influence on our results due to its large frequency separations."," We started with the $k= 22$ complex, that naively could be thought to have the strongest influence on our results due to its large frequency separations."
" Indeed, with an average frequency spacing of 4.53 Hz, it is significantly larger than the other spacings (see Table 2))."," Indeed, with an average frequency spacing of $4.53\, \mu$ Hz, it is significantly larger than the other spacings (see Table \ref{table2}) )."
" We again find differential rotation with the core spinning more than twice faster than the surface, being the variations in the values of Ως and (X. negligible when compared with the case in which all the seven splittings are taken into account."," We again find differential rotation with the core spinning more than twice faster than the surface, being the variations in the values of $\Omega_{\rm
s}$ and $\Omega_{\rm c}$ negligible when compared with the case in which all the seven splittings are taken into account."
" In fact, in this case, we have Ως=10.75 wHz and €,=4.40 wHz."," In fact, in this case, we have $\Omega_{\rm c}= 10.75 \, \mu$ Hz and $\Omega_{\rm
s}= 4.40 \, \mu$ Hz."
" THe similarity of the solutions is due to its large observational uncertainty, which, according to Eq. (1)),"," THe similarity of the solutions is due to its large observational uncertainty, which, according to Eq. \ref{xi}) ),"
 strongly attenuates its impact., strongly attenuates its impact.
 A similar experiment in which we remove the k—24 triplet leads to very similar results., A similar experiment in which we remove the $k= 24$ triplet leads to very similar results.
" Clearly, these two triplets, which suffer from the largest uncertainties in the list, have no appreciable influence in our results."," Clearly, these two triplets, which suffer from the largest uncertainties in the list, have no appreciable influence in our results."
 We repeated this experiment with the rest of the splittings., We repeated this experiment with the rest of the splittings.
" We found that the most critical triplets are, in order of decreasing importance, k—18,17,14 and 15."," We found that the most critical triplets are, in order of decreasing importance, $k= 18, 17, 14$ and $15$."
 The frequency separations for these triplets are accurately known., The frequency separations for these triplets are accurately known.
" In particular, if we discard the k=18 triplet, the solution becomes €),=16.45 uHz and Ὡς=0.95 Hz, that is, strong differential rotation."," In particular, if we discard the $k= 18$ triplet, the solution becomes $\Omega_{\rm c}= 16.45 \, \mu$ Hz and $\Omega_{\rm s}= 0.95 \,
\mu$ Hz, that is, strong differential rotation."
" On the other hand, if the k—17 triplet is not considered, the solution turns out be compatible with rigid rotation, with Ως=7.70 wHz and Ὡς=6.55 Hz."," On the other hand, if the $k= 17$ triplet is not considered, the solution turns out be compatible with rigid rotation, with $\Omega_{\rm c}= 7.70 \, \mu$ Hz and $\Omega_{\rm
 s}= 6.55 \, \mu$ Hz."
" Hence, our results rely mostly on these two triplets, for which the frequency spacings are well determined (although see Sect. 3.4))."," Hence, our results rely mostly on these two triplets, for which the frequency spacings are well determined (although see Sect. \ref{drift}) )."
" Finally, we examined the quality of the match between the observed and theoretical frequency splittings."," Finally, we examined the quality of the match between the observed and theoretical frequency splittings."
" We find that the best fit solution reproduces the observed frequency splittings with a mean difference of z0.25wHz, which is reduced to z0.1uLHz when we do not consider the difference corresponding the k=22 triplet."," We find that the best fit solution reproduces the observed frequency splittings with a mean difference of $\approx 0.25\ \mu$ Hz, which is reduced to $\approx 0.1\ \mu$ Hz when we do not consider the difference corresponding the $k= 22$ triplet."
" In order to assess the uncertainties in the derived parameters, we repeated our optimization procedure adding artificial, normally (Gaussian) distributed uncertainties to the set of 14 observed splittings, with a standard deviation Of Onoise=0.08 «Hz, which is comparable with the best data available for pulsating white dwarf and pre-white dwarf stars observed with the Whole Earth Telescope (Kawaler et al."," In order to assess the uncertainties in the derived parameters, we repeated our optimization procedure adding artificial, normally (Gaussian) distributed uncertainties to the set of 14 observed splittings, with a standard deviation of $\sigma_{\rm noise}= 0.08$ $\mu$ Hz, which is comparable with the best data available for pulsating white dwarf and pre-white dwarf stars observed with the Whole Earth Telescope (Kawaler et al."
 1999)., 1999).
 We performed about 400 realizations of this type., We performed about 400 realizations of this type.
 The resulting solutions are shown in the left panel of Fig., The resulting solutions are shown in the left panel of Fig.
 5 with blue dots., \ref{errores} with blue dots.
" Assuming that the distribution of solutions (Qe, Qs) is also Gaussian, we can estimate its dispersion, c."," Assuming that the distribution of solutions $\Omega_{\rm c}$, $\Omega_{\rm s}$ ) is also Gaussian, we can estimate its dispersion, $\sigma$."
" Wefind that the solutions that deviate more from the best-fit solution (free of uncertainties, black circle) are located at more than 2c from the line that defines the solutions of rigid- rotation (green line)."," Wefind that the solutions that deviate more from the best-fit solution (free of uncertainties, black circle) are located at more than $2 \sigma$ from the line that defines the solutions of rigid-body rotation (green line)."
" In other words, rigid rotation can be discarded at a level of confidence of more than 2c."," In other words, rigid rotation can be discarded at a level of confidence of more than $2 \sigma$."
 We find Ως=10.62+ 1.8yHz and Ὡς=4.41+1.14 Hz.," We find $\Omega_{\rm c}= 10.62 \pm 1.8 \, \mu$ Hz and $\Omega_{\rm s}= 4.41 \pm 1.1 \, \mu$ Hz."
 We have repeated the above analysis adopting a deviation Of Onoise=0.12 uHz for the distribution of uncertainties in the frequency splittings (right panel in Fig. 5)).," We have repeated the above analysis adopting a deviation of $\sigma_{\rm noise}= 0.12 \, \mu$ Hz for the distribution of uncertainties in the frequency splittings (right panel in Fig. \ref{errores}) )."
 These artificial uncertainties are by far larger than the average of the uncertainties in frequencies quoted by Fu et al. (, These artificial uncertainties are by far larger than the average of the uncertainties in frequencies quoted by Fu et al. (
2007) in their table 4.,2007) in their table 4.
" Even in this extreme case, we can discard uniform rotation at a level of confidence of more than 1.5c."," Even in this extreme case, we can discard uniform rotation at a level of confidence of more than $1.5\sigma$."
" Thus, the conclusion of differential rotation for remains unchanged even when we consider exaggerated uncertainties in the measured frequency splittings."," Thus, the conclusion of differential rotation for remains unchanged even when we consider exaggerated uncertainties in the measured frequency splittings."
" Finally, we investigated the effects of uncertainties in the asteroseismological model on our results, and we found that they are not relevant."," Finally, we investigated the effects of uncertainties in the asteroseismological model on our results, and we found that they are not relevant."
 Vauclair et al. (, Vauclair et al. (
2011) have published the results of a comprehensive monitoring of seven oscillation frequencies of this star.,2011) have published the results of a comprehensive monitoring of seven oscillation frequencies of this star.
" They report changes of these oscillation frequencies over time, with much larger amplitudes and shorter time scales than those expected by cooling, although the data resolution is rather low, of the order of 1 µΗΖ."," They report changes of these oscillation frequencies over time, with much larger amplitudes and shorter time scales than those expected by cooling, although the data resolution is rather low, of the order of $1\,
\mu$ Hz."
 We focus on Table 4 of Vauclair et al. (, We focus on Table 4 of Vauclair et al. (
"2011), which shows the frequency and amplitude variations of the 7 largest amplitude modes of ccorresponding to the triplets centered at 2224 wHz (k= 17) and 2493 Hz (k= 15).","2011), which shows the frequency and amplitude variations of the 7 largest amplitude modes of corresponding to the triplets centered at $2224\ \mu$ Hz $k= 17$ ) and $2493\ \mu$ Hz $k= 15$ )."
 These two triplets are the only ones (out of seven triplets present in the star) which have been well documented to exhibit changes over time., These two triplets are the only ones (out of seven triplets present in the star) which have been well documented to exhibit changes over time.
 From Table 4 of Vauclair et al. (, From Table 4 of Vauclair et al. (
"2011) we have computed the average value of the rotational shifts for each triplet as: where 64(4)=VUm=+1)—V(m-0) and ὃν)= at a given epoch t, of observation, where v(m-o)n= Nobs, being Nops=6 for the triplet centered at","2011) we have computed the average value of the rotational shifts for each triplet as: where $\delta \nu_{(+)}= \nu_{(m= +1)}-\nu_{(m= 0)}$ and $\delta \nu_{(-)}= \nu_{(m= 0)}-\nu_{(m= -1)}$ at a given epoch $t_n$ of observation, where $n=
1,\cdots, N_{\rm obs}$ , being $N_{\rm obs}= 6$ for the triplet centered at"
solar Zr abundance.,solar Zr abundance.
 Figure 18 shows synthetic spectra fits to two lines., Figure \ref{f_zr_synthesis} shows synthetic spectra fits to two lines.
" 0.2 cm The Ba abundance is based on the line, the weakest line of RMT 2."," 0.2 cm The Ba abundance is based on the line, the weakest line of RMT 2."
 The gf-value adopted is the mean of the experimental values from Gallagher (1967) and Davidson et al. (, The $gf$ -value adopted is the mean of the experimental values from Gallagher (1967) and Davidson et al. (
1992).,1992).
 Hyperfine and isotopic splittings are taken into account from McWilliam (1998)., Hyperfine and isotopic splittings are taken into account from McWilliam (1998).
 The stronger lines from RMT 2 and the resonance lines (RMT 1) give a roughly 0.7 dex higher Ba abundance and are not considered further., The stronger lines from RMT 2 and the resonance lines (RMT 1) give a roughly 0.7 dex higher Ba abundance and are not considered further.
 0.2 cm The line at 5114 rregion was too weak to measure but used to set the upper limit: log e(La) < 0.31., 0.2 cm The line at 5114 region was too weak to measure but used to set the upper limit: $\log\epsilon$ (La) $\leq$ 0.31.
 The g f-value for the line is taken from Lawler et al. (, The $gf$ -value for the line is taken from Lawler et al. (
20014).,2001a).
 0.2 cm The absent line at 4303 iis used to set an upper limit: log e(Nd) < 0.44., 0.2 cm The absent line at 4303 is used to set an upper limit: $\log\epsilon$ (Nd) $\leq$ 0.44.
 The g f-value for the line is taken from Den Hartog et al. (, The $gf$ -value for the line is taken from Den Hartog et al. (
2003).,2003).
 0.2 cm The resonance lines are aand wwere searched for in the spectra., 0.2 cm The resonance lines are and were searched for in the spectra.
" The g f-values, hyperfine and isotopic structure are taken from Lawler et al. ("," The $gf$ -values, hyperfine and isotopic structure are taken from Lawler et al. ("
2001b).,2001b).
 Spectrum synthesis of the lline gives the abundance in Table 5., Spectrum synthesis of the line gives the abundance in Table 5.
 The Iline is too seriously blended to yield a useful a Eu abundance., The line is too seriously blended to yield a useful a Eu abundance.
 Mean abundances are summarized in Table 5., Mean abundances are summarized in Table 5.
" Absolute uncertainties for the abundances arising from uncertainties of the atmospheric parameters Te, logg, and € are summarized in Table 6 for changes with respect to the model of +150 K, 40.5 cm ?, and +1.0 km s! for representative lines (i.e., the Co entries are based on the EW upper limits for the 3995 aand 4121 llines)."," Absolute uncertainties for the abundances arising from uncertainties of the atmospheric parameters $T_{\rm eff}$, $\log\,g$, and $\xi$ are summarized in Table 6 for changes with respect to the model of +150 K, +0.5 cm $^{-2}$, and $\pm$ 1.0 km $^{-1}$ for representative lines (i.e., the Co entries are based on the EW upper limits for the 3995 and 4121 lines)."
" From the uncertainties listed in Table 6, we find the total absolute uncertainty to be ranging from 0.08 for to 0.19 for by taking the square root of the sum of the square of individual errors (for each species) associated with uncertainties in temperature, gravity, and microturbulent velocity."," From the uncertainties listed in Table 6, we find the total absolute uncertainty to be ranging from 0.08 for to 0.19 for by taking the square root of the sum of the square of individual errors (for each species) associated with uncertainties in temperature, gravity, and microturbulent velocity."
" In light of the line-to-line scatter of abundances, the absolute uncertainties, and, more importantly, the probability that"," In light of the line-to-line scatter of abundances, the absolute uncertainties, and, more importantly, the probability that"
open field lines firmly anchored on a rapidly rotating disk. and centrifugal acceleration of (hud parcels along the field lines into a higli-speed wind.,"open field lines firmly anchored on a rapidly rotating disk, and centrifugal acceleration of fluid parcels along the field lines into a high-speed wind."
 The wind trails behind the clisk in rotation. generating a toroidal component of magnetic field. which exerts a braking torque on the disk.," The wind trails behind the disk in rotation, generating a toroidal component of magnetic field, which exerts a braking torque on the disk."
 It is this magnetic torque Chat is responsible for extracting both energv. and angular monmentun trom (he disk and for powering the wind., It is this magnetic torque that is responsible for extracting both energy and angular momentum from the disk and for powering the wind.
 Since (he energv extracted is simply (he work done by the rotating disk against the magnetic torque jitealts96)). the rate of enerev extraction is directly. proportional {ο the rate of angular --noment(un extraction. with the proporGonality constant being the angular speed O of je disk rotation.," Since the energy extracted is simply the work done by the rotating disk against the magnetic torque \\citealt{s96}) ), the rate of energy extraction is directly proportional to the rate of angular momentum extraction, with the proportionality constant being the angular speed $\Omega_0$ of the disk rotation."
 The extracted οποιον and. angular momentum are initially stored in an electromagnetic form., The extracted energy and angular momentum are initially stored in an electromagnetic form.
 Thev are gradually converted. into a kinetic form as the flow accelerates., They are gradually converted into a kinetic form as the flow accelerates.
 At large observable distances. (he conversion is nearly complete. ancl the wind becomes kinetically dominated.," At large observable distances, the conversion is nearly complete, and the wind becomes kinetically dominated."
 Along any given field line. the kinetic energy will then be proportional to the fIuid angular momentum. with the same proportionality constant Oy because of energy and angular momentum conservation.," Along any given field line, the kinetic energy will then be proportional to the fluid angular momentum, with the same proportionality constant $\Omega_0$ because of energy and angular momentum conservation."
" At any observable location (sav of a distance σας. from the rotation axis. where the subscript denotes a quantity far from (he launching reeion). the specilic kinetic energy. ancl angular momentum of the wind are given by the poloidal ancl toroidal components of the[lid velocity. 0,4 and 0,4. through DN+TM)/2 and ce,szm respectively."," At any observable location (say of a distance $\varpi_\infty$ from the rotation axis, where the subscript denotes a quantity far from the launching region), the specific kinetic energy and angular momentum of the wind are given by the poloidal and toroidal components of thefluid velocity, $v_{p,\infty}$ and $v_{\phi,\infty}$, through $(v_{p,
\infty}^2+v_{\phi,\infty}^2)/2$ and $v_{\phi,\infty}\varpi_\infty$ respectively."
 Both velocity components can in principle be derived empirically from spatially resolved spectra ancl proper motion observations. as has been done for DG Tau (Bacciottietal.2002).," Both velocity components can in principle be derived empirically from spatially resolved spectra and proper motion observations, as has been done for DG Tau \citep{bac02}."
". Once measured. they can be used to deduce (he rate of disk rotation at the foot point of the field line passing through that location. through Where &,4<Up Is assumed. as is generally (rue at large distances."," Once measured, they can be used to deduce the rate of disk rotation at the foot point of the field line passing through that location, through where $v_{\phi,\infty}\ll v_{p,\infty}$ is assumed, as is generally true at large distances."
 The above relation provides a simple wax to infer (he rotation rate of the wind-Izaunching region on the Keplerian disk aud. if the central stellar mass is known independently. (he distance Irom the star.," The above relation provides a simple way to infer the rotation rate of the wind-launching region on the Keplerian disk and, if the central stellar mass is known independently, the distance from the star."
 In the next subsection. we will use the steady magnetohvdrodyvnamie (MIID) wind theory to derive a refined. version of this relation. taking into account of the energv and angular momentum associated with fluid rotation at the base of the wind.," In the next subsection, we will use the steady magnetohydrodynamic (MHD) wind theory to derive a refined version of this relation, taking into account of the energy and angular momentum associated with fluid rotation at the base of the wind."
 It is well known that in a steady. axisvnunetric MIID wind. several quantities are conserved along a given field line (Mestel 1963)..," It is well known that in a steady, axisymmetric MHD wind, several quantities are conserved along a given field line \citep{m68}. ."
 These include (he total specilic energy, These include the total specific energy
dis work deseribes the lydrocwuamic loop modeling of ji N-rav flare observed ou Proxima €'entaur with μα...ewton.,This work describes the hydrodynamic loop modeling of an X-ray flare observed on Proxima Centauri with XMM-Newton.
 The «ata are very detailed: we can distinguish six cl-defined phases in the light curve aud a well-defiued atl in the deusitv-teiiperature diagram., The data are very detailed: we can distinguish six well-defined phases in the light curve and a well-defined path in the density-temperature diagram.
 Our approach ns] con to model cach plase in detail takine the time-Osolved. deusitv/temperature information iuto account., Our approach has been to model each phase in detail taking the time-resolved density/temperature information into account.
 We find hat this flare is best-deseribed with the fowing combination of compoucuts: Table 1 sumnuananzes the parameters of the model flaring loops., We find that this flare is best-described with the following combination of components: Table \ref{tab:loopab} summarizes the parameters of the model flaring loops.
 Ibvcdrodyuuuic modchug allowed us to derive a very ctailed scenario with qualitative aud quantitative coustraints on the loop morphology and on the heating function., Hydrodynamic modeling allowed us to derive a very detailed scenario with qualitative and quantitative constraints on the loop morphology and on the heating function.
 The modeling may not hx nique: we cannot exclude that other combinations of he parameters may be found with a significantly higher uodeling effort. aud a few parameters are not otally coustrained.," The modeling may not be unique: we cannot exclude that other combinations of the parameters may be found with a significantly higher modeling effort, and a few parameters are not totally constrained."
 On the other haud. we notice that such a complex and detailed eveut is reasonably well-described iu terms of only two dominant oop colmpoucnts and a well-defined heating function. vali for both faring loops.," On the other hand, we notice that such a complex and detailed event is reasonably well-described in terms of only two dominant loop components and a well-defined heating function, valid for both flaring loops."
 This resul 11lay xovide a ecneral pattern for the interpretation of stellar flares. evei those which show lieht curves more complex than a smile | decay.," This result may provide a general pattern for the interpretation of stellar flares, even those which show light curves more complex than a simple $+$ decay."
 The complex ight curve of this flare ma be in par explained by the OYSOY Collectiug area of NATAL together with its capability of long uuinterrupted observations witli respect to previots satelites. and we may have nuüssed 1 in other stellar flares just because of insufficient S/N ratio and time coverage.," The complex light curve of this flare may be in part explained by the larger collecting area of XMM together with its capability of long uninterrupted observations with respect to previous satellites, and we may have missed it in other stellar flares just because of insufficient S/N ratio and time coverage."
 Ou the other haud. there are stellar flares observed wih euoeh time resolution. coverage and statistics. which are less complex (e.g. vau deu Oord Alewe 1989. Pallavicini et al.," On the other hand, there are stellar flares observed with enough time resolution, coverage and statistics, which are less complex (e.g. van den Oord Mewe 1989, Pallavicini et al."
 1990). and also many solar Ones (0.2. Sato ot al.," 1990), and also many solar ones (e.g. Sato et al."
 2003)., 2003).
spectral index does not change significantly as a function of the orbital phase (Abdoetal.2009).,spectral index does not change significantly as a function of the orbital phase \citep{Fermi}.
. Therefore. we asstume a constaut spectrum as a function of the phase at TeV energies.," Therefore, we assume a constant spectrum as a function of the phase at TeV energies."
 The integrated flux is then where Ej depeuds ou the detection threshold energy of the detector. aud £j is a uormalizatiou [actor that is taken as a free parameter.," The integrated flux is then where $E_0$ depends on the detection threshold energy of the detector, and $F_0$ is a normalization factor that is taken as a free parameter."
 Figure 2 shows tlie attenuation as a fuuctiou of the orbital phase for several different energies or the case of the compact object having the canonical neutron star mass (2zzOL or ALzzLAL. )., Figure \ref{att_gamma_gamma} shows the attenuation as a function of the orbital phase for several different energies for the case of the compact object having the canonical neutron star mass $i\approx 64^{\circ}$ or $M\approx 1.5M_\odot$ ).
 In figure 2 we essentially reproduce oue of the results obtained by Dubus. except that the orbital j»arameters used are the newer set obtained by Aragonaetal.(2009).," In figure \ref{att_gamma_gamma} we essentially reproduce one of the results obtained by Dubus, except that the orbital parameters used are the newer set obtained by \citet{Aragona}."
. When only interactions with he backgrouud black body photous are takeu iuto account. and the orbital plane is closer to being seen edge-on. the optical depth approaches a ΕΠΙ when the compact object is close to the vain sequence star.," When only interactions with the background black body photons are taken into account, and the orbital plane is closer to being seen edge-on, the optical depth approaches a minimum when the compact object is close to the main sequence star."
 This is especially the case lor very high inclination augles. correspoudiug to he mass of the compact object being small. aud close to the Clhaudrasekliar mass.," This is especially the case for very high inclination angles, corresponding to the mass of the compact object being small, and close to the Chandrasekhar mass."
 This behavior can be understood from the angular depeudeuce of the threshold euergy in addition to the relative velocity of the incident and target photons approaching a minimum., This behavior can be understood from the angular dependence of the threshold energy in addition to the relative velocity of the incident and target photons approaching a minimum.
 Also. at high energies. the cross section for pair creation decreases as the inverse square of the center of mass euergy. decreasing the optical depth even more.," Also, at high energies, the cross section for pair creation decreases as the inverse square of the center of mass energy, decreasing the optical depth even more."
" That is. even though the total deusity of backgrouud photons increases (as qo l/r) when the compact objectH approaches the Be star. a combination. of+ the previously"" mentioned"" factors dominates as Can seen in figure 2.."," That is, even though the total density of background photons increases (as $1/r^2$ ) when the compact object approaches the Be star, a combination of the previously mentioned factors dominates as can seen in figure \ref{att_gamma_gamma}."
 Figure 3. shows the normalized. integrated [lux assuming differeut inclination angles aud corresponding compact object masses., Figure \ref{light_curve_gamma_gamma} shows the normalized integrated flux assuming different inclination angles and corresponding compact object masses.
 The VERITAS data shown iu figure )3. were binued to show a siugle light-curve as opposed to monthly data., The VERITAS data shown in figure \ref{light_curve_gamma_gamma} \citep{Veritas2} were binned to show a single light-curve as opposed to monthly data.
 IC we assume that the enission comes [roi the vicinity of the compact object. aud is isotropic. and coustant as a fuuctiou of the orbital phase. then these results lead. us to conclude that there must be an additional allenuation mechanism at," If we assume that the emission comes from the vicinity of the compact object, and is isotropic, and constant as a function of the orbital phase, then these results lead us to conclude that there must be an additional attenuation mechanism at"
 IC we assume that the enission comes [roi the vicinity of the compact object. aud is isotropic. and coustant as a fuuctiou of the orbital phase. then these results lead. us to conclude that there must be an additional allenuation mechanism at.," If we assume that the emission comes from the vicinity of the compact object, and is isotropic, and constant as a function of the orbital phase, then these results lead us to conclude that there must be an additional attenuation mechanism at"
emission lines but no redshift was detected in this range using absorption lines.,emission lines but no redshift was detected in this range using absorption lines.
 These three bursts yield a rate that is significantly higher than predicted by the model., These three bursts yield a rate that is significantly higher than predicted by the model.
 This cannot be explained by a misidentification of the host galaxies because the probability for such an effect is too small 2008).., This cannot be explained by a misidentification of the host galaxies because the probability for such an effect is too small \citep{Cobb(2008)}.
 These low redshift bursts have low luminosities and they could not have been detected at a much higher redshift., These low redshift bursts have low luminosities and they could not have been detected at a much higher redshift.
 We conclude that they possibly represent a low luminosity population that is different from the majority of the bursts (see a discussion in refsec:lowlumpop))., We conclude that they possibly represent a low luminosity population that is different from the majority of the bursts (see a discussion in \\ref{sec:lowlumpop}) ).
 A particularly interesting question is how many high redshift bursts are expected to be observed., A particularly interesting question is how many high redshift bursts are expected to be observed.
 Such bursts are of great interest as they may shed new light on the very early universe., Such bursts are of great interest as they may shed new light on the very early universe.
 Already now GRB090423 is amongst the most distant and hence the earliest objects observed so far., Already now GRB090423 is amongst the most distant and hence the earliest objects observed so far.
" Figure 8 depicts the observed cumulative redshift distribution and the predicted onc, for high-redshift bursts forSwift and for past and future missions: BATSE, EXIST (Bandetal.2008) and SVOM (Schanneetal. 2009).."," Figure \ref{fig:cumsum_redshift} depicts the observed cumulative redshift distribution and the predicted one, for high-redshift bursts for and for past and future missions: BATSE, EXIST \citep{Band(2008)} and SVOM \citep{Schanne(2008),Gotz(2009)}."
 The fraction of bursts with z>7.3.9.10.15.20 respectively is shown in Table 2..," The fraction of bursts with $z > 7, 8, 9, 10, 15, 20$ respectively is shown in Table \ref{tbl:redshift_frac}."
 Detection of high redshift bursts is one of the main objectives of EXIST., Detection of high redshift bursts is one of the main objectives of EXIST.
 The number of high redshift bursts expected to be detected by EXIST is presented in Table 3.., The number of high redshift bursts expected to be detected by EXIST is presented in Table \ref{tbl:redshift_count_EXIST}.
 During a five year mission EXIST will detect many high redshift (2> 10) (2 30) bursts., During a five year mission EXIST will detect many high redshift $z > 10$ ) $\approx$ 30) bursts.
 In the best-fit model there is a good probability for EXIST to detect even a 2>20 burst during mission., In the best-fit model there is a good probability for EXIST to detect even a $z > 20$ burst during mission.
" This is. of course, provided that such early bursts exist and that the rate at 5zz—3 can be extrapolated to such high redshifts."," This is, of course, provided that such early bursts exist and that the rate at $z \approx 5-8$ can be extrapolated to such high redshifts."
Although our new algorithm: provides a dramatically iniproved: satellite mass. loss history. discrepancies remain owing to the many complicated: physical processes involved in satellite disruption.,"Although our new algorithm provides a dramatically improved satellite mass loss history, discrepancies remain owing to the many complicated physical processes involved in satellite disruption."
 First. the compressive gravitational shock is a strongly. non-linear. time-dependent perturbation.," First, the compressive gravitational shock is a strongly non-linear, time-dependent perturbation."
 Although the adiabatic correction includes some of this time dependence by including the work done at the frecqucney »eak. many channels of possible coupling are. ignored.," Although the adiabatic correction includes some of this time dependence by including the work done at the frequency peak, many channels of possible coupling are ignored."
 Second. satellite heating by resonant interactions is too complicated to be accurately represented. by our. simple »xwametrisation.," Second, satellite heating by resonant interactions is too complicated to be accurately represented by our simple parametrisation."
 For example. the aciahatie correction extends the coupling to. lower energies by noting that he resonant coupling will have a power law rather than an exponential scaling with frequency.," For example, the adiabatic correction extends the coupling to lower energies by noting that the resonant coupling will have a power law rather than an exponential scaling with frequency."
 However. for a xutieular interaction. individual resonances may dominate he response: this is analogous to the dilference between a and. spectrum.," However, for a particular interaction, individual resonances may dominate the response; this is analogous to the difference between a and spectrum."
 Third. a real satellite halo iàs à wide range of orbits. [rom circular to radial.," Third, a real satellite halo has a wide range of orbits, from circular to radial."
 The mass shell scheme cannot accurately. capture the dynamics of these dillerent orbits., The mass shell scheme cannot accurately capture the dynamics of these different orbits.
 Moreover. the entire. satellite structure readjusts its evolution in the course of satellite disruption.," Moreover, the entire satellite structure readjusts its evolution in the course of satellite disruption."
 Εις reacljustment process is also not included., This readjustment process is also not included.
 ‘These arguments suggest that although our new algorithm provides improved satellite mass loss histories. high qualiE simulations are still necessary for an accurate prediction of mass loss.," These arguments suggest that although our new algorithm provides improved satellite mass loss histories, high quality simulations are still necessary for an accurate prediction of mass loss."
 Using high resolution simulations with cosmologically motivated. initial conditions. we investigate the physical processes responsible for the evolution of satellite galaxies in their host halo.," Using high resolution simulations with cosmologically motivated initial conditions, we investigate the physical processes responsible for the evolution of satellite galaxies in their host halo."
 We identified and explored the following important physical mechanisms that result. in. satellite galaxy disruption., We identified and explored the following important physical mechanisms that result in satellite galaxy disruption.
 Our main results are as follows: The satellite evolution. using N-body simulation has xen studied by a number of authors recently (e.g.Llavashi2007)..," Our main results are as follows: The satellite evolution using N-body simulation has been studied by a number of authors recently \citep[e.g.][]{Hayashi.etal:03,Kazantzidis.etal:04,
 Boylan-Kolchin.Ma:07}."
 Their approach differs ours in several ways., Their approach differs ours in several ways.
 We use a highly. idealised simulation configuration to better discern the details of the dynamical mechanisms. focussing on the detailed. physical processes allecting satellite evolution such as resonant dvnamics that have not been rigorously accdressed: elsewhere.," We use a highly idealised simulation configuration to better discern the details of the dynamical mechanisms, focussing on the detailed physical processes affecting satellite evolution such as resonant dynamics that have not been rigorously addressed elsewhere."
 For example. we cquilibratecl the initially tically truncated satellite to reduce the artificial perturbation from. sudden. introduction. of the host. halo potential.," For example, we equilibrated the initially tidally truncated satellite to reduce the artificial perturbation from sudden introduction of the host halo potential."
 We also explored. a circular. orbit. simulation to clearly demonstrate the elfects of resonant. torque., We also explored a circular orbit simulation to clearly demonstrate the effects of resonant torque.
 We performed. a. linear. perturbation calculation to. estimate the required. resolution for resonant effect. ancl made the simulations satisfv this requirement.," We performed a linear perturbation calculation to estimate the required resolution for resonant effect, and made the simulations satisfy this requirement."
 In addition. we use an expansion code to reduce the force [uctuations on. small scales.," In addition, we use an expansion code to reduce the force fluctuations on small scales."
 Owing to these ellorts. we were able to demonstrate the iniportance of the resonant torque in satellite disruption.," Owing to these efforts, we were able to demonstrate the importance of the resonant torque in satellite disruption."
 llavashietal.(2003) used an NEW. halo as a satellite initial conditions and a Tree code for a potential solver., \citet{Hayashi.etal:03} used an NFW halo as a satellite initial conditions and a Tree code for a potential solver.
 They found that the simple tidal-Iimit approximation unclerestimates the mass loss. as we do. and found structural evolution. also similar to our findings.," They found that the simple tidal-limit approximation underestimates the mass loss, as we do, and found structural evolution, also similar to our findings."
 lxazantzidisctal.(2004) and DBovlan-Ixolchin.&Aa(2007) also. found wt similar internal structure evolution using trec-cocle simulations., \citet{Kazantzidis.etal:04} and \citet{Boylan-Kolchin.Ma:07} also found that similar internal structure evolution using tree-code simulations.
 This consensus suggests that the inner cusp of satellite halo is not strongly allectec by the tides from 10 host halo., This consensus suggests that the inner cusp of satellite halo is not strongly affected by the tides from the host halo.
 However. owing to their rather complicated configuration. they could. not. discern. the οσο of. the resonant torque. although some authors (e.g.Hayashictal.103) has noticed that analytic formulae underestimate the mass loss.," However, owing to their rather complicated configuration, they could not discern the effect of the resonant torque, although some authors \citep[e.g.][] {Hayashi.etal:03} has noticed that analytic formulae underestimate the mass loss."
 Although we have improved our understanding of )0 detailed. physical processes responsible for satellite isruption. some issues remain.," Although we have improved our understanding of the detailed physical processes responsible for satellite disruption, some issues remain."
 By separating the heating mechanisms into two distinct regimes. we achieved an improved. understanding of resonant dynamics for an eccentric orbit.," By separating the heating mechanisms into two distinct regimes, we achieved an improved understanding of resonant dynamics for an eccentric orbit."
 However. we have not compared this approximation with a comprehensive perturbation theory calculation.," However, we have not compared this approximation with a comprehensive perturbation theory calculation."
 The resonant heating of satellites on eccentric orbits is similar to heating by other subhaloes., The resonant heating of satellites on eccentric orbits is similar to heating by other subhaloes.
" ""This interaction is an important source of satellite evolution in addition to the interaction with the smooth host halo.", This interaction is an important source of satellite evolution in addition to the interaction with the smooth host halo.
 In addition. the initially spherical satellite is. deformed," In addition, the initially spherical satellite is deformed"
Aside from the possibility mentioned at the outset that the IIubble constant is too large aud the universe too old for any O=1 model to be viable. the main potential problem for CIIDME appears to be forming cuough structure at veh vedshift.,"Aside from the possibility mentioned at the outset that the Hubble constant is too large and the universe too old for any $\Omega=1$ model to be viable, the main potential problem for CHDM appears to be forming enough structure at high redshift."
 Although. as I meutiored above. the prediction o: CIIDM that the amount of eas in caped Lynido svstenas ds starting to decrease at Inm1 redshift +A3 secus to be in accord with the available data. he large velmd vospread of the associaed iietal-Πιο svstenis305 midieae that these SVSelus ae OTe nüiassive fvan CIIDA would predict (see e.e.. [65.69)).," Although, as I mentioned above, the prediction of CHDM that the amount of gas in damped Lyman $\alpha$ systems is starting to decrease at high redshift $z \gsim 3$ seems to be in accord with the available data, the large velocity spread of the associated metal-line systems indicate that these systems are more massive than CHDM would predict (see e.g., \cite{Lu,Wolfe}) )."
 Also. results from a receut CDM liverodyvnanic sinulatiou [70] iu whic1 the amount of nuitral hydrogen in protogzdaxies ποσοστο. cousisteut with that observed in darmixb πα o systems le the authors to specilate that CIIDN models would produce less than οποιο]i however. since the regiois identified as damped Lylia jio systems in the nulations were not actualv resolved. lis will need to X! addressed by higher resoution simulations for all he models cousidered.," Also, results from a recent CDM hydrodynamic simulation \cite{HernDLAS} in which the amount of neutral hydrogen in protogalaxies seemed consistent with that observed in damped Lyman $\alpha$ systems led the authors to speculate that CHDM models would produce less than enough; however, since the regions identified as damped Lyman $\alpha$ systems in the simulations were not actually resolved, this will need to be addressed by higher resolution simulations for all the models considered."
 Firally. Steidel et al.," Finally, Steidel et al."
 [T1] have found objects by their cuutted light at redsuifts +=33.5 apparcutly with relatively high velociyv dispersions (indicated by tl106 equivalent widlis of absorption lines). which trey tentatively identify as the progenitors of giaat elliptical galaxies.," \cite{Steidel} have found objects by their emitted light at redshifts $z=3-3.5$ apparently with relatively high velocity dispersions (indicated by the equivalent widths of absorption lines), which they tentatively identify as the progenitors of giant elliptical galaxies."
"dssining tha the indicated velocity dispESIOUS are 1ideed gravitational velocities. Mo Fusugita (ME) |72]. have argued that the abuidance of these objects is higher than expected for the COBE-normaliz dO— DAL-type models tmat cau fit the low-redslüft data. including CTIDA. bu in accord with predictions of the numodel considered Lieπο, ("," that the indicated velocity dispersions are indeed gravitational velocities, Mo Fukugita (MF) \cite{MoF} have argued that the abundance of these objects is higher than expected for the COBE-normalized $\Omega=1$ CDM-type models that can fit the low-redshift data, including CHDM, but in accord with predictions of the model considered here. ("
"Dui iore detail. he 3F analvsis distavors CIIDM with h15 and QO,20.2 in a sinele species of τςntrüuos.","In more detail, the MF analysis disfavors CHDM with $h=0.5$ and $\Omega_\nu
\gsim 0.2$ in a single species of neutrinos."
 They appareutlv would arene tha this model is then iu diffieutv SLC (dt overproduces rich clusters aud if tha problemi were solve| with a ittle til nyzz0.9. the resuline decrease in fhctuation power on sumiall «les wotId not cad to formation of choueh carly OhFCC‘ts.," They apparently would argue that this model is then in difficulty since it overproduces rich clusters — and if that problem were solved with a little tilt $n_p
\approx 0.9$, the resulting decrease in fluctuation power on small scales would not lead to formation of enough early objects."
" Hosvever. if QO,τε0.2 is sluurel )etween two species of neutrinos. the resultine uodel apours to be at least uarginallv consistent witli thi clusters aud he 5cidel objects even with the asstunptious of ME."," However, if $\Omega_\nu \approx 0.2$ is shared between two species of neutrinos, the resulting model appears to be at least marginally consistent with both clusters and the Steidel objects even with the assumptions of MF."
 The TIO¢cl with h=thi consscut with tie most restrictive ATF assiniptious has mA. hence ty<12 Cyr.," The model with $h=0.7$ consistent with the most restrictive MF assumptions has $\Omega_0 \gsim 0.5$, hence $t_0
\lsim 12$ Gyr."
 TIO¢els inviug tit and lower h. aud trerefore more consistent wih the ower constraint ciscissed above. may also be in trouble with the ME analysis.)," models having tilt and lower $h$, and therefore more consistent with the small-scale power constraint discussed above, may also be in trouble with the MF analysis.)"
 But inadditiou to ταncertaiuties about the actial velocity dispersion aud ]plivsical size of he Steide et al., But in addition to uncertainties about the actual velocity dispersion and physical size of the Steidel et al.
" objects. the couclusious of the ME analysis Call:dso be significaitly weaNned if the| eravitational velocities of the observed birvoIs are svsteiaically higler than t16 eravitational veocitics in the surrounding dark matter halos. as is perhaps the case at low redshift for arge spiral galaxies [73]. and even more so for elliptica ealaxies which are lavecly selt-exavitatiug stellar systenis in their central regionIs,"," objects, the conclusions of the MF analysis can also be significantly weakened if the gravitational velocities of the observed baryons are systematically higher than the gravitational velocities in the surrounding dark matter halos, as is perhaps the case at low redshift for large spiral galaxies \cite{NFW}, and even more so for elliptical galaxies which are largely self-gravitating stellar systems in their central regions."
 (dven the regular morprologies of the high-redshift objects seeu iu the Itubble Deep Field [T1]. aud other deep HST images. it scenis more likely that they are relatively low iiass objects uuidergoiug starbursts. possibly triggered by," Given the irregular morphologies of the high-redshift objects seen in the Hubble Deep Field \cite{vdB96} and other deep HST images, it seems more likely that they are relatively low mass objects undergoing starbursts, possibly triggered by"
considered m sinmimlatious).,considered in simulations).
" Then. the radial ficld yy inside the disk exactly writes (0.6. Durand 1961) where K and E are the complete elliptic integrals of the first aud second kinds respectively. (dj,infROX Ἐν tor=PdourOl. where αμ=O is the disk inner edge and dourD>604, is the outer edges"," Then, the radial field $g_R$ inside the disk exactly writes (e.g. Durand 1964) where $\elik$ and $\elie$ are the complete elliptic integrals of the first and second kinds respectively, $\uin =\ain/R \le 1$ , $\vout = R/\aout \le 1$, where $\ain \ge 0$ is the disk inner edge and $\aout > \ain$ is the outer edge."
 A more tractable expressions for gy cau be derived from truucated expansions of the complete elliptic iutegrals., A more tractable expressions for $g_R$ can be derived from truncated expansions of the complete elliptic integrals.
 For instance. with a classical secoud-order expausious over the modulus weed (es Gradshtevn Ryzhik 1991). namely the field) far from cdges ds given du a good approximation bv Note that contrary to Eq.(12)). that formmla remains finite at the edges. aud then is more realistic although approximated.," For instance, with a classical second-order expansions over the modulus $x < 1$ (e.g. Gradshteyn Ryzhik 1994), namely the field far from edges is given in a good approximation by Note that contrary to \ref{eq:gr}) ), that formula remains finite at the edges, and then is more realistic although approximated."
" Since the eravitational potential of the disk is nünninuuni very close to the immer edge (like iu most astroplivsical disks). we can simply Eq.(11)) by considering oulv the terii linear in A. aud so Iu the homogeneous disk model. we thus have where gp=ALGο Ma being the disk mass. and assuage: ez,D<<a5."," Since the gravitational potential of the disk is minimum very close to the inner edge (like in most astrophysical disks), we can simplify \ref{eq:grapprox}) ) by considering only the term linear in $R$ , and so In the homogeneous disk model, we thus have where $\mu=\md/M$, $\md$ being the disk mass, and assuming $\ain^2 \ll \aout^2$."
 Deuce. 5=0.," Hence, $\gamma=0$."
" UsingToi . Eq.(3)). our formula for the relative shift becomes We see that the relative shift is strongly sensitive to the location of the planet with respect to the disk outer οσο,"," Using \ref{eq:reff}) ), our formula for the relative shift becomes We see that the relative shift is strongly sensitive to the location of the planet with respect to the disk outer edge."
 Figure 1 shows d6R/R. as predicted by Eq.(17))} for typical parameters., Figure \ref{fig:anavsnum.eps} shows $\delta R/\re$ as predicted by \ref{eq:shifthomo}) ) for typical parameters.
 The agreciieut ο the above formmla aud values determined muucericallv from Eq.(5)) by standard root finding techniques is excellent., The agreement between the above formula and values determined numerically from \ref{eq:dsg}) ) by standard root finding techniques is excellent.
" Both Liudblad resonances ect closer to the λαο,", Both Lindblad resonances get closer to the planet.
 The OLRs are slightly more shifted than the ILRs. specially for simall m (depending on the disk aspect ratio).," The OLRs are slightly more shifted than the ILRs, specially for small $m$ (depending on the disk aspect ratio)."
 This can be understood by computing the relative displacement Figure 2. displavs the Lindblad torques computed following the procedure deseribed in Ward (1997)., This can be understood by computing the relative displacement Figure \ref{fig:torques.eps} displays the Lindblad torques computed following the procedure described in Ward (1997).
 It confinis that the outer torques have larger amplitiucle than inner ones when the disk eravitv is accounted for., It confirms that the outer torques have larger amplitude than inner ones when the disk gravity is accounted for.
 Differential torques are larger. and iuwards nüeration should be accelerated.," Differential torques are larger, and inwards migration should be accelerated."
 We have computed the resonance shifts aud associated Lindblad torques for disks with power law surface density profiles (o. Sox RW). typical from disk models and observations.," We have computed the resonance shifts and associated Lindblad torques for disks with power law surface density profiles (i.e. $\Sigma \propto R^{-s}$ ), typical from disk models and observations."
 The disk eravity field gj has been determinedmuuencallv from the splitting method described in Iburé Pierens (2001)., The disk gravity field $g_R$ has been determinednumerically from the splitting method described in Huré Pierens (2004).
 As shown in Fig. 3..," As shown in Fig. \ref{fig:vss.xfig.eps}, ,"
"BL Lacs and FSRQs, intrinsically relate to their physical properties during pre-flare conditions, i.e., to dearth or plenty of gas surrounding the central SMBH.","BL Lacs and FSRQs, intrinsically relate to their physical properties during pre-flare conditions, i.e., to dearth or plenty of gas surrounding the central SMBH."
 we have studied how this physical dichotomy is reflected in the flaring activity., we have studied how this physical dichotomy is reflected in the flaring activity.
" We find that during the flares the BL Lac and the FSRQ follow different patterns, ones on the - We ga"," We find that during the flares the BL Lac and the FSRQ follow different patterns, ones on the - plane."
"mma,focused plane.on the pattern derived from recent )observationsofthetwoFSRQs, 3C279and3C454.3."," We focused on the pattern derived from recent observations of the two FSRQs, 3C 279 and 3C 454.3."
"We findthese f sholoesd lyowatidahenthe plane to be driven mainly by an increasing flux of the gamma,external photons, under further growth of the already substantial rates."," We find flares to closely on the - plane to be driven mainly by an increasing flux of the external photons, under further growth of the already substantial rates."
" Meanwhile the electron energies hardly increase, being constrained by the shortening cooling times in the denser radiative field."," Meanwhile the electron energies hardly increase, being constrained by the shortening cooling times in the denser radiative field."
" These sources add to three BL Lacs that we previously studied (S50716+714, Mrk421 and Mrk501) and complete our picture."," These sources add to three BL Lacs that we previously studied (S50716+714, Mrk421 and Mrk501) and complete our picture."
" We interpret their flares in terms of increasing under furtheracceleration, to yield gamma,patestns"," We interpret their flares in terms of increasing under further, to yield patterns toward the"
scales limited on small scales by thermal pressure. (ad (he. Jeans length. Aj=C?/GX) and on large scales by rotation (at the critical length set by rotation. Αμ=4226/57).,"scales limited on small scales by thermal pressure (at the Jeans length $\lambda_{\rm Jeans} = C_{\rm s}^2 / G\Sigma $ ) and on large scales by rotation (at the critical length set by rotation, $\lambda_{\rm rot}= 4\pi^2 G \Sigma / \kappa^2$ )."
 All intermediate leneth scales are unstable. and the most rapidly growing mode has a wavelength 2A\Jeans:," All intermediate length scales are unstable, and the most rapidly growing mode has a wavelength $2 \, \lambda_{\rm Jeans}$."
 The maxinum unstable length scale in a disk. A44. is a robust quantity because il depends only on (he surface density aud epicvelic frequency of the disk and not on the smaller scale physics.," The maximum unstable length scale in a disk, $\lambda_{\rm rot}$, is a robust quantity because it depends only on the surface density and epicyclic frequency of the disk and not on the smaller scale physics."
" Such scale has a characteristic mass associated. defined equals to X,,,(A,,/2)7"" and (hat can be expressed as: Coultrarily. due to (he complex structure and dynamics of real interstellar medium (ISM) in galaxies. which cannot be described by a simple equation of state. there is not a well-defined Jeans length at intermediate scales."," Such scale has a characteristic mass associated, defined equals to $ \Sigma_{gas} (\lambda_{\rm rot}/2)^{2}$ and that can be expressed as: Contrarily, due to the complex structure and dynamics of real interstellar medium (ISM) in galaxies, which cannot be described by a simple equation of state, there is not a well-defined Jeans length at intermediate scales."
 Therefore. no real lower limit on the sizes of the structures that can form until the thermal Jeans scale is reached in molecular ‘loud cores (Escala Larson 2008).," Therefore, no real lower limit on the sizes of the self-gravitating structures that can form until the thermal Jeans scale is reached in molecular cloud cores (Escala Larson 2008)."
 One of the intermediate scales usually claimed (o. be a relevant in the stability of je ISM. is the Jeans length defined by the turbulent. pressure in the medium vOSH/GM).," One of the intermediate scales usually claimed to be a relevant in the stability of the ISM, is the Jeans length defined by the turbulent pressure in the medium $\lambda_{J}^{turb}= v_{\rm turb}^2 / G\Sigma $ )."
 Besides the convenience of the idea of a turbulent pressure term that oSgeneralize 1¢ eravitational instability analvsis. this term is however not a well defined quantity in ISM (Elneereen Sealo 2004).," Besides the convenience of the idea of a turbulent pressure term that generalize the gravitational instability analysis, this term is however not a well defined quantity in ISM (Elmegreen Scalo 2004)."
 This comes from the fact that this pressure term could only be defined in the case where the dominant turbulent scale is much smaller than the region under consideration Cmieroturbulence). which is in fact is not the ease of ISM.," This comes from the fact that this pressure term could only be defined in the case where the dominant turbulent scale is much smaller than the region under consideration (`microturbulence'), which is in fact is not the case of ISM."
 Rigorous analysis endeed shows that turbulence can be represented as a pressure only in (he microturbulence case (Donazzola et al., Rigorous analysis endeed shows that turbulence can be represented as a pressure only in the microturbulence case (Bonazzola et al.
 1992)., 1992).
 Desides turbulent motions could help to prevent collapse on some scales. those complex motions cannot be represented as a pressure.," Besides turbulent motions could help to prevent collapse on some scales, those complex motions cannot be represented as a pressure."
" Therefore the eravitational instability analysis is not strictly applicable with a turbulent pressure term that could stabilize and damp all the substructure below A/"".", Therefore the gravitational instability analysis is not strictly applicable with a turbulent pressure term that could stabilize and damp all the substructure below $\lambda_{J}^{turb}$.
 Although turbulence cannot be modelled as simple pressure supporting term that will prevent all gravitational instabilities below some scale. turbulence could prevent the growth ol instabilites by heating the gas trough compression.," Although turbulence cannot be modelled as simple pressure supporting term that will prevent all gravitational instabilities below some scale, turbulence could prevent the growth of instabilities by heating the gas trough compression."
 However. compression [rom turbulence nol only heat the gas. but also may enhance collapse and star formation on smaller scales (Elmegreen 2002: Krumbholty Mekee. 2005).," However, compression from turbulence not only heat the gas, but also may enhance collapse and star formation on smaller scales (Elmegreen 2002; Krumholtz Mckee 2005)."
 Only when the turbulent heating becomes faster than the gas cooling. the net effect of turbulence is to effectively heat the gas and therefore prevent gravitational instabilities.," Only when the turbulent heating becomes faster than the gas cooling, the net effect of turbulence is to effectively heat the gas and therefore prevent gravitational instabilities."
 Once this condition is satisfied (he gas is heated towards c47viu. the thermal pressure will stabilize aud damp all (he substructure below Ay0Alt," Once this condition is satisfied the gas is heated towards $\rm c_{S} \sim v_{turb}$, the thermal pressure will stabilize and damp all the substructure below $\lambda_{J} \sim \lambda_{J}^{turb}$."
" If the condition Ay~Apo, (equivalent to Q ~ 1) is also satisfied. the svstem will"," If the condition $\lambda_{J} \sim \lambda_{rot}$ (equivalent to Q $\sim$ 1) is also satisfied, the system will"
"As more distant galaxy populations become accessible to modern surveys, astronomers are striving to estimate their physical properties, despite the challenges inherent in such pioneering tasks.","As more distant galaxy populations become accessible to modern surveys, astronomers are striving to estimate their physical properties, despite the challenges inherent in such pioneering tasks."
" Light which barely registers on our instruments is analysed to infer the stellar mass and star formation activity of its source, providing valuable stepping stones on which our physical picture of structure formation can progress."," Light which barely registers on our instruments is analysed to infer the stellar mass and star formation activity of its source, providing valuable stepping stones on which our physical picture of structure formation can progress."
" For example, produced estimates of stellar mass for 1038 galaxies from the the GOODS survey, grouped into three populations by redshift: z7 4, 5 and 6."," For example, produced estimates of stellar mass for 1038 galaxies from the the GOODS survey, grouped into three populations by redshift: $z\approx$ 4, 5 and 6."
 These stellar masses were estimated using a population synthesis model which searches for the stellar population which best fits the observed spectral energy distribution of each galaxy (see refStellarMass))., These stellar masses were estimated using a population synthesis model which searches for the stellar population which best fits the observed spectral energy distribution of each galaxy (see \\ref{StellarMass}) ).
" Star formation rates were specifically not derived for this sample, because of uncertainties in the extinction correction."," Star formation rates were specifically not derived for this sample, because of uncertainties in the extinction correction."
" In lieu of this, the galaxies’ “emerging” UV luminosities were computed (the luminosity at without any dust correction)."," In lieu of this, the galaxies' “emerging” UV luminosities were computed (the luminosity at without any dust correction)."
" However, figure 9 of does include the star formation rates that would be inferred if a standard proportionality between UV luminosity and star formation rate were assumed??)mainBodyCitationEnd182]Madau98."," However, figure 9 of does include the star formation rates that would be inferred if a standard proportionality between UV luminosity and star formation rate were assumed."
" This figure, for the galaxies in the nearest of the three samples, is reproduced for reference in the upper panel of our Fig ??.."," This figure, for the galaxies in the nearest of the three samples, is reproduced for reference in the upper panel of our Fig \ref{obs}."
" Despite only a fleeting appearance in the observational paper, these star formation rate estimates have since been the subject of quite detailed theoretical analysis."," Despite only a fleeting appearance in the observational paper, these star formation rate estimates have since been the subject of quite detailed theoretical analysis."
 summarise the trend given by the sample in Fig.?? as:, summarise the trend given by the sample in \ref{obs} as:
The present work was supported by Russian Foundation of Basic Researches (Grant 96-02-16140) and by INTEAS (Crant 93-03-ENT).,The present work was supported by Russian Foundation of Basic Researches (Grant 96-02-16140) and by INTAS (Grant 93-93-EXT).
"Water vapour plays a sev role in the Martian atinosplieric chemistry and psics,",Water vapour plays a key role in the Martian atmospheric chemistry and physics.
 As source of hvdrosen radicals it has an nmuporaut Πλας ou the ozone (Lefevre ο al. 2001))," As source of hydrogen radicals it has an important impact on the ozone (Lefevre et al., \cite{le04}) )"
 conceitration. f1ο stability. of the Martian atinosphere (Parlsinson and Tuten. 1972)) and plays au nuportaut role i the comparative planctologey coutex (e.g. Tartogh et a. 200L.. 20] 0)).," concentration, the stability of the Martian atmosphere (Parkinson and Hunten, \cite{pa72}) ) and plays an important role in the comparative planetology context (e.g. Hartogh et al. \cite{ha04}, \cite{ha10}) )."
 The water cvele aud the associated: variabe hveropatwe (Clancy 1996... Eucreuaz 2001)) with a sLOU asviunetry of ifs distribution a perihelion aud apicliou is one of the principal probleuis of the present Martian climate.," The water cycle and the associated variable hygropause (Clancy \cite{cl96}, Encrenaz \cite{en01}) ) with a strong asymmetry of its distribution at perihelion and aphelion is one of the principal problems of the present Martian climate."
 Tf confined. the asviuimetiv is expected to drive a stroug unbalance in the meridiona transport.," If confirmed, the asymmetry is expected to drive a strong unbalance in the meridional transport."
 A better observajonal characterization will be a stroug constraint for general circulation iodes of the Martian. atinosphere anL particular its meridional transport (c.g. Moedvedev aid Wartogh. 2007)).," A better observational characterization will be a strong constraint for general circulation models of the Martian atmosphere and particular its meridional transport (e.g. Medvedev and Hartogh, \cite{me07}) )."
 Since its first detection im the near infrared (Spiurad et al. 1963)) water vapour iu the ALartian atinosphere has beeu exteusively studied by growid-hased and space borne observations (6.8. Sprague. e al. 1996..," Since its first detection in the near infrared (Spinrad et al, \cite{sp63}) ) water vapour in the Martian atmosphere has been extensively studied by ground-based and space borne observations (e.g. Sprague, et al. \cite{sp96},"
 Duredorf. et al. 2000.. ," Burgdorf, et al. \cite{bu00}, ,"
Camwell. et al. 2000..," Gurwell, et al. \cite{gu00},"
 Suüth 2002: Diver. et al. 2005..," Smith \cite{sm02}; Biver, et al. \cite{bi05},"
 Fouchet. 2007.. Eucreiz et al. 9005. and σαι et al. 20093).," Fouchet, \cite{fo07}, Encrenaz et al, \cite{en08} and Smith et al, \cite{sm09}) )."
 Compared to water. there is oulv a sparse database for carbon monoxide at \ars.," Compared to water, there is only a sparse database for carbon monoxide at Mars."
 Space borne (Encrenaz. 2006.. and Suaith. 2009)) and erouud-based (e.g. Lellouch et al. 1991.. Ίναasnopolsky. 2007)) show an average abundance of 700 ppm.," Space borne (Encrenaz, \cite{en06}, and Smith, \cite{sm09}) ) and ground-based (e.g. Lellouch et al, \cite{ll91}, Krasnopolsky, \cite{kr07}) ) show an average abundance of 700 ppm."
 One of the major goals of the is to tracethe origin. evoliition and chemistry of water," One of the major goals of the is to tracethe origin, evolution and chemistry of water"
south of the SDSS position (Creecne&IIo2001) to avoid any phase-ceuter artifacts;,south of the SDSS position \citep{gre04} to avoid any phase-center artifacts.
 Observations were lmade assuniue a coordinate equinox of 2000., Observations were made assuming a coordinate equinox of 2000.
" Data were acquired with a bandwidth of 100 MITIz for cach circular polarization,", Data were acquired with a bandwidth of 100 MHz for each circular polarization.
 Observations of 2286 were used to set the amplitude scale to an accuracy of about354., Observations of 286 were used to set the amplitude scale to an accuracy of about.
 Ou 2007 Jun 1δ. net exposure times were 1500 s. του s; and 8300-8 at Ll. LO. and 8.5 GIIz. respectively.," On 2007 Jun 18, net exposure times were 1500 s, 7800 s, and 8300 s at 1.4, 4.9, and 8.5 GHz, respectively."
 Twenty-five to 27 antennas provided data of acceptable quality. with most of the data loss attributable to EVLA retrofitting activities.," Twenty-five to 27 antennas provided data of acceptable quality, with most of the data loss attributable to EVLA retrofitting activities."
 The obscrvatious at 1.1 αν were made at a lage hour anele aud were followed wi alteruatiug observations at L9 and 8.5 CGIIz to suuple similar hour ausles., The observations at 1.4 GHz were made at a large hour angle and were followed by alternating observations at 4.9 and 8.5 GHz to sample similar hour angles.
 Ou 2007 Jul 15. L.1-CHIIz observations were made at transit aud resulted iu a iet exposure time of 2800 s. Tweuty-four autennas xovided data of acceptable quality. with most of the data loss attributable to EVLA retrofitting activitics and software failures.," On 2007 Jul 15, 1.4-GHz observations were made at transit and resulted in a net exposure time of 2800 s. Twenty-four antennas provided data of acceptable quality, with most of the data loss attributable to EVLA retrofitting activities and software failures."
 The data were calibrated usine he 2007 Dec 31 release of the NRAO AIPS software., The data were calibrated using the 2007 Dec 31 release of the NRAO AIPS software.
 No selfcalibrations were performed., No self-calibrations were performed.
 No polarization calibration was performed. as onlv upper Ματς to the iuear polarization perceutages were sought.," No polarization calibration was performed, as only upper limits to the linear polarization percentages were sought."
 The AIPS ask was used to form: aud deconvolye images of he cussion from 110., The AIPS task was used to form and deconvolve images of the emission from 10.
 Following Wrobel(2000)..imagr was used to form and decouvolve Stokes J images at LO and 8.5 GIIz at a matched resolution at FWIIAE of0.," Following \citet{wro00}, was used to form and deconvolve Stokes $I\/$ images at 4.9 and 8.5 GHz at a matched resolution at FWHM of."
"12""... This led to detections of a compact source at both frequencies. with no evidence for adjacent cussion on larger scales."," This led to detections of a compact source at both frequencies, with no evidence for adjacent emission on larger scales."
 The 19 GIIz image appears iu Figure 1., The 4.9 GHz image appears in Figure 1.
 Quadratic fits in the image plane vielded the integrated flux densities aud the two-point spectral index eiven in Table 1., Quadratic fits in the image plane yielded the integrated flux densities and the two-point spectral index given in Table 1.
 These fiux densities frour 2007 Jun 18 are plotted in Figure 2., These flux densities from 2007 Jun 18 are plotted in Figure 2.
" The nuage fits also vielded ideutical decouvolved diameters of less than0.21""., as well as almost ideutical positions."," The image fits also yielded identical deconvolved diameters of less than, as well as almost identical positions."
" The L9-GIIz position is α20000=12?[0355.825 and ó(72000)—00729/19"".53. with a 1-D error of0.1"".. the quadratic sum of teris due to the signal-to-noise ratio (SNR) of the detection (less than 0.019). the phase-calibrator position error (0.011). and the pliase-referenciug strategies (estimated to be 0.17))."," The 4.9-GHz position is $\alpha(J2000) = 12^{h}
40^{m} 35^{s}.825$ and $\delta(J2000) = -00^{\circ} 29' 19''.53$, with a 1-D error of, the quadratic sum of terms due to the signal-to-noise ratio (SNR) of the detection (less than ), the phase-calibrator position error ), and the phase-referencing strategies (estimated to be )."
 This 1.9-GIIz position is consistent with the SDSS position (Creeue&Πο2001). marked in Figure 1.," This 4.9-GHz position is consistent with the SDSS position \citep{gre04}
 marked in Figure 1."
imagr was also used to formu matched-resolution images of Stokes Q and (oat L9 and 8.5 GIIz., was also used to form matched-resolution images of Stokes $Q\/$ and $U\/$ at 4.9 and 8.5 GHz.
" Those images led to no detections, with the most constraiiung upper hit to the linear polarization percentage being less tha at 1.9 CIIz."," Those images led to no detections, with the most constraining upper limit to the linear polarization percentage being less than at 4.9 GHz."
 Following Wo&Ulvestad(2001).. was used in a two-step approach to form aud deconvolve Stokes f iuages at Lt Cz.," Following \citet{ho01b}, was used in a two-step approach to form and deconvolve Stokes $I\/$ images at 1.4 GHz."
 First. a search for confusing sources was done by using the iuner array baselines to iniage. at low resolution. the primary autenua beam to the halfpower point.," First, a search for confusing sources was done by using the inner array baselines to image, at low resolution, the primary antenna beam to the half-power point."
 Second. all array baselimes were used to nuaee 110 and its confusiug sources. sinultancously aud at a lich resolution of about af ΕΛΠ," Second, all array baselines were used to image 10 and its confusing sources, simultaneously and at a high resolution of about at FWHM."
 From the (11110 data acquired at transit ou 2007 Jul 15. this approach led to a high-resolution nuaee mildly affected by confusion. showing a compact detection plus uo evidence for adjacent cussion ou larecr scales: a quadratic fit im the image pluie vielded the integrated flux deusity in Table 1 and a decouvolved diameter of less than0.," From the 10 data acquired at transit on 2007 Jul 15, this approach led to a high-resolution image mildly affected by confusion, showing a compact detection plus no evidence for adjacent emission on larger scales; a quadratic fit in the image plane yielded the integrated flux density in Table 1 and a deconvolved diameter of less than."
72”.. Frou the GIII10 data acquired at a large hour angle on 2007 Jun Is. the hieh-resolutiou image was more degraded bv confusion so onlv a parabolic fit in the image plane was made.," From the 10 data acquired at a large hour angle on 2007 Jun 18, the high-resolution image was more degraded by confusion so only a parabolic fit in the image plane was made."
 This vielded the peak flux deusitv in Table 1 that is plotted in Figure 2., This yielded the peak flux density in Table 1 that is plotted in Figure 2.
 That figure. also shows the spectral iudex a=0.76£0.05 derived from all the photometry ou 2007 Jun 18., That figure also shows the spectral index $\alpha = -0.76\pm0.05$ derived from all the photometry on 2007 Jun 18.
 Finally. the L.9-CIIz data were used to produce au nuage at the EWIIM resolution of the 1.I-GIIz data ou 2007 Jun Ls.," Finally, the 4.9-GHz data were used to produce an image at the FWHM resolution of the 1.4-GHz data on 2007 Jun 18."
 A parabolic fit to that L9-GIIz image eave the fiux density listed in Table 1. leading in turi to the tabulated spectral index between 1.1 and 1.9 GIIz.," A parabolic fit to that 4.9-GHz image gave the flux density listed in Table 1, leading in turn to the tabulated spectral index between 1.4 and 4.9 GHz."
" Also. that [.9-CHIIz fux density is consistent with the value obtained at a EWIIM resolution of Huplvine no additional cussion on scales between ind about1""."," Also, that 4.9-GHz flux density is consistent with the value obtained at a FWHM resolution of implying no additional emission on scales between and about."
. The new VLA data imply that the cimission from 110 is compact. with a diameter of less than (320 pc) at LO Giz (Fie.," The new VLA data imply that the emission from 10 is compact, with a diameter of less than (320 pc) at 4.9 GHz (Fig."
 1) aud 8.5 GIIz. aud less than (1.1 pe) at 1.4 GIIz.," 1) and 8.5 GHz, and less than (1.1 kpc) at 1.4 GHz."
 These upper lits are derived from data with SNRs of 15 or more., These upper limits are derived from data with SNRs of 15 or more.
 The new 1.1-CIIz nuages show only a compact component and do nof confirm the prior hint frou low-SNR data of mareiually resolved cussion with a FEWITIM diameter of (0.5kpe:Whiteetal.1997)., The new 1.4-GHz images show only a compact component and do not confirm the prior hint from low-SNR data of marginally resolved emission with a FWHM diameter of \citep[7.5~kpc;][]{whi97}.
. Stilarly. the new L.9-CdIz nage does not coufirm the prior suggestion from low SNR-data of mareiualh-vesolved cussion with a FWIAL diameter of (530pe:Greeneetal.2006).," Similarly, the new 4.9-GHz image does not confirm the prior suggestion from low SNR-data of marginally-resolved emission with a FWHM diameter of \citep[530~pc;][]{gre06}."
. The emission from GITITO has a steep radio spectrum. with a three-point spectral iudex of a=0.7640.05 measured on 2007 Jun 18 aud thus free from concerns about time variability (Fig.," The emission from 10 has a steep radio spectrum, with a three-point spectral index of $\alpha = -0.76\pm0.05$ measured on 2007 Jun 18 and thus free from concerns about time variability (Fig."
 2)., 2).
 This spectral index agrees with the two-point. matched-resolution spectral iudices even in Table 1. as is expected for compact. unresolved cluission.," This spectral index agrees with the two-point, matched-resolution spectral indices given in Table 1, as is expected for compact, unresolved emission."
 The sinele-cpoch radio spectrum shown for 110 in Fieure 2 - the first ever measured for ay of the 19 low-mass ACNs discovered bv Creene aud Πο (Cirecne&Πο2001:Crecneetal.2006). - is consistent with optically thin svuchrotrou cussion.," The single-epoch radio spectrum shown for 10 in Figure 2 - the first ever measured for any of the 19 low-mass AGNs discovered by Greene and Ho \citep{gre04,gre06} - is consistent with optically thin synchrotron emission."
 This emission is less than nearly polarized at LO GIIz. well below the expected for electrons with a power-law cuerey distribution. with iudex p=1 26.1n a uniforii maguetic field (Shu1991.page197)..," This emission is less than linearly polarized at 4.9 GHz, well below the expected for electrons with a power-law energy distribution, with index $p = 1 -
2\alpha$, in a uniform magnetic field \citep[][page 197]{shu91}."
 Faraday depth effects could also be at play., Faraday depth effects could also be at play.
 Variability timescales for AGNs are expected to scale with black-hole mass (e.g...Edelson&Naudra1999).. and have not vet been explored for auv of the 19 low-mass Greene aud Πο objects.," Variability timescales for AGNs are expected to scale with black-hole mass \citep[e.g., ][]{ede99}, and have not yet been explored for any of the 19 low-mass Greene and Ho objects."
 The new VLA plotometry can be used to assess LLO's time variabilty. for future conrparison with other CULLLO timescales to search for plivsicallinkages.," The new VLA photometry can be used to assess 10's time variabilty, for future comparison with other 10 timescales to search for physicallinkages."
" Using only A-coufieuration data. the ratio of the L9-CGIIZ flux density on 2007 Jun 15 (Q.7540.05 ude) to that on 2001 Oct 9 (0.70.1ινCreeneetal.9006) is τον, cousistent with steady Cluission on timescales of vears."," Using only A-configuration data, the ratio of the 4.9-GHz flux density on 2007 Jun 18 $\pm$ 0.05 mJy) to that on 2004 Oct 9 \citep[0.7$\pm$0.1~mJy;][]{gre06} is $\pm$, consistent with steady emission on timescales of years."
 The ratio of the 1.l- flux density ou 2007 Jul 15 (1.880.183 mJy) to that on 2007 Jun 18 (1.8040.09 inv) is Lt9%.. also," The ratio of the 1.4-GHz flux density on 2007 Jul 15 $\pm$ 0.13 mJy) to that on 2007 Jun 18 $\pm$ 0.09 mJy) is $\pm$ , also"
We have shown iu lis article how to extend to nou-Ciassieui initial concitious the computation of he miss unction aux of the bias of dark matter halos presened in Valaecas (20095) or the Caussian case.,We have shown in this article how to extend to non-Gaussian initial conditions the computation of the mass function and of the bias of dark matter halos presented in Valageas (2009b) for the Gaussian case.
 This relics Ol v saclee-poiut approaci that allows to derive the larec-nass asvinptotic tails Q ‘the quantities of interest from he y.atisticul weight of he initial conditions. suppeniened we adiitional nonlinear constraints.," This relies on a saddle-point approach, that allows to derive the large-mass asymptotic tails of the quantities of interest from the statistical weight of the initial conditions, supplemented by additional nonlinear constraints."
" Then. focisues O1 je case of ""fxp-tvpe xunaordial nou-Ciuissiauity. where je linear gravitationa potential can be written as the ""Hu o: Linear aud qiadratic terms over an anxiliarv Gaussian field x. we exλα1ος] how to obtain the relevant saddle-poiuts as a perturbative series over the nonlinear paraueter νε,"," Then, focussing on the case of $\fNL$ -type” primordial non-Gaussianity, where the linear gravitational potential can be written as the sum of linear and quadratic terms over an auxiliary Gaussian field $\chi$, we explained how to obtain the relevant saddle-points as a perturbative series over the nonlinear parameter $\fNL$."
 This ucthod is very ecucral aud it applies to any case of small primordial nou-Caussianity. where Bardecu’s poteutial (or equivalently. below the Παολο Tacius. the exavitational poteutial or the deusity feld) cau be written as a polvuonidal over a Gaussian field X as E ( ) where the nonlinear parameters f; are μηνα.," This method is very general and it applies to any case of small primordial non-Gaussianity, where Bardeen's potential (or equivalently, below the Hubble radius, the gravitational potential or the density field) can be written as a polynomial over a Gaussian field $\chi$ as + f_i ( ^i - ), where the nonlinear parameters $f_i$ are small."
 Then. olOwlis the method preseuted iu sections 5. Dd. the αμ...] aud two-cell sacdle-poits (associated with the one-point and two-poiu density. distributions. whence he haο lnass function and bias) can be computed as a )orturtivo series OVCL the coefficients f£; (which ueed not jc oftre sune order).," Then, following the method presented in sections \ref{Mass-function-halos}, \ref{Bias-of-halos}, the one-cell and two-cell saddle-points (associated with the one-point and two-point density distributions, whence the halo mass function and bias) can be computed as a perturbative series over the coefficients $f_i$ (which need not be of the same order)."
 This includes in particular the case of à οιbic term ονκ”., This includes in particular the case of a cubic term $g_{\rm NL} \chi^3$.
 Of course. this also extends to he cases where the coefficients f; are not mere nuuIOLS ut cowolution kernels and to several Gaussian fields Y; (which cau have a nouzero cross-correlation).," Of course, this also extends to the cases where the coefficients $f_i$ are not mere numbers but convolution kernels and to several Gaussian fields $\chi_j$ (which can have a nonzero cross-correlation)."
 Iu all such cases. the large-niass asyuptotics are set by the statistical weight (ο...ol⋅ of these Camussian fields. taken at the saddle-poiut associated with he maxinization of this weight uuder appropriate nonilear constraints. hat OXpress he mapping fic1n y; to the relevaut quautitv. such as the nonlinear density within a spicrical cell.," In all such cases, the large-mass asymptotics are set by the statistical weight $e^{-(\chi_i.C_{ij}^{-1}.\chi_j)/2}$ of these Gaussian fields, taken at the saddle-point associated with the maximization of this weight under appropriate nonlinear constraints, that express the mapping from $\chi_j$ to the relevant quantity, such as the nonlinear density within a spherical cell."
 Focussing on the ¢ase of local-type primordial nou-Ciassieuty. we described how to obtain. up to linear order over fxp. the onc-cell saddle-Olif associated with the probability dcistribitions Pp(p) arLP (6) ofthe linear and noulinear deusitv coutrasts within syherical cells.," Focussing on the case of local-type primordial non-Gaussianity, we described how to obtain, up to linear order over $\fNL$, the one-cell saddle-point associated with the probability distributions $\cP_L(\delta_L)$ and $\cP(\delta)$ of the linear and nonlinear density contrasts within spherical cells."
 This eives the quasi-ucar nuit of these cistributions as well as the arge-uass exponeutial falloff of the alo Lass ‘ction., This gives the quasi-linear limit of these distributions as well as the large-mass exponential falloff of the halo mass function.
 An advantagee of our methoc is that it alows ns to exMicitlv check. tat realistic amounts of onordialpri 10Causslanityv have 10 Significant effect ou f clesity DEile «ot this saddle-poiut., An advantage of our method is that it allows us to explicitly check that realistic amounts of primordial non-Gaussianity have no significant effect on the density profile of this saddle-point.
 This eusures that shell«TOosslue apyears for (almost) the same nonlinear deusitv ὃz200 (seo. Valageas. 20095). so that the large-niass tail of the laο. lnass fiction can be derived provided halos are defined by à nouliuear density threshold that is beow this upxx bound (which is indeed the case).," This ensures that shell-crossing appears for (almost) the same nonlinear density $\delta_+ \ga 200$ (see Valageas 2009b), so that the large-mass tail of the halo mass function can be derived provided halos are defined by a nonlinear density threshold that is below this upper bound (which is indeed the case)."
 Althoieh this WOCCEure only gives the largeanass tail. we have proposed a snplec mange of variable. applied to the mass function fitted ο Cassial nuinerical suilations. that ohevs this ALEC-Mass asviuptotie while keevine the mass function LOLAlze to uuitv.," Although this procedure only gives the large-mass tail, we have proposed a simple change of variable, applied to the mass function fitted to Gaussian numerical simulations, that obeys this large-mass asymptotic while keeping the mass function normalized to unity."
 If one uses tl1ο Press-Schechter mass Tactient us gives back the result of Matarrese et al. (, If one uses the Press-Schechter mass function this gives back the result of Matarrese et al. (
"2000) but we aYSue that this srocedure js somewlha LOLre vatral if one wishes to τίCOVOP a lore accurate nass function iu ΓΙ casο,",2000) but we argue that this procedure is somewhat more natural if one wishes to recover a more accurate mass function in the Gaussian case.
 We also checked tha his aerees well with results from παο miuerical sinulations., We also checked that this agrees well with results from non-Gaussian numerical simulations.
 Next. we applied this luctrod to the two-poiu correlaion of massive halos. following the apmroach of INaiser (1081).," Next, we applied this method to the two-point correlation of massive halos, following the approach of Kaiser (1984)."
 As in Valageas (2009b) we take iito accoull the displacement of halo pairs under their mutua eyavitaional attraction., As in Valageas (2009b) we take into account the displacement of halo pairs under their mutual gravitational attraction.
 This gives the realsowe halo correlaion Sa;Gc). whence the real-space bias ba;Ce).," This gives the real-space halo correlation $\xi_M(x)$, whence the real-space bias $b_M(x)$."
 $iice his apxoach does not assunue that t1ο halo correlation Is sna1. the uoulinear formula it vields can be used o check whether the “lnuearized” fornt where oue oulv =keeps tjio linear term over the matter correlation €(.0)) is a good approximation iu the reei11ο of interest.," Since this approach does not assume that the halo correlation is small, the nonlinear formula it yields can be used to check whether the “linearized” form (where one only keeps the linear term over the matter correlation $\xi(x)$ ) is a good approximation in the regime of interest."
" As expected. we find that the correction Abss(e.fxL) to je Gaussian bias grows with ba;6e.0)ag uid with scale. ‘oughly as fxpbar(GeOe. up tow ~1TOEΊλῃος, "," As expected, we find that the correction $\Delta b_M(x,\fNL)$ to the Gaussian bias grows with $b_M(x,0)$ and with scale, roughly as $\fNL b_M(x,0) x^2$, up to $x \sim 100 h^{-1}$ Mpc."
Bevoud us scale. the barvon acoustic oscilliution and the fact iat the halo aud nujatter correlations do not chanec seu at tie same point lead to stro1 (scillations aud civerexut spikes for Όλα).," Beyond this scale, the baryon acoustic oscillation and the fact that the halo and matter correlations do not change sign at the same point lead to strong oscillations and divergent spikes for $b_M(x)$."
 This ieals that for ur1007; ANx: the bias is no longer a 1sefu quantity and Qιο should directly work with the halo ando latter correlations., This means that for $x > 100 h^{-1}$ Mpc the bias is no longer a useful quantity and one should directly work with the halo and matter correlations.
 Iu agrecinent with Desjacques (2008) we fixd that he two-point correlation of nassive halos. hat have a lareeoO bias. stronely[m amplifies tjio barvon acoistic oscillaion.," In agreement with Desjacques (2008) we find that the two-point correlation of massive halos, that have a large bias, strongly amplifies the baryon acoustic oscillation."
 Iu additioji we also obtain the meodificatloli assoclatec with prinordial uou-Caussianitv., In addition we also obtain the modifications associated with primordial non-Gaussianity.
 The barvon oscillalon renmius strouglv amplified. with a s1iall sift. but sonewhat less so Or positive fx," The baryon oscillation remains strongly amplified, with a small shift, but somewhat less so for positive $\fNL$."
" Finally. we used he “linearized” form, of the alo two-poiut correlation to derive the 1ido power spectruni and tje halo bias di Fourier space."," Finally, we used the “linearized” form of the halo two-point correlation to derive the halo power spectrum and the halo bias in Fourier space."
 We also give a siniple recipe that cusures that the ido power spectruni alwavs ronidus positive (tus onlv differs from the direct medicjon bw terms of order feye) and higher)., We also give a simple recipe that ensures that the halo power spectrum always remains positive (this only differs from the direct prediction by terms of order $\fNL^2$ and higher).
 We obtain a good agreement with numerical simulations. witlott introducimg :my yee parameter.," We obtain a good agreement with numerical simulations, without introducing any free parameter."
" Moreover. the wo formulae described al)ove allow oue to estimate the rage over which linCar approxinunatlons over fx, are sufficient."," Moreover, the two formulae described above allow one to estimate the range over which linear approximations over $\fNL$ are sufficient."
 Thus. we fiud hat terms of order I start dlaving a role at low & (he«0.012 1) for larec reeative [νι Cfxr1 KD. where the direct formula would give a negative power spectra.," Thus, we find that terms of order $\fNL^2$ start playing a role at low $k$ $k < 0.01 h$ $^{-1}$ ) for large negative $\fNL$ $\fNL < -100$ ), where the direct formula would give a negative power spectrum."
 These results. which «o not involve free parameters (except for the mass function. if one requires its full shape. where one needs the fit to ummerical simulations for Gaussian initial conditious) should be useful to coustraiu," These results, which do not involve free parameters (except for the mass function, if one requires its full shape, where one needs the fit to numerical simulations for Gaussian initial conditions) should be useful to constrain"
spectrum shown in retft 162316..,spectrum shown in \\ref{ft162316}.
 This star shows a peculiar spectrum. but most of the metal ines are weak and shallow with rotational and magnetic xoadening.," This star shows a peculiar spectrum, but most of the metal lines are weak and shallow with rotational and magnetic broadening."
 The spectrum has rather weak lines ofextsclii.. andextscii.," The spectrum has rather weak lines of, and."
. Phe magnetic field can be recognised from splitting of the line. although it was not obvious ancl required à synthetic spectrum for comparison as shown in relsv 168767..," The magnetic field can be recognised from splitting of the line, although it was not obvious and required a synthetic spectrum for comparison as shown in \\ref{sy168767}."
. The απο also shows doublet Zeeman splitting. which fits well when compared to a synthetic spectrum. calculated. for a magnetic field of κα. Despite this strong field. most other lines do not show splitting because of relatively rapid rotation with esiné=LEO+ 1.5kkmss+.," The line also shows doublet Zeeman splitting, which fits well when compared to a synthetic spectrum calculated for a magnetic field of kG. Despite this strong field, most other lines do not show splitting because of relatively rapid rotation with $v \sin i = 14.0 \pm 1.5$ $^{-1}$."
 With a strong magnetic field ancl relatively short rotational period. (not more than several davs given the relatively high. esin/). this star is one of the most interesting targets for future observations.," With a strong magnetic field and relatively short rotational period (not more than several days given the relatively high $v \sin i$ ), this star is one of the most interesting targets for future observations."
 We have only one spectrum: observations at other rotational phases may reveal an even stronger magnetic field., We have only one spectrum; observations at other rotational phases may reveal an even stronger magnetic field.
 This star was observed with FEROS three times as H was not clear whether it shows magnetic splitting., This star was observed with FEROS three times as it was not clear whether it shows magnetic splitting.
 Moderate intensity spectral lines of rare earth. elements are. present in the spectrum., Moderate intensity spectral lines of rare earth elements are present in the spectrum.
 The partial Zeeman splitting of the line is visible for two spectra ancl shows some hint of splitting for the third., The partial Zeeman splitting of the line is visible for two spectra and shows some hint of splitting for the third.
 his means that the magnetic ield is variable with an unknown rotation period., This means that the magnetic field is variable with an unknown rotation period.
 ASAS ohotometry did not give a clue to à possible period., ASAS photometry did not give a clue to a possible period.
 refsv177268. demonstrates a portion of one of the spectra ogether with a synthetic spectrum., \\ref{sy177268} demonstrates a portion of one of the spectra together with a synthetic spectrum.
 Another line also shows partial Zeeman splitting. thus supporting he discovery of à magnetic field in this star.," Another line also shows partial Zeeman splitting, thus supporting the discovery of a magnetic field in this star."
 The star has ohvsical parameters similar to known roAp stars. which led Abutinez&Ixurtz(1904). to search photometrically for rapid oscillations. but no evidence for pulsation was found.," The star has physical parameters similar to known roAp stars, which led \citet{Mart94} to search photometrically for rapid oscillations, but no evidence for pulsation was found."
 ‘This is another star with very strong lines oftextscii., This is another star with very strong lines of.
. Other strong rare earth. clement lines are also [ound in the spectrum. including and," Other strong rare earth element lines are also found in the spectrum, including and"
The dominant thermal ionisation ions in. protostellar discs are and lx (Umoebayashi&Nakano 1981).,The dominant thermal ionisation ions in protostellar discs are $^+$ and $^+$ \citep{umebayashi88}. .
. However the ion is more important at the onset of dynamically interesting ionisation levels because of its smaller ionisation potential., However the $^+$ ion is more important at the onset of dynamically interesting ionisation levels because of its smaller ionisation potential.
 Phe Saha equation for the electron fraction fron thermal ionisation can be approximated. by (Dalbus&Hawley2000). where ΑΗ —logGfI)JC{να is the potassium abundance relative to. hvdrogen. and. log(Av/HM)=7.We consider the elfeet of varving AI] in. Sections. 3.1.," The Saha equation for the electron fraction from thermal ionisation can be approximated by \citep{balbus00}, where $[K/H]=\log_{10}(K/H)-\log_{10}(K/H)_{\rm
 solar}$ is the potassium abundance relative to hydrogen and $\log_{10}(K/H)_{\rm solar}=-7$.We consider the effect of varying $[K/H]$ in Sections \ref{a}."
 Thermal ionisation is most important in the inner parts of the dise where the temperatures are higher and it falls olf exponentially., Thermal ionisation is most important in the inner parts of the disc where the temperatures are higher and it falls off exponentially.
 When thermal ionisation becomes negligible. the ionisation fraction is found with the balance of cosmic ray ionisation and recombination effects.," When thermal ionisation becomes negligible, the ionisation fraction is found with the balance of cosmic ray ionisation and recombination effects."
 Electrons may be captured: by dissociative recombination with molecular ions (of density napi) and raciative recombination with heavy-metal ions (of density Όρια)., Electrons may be captured by dissociative recombination with molecular ions (of density $n_{\rm M^+}$ ) and radiative recombination with heavy-metal ions (of density $n_{\rm m^+}$ ).
" Charge will be transferred from molecular ions to metal atoms and so the electron. fraction depends also upon the metal fraction where ny, is the number density of metals.", Charge will be transferred from molecular ions to metal atoms and so the electron fraction depends also upon the metal fraction where $n_{\rm m}$ is the number density of metals.
 Phe stancdare notation for metallicitv is by mass fraction. Z.," The standard notation for metallicity is by mass fraction, $Z$."
 This is related to the metal fraction with where gris the average mass of a particle in units of the mass of hydrogen., This is related to the metal fraction with where $\mu$ is the average mass of a particle in units of the mass of hydrogen.
 For example. the molecule of highest mass taking part in the reactions may be CO. which has a mass of 28 units.," For example, the molecule of highest mass taking part in the reactions may be CO, which has a mass of 28 units."
 In our galaxy Z—0.02 so ry2710L, In our galaxy $Z=0.02$ so $x_{\rm m}\gtrsim 7\times 10^{-4}$.
" In the LMC Z=0.008 is equivalent to ayz310 land in the SAIC Z=0.002 is equivalent tory,το107.", In the LMC $Z=0.008$ is equivalent to $x_{\rm m}\gtrsim 3\times 10^{-4}$ and in the SMC $Z=0.002$ is equivalent to $x_{\rm m}\gtrsim 7\times 10^{-5}$.
 Assuming the rates are the same for all species. the rate equation for the electron density ts and for the molecular ion density The ionisation rate. C. is discussed in the next section.," Assuming the rates are the same for all species, the rate equation for the electron density is and for the molecular ion density The ionisation rate, $\zeta$, is discussed in the next section."
 The recombination rate coefficients are for the radiative recombination of electrons with metal jons. for the cdissociative recombination of electrons with molecular ions and for the charge transfer from molecular ions to metal atoms.," The recombination rate coefficients are for the radiative recombination of electrons with metal ions, for the dissociative recombination of electrons with molecular ions and for the charge transfer from molecular ions to metal atoms."
 Conservation of charge tells us that Solving equations (11)) and (12)) in steady state with equation (16)) gives the cubic equation (Oppenheimer&Dalgarno1974:?:MillarFarquharWillacy) that can be solved to find the electron density.," Conservation of charge tells us that Solving equations \ref{ne}) ) and \ref{nions}) ) in steady state with equation \ref{charge}) ) gives the cubic equation \citep{oppenheimer74,spitzer98,millar97} that can be solved to find the electron density."
 There are three solutions to equation (17)) but. only one that is both real ancl positive., There are three solutions to equation \ref{cubic}) ) but only one that is both real and positive.
" The solution to equation (17)) can be simplified for the case where there is zero metallicity. ay,0. and find and also when the metals dominate so that ry,Sa. and (ο.Matsumura&Puelritz2003)."," The solution to equation \ref{cubic}) ) can be simplified for the case where there is zero metallicity, $x_{\rm m}=0$, and find and also when the metals dominate so that $x_{\rm m}\gg x_{\rm e}$ and \citep[e.g.][]{matsumura03}."
. We use these limits in Section 3.., We use these limits in Section \ref{disc}.
 We note that this model does not take the ellects of cust into account that may increase the size of the dead zone (e.g.Turner&Sano2008:Okuzumi 2011).," We note that this model does not take the effects of dust into account that may increase the size of the dead zone \citep[e.g.][]{turner08,okuzumi11}."
. We consider the only external ionisation source to be cosmic pavs with ionisation rate (wherewehaveneglectedthesecondtermtionbvSanoetal. 2000).. à)=10ts+ Vomasko1968) and x4=l100gcm (Umebavashi&Nakano 1981).," We consider the only external ionisation source to be cosmic rays with ionisation rate \cite[where we have neglected the second term in the equation
 by][]{sano00}, $\zeta_0=10^{-17}\,\rm s^{-1}$ \citep{spitzer68} and $\chi_{\rm cr}=100\,\rm g\,cm^{-2}$ \citep{umebayashi81}."
. The total surface censity in the active lavers is where zo; is the height of the disc above the mid-plane where the dead zone ends which we find in Section 2.5.., The total surface density in the active layers is where $z_{\rm crit}$ is the height of the disc above the mid-plane where the dead zone ends which we find in Section \ref{dz}.
 Young stellar. objects may be active N-ràav. sources. (e.glxovamaetal. 1994)., Young stellar objects may be active X-ray sources \citep[e.g][]{koyama94}.
. llowever. Alatsumura&Puclritz(2003) lind that cosmic rav ionisation clominates X-ray ionisation.," However, \cite{matsumura03} find that cosmic ray ionisation dominates X-ray ionisation."
 For example. comparing their figures 2 and 3 the dead: zone is muchlarger in both height and. radius with only X-rays.," For example, comparing their figures 2 and 3 the dead zone is muchlarger in both height and radius with only X-rays."
 X-rays may dominate only if the energy is very high. ος=510 keV. but this is much higher than for most observed. sources.," X-rays may dominate only if the energy is very high, $k T_{\rm X}=5-10\,\rm keV$ , but this is much higher than for most observed sources."
 In this work we neglect. X-ray ionisation and concentrate on the cllects of cosmic ray ionisation., In this work we neglect X-ray ionisation and concentrate on the effects of cosmic ray ionisation.
 Kibcseb5kms !kpe !/,K if $<\kappa> \simeq 5$ km $^{-1}$ $^{-1}$.
 Note that we have not had to assume a value of AV for this calculation. it can be anything between the typical absorption linewidth of 5 kms 1 and the typical warm phase linewidth Z10 kms 1.," Note that we have not had to assume a value of $\Delta V$ for this calculation, it can be anything between the typical absorption linewidth of $\lesssim$ 5 km $^{-1}$ and the typical warm phase linewidth $\gtrsim$ 10 km $^{-1}$."
" The brightness temperature we expect from regions with p=1 is different from this. however. since Z,., is partially absorbed by the cool clouds before it gets to us. ancl the cool clouds themselves contribute some emission. given bv T,(1—€7)."," The brightness temperature we expect from regions with $\rho \gtrsim 1$ is different from this, however, since $T_{w,u}$ is partially absorbed by the cool clouds before it gets to us, and the cool clouds themselves contribute some emission, given by $T_{cool} \ (1 \ - \ e^{-\tau})$."
" If the warn phase gas is well mixed around ancl among the cool clouds. then its attenuated brightness temperature is given by Aclding on the emission [rom the cool cloud(s). 744,(1—€ 7). gives the total brightness lemperature seen in emission For our assumed 7| this gives ~0.6 (Z4+ Diy). or about 0.6 (651x + 80IN) = 87x. In the alternative geometry. assumed in section 4. where (he warm phase gas is partly in front. of the absorbing cloud and partly.behind. we get or. in (he notation of equations 10 - 14. For εξ. and 7=0.63 (7-1) this gives 102 Ix. In both cases the contributions from the warm and cool phases are about equal al about. 50 Ix each."," If the warm phase gas is well mixed around and among the cool clouds, then its attenuated brightness temperature is given by Adding on the emission from the cool cloud(s), $T_{cool} \ (\ 1\ - \ e^{-\tau} )$ , gives the total brightness temperature seen in emission For our assumed $\tau \simeq 1$ this gives $\sim$ 0.6 $T_{cool}
+ T_{w,u}$ ), or about 0.6 (65K + 80K) = 87K. In the alternative geometry assumed in section 4, where the warm phase gas is partly in front of the absorbing cloud and partlybehind, we get or, in the notation of equations 10 - 14, For $\epsilon$ =0.5 and $x$ =0.63 $\tau$ =1) this gives 102 K. In both cases the contributions from the warm and cool phases are about equal at about 50 K each."
 Going to higher optical depths does not change these numbers much., Going to higher optical depths does not change these numbers much.
" For 7—o€ Gr— 1) equation 26 predicts 754115 Ix and in equation 24 the peak brightness actually decreases to approach ο,& 65Ix. Of course the actual value of (he peak brightness temperature will fluctuate around (these values due to the relatively wide distribution ofvalues of ρω.", For $\tau \rightarrow \infty$ $x \rightarrow 1$ ) equation 26 predicts $T_B \rightarrow 115$ K and in equation 24 the peak brightness actually decreases to approach $T_{cool}\simeq$ 65K. Of course the actual value of the peak brightness temperature will fluctuate around these values due to the relatively wide distribution ofvalues of $T_{cool}$ .
 The highest values we see presumably, The highest values we see presumably
D. Teyssier**. Observatoire de Paris-Meudon. Meudon. France starr-forming regions Is not well known.,"D. Teyssier Observatoire de Paris-Meudon, Meudon, France r-forming regions is not well known."
might be due to the SPH artificial viscosity or the artificial smoothing of density enhancements.,might be due to the SPH artificial viscosity or the artificial smoothing of density enhancements.
" In this Letter, we present results on the convergence of the critical value of 6 using the grid-based hydrodynamics codeFARGO2)mainBodyCitationEnd437|masset00a,masset0Ob."," In this Letter, we present results on the convergence of the critical value of $\beta$ using the grid-based hydrodynamics code."
" First of all, with a smooth initial setup as in many previous simulations of gravitationally unstable discs, we find similar non-convergence withFARGO, indicating that the problem is not specific to SPH."," First of all, with a smooth initial setup as in many previous simulations of gravitationally unstable discs, we find similar non-convergence with, indicating that the problem is not specific to SPH."
" We then show that numerical convergence can be reached by avoiding a sharp boundary between the gravito-turbulent inner part of the disc and the still laminar outer part, which arises when using smooth initial conditions during the initial stage of cooling towards Qz1."," We then show that numerical convergence can be reached by avoiding a sharp boundary between the gravito-turbulent inner part of the disc and the still laminar outer part, which arises when using smooth initial conditions during the initial stage of cooling towards $Q\approx 1$."
" If smooth initial conditions are not used, we find a value of f close to that found by?."," If smooth initial conditions are not used, we find a value of $\betac$ close to that found by."
". We use the grid-based hydrodynamics codeFARGO(7??),, for which a self-gravity solver was presented in?."," We use the grid-based hydrodynamics code, for which a self-gravity solver was presented in."
. We have implemented the simple cooling law where c is the internal energy and too is given by equation (2))., We have implemented the simple cooling law where $\epsilon$ is the internal energy and $t_\mathrm{cool}$ is given by equation \ref{eqtcool}) ).
 A standard prescription for artificial viscosity is used to handle shocks., A standard prescription for artificial viscosity is used to handle shocks.
" Following ?,, we choose our disc to lie between R=0.25 and R=25."," Following , we choose our disc to lie between $R=0.25$ and $R=25$."
" The self-gravity module of requires a logarithmic grid in the radial direction, and we choose our grid so as to give square cells everywhere (AR/R= Ao)."," The self-gravity module of requires a logarithmic grid in the radial direction, and we choose our grid so as to give square cells everywhere $\Delta R/R \approx \Delta \phi$ )."
" Our lowest resolution has 512 cells in the radial direction and 768 cells in the azimuthal direction, and we increase the resolution by a factor of 2 and 4."," Our lowest resolution has 512 cells in the radial direction and 768 cells in the azimuthal direction, and we increase the resolution by a factor of 2 and 4."
 We use outflow boundary conditions at the inner and outer grid edges., We use outflow boundary conditions at the inner and outer grid edges.
" Following?,, we choose our disc to have surface mass density XοςR! and temperature ΤοςR/?."," Following, we choose our disc to have surface mass density $\Sigma \propto R^{-1}$ and temperature $T \propto R^{-1/2}$."
" The angular velocity is Keplerian, corrected for the mass and pressure of the disc."," The angular velocity is Keplerian, corrected for the mass and pressure of the disc."
 A small level (~ 0.1%) of white noise is put on top to break the axisymmetry and allow spiral waves to form., A small level $\sim 0.1\%$ ) of white noise is put on top to break the axisymmetry and allow spiral waves to form.
" The surface density and aspect ratio H/R at R=1 are chosen so as to give a total disc mass of 0.25 M., and so that the minimum value of Qz2 at the outeredge of the disc."," The surface density and aspect ratio $H/R$ at $R=1$ are chosen so as to give a total disc mass of $0.25$ $M_*$, and so that the minimum value of $Q\approx 2$ at the outeredge of the disc."
 We use y=5/3 throughout., We use $\gamma=5/3$ throughout.
" We measure the total shear stress in the simulations by calculating where gr and gy, are the gravitational accelerations, dur and óu, are velocity fluctuations, and (-) denotes an azimuthal average."," We measure the total shear stress in the simulations by calculating where $g_R$ and $g_\varphi$ are the gravitational accelerations, $\delta v_R$ and $\delta v_\varphi$ are velocity fluctuations, and $\left< \cdot \right>$ denotes an azimuthal average."
 The vertical integral iscalculated by computing gg and gy at different values of z details)., The vertical integral iscalculated by computing $g_R$ and $g_\varphi$ at different values of $z$ .
". To compare with equation (3)), we can use an a-parametrisation:"," To compare with equation \ref{eqabc}) ), we can use an $\alpha$ -parametrisation:"
Type Ia supernovae (SNe la) play an important role in the study of cosmic evolution. especially in cosmology.,"Type Ia supernovae (SNe Ia) play an important role in the study of cosmic evolution, especially in cosmology."
 They have been applied successfully in determining cosmological parameters (e.g.. O and A: Riess et al.," They have been applied successfully in determining cosmological parameters (e.g., $\Omega$ and $\Lambda$; Riess et al."
 1998: Perlmutter et al., 1998; Perlmutter et al.
 1999)., 1999).
 It is generally believed that SNe Ia are thermonuclear explosions of carbon-oxygen white dwarfs (CO WDs) in binaries (for the review see Nomoto et al., It is generally believed that SNe Ia are thermonuclear explosions of carbon–oxygen white dwarfs (CO WDs) in binaries (for the review see Nomoto et al.
 1997)., 1997).
 However. there is still no agreement on the nature of their progenitors (Hillebrandt Niemeyer 2000; Podsiadlowski et al.," However, there is still no agreement on the nature of their progenitors (Hillebrandt Niemeyer 2000; Podsiadlowski et al."
 2008; Wang et al., 2008; Wang et al.
 2008). and no 5 la progenitor system has been conclusively identified from before the explosion.," 2008), and no SN Ia progenitor system has been conclusively identified from before the explosion."
 Over the past few decades. two families of SN Ia progenitor models have been proposed. t-e.. the double-degenerate (DD) and single-degenerate (SD) models.," Over the past few decades, two families of SN Ia progenitor models have been proposed, i.e., the double-degenerate (DD) and single-degenerate (SD) models."
 Of these two models. the SD model is widely accepted at present (Nomoto et al.," Of these two models, the SD model is widely accepted at present (Nomoto et al."
 1984)., 1984).
 It is suggested that the DD model. which involves the merger of two CO WDs (Iben Tutukov 1984: Webbink 1984: Han 1998). likely leads to an accretion-induced collapse rather than to an SN Ia (Nomoto Iben 1985).," It is suggested that the DD model, which involves the merger of two CO WDs (Iben Tutukov 1984; Webbink 1984; Han 1998), likely leads to an accretion-induced collapse rather than to an SN Ia (Nomoto Iben 1985)."
 For the SD model. the companion is probably a main-sequence (MS) star. a slightly evolved subgiant star (WD + MS channel). or a red-giant star (WD + RG channel) (e.g.. Hachisu et al.," For the SD model, the companion is probably a main-sequence (MS) star, a slightly evolved subgiant star (WD + MS channel), or a red-giant star (WD + RG channel) (e.g., Hachisu et al."
 1996. 1999a.b: Li & van den Heuvel 1997; Langer et al.," 1996, 1999a,b; Li $\&$ van den Heuvel 1997; Langer et al."
 2000: Han & Podsiadlowski 2004. 2006: Che & Li 2007. 2009; Meng et al.," 2000; Han $\&$ Podsiadlowski 2004, 2006; Chen $\&$ Li 2007, 2009; Meng et al."
 2009: Lü et al., 2009; Lü et al.
 2009; Wang. Li Han 2009).," 2009; Wang, Li Han 2009)."
 An explosion following the merger of two WDs would leave no remnant. while the companion star 1 the SD model would survive and be potentially identifiable (Podsiadlowski 2003).," An explosion following the merger of two WDs would leave no remnant, while the companion star in the SD model would survive and be potentially identifiable (Podsiadlowski 2003)."
 There has been no conclusive proof yet that any individual object is the surviving companion star of a SN Ia. It will be a promising method to test SN Ia progenitor models by identifying their surviving companion stars., There has been no conclusive proof yet that any individual object is the surviving companion star of an SN Ia. It will be a promising method to test SN Ia progenitor models by identifying their surviving companion stars.
 Yoon & Langer (2003) followed the evolution of a CO WD + He star system with a 1.044. CO WD and a 1.6M. He star in a dd orbit.," Yoon $\&$ Langer (2003) followed the evolution of a CO WD + He star system with a $1.0\,M_{\odot}$ CO WD and a $1.6\,M_{\odot}$ He star in a d orbit."
 In this binary. the WD accretes He from the He star and grows in mass to the Chandrasekhar (Ch) mass.," In this binary, the WD accretes He from the He star and grows in mass to the Chandrasekhar (Ch) mass."
 SNe Ia from this binary channel can neatly avoid H lines., SNe Ia from this binary channel can neatly avoid H lines.
 Recently. Wang et al. (," Recently, Wang et al. ("
20092) systematically studied the WD + He star channel of SNe Ia. In the study. they carried out binary evolution calculations of this channel for about 2600 close WD binaries. in which a CO WD accretes material from an He MS star or an He subgiant to increase its mass to the Ch mass.,"2009a) systematically studied the WD + He star channel of SNe Ia. In the study, they carried out binary evolution calculations of this channel for about 2600 close WD binaries, in which a CO WD accretes material from an He MS star or an He subgiant to increase its mass to the Ch mass."
 The study shows the parameter spaces for the progenitors of SNe la. By using a detailed binary population synthesis (BPS) approach. Wang et al. (," The study shows the parameter spaces for the progenitors of SNe Ia. By using a detailed binary population synthesis (BPS) approach, Wang et al. ("
"2009b) find that the Galactic SN la birthrate from this channel is «0.3x107yr! and that this channel can produce SNe la with short delay times (45-140 MMyr) from the star formation to SN explosion,","2009b) find that the Galactic SN Ia birthrate from this channel is $\sim$$0.3\times 10^{-3}\ {\rm
yr}^{-1}$ and that this channel can produce SNe Ia with short delay times $\sim$ $-$ Myr) from the star formation to SN explosion."
 The companion star in this channel would survive and show distinguishing properties., The companion star in this channel would survive and show distinguishing properties.
 In recent years hypervelocity stars (HVSs) have been observed in the halo of the Galaxy., In recent years hypervelocity stars (HVSs) have been observed in the halo of the Galaxy.
 HVSs are stars with a velocity so great that they are able to escape the gravitational pull of the Galaxy., HVSs are stars with a velocity so great that they are able to escape the gravitational pull of the Galaxy.
 However. the formation of HVSs is still unclear (for a recent review see Tutukov Fedorova 2009).," However, the formation of HVSs is still unclear (for a recent review see Tutukov Fedorova 2009)."
 It has been suggested that such HVSs can be formed by the tidal disruption of a binary through interaction with the super-massive black hole (SMBH) at the Galactic center (GC) (Hills 1988; Yu Tremaine 2003)., It has been suggested that such HVSs can be formed by the tidal disruption of a binary through interaction with the super-massive black hole (SMBH) at the Galactic center (GC) (Hills 1988; Yu Tremaine 2003).
 The first three HVSs have only recently. been discovered. serendipitously (e.g.. Brown et al.," The first three HVSs have only recently been discovered serendipitously (e.g., Brown et al."
 2005: Hirsch et al., 2005; Hirsch et al.
 2005: Edelmann et al., 2005; Edelmann et al.
 2005)., 2005).
 Up to now. about 17 HVSs have been discovered in the Galaxy (Brown et al.," Up to now, about 17 HVSs have been discovered in the Galaxy (Brown et al."
 2009; Tillich et al., 2009; Tillich et al.
 2009). most of which are B-type stars. probably with masses ranging from 3 to 5M. (Brown et al.," 2009), most of which are B-type stars, probably with masses ranging from 3 to $M_\odot$ (Brown et al."
 2005. 2009: Edelmann et al.," 2005, 2009; Edelmann et al."
 2005)., 2005).
 One HVS. HE 0437-5439. is known to be an apparently normal early B-type star.," One HVS, HE 0437-5439, is known to be an apparently normal early B-type star."
 Edelmann et al. (, Edelmann et al. (
2005) suggests that the star could have originated in the Large Magellanie Cloud. because it is much closer to this galaxy (18 kpe) than to the GC (see also Przybilla et al.,"2005) suggests that the star could have originated in the Large Magellanic Cloud, because it is much closer to this galaxy (18 kpc) than to the GC (see also Przybilla et al."
 2008)., 2008).
 At present. only one HVS. US 708. is a subdwarf O (sdO) star. and Hirsch et al. (," At present, only one HVS, US 708, is a subdwarf O (sdO) star, and Hirsch et al. ("
2005) speculatS that US 708 is formed by the merger of two He WDs in a close binary induced by the interaction with the SMBH in the GC and then escaped.,2005) speculatS that US 708 is formed by the merger of two He WDs in a close binary induced by the interaction with the SMBH in the GC and then escaped.
 Recently. Perets (2009) has suggested that US 708 may have been ejected as a binary from a triple disruption by the SMBH. which later on evolved and merged to form an sdO star.," Recently, Perets (2009) has suggested that US 708 may have been ejected as a binary from a triple disruption by the SMBH, which later on evolved and merged to form an sdO star."
 Because of the existence of the short orbital periods ( hh) for the WD + He star systems. Justham et al. (," Because of the existence of the short orbital periods $\sim$ h) for the WD + He star systems, Justham et al. ("
2009) argues that the WD + He star channel of SNe Ia may provide a natural explanation for stars like US 708.,2009) argues that the WD + He star channel of SNe Ia may provide a natural explanation for stars like US 708.
The implemented method (WASIET) for warped galaxies can be applied to galaxies where the sampling of the velocity. field. is poor. ancl it enables us to trace the rotation curve along any possible warp instead of keeping a fixed position angle. like in methods based on the analysis of the position-velocity diagram: the shortcomings of keeping a fixed position angle are discussed in Verganietal.(2004).,"The implemented method (WAMET) for warped galaxies can be applied to galaxies where the sampling of the velocity field is poor, and it enables us to trace the rotation curve along any possible warp instead of keeping a fixed position angle, like in methods based on the analysis of the position-velocity diagram; the shortcomings of keeping a fixed position angle are discussed in \citet{V:04}."
.. In order to derive the rotation curves. we applied the tilted-ring mocdelling of the velocity Ποια for NGC 1090.," In order to derive the rotation curves, we applied the tilted-ring modelling of the velocity field for NGC 1090."
 For the other galaxies we determined. the kinematical centre and the systemic velocity by minimising the dillerences between the two sides., For the other galaxies we determined the kinematical centre and the systemic velocity by minimising the differences between the two sides.
" The errors are the maximum of the three Following values: the dillerence between the velocities of the approaching and the receding sides. our correction for beam broadening of the profiles. and a ""minimum error"" equal to (2/sin i) kms ¢."," The errors are the maximum of the three following values: the difference between the velocities of the approaching and the receding sides, our correction for beam broadening of the profiles, and a “minimum error” equal to $2/sin~i$ ) km $^{-1}$."
 The METE/WAAMIZT rotation curve served as an input to construct model cata cubes with the task GALALOD within GIPSY (vanderHulst.etal.1992):: these models have been then compared to the observed data cubes., The MET/WAMET rotation curve served as an input to construct model data cubes with the task GALMOD within GIPSY \citep{vdH:92}; these models have been then compared to the observed data cubes.
 Such models are built with the assumption that the neutral hydrogen moves along circular orbits: a series of ecometrical ancl physical parameters. among which the rotation curve. allows us to create. model. observations of the disces.," Such models are built with the assumption that the neutral hydrogen moves along circular orbits; a series of geometrical and physical parameters, among which the rotation curve, allows us to create model observations of the discs."
 These models. when compared to the real data cubes. verily whether they are a fair representation of the dise.," These models, when compared to the real data cubes, verify whether they are a fair representation of the disc."
 Some input parameters are inferred directly from the observations ancl kept fixed (e.g. the surface density profile). while others (e.g.. the kinematical centre and the svstemic velocity) are first estimated. from the observations and then adjusted in order to optimise the mateh between the observed. and the model cubes.," Some input parameters are inferred directly from the observations and kept fixed (e.g., the surface density profile), while others (e.g., the kinematical centre and the systemic velocity) are first estimated from the observations and then adjusted in order to optimise the match between the observed and the model cubes."
 In two cases (ESO 287-C:13 and NCC 1090. the galaxies with the steepest rotation curves) the rotation curve derived with AIF/WAAIE'D had to be slightly moclilied in order to better reproduce the observed. data cubes.," In two cases (ESO 287-G13 and NGC 1090, the galaxies with the steepest rotation curves) the rotation curve derived with MET/WAMET had to be slightly modified in order to better reproduce the observed data cubes."
 We found. out that it was necessary to increase the value of the rotation velocity of the 23 innermost points by values of the order of their 1@ error: this is probably due to the fact. that. even with the MET/WASMIZT method. when the rotation curve rises steeply we are not able to retrieve precisely the true rotation velocity: this is however possible by mocelling the data cube.," We found out that it was necessary to increase the value of the rotation velocity of the 2–3 innermost points by values of the order of their $1~ \sigma$ error; this is probably due to the fact that, even with the MET/WAMET method, when the rotation curve rises steeply we are not able to retrieve precisely the true rotation velocity; this is however possible by modelling the data cube."
 In Fig. 5.," In Fig. \ref{channels2},"
 we display some channel maps (observed and modelled) to show that with our final rotation curves we are able to provide a very good reproduction of the observations., we display some channel maps (observed and modelled) to show that with our final rotation curves we are able to provide a very good reproduction of the observations.
 All the galaxies of our sample were successfully, All the galaxies of our sample were successfully
use Doppler imaging to map the acerction disc.,use Doppler imaging to map the accretion disc.
 Observations are presented. in section 2 followed by the analysis of the tomograms in section 3., Observations are presented in section 2 followed by the analysis of the tomograms in section 3.
 The tidal origin of the spirals is discussed in section 4., The tidal origin of the spirals is discussed in section 4.
 The data we present here are part of a long termi service orogram to study LP Peg throughout its outburst. evcle., The data we present here are part of a long term service program to study IP Peg throughout its outburst cycle.
 Time-resolved CCD spectrophotometry with the 2.5m Isaac ewton Telescope on La Palma was used to study the strong emission lines originating in the accretion disc both during quiescence and outburst., Time-resolved CCD spectrophotometry with the 2.5m Isaac Newton Telescope on La Palma was used to study the strong emission lines originating in the accretion disc both during quiescence and outburst.
 Here we will focus our attention on he data obtained during the night of 19 August. 1993.," Here we will focus our attention on the data obtained during the night of 19 August, 1993."
 LP ee had just gone into outburst a day before and was close to its maximum brightness level., IP Peg had just gone into outburst a day before and was close to its maximum brightness level.
 Phe Intermediate Dispersion Spectrograph was used to obtain spectra between 6300 and 6800... covering La and HoL(A6G678) at à mean dispersion of 0.56 ppixel or 38 km “pixel 7.," The Intermediate Dispersion Spectrograph was used to obtain spectra between 6300 and 6800, covering $\alpha$ and $\lambda 6678$ ) at a mean dispersion of 0.56 $^{-1}$ or 38 km $^{-1}$ $^{-1}$."
 A lO24pixel TEI CCD chip recorded long slit. spectra of LP Peg and a Comparison star to account for slit-losses., A $\times$ 1024-pixel TEK CCD chip recorded long slit spectra of IP Peg and a comparison star to account for slit-losses.
 Neon are spectra were regularly recorded Lor wavelength calibration ane the lux standard 2874211. was used. for Dux. calibration., Neon arc spectra were regularly recorded for wavelength calibration and the flux standard $^{\circ}$ 4211 was used for flux calibration.
 This setup allowed us to optimally extract spectra with an absolute Dux scale., This setup allowed us to optimally extract spectra with an absolute flux scale.
 A total of 15 spectra with an exposure ime of 360 s where obtained sampling of the 3.8 hour xunarv orbit., A total of 15 spectra with an exposure time of 360 s where obtained sampling of the 3.8 hour binary orbit.
 The top panels of Figure 1 show the Ho and Hel(6678) ine profiles as a function of binary phase after subtracting a low order spline fit to the continuum. of the individual spectra., The top panels of Figure 1 show the $\alpha$ and HeI(6678) line profiles as a function of binary phase after subtracting a low order spline fit to the continuum of the individual spectra.
 Orbital phases were caleulated using the Wolfet al. (, Orbital phases were calculated using the Wolf et al. (
1993) ephemeris without their quadratic term: with Z5 corresponding to mid-eclipse.,1993) ephemeris without their quadratic term; with $T_0$ corresponding to mid-eclipse.
 Phe AB Y12.6 mag continuum increasing by ~7% during the 2 hour observing window. shows that LP Peg was near the top of its rise to outburst. which tvpically lasts 1.1.5 d. To interpret. the phase dependent. line profiles flr.) (Fig.," The $\sim$ 12.6 mag continuum increasing by $\sim 7\%$ during the 2 hour observing window, shows that IP Peg was near the top of its rise to outburst, which typically lasts 1–1.5 d. To interpret the phase dependent line profiles $f(v,\phi)$ (Fig."
 1). we use Doppler tomography. (Marsh. Horne 1988). an indirect de-projection technique very similar to CAT scanning used in medical imaging.," 1), we use Doppler tomography (Marsh Horne 1988), an indirect de-projection technique very similar to CAT scanning used in medical imaging."
" Phe Doppler map I(V,.N 4) gives the emission line lux of gas moving with velocity vector V—(VV) in the rotating frame of the binary."," The Doppler map $I$ $_x$ $_y)$ gives the emission line flux of gas moving with velocity vector $V=(V_x,V_y)$ in the rotating frame of the binary."
 As the binary rotates. projections of the rotating velocity. vector onto the line of sight traces the sinusoidal radial velocity curve: The observed line profiles οιὁ) can therefore be niodelled as projections of the map Z(V V) without making specific assumptions about the form of the velocity. field. of the accretion How (see also Robinson. Marsh Smak 1993 and Horne 1991).," As the binary rotates, projections of the rotating velocity vector onto the line of sight traces the sinusoidal radial velocity curve; The observed line profiles $f(v,\phi)$ can therefore be modelled as projections of the map $I$ $_x$ $_y)$ without making specific assumptions about the form of the velocity field of the accretion flow (see also Robinson, Marsh Smak 1993 and Horne 1991)."
 A maximum entropy implementation was used where the Doppler image is built up iteratively., A maximum entropy implementation was used where the Doppler image is built up iteratively.
 Any given map is projected to produce the predicted line profiles for the particular map., Any given map is projected to produce the predicted line profiles for the particular map.
 x72 statistic is used to determine goodness of fit while the entropy is niaximised to select the simplest image that can fit the data to the required X7 value., $\chi^2$ statistic is used to determine goodness of fit while the entropy is maximised to select the simplest image that can fit the data to the required $\chi^2$ value.
 This technique assumes that the dise pattern Is constant hroughout the data set (in the co-rotating [rame of he binary) so that the line variations can be mocleled w projection effects., This technique assumes that the disc pattern is constant throughout the data set (in the co-rotating frame of the binary) so that the line variations can be modeled by projection effects.
 Transient features will therefore. be averaged oul over the map so that the average co-rotating xutern is recovered., Transient features will therefore be averaged out over the map so that the average co-rotating pattern is recovered.
 Tidal distortions co-rotate in the binary rame and therefore do not suller from this restriction and are ideally recovered. by Doppler. tomography., Tidal distortions co-rotate in the binary frame and therefore do not suffer from this restriction and are ideally recovered by Doppler tomography.
 (X. second xoblem can be secular variability of the svstem within the data set used. for tomography., A second problem can be secular variability of the system within the data set used for tomography.
 In: our case the continuum showed little increase during the course of our observations (i.c. outburst. was developed) and as our observations cover only ~ 2h. which is sullicient to caleulate a Doppler image as more than hall of the orbital period is covered. secular changes were negligible.," In our case the continuum showed little increase during the course of our observations (i.e. outburst was developed) and as our observations cover only $\sim$ 2h, which is sufficient to calculate a Doppler image as more than half of the orbital period is covered, secular changes were negligible."
 Furthermore. line flux. variations were compatible with the changing contribution of the companion star as the illuminated. inner face comes into view. while the dise contribution was stable.," Furthermore, line flux variations were compatible with the changing contribution of the companion star as the illuminated inner face comes into view, while the disc contribution was stable."
 Middle panels of Figure 1. show the two maps constructed from the observed Lo. ancl Hel(6673) line Dux., Middle panels of Figure 1 show the two maps constructed from the observed $\alpha$ and HeI(6678) line flux.
 As à comparison. bottom panels show preclicted data. anc can be used to check how well the Doppler image reproduces the observed. line emission.," As a comparison, bottom panels show predicted data and can be used to check how well the Doppler image reproduces the observed line emission."
 The gas strezum trajectory anc position of the companion star's Roche lobe is plotted basec on the known system parameters (Marsh. Horne 1990)., The gas stream trajectory and position of the companion star's Roche lobe is plotted based on the known system parameters (Marsh Horne 1990).
 Strong secondary star emission (Ix3—300 Emi 1) is visible in both lines. à common feature of cwarf novae in outburs and is thought to be due to irradiation of the inner face of the star.," Strong secondary star emission $_2$ =300 km $^{-1}$ ) is visible in both lines, a common feature of dwarf novae in outburst and is thought to be due to irradiation of the inner face of the star."
 However. emission from the companion has also been observed. during quiescence (Llarlaftis et al.," However, emission from the companion has also been observed during quiescence (Harlaftis et al."
 1994) anc can be related to intrinsic activity of the late type star as the secondary star is co-rotating in a binary with a period of only several hours., 1994) and can be related to intrinsic activity of the late type star as the secondary star is co-rotating in a binary with a period of only several hours.
 There is also à weak low velocity componen in the Ho image. which was observed a week later by Steeghs et.," There is also a weak low velocity component in the $\alpha$ image, which was observed a week later by Steeghs et."
 al (1996) who propose prominence like structures to be responsible for this feature., al (1996) who propose prominence like structures to be responsible for this feature.
 This emission is thus already present early in outburst. even though it is more pronounced a week Later.," This emission is thus already present early in outburst, even though it is more pronounced a week later."
 Disc emission is centered on the white cwarl (Ix; —147 km s 5) and has a strong azimuthal asvmimetrv in the form of à two armed spiral pattern., Disc emission is centered on the white dwarf $_1$ =147 km $^{-1}$ ) and has a strong azimuthal asymmetry in the form of a two armed spiral pattern.
 Both lines show similar structure but the arms are more sharply defined. in the Ucl map., Both lines show similar structure but the arms are more sharply defined in the HeI map.
 The line llux in these spirals is about a factor of ~4 stronger than that of the disc emission. outside the spirals pointing to considerable heating and density enhancement., The line flux in these spirals is about a factor of $\sim$ 4 stronger than that of the disc emission outside the spirals pointing to considerable heating and density enhancement.
 The velocities of the dise material in the two arms decrease [rom ~700 kms 1 down to ~500 kn st with increasing azimuth. suggesting a highly non Ixeplerian Low.," The velocities of the disc material in the two arms decrease from $\sim$ 700 km $^{-1}$ down to $\sim$ 500 km $^{-1}$ with increasing azimuth, suggesting a highly non Keplerian flow."
 A Ixeplerian accretion disc on the other hand would produce circular rings of emission. cach velocity corresponding to," A Keplerian accretion disc on the other hand would produce circular rings of emission, each velocity corresponding to"
"equation (2) becomes To simplify Che equations. we make the following substitutions: where Under these (ransformations. equations (1) and (10) do not change. but (he other equations become Nolice the equations are not invariant under (he transformation /— —/. e,— τν, and e.—στον, so the solutions can not be Gime-reversible.","equation (2) becomes To simplify the equations, we make the following substitutions: where Under these transformations, equations (1) and (10) do not change, but the other equations become Notice the equations are not invariant under the transformation $t \rightarrow -t$ , $v_{r} \rightarrow -v_{r}$ , and $v_{\varphi} \rightarrow -v_{\varphi}$, so the solutions can not be time-reversible."
 Thus. these equations describe a nonlinear evolving svstem.," Thus, these equations describe a nonlinear evolving system."
 In following the nonlinear evolution of dynamically evolving systems. the technique of sell-similar analvsis is useful. as it allows a set of partial differential equations. such as those above. to be transformed into a set of ordinary differential equations.," In following the nonlinear evolution of dynamically evolving systems, the technique of self-similar analysis is useful, as it allows a set of partial differential equations, such as those above, to be transformed into a set of ordinary differential equations."
 A similarity solution. although constituting only a limited part of the problem. is often useful in understandingthe," A similarity solution, although constituting only a limited part of the problem, is often useful in understandingthe"
We thank DDavis for his many vears of service maintaining the database for the CLA dieital speeclometers. and CCaruso. ZZajac. DDBerlind. aud TTorres lor making many of the observations.,"We thank Davis for his many years of service maintaining the database for the CfA digital speedometers, and Caruso, Zajac, Berlind, and Torres for making many of the observations."
 BWC! thanks the National Science Foundation for grants and AST-0305431 to the University of North Carolina., BWC thanks the National Science Foundation for grants AST-9988156 and AST-0305431 to the University of North Carolina.
 JBL thanks the National Science Foundation for grants AST-9988247 and AST-0307340 to Bowling Green State University., JBL thanks the National Science Foundation for grants AST-9988247 and AST-0307340 to Bowling Green State University.
 We especially acknowledge the great utility of the 2\TASS database. as well as (he SIMDBAD database. maintained by the CDS in Strasbourg. France.," We especially acknowledge the great utility of the 2MASS database, as well as the SIMBAD database, maintained by the CDS in Strasbourg, France."
contributions (see the bottom panels of Figures 2a.b) and can be neglected for most. purposes.,"contributions (see the bottom panels of Figures 2a,b) and can be neglected for most purposes."
 Recently Zaldarriaga and. Seljak (1998). released a version of CAIBPAST that includes gravitational lensing., Recently Zaldarriaga and Seljak (1998) released a version of CMBFAST that includes gravitational lensing.
 A comparison between our code ancl theirs shows good agreement with differences smaller than 5% of a lensing eencrated contribution., A comparison between our code and theirs shows good agreement with differences smaller than $5\%$ of a lensing generated contribution.
 To assess whether the lensing cllects are detectable. we analyse the two classes of degenerate models illustrated in Figures d. and 2.," To assess whether the lensing effects are detectable, we analyse the two classes of degenerate models illustrated in Figures \ref{fig1}
 and 2."
" Models of the first class (hereafter Class 1) have power spectra which are indistinguishable from that of a spatially Dat cosmological model with ay,=0.1. a,=0.0125 and ων=0.15."," Models of the first class (hereafter Class 1) have power spectra which are indistinguishable from that of a spatially flat cosmological model with $\omega_{\rm m}=0.1,$ $\omega_{\rm b}=0.0125$ and $\omega_\Lambda=0.15$."
" Class 2 models have power spectra that are indistinguishable with that of an open universe with zero cosmological constant. wy,=OL. a,=0.0125 and wy=0.15."," Class 2 models have power spectra that are indistinguishable with that of an open universe with zero cosmological constant, $\omega_{\rm m}=0.1$, $\omega_{\rm b}=0.0125$ and $\omega_K=0.15$."
 The models in each class are labelled with a letter. (a. b. c. d) in ascending order of the value of the Hubble constant (&= 0.5.0.6.0.7.0.8).," The models in each class are labelled with a letter, (a, b, c, d) in ascending order of the value of the Hubble constant $h=0.5,0.6,0.7, 0.8$ )."
 For all models we assume a precisely scalo-invariant spectrum. of scalar adiabatie perturbations and no contribution from. tensor modes., For all models we assume a precisely scale-invariant spectrum of scalar adiabatic perturbations and no contribution from tensor modes.
 The parameters for the two families of models are specified in Table 1., The parameters for the two families of models are specified in Table 1.
" For comparison we have also computed results for the standard’ CDM. model ay,= 0.25."," For comparison we have also computed results for the `standard' CDM model $\omega_{\rm m}=0.25$ ,"
between interacting particles is only limited by the system's inhomogencity.,between interacting particles is only limited by the system's inhomogeneity.
 The spatial structure of the svstem matters as well as the details of the particle orbits., The spatial structure of the system matters as well as the details of the particle orbits.
 The consistent inclusion of collective screening cllects in a kinetic equation for electrically interacting weakly coupled xwticles has been one of the major theoretical achievements in plasma physics when Baleseu(1960) and. Lenarel(1960) could derive an equation surpassing in consistency the simple Fokker-Planck equation (Spitzer1962)., The consistent inclusion of collective screening effects in a kinetic equation for electrically interacting weakly coupled particles has been one of the major theoretical achievements in plasma physics when \citet{Balescu} and \citet{Lenard} could derive an equation surpassing in consistency the simple Fokker-Planck equation \citep{petitSpitzer}.
. It is the aim of this per to derive a similar equation for self-gravitating svstems., It is the aim of this paper to derive a similar equation for self-gravitating systems.
 The task is slightly more difficult because the screening of the electrical interaction at the. usually small. Debye Length allows. in electrically interacting svstems. to take the homogeneous and uniform motion limits.," The task is slightly more difficult because the screening of the electrical interaction at the, usually small, Debye length allows, in electrically interacting systems, to take the homogeneous and uniform motion limits."
 These limits cannot be taken in a self-gravitating system., These limits cannot be taken in a self-gravitating system.
 We overcome this dillicultv by expressing he kinetic equation in action angle space rather than in position momentum space., We overcome this difficulty by expressing the kinetic equation in action angle space rather than in position momentum space.
 This is possible when the Hamiltonian corresponding to the average potential C(r) of the system is integrable., This is possible when the Hamiltonian corresponding to the average potential $U({\mathbf{r}})$ of the system is integrable.
 It is nevertheless uneasy in &eneral to togele from one o the other space. although this is certainly possible for spherically svniumetric potentials. for Hat svstems (which may however x6 unstable) and for special thick disk potentials.," It is nevertheless uneasy in general to toggle from one to the other space, although this is certainly possible for spherically symmetric potentials, for flat systems (which may however be unstable) and for special thick disk potentials."
 Numerical methods could be used to achieve the necessary transformation (Pichon&Cannon1997:MeMillainBinney2008).," Numerical methods could be used to achieve the necessary transformation \citep{PichonCannon, McMillanBinney}."
.. As an illustrative example. we shall give special attention to spherically svmametrie potentials. expanding their kinetic equation into à system which almost entirely avoids any calculation in the »osition-momentunmr space.," As an illustrative example, we shall give special attention to spherically symmetric potentials, expanding their kinetic equation into a system which almost entirely avoids any calculation in the position-momentum space."
 The system's inhomogeneity requires that solutions to the Poisson equation are easily. found. for any inhomogeneous mass distributions., The system's inhomogeneity requires that solutions to the Poisson equation are easily found for any inhomogeneous mass distributions.
" This is achieved. by projecting on a biorthogonal basis of pairs of density-potential '""unctions.", This is achieved by projecting on a biorthogonal basis of pairs of density-potential functions.
 Alany astrophysical systems which have evolved to à quasi-stationary collisionless equilibrium still keep evolving on time scales longer than the dynamical time as a result of gravitational noise induced by their own constituents or by external ones., Many astrophysical systems which have evolved to a quasi-stationary collisionless equilibrium still keep evolving on time scales longer than the dynamical time as a result of gravitational noise induced by their own constituents or by external ones.
 We disregard. external. perturbators. which we define as unbound to the system. although. as did. Weinberg(2001b).. these could be treated. if numerous and frequent enough. as a given. non-evolving. population providing a source of gravitational noise for other populations.," We disregard external perturbators, which we define as unbound to the system, although, as did \citet{Weinberg2001II}, these could be treated, if numerous and frequent enough, as a given, non-evolving, population providing a source of gravitational noise for other populations."
 Loosely bound satellites or remote star populations are regarded as internal to the svstem., Loosely bound satellites or remote star populations are regarded as internal to the system.
 This is possible because our set of kinetic equations allows to simultaneously follow different mass populations., This is possible because our set of kinetic equations allows to simultaneously follow different mass populations.
 Dwarf satellite galaxies could be regarded for example as one such mass population., Dwarf satellite galaxies could be regarded for example as one such mass population.
 Clobular clusters. dwarl galaxies. disk galaxies and their haloes are examples of bound systems still evolving as a result of internal noise caused by particle discreteness.," Globular clusters, dwarf galaxies, disk galaxies and their haloes are examples of bound systems still evolving as a result of internal noise caused by particle discreteness."
 Such svstems are the object of our study., Such systems are the object of our study.
 As in any weakly coupled system. the particles sullering collisions are ciressed by the polarization clouds caused by their own influence on other particles.," As in any weakly coupled system, the particles suffering collisions are dressed by the polarization clouds caused by their own influence on other particles."
 Collisions between dressed. particles have quantitatively different outcomes than collisions between naked ones (Weinberg1998)., Collisions between dressed particles have quantitatively different outcomes than collisions between naked ones \citep{Weinberg98}.
. Fhis may rellect in significant dillerences in calculated effective relaxation times and braking or cillusion cocllicicnts. especially when the svstem. though stable. is not too far from instability 1993).," This may reflect in significant differences in calculated effective relaxation times and braking or diffusion coefficients, especially when the system, though stable, is not too far from instability \citep{Weinberg93}."
. It is therefore useful to account for collective dressing when caleulating such processes as secular thick disk evolution. mass segregation in galaxies or in star clusters. or the damping by dynamical friction of galactic populations on high energy orbits.," It is therefore useful to account for collective dressing when calculating such processes as secular thick disk evolution, mass segregation in galaxies or in star clusters, or the damping by dynamical friction of galactic populations on high energy orbits."
 For simplicity. the kinetic equations to be derived. below assume that the system is stationary on a dynamical time scale.," For simplicity, the kinetic equations to be derived below assume that the system is stationary on a dynamical time scale."
 They thus cannot address questions in which the distribution in angle variable matters. such as the dissolution of freshly accreted satellites. although a simple extension of the theory could.," They thus cannot address questions in which the distribution in angle variable matters, such as the dissolution of freshly accreted satellites, although a simple extension of the theory could."
 Since however our equations describe the coupled evolution of all populations present in the system. they are well suited to study. for example. the simultaneous evolution by dynamical friction ancl dillusion of a stellar population and the population of molecular clouds olf which these stars scatter.," Since however our equations describe the coupled evolution of all populations present in the system, they are well suited to study, for example, the simultaneous evolution by dynamical friction and diffusion of a stellar population and the population of molecular clouds off which these stars scatter."
 The collective response of a self &ravitating svstem to the presence of a perturbing body has been considered. by a number of authors. analytically (Weinberg1989.1995:Murali&Tremaine1998:SahaJog2006) or numerically &Wolll1984:GrecdinOstriker 1999).," The collective response of a self gravitating system to the presence of a perturbing body has been considered by a number of authors, analytically \citep{Weinberg89, Weinberg95,
MuraliTremaine, SahaJog} or numerically \citep{ThielheimWolff, GnedinOstriker}."
. Sometimes. the reaction of this perturbation on the perturbing body. itself is calculated. as did Ixalnajs(1972).. who computed the drag on a large body moving in an homogencous medium. taking the collective response of this medium into account. and Tremaine&Weinberg(1984).. who considered the global. self-consistent. perturbation caused by a satellite or a barred structure in a spherically symmetric svstem and its reaction on the perturbator object bv the effect of dynamical friction.," Sometimes, the reaction of this perturbation on the perturbing body itself is calculated, as did \citet{Kalnajs72}, who computed the drag on a large body moving in an homogeneous medium, taking the collective response of this medium into account, and \citet{WeinbergTremaine}, who considered the global, self-consistent, perturbation caused by a satellite or a barred structure in a spherically symmetric system and its reaction on the perturbator object by the effect of dynamical friction."
 The secular evolution of the svstem in response to such perturbations has been considered by Weinberg(2001a).. who considered &eneral types of perturbations on a galaxy. and by who considered. perturbations caused by the cosmological environment on dark matter haloes.," The secular evolution of the system in response to such perturbations has been considered by \citet{Weinberg2001}, who considered general types of perturbations on a galaxy, and by \citet{PichonAubert} who considered perturbations caused by the cosmological environment on dark matter haloes."
 This evolution is of course in xinciple observable in N-body simulations. which however have their own dillieulties in calculating the long term evolution of such svstems (Binney2004).," This evolution is of course in principle observable in N-body simulations, which however have their own difficulties in calculating the long term evolution of such systems \citep{Binney2004}."
.. A number of authors (Murali1999:Weinberg2001a:Pichon&Aubert2006) have studie he collective perturbations caused in a massive spherical galactie halo by its environment.," A number of authors \citep{Murali, Weinberg2001, PichonAubert} have studied the collective perturbations caused in a massive spherical galactic halo by its environment."
 They could. calculate the response of this system by resorting to a representation of the particles motion in action and angle variables. a method first used by IxalnajsC1977).," They could calculate the response of this system by resorting to a representation of the particle's motion in action and angle variables, a method first used by \citet{Kalnajs77}."
. We follow them on this road., We follow them on this road.
" They. also made good. use of a basis of biorthogonal pairs of density-potentia ""unctions.", They also made good use of a basis of biorthogonal pairs of density-potential functions.
 Weinberg(1993). first. derived a kinetic equation for the collisional relaxation of a self eravitating system along hese lines., \citet{Weinberg93} first derived a kinetic equation for the collisional relaxation of a self gravitating system along these lines.
 His equation accounts for the self-consistent gravitational dressing of the particles. but is otherwise simplified. the geometry supposedly being that of an homogeneously. filled. periodic cube.," His equation accounts for the self-consistent gravitational dressing of the particles, but is otherwise simplified, the geometry supposedly being that of an homogeneously filled periodic cube."
 The inhomogeneous nature of the svstem: shoulc »e described more accurately. still accounting for collective gravitational dressing elfects.," The inhomogeneous nature of the system should be described more accurately, still accounting for collective gravitational dressing effects."
 his is specifically the aim of this, This is specifically the aim of this
the svuthetic xofiles to eet a reasonale match (by eve) between the calculated and measured radii of saturation.,the synthetic profiles to get a reasonable match (by eye) between the calculated and measured radii of saturation.
 As uoted. this is the only parameter tha we adjusted. after fixing a zeropolu for t1e svuthetic magnitudes.," As noted, this is the only parameter that we adjusted, after fixing a zero–point for the synthetic magnitudes."
 The agreement is reasonable for the larger radii. bu is poor for the fainter stars.," The agreement is reasonable for the larger radii, but is poor for the fainter stars."
 Iu xut this may be a consequeuce of tje lamited resolution of the plate scaus: The svuthetic profiles were calculatec Ol a Loare second axis. then resampled Oa 5.25 are second pixel scale to match t16 plate scan resolution.," In part this may be a consequence of the limited resolution of the plate scans: The synthetic profiles were calculated on a 1–arc second axis, then resampled to a 5.25 arc second pixel scale to match the plate scan resolution."
 When the area of saturation of the profile on the! plate is comparable to the «5.25r arc sec resolution of the scan. uudersauiyne starts to reduce the observed saturalon radius. heuce Fig.," When the area of saturation of the profile on the plate is comparable to the $\times$ 5.25 arc sec resolution of the scan, undersampling starts to reduce the observed saturation radius, hence Fig."
 5. shows poorer agrecuent for the fainter stars., \ref{radius_comp} shows poorer agreement for the fainter stars.
(Eq. 4)).,(Eq. \ref{evaporation_rate}) ).
 The species that arrive at a location on the surface can move randomly across the surface by means of tunneling effects and thermal hopping (Eq. 5))., The species that arrive at a location on the surface can move randomly across the surface by means of tunneling effects and thermal hopping (Eq. \ref{hopping_rate}) ).
 Species present on the surface can also receive a UV photon and become dissociated at the rates shown in table 2.., Species present on the surface can also receive a UV photon and become dissociated at the rates shown in table \ref{photorates}.
 Once a chemical species present oi the dust moves. or adsorbs a photon. or meets another species. the next event that occurs to this species is determined as well as its time of occurrence.," Once a chemical species present on the dust moves, or adsorbs a photon, or meets another species, the next event that occurs to this species is determined as well as its time of occurrence."
 Therefore. the events that concert every species on the dust are ordered by time of occurrence. and for each event that occurs. a next event for the concerned species is determined.," Therefore, the events that concern every species on the dust are ordered by time of occurrence, and for each event that occurs, a next event for the concerned species is determined."
 In figure. 1.. we follow the fraction of oxygen that is released in the gas under the form of OH anc H:O for environments |. 2. 4 and 5 (see Table 2. for adoptec parameters).," In figure \ref{MC}, we follow the fraction of oxygen that is released in the gas under the form of OH and $_2$ O for environments 1, 2, 4 and 5 (see Table \ref{params_env} for adopted parameters)."
 The parameters of these environments represent typical ambient low and high densities and radiation fields iu the ISM of active galaxy centers. and are similar to those i1(2005).," The parameters of these environments represent typical ambient low and high densities and radiation fields in the ISM of active galaxy centers, and are similar to those in."
. The solid and dashed lines show our results for Ty. = 20 and 30 K respectively., The solid and dashed lines show our results for $T_{\rm{dust}}$ = 20 and 30 K respectively.
 For each environments. 30-60 % (30-40%) of the oxygen arriving on the dust goes back in the gas phase under the form of H»O (OH).," For each environments, 50-60 $\%$ $\%$ ) of the oxygen arriving on the dust goes back in the gas phase under the form of $_2$ O (OH)."
 At higher dust temperatures (~35—40 K). on the other hand. species coming on the dust evaporate very fast.," At higher dust temperatures $\sim 35-40$ K), on the other hand, species coming on the dust evaporate very fast."
 Therefore. the efficiencies decrease dramatically since the grain surfaces are only sporadically covered by species.," Therefore, the efficiencies decrease dramatically since the grain surfaces are only sporadically covered by species."
 The rate equations method can be applied when there is a minimum coverage of | atom on the dust., The rate equations method can be applied when there is a minimum coverage of 1 atom on the dust.
 This condition is satisfied if the grains are not too warm. that species can encounter each other before evaporating.," This condition is satisfied if the grains are not too warm, that species can encounter each other before evaporating."
 By comparing our rate equations results with Monte Carlo simulations. we will define a range of dust temperatures for which rate equations are valid.," By comparing our rate equations results with Monte Carlo simulations, we will define a range of dust temperatures for which rate equations are valid."
 The rate equation method couples the different equations for the different species on the dust., The rate equation method couples the different equations for the different species on the dust.
 The processes described in the previous subsection are dependent on the dust and gas temperature. UV field. and exo-thermicity of the reaction occurring on the dust.," The processes described in the previous subsection are dependent on the dust and gas temperature, UV field, and exo-thermicity of the reaction occurring on the dust."
" The equation for species X can thus be written às follows: where «,,CX) and «,,(Y) are the mobilities of the species X and Y (see Eq. 5)).", The equation for species $X$ can thus be written as follows: where $\alpha_{pp}(X)$ and $\alpha_{pp}(Y)$ are the mobilities of the species $X$ and $Y$ (see Eq. \ref{hopping_rate}) ).
 The different terms of this equation present (D) accretion of species X from the gas phase. (2) the formation of the new species that involve species X. (3) creation of the species X by the encountering of spectes Y and Z that stay on the surface. (4) evaporation of species X. (5) photo-dissociation of X on the dust. and (6) the creation of X by photo-dissociation of species Y.," The different terms of this equation present (1) accretion of species $X$ from the gas phase, (2) the formation of the new species that involve species $X$, (3) creation of the species $X$ by the encountering of species $Y$ and $Z$ that stay on the surface, (4) evaporation of species $X$, (5) photo-dissociation of $X$ on the dust, and (6) the creation of $X$ by photo-dissociation of species $Y$."
 The formation efficiency is defined as the fraction of oxygen 10) that is released in the gas phase under the form of another species.," The formation efficiency is defined as the fraction of oxygen $f({\rm
 O})$ that is released in the gas phase under the form of another species."
 We studied this in a few different environments. chosen such that they resemble the conditions as found in the XDR models of Meterink Spaans (2005).," We studied this in a few different environments, chosen such that they resemble the conditions as found in the XDR models of Meijerink Spaans (2005)."
 The gas temperature is fixed at à constant temperature T=100 K. and dust temperatures ranging from 7410—60 K. Four different radiation fields were considered at two fixed densities. 2=10? and 10? em. in both the atomie and molecular regime.," The gas temperature is fixed at a constant temperature $T=100$ K, and dust temperatures ranging from $T_{\rm{dust}}=10-60$ K. Four different radiation fields were considered at two fixed densities, $n=10^3$ and $10^{5.5}$ $^{-3}$, in both the atomic and molecular regime."
 The parameters used are summarized in Table 2.., The parameters used are summarized in Table \ref{params_env}.
 In Figs., In Figs.
 2. and 3.. we show the obtained efficiencies for OH and H:O. The figures show that between dust temperatures Ta~15—35 K. the efficiencies for OH and HO are roughly constant.," \ref{low_dens_eff} and \ref{high_dens_eff}, we show the obtained efficiencies for OH and $_2$ O. The figures show that between dust temperatures $T_{\rm dust}\sim 15-35$ K, the efficiencies for OH and $_2$ O are roughly constant."
 In this regime most of the oxygen goes to H;O. For dust temperatures Zi.235—45 K. the total efficiency for forming molecules is rapidly decreasing. and most of the atomic oxygen will not react to another species. before it evaporates from the grain.," In this regime most of the oxygen goes to $_2$ O. For dust temperatures $T_{\rm{dust}} \gtrsim 35-45$ K, the total efficiency for forming molecules is rapidly decreasing, and most of the atomic oxygen will not react to another species, before it evaporates from the grain."
 Therefore. the dust surface chemistry can only be effective. when the dust temperatures are not too high.," Therefore, the dust surface chemistry can only be effective, when the dust temperatures are not too high."
 Recall that the way to form OH and H?O is through O-H—OHand OH-H—H:O.," Recall that the way to form OH and $_2$ O is through $\rm O + H \rightarrow
OH$and $\rm OH + H \rightarrow H_2O$."
 For slightly higher dust temperatures. the chance of an OH molecule to evaporate from the grain becomes higher than that for meeting another H atom.," For slightly higher dust temperatures, the chance of an OH molecule to evaporate from the grain becomes higher than that for meeting another H atom."
 Therefore at temperatures Ty.z35—40 K. most of the oxygen atoms are incorporated into OH. and then εοη Is highest.," Therefore at temperatures $T_{\rm{dust}} \gtrsim 35 - 40$ K, most of the oxygen atoms are incorporated into OH, and then $\epsilon_{\rm OH}$ is highest."
 Increasing the radiation field. has a similar effect as increasing the dust temperature. Le. it results in a lower €4.0/€oH ratio. which is caused by photo-dissociation of H:O on the surface.," Increasing the radiation field, has a similar effect as increasing the dust temperature, i.e., it results in a lower $\epsilon_{\rm H_2O}/
\epsilon_{\rm OH}$ ratio, which is caused by photo-dissociation of $_2$ O on the surface."
 The strongest effect is seen in the environments 3 and 6. where the dominant hydrogen fraction is contained in H».," The strongest effect is seen in the environments 3 and 6, where the dominant hydrogen fraction is contained in $_2$."
 In this environment. the hydrogen accretion rates are lowest. and thus also have the lowest atomic surface density on grains.," In this environment, the hydrogen accretion rates are lowest, and thus also have the lowest atomic surface density on grains."
 When most of the gas is in atomic form. the hydrogen aceretion is such that at high dust temperatures most of the oxygen Is converted into OH.," When most of the gas is in atomic form, the hydrogen accretion is such that at high dust temperatures most of the oxygen is converted into OH."
 However. when the gas become more and more molecular. the atomic hydrogen accretion rate onto the dust decreases. and a larger fraction of atomic oxygen ends up in Os.," However, when the gas become more and more molecular, the atomic hydrogen accretion rate onto the dust decreases, and a larger fraction of atomic oxygen ends up in $_2$ ."
" In environment 3. this is even the most likely product at temperature Z4,>35 K."," In environment 3, this is even the most likely product at temperature $T_{\rm{dust}} > 35$ K."
The similar process ¢|έI5 du uuparticle plivsies has been ]5usued in Ref.|0].. aud we can straightforwardly obtain the spi-averaged amplitude squared for €clos[a| as where E is from the SAL ID is from the teusor uuparticle exchanec. arc the interference between the SM and wnparticle amplitucles is provided as IV. which are all cousisteut with Ref.[0]..,"The similar process $e^-+e^+\to \gamma+\gamma$ in unparticle physics has been pursued in \cite{2}, and we can straightforwardly obtain the spin-averaged amplitude squared for $e^-+e^+\to
\gamma+\gamma$ as where ${\bf I}$ is from the SM, ${\bf II}$ is from the tensor unparticle exchange, and the interference between the SM and unparticle amplitudes is provided as ${\bf IV}$ , which are all consistent with \cite{2}."
 The term III is the contribution from the scalar uuparticle exchange. which is absent in Ref.[O]..," The term ${\bf III}$ is the contribution from the scalar unparticle exchange, which is absent in \cite{2}."
 We investigate further the process in the ceuter-ofmementium svstem of the two initial ploous. and the Manudelstaui variables can be written as [f|2s(1cos@)/2 aud |u|=s(1|cos 0)/2. iux Mis the scattering angle.," We investigate further the process in the center-of-momentum system of the two initial photons, and the Mandelstam variables can be written as $|t|= s(1-\cos\theta)/2$ and $|u|= s(1+\cos\theta)/2$ , and $\theta$ is the scattering angle."
 Plotting the terms E. EE. EE. and IV. versus 0 in Fig... we find out that coiribution TIL from scalar uuparticle exchange to the process ο|60!»5 is consequential.," Plotting the terms ${\bf I}$, ${\bf II}$, ${\bf III}$, and ${\bf
IV}$ versus $\theta$ in Fig.1, we find out that contribution ${\bf
III}$ from scalar unparticle exchange to the process $e^-+e^+\to
\gamma+\gamma$ is consequential."
" The terms E. IL. IHE aud IV in M? ofc|¢>> versus ϐ in the case (4=d,—1.1. Ay=AxLO aud Aj,= LOTeV at 5= 0.5TeV. The diphoton can ouly clastically scatters via the loop-level iu the SM and thus is highly suppressed. however. it can take place via scalar aud tensor uuparticle exchanges in al ]s- ἐν snd e-chanuels at the trec-level iu the uupiirticle plivsies."," The terms ${\bf I}$, ${\bf II}$, ${\bf III}$ and ${\bf
IV}$ in $\overline{|{\cal M}|^2}$ of $e^-+e^+\to \gamma+\gamma$ versus $\theta$ in the case $d_{\rm t}=d_{\rm s}=1.1$, $\lambda_0=\lambda_2=1.0$ and $\Lambda_{\cal U}=1.0$ TeV at $\sqrt{s}=0.5$ TeV. The diphoton can only elastically scatters via the loop-level in the SM and thus is highly suppressed, however, it can take place via scalar and tensor unparticle exchanges in all$s$ -, $t$ -, and $u$ -channels at the tree-level in the unparticle physics."
 The scattering amplitude through scalar uuparticleexchauege, The scattering amplitude through scalar unparticleexchange
cluster would be tidallv disrupted.,cluster would be tidally disrupted.
 This induces only a slight. (~5.4%) reduction in the enerev requirements of the collision., This induces only a slight $\sim 5.4\%$ ) reduction in the energy requirements of the collision.
 Within the lormalism of 2.. the effect is equivalent to decreasing (1—€) bv 5.4%.. or increasing e [rom 0.87 to 0.877.," Within the formalism of \ref{sec:anal}, the effect is equivalent to decreasing $(1-e)$ by , or increasing $e$ from 0.87 to 0.877."
" This reduces the overall mass scale for mg=m3/mz, by a factor (1—0.054)*~0.35 (see eq. (6]]).", This reduces the overall mass scale for $m_f=m_2^3/m_\tot^2$ by a factor $(1-0.054)^3\sim 0.85$ (see eq. \ref{eqn:meq}] ]).
 Given the high nass scale of this equation and of Figure 2.. (his reduction is welcome. but modest.," Given the high mass scale of this equation and of Figure \ref{fig:massdis}, this reduction is welcome, but modest."
 llowever. (his finite semi-major axis has another important effect: the binary re-encounters Ser A* on timescales of ~107 ves.," However, this finite semi-major axis has another important effect: the binary re-encounters Sgr A* on timescales of $\sim 10^3$ yrs."
 When it does so. it preserves all the binary parameters ib acquired in the last collision. save the internal orbital phase relative to (he phase of the orbit around Ser À*. which is effectively randomized.," When it does so, it preserves all the binary parameters it acquired in the last collision, save the internal orbital phase relative to the phase of the orbit around Sgr A*, which is effectively randomized."
 Hence. thousands of such encounters nav (ake place. which increase the chance that (he binary will be disrupted.," Hence, thousands of such encounters may take place, which increase the chance that the binary will be disrupted."
 Partly. the infernal binary orbit changes each time.," Partly, the internal binary orbit changes each time."
 ILowever. even if the orbit is left unchanged. the mere randomizing of the internal orbital phase permits a new opportunity. lor a disruptive encounter.," However, even if the orbit is left unchanged, the mere randomizing of the internal orbital phase permits a new opportunity for a disruptive encounter."
 To explore this effect. we simulate repeat encounters by binaries in orbit around Ser A*. continue the simulation until the binary is disrupted or until it completes 1000 orbits.," To explore this effect, we simulate repeat encounters by binaries in orbit around Sgr A*, continuing the simulation until the binary is disrupted or until it completes 1000 orbits."
 In order (ο permit efficient computation. we place the binaries in orbits with semi-major axis a=3000AU.," In order to permit efficient computation, we place the binaries in orbits with semi-major axis $a=3000\,\au$."
 Then. in order to be able to compare directly with the previous simulations we consider as a success. Collisions with the same energy change. that is. with final orbits having eccentricitv e=0.847.," Then, in order to be able to compare directly with the previous simulations we consider as a success, collisions with the same energy change, that is, with final orbits having eccentricity $e=0.847$."
 Finally. to minimize the role of Poisson fluctuations. we focus on (he relative number successes Irom (he entire (up to 1000 orbit)) simulation with the number coming [rom the first encounter.," Finally, to minimize the role of Poisson fluctuations, we focus on the relative number successes from the entire (up to 1000 orbit) simulation with the number coming from the first encounter."
 For mq=15M. and m»=15.M.. the 1000-0orbit simulations doubles the number of successes. while for i=100M. the improvement rises io about a factor 2.5.," For $m_1=15\,M_\odot$ and $m_2=75\,M_\odot$, the 1000-orbit simulations doubles the number of successes, while for $m_2=100\,M_\odot$ the improvement rises to about a factor 2.5."
 About of the post-first-orbit successes occur within the first 100 orbits. and none past 200 orbits. indicating that our 1000-0rbit simulations were quite adequate.," About of the post-first-orbit successes occur within the first 100 orbits, and none past 200 orbits, indicating that our 1000-orbit simulations were quite adequate."
 The great majority of additional successes have initial semi-major axes LAUαμ«2.5AU.," The great majority of additional successes have initial semi-major axes $1\,\au<a_\bin<2.5\,\au$."
 Of those that do not. all but one are in retrograde orbits.," Of those that do not, all but one are in retrograde orbits."
 That is. the additional successes. arise because binaries with unfavorable parameters are gradually perturbed into a more favorable reeime.," That is, the additional successes arise because binaries with unfavorable parameters are gradually perturbed into a more favorable regime."
" Given the enhanced probability for binaries with ej; in the lowest range simulated. we conduct an additional simulation with io=LOOM. and 0.1AU<ay,1AU."," Given the enhanced probability for binaries with $a_\bin$ in the lowest range simulated, we conduct an additional simulation with $m_2=100\,M_\odot$ and $0.1\,\au<a_\bin<1\,\au$."
 Lowever. we find. very [ew successful collisions from this range.," However, we find very few successful collisions from this range."
 In brief. allowing for multiple encounters in repeated orbits roughly doubles thenumber ol successful encounters. Le.. those leaving the secondary in an orbit similar to SO-2.," In brief, allowing for multiple encounters in repeated orbits roughly doubles thenumber of successful encounters, i.e., those leaving the secondary in an orbit similar to SO-2."
 , 
the signal-to-noise ratio in the line data cubes is larger than δσ (corresponding to 0.5 Jy beam”! ).,the signal-to-noise ratio in the line data cubes is larger than $5\sigma$ (corresponding to 0.5 Jy $^{-1}$ ).
 As à first step to investigate the association. of given transitions with the molecular outflows in the region. with potential accretion disks or with the envelope around protostars. we derived the integrated intensity of different lines to study their spatial distribution.," As a first step to investigate the association of given transitions with the molecular outflows in the region, with potential accretion disks or with the envelope around protostars, we derived the integrated intensity of different lines to study their spatial distribution."
 In Figs. 4.. 5.. 6..," In Figs. \ref{ch3ohmom0}, \ref{ocsmom0}, \ref{ch3cnmom0},"
 and 7 we present the integrated intensity distribution of selected, and \ref{hncomom0} we present the integrated intensity distribution of selected
spectrum with an absorbed multi-color disk (MCD) blackbody model. but the fit was unacceptable(y7/DOF 238/28).,"spectrum with an absorbed multi-color disk (MCD) blackbody model, but the fit was unacceptable$\chi^2$ /DOF =38/28)."
 We then tried an absorbed power-law model and obtained an acceptable fit (Z7/DOF =27.8/28) with D=1.60.3 confidence limit) and galactic absorption (see Fig., We then tried an absorbed power-law model and obtained an acceptable fit $\chi^2$ /DOF =27.8/28) with $\Gamma=1.6 \pm 0.3$ confidence limit) and galactic absorption (see Fig.
 7)., 7).
 Adding a disk component to the power-law model did not tmprove the fit further., Adding a disk component to the power-law model did not improve the fit further.
 The unabsorbed 0.3-10.0 keV flux is 2.0«107 eres em-7 s7!. which corresponds to L(0.3-10.0 keV)2 2.7«IQ? ergs s7!.," The unabsorbed 0.3-10.0 keV flux is $2.0 \times 10^{-13}$ ergs $^{-2}$ $^{-1}$ , which corresponds to L(0.3-10.0 keV)= $2.7 \times 10^{39}$ ergs $^{-1}$."
" The probability of detecting a background source with a flux equal to. or greater than this value within the central 7"" of the galaxy is less than 1077."," The probability of detecting a background source with a flux equal to, or greater than this value within the central $7^{\prime\prime}$ of the galaxy is less than $10^{-4}$."
 The intensity. of this source varies smoothly by a factor of two during the course of the 30 ksee observation (see Fig., The intensity of this source varies smoothly by a factor of two during the course of the 30 ksec observation (see Fig.
 8)., 8).
 The peak luminosity of the source is 3.5«10? ergs s7! which places it among the ULX class of objects and corresponds to the Eddington luminosity of a 30M. black hole., The peak luminosity of the source is $3.5 \times 10^{39}$ ergs $^{-1}$ which places it among the ULX class of objects and corresponds to the Eddington luminosity of a $30 \Mo$ black hole.
 We also generated a hardness ratio light curve. but did not find significant variations during the observation.," We also generated a hardness ratio light curve, but did not find significant variations during the observation."
 To check for long-term variability in the ULX. we examined the 24 ksee ROSAT HRI observation of NGC 3379 taken on Aug. 2. 1996 and found that the X-ray emission was centered on the ULX. not the center of the galaxy.," To check for long-term variability in the ULX, we examined the 24 ksec ROSAT HRI observation of NGC 3379 taken on Aug. 2, 1996 and found that the X-ray emission was centered on the ULX, not the center of the galaxy."
 The net HRI count rate within a 10” radius aperture (corresponding to the encircled energy radius) centered on the J2000 coordinates of the ULX as determined from the S3 image i5 3.3.107 et s7!., The net HRI count rate within a $10^{\prime\prime}$ radius aperture (corresponding to the encircled energy radius) centered on the J2000 coordinates of the ULX as determined from the S3 image is $3.3 \times 10^{-3}$ ct $^{-1}$.
 Assuming an absorbed power-law model with the best-fit parameters listed above gives an unabsorbed 0.3-10.0 keV flux of 3.0«107 eres em s7!. which is higher than the ACIS measured flux of the ULX in 2001.," Assuming an absorbed power-law model with the best-fit parameters listed above gives an unabsorbed 0.3-10.0 keV flux of $3.0 \times 10^{-13}$ ergs $^{-2}$ $^{-1}$, which is higher than the ACIS measured flux of the ULX in 2001."
 The HRI flux measurement is undoubtedly affected by contamination from other point sources., The HRI flux measurement is undoubtedly affected by contamination from other point sources.
 Based on the $3 image. the ULX contributes of the flux within the same apperature used to compute the HRI flux.," Based on the S3 image, the ULX contributes of the flux within the same apperature used to compute the HRI flux."
 Thus. there is strong evidence that the ULX was at a comparable luminosity 5 years prior to the Chandra observation.," Thus, there is strong evidence that the ULX was at a comparable luminosity 5 years prior to the Chandra observation."
 There are insufficient counts in the HRI data to measure short term variability., There are insufficient counts in the HRI data to measure short term variability.
" After excising the emission from all detected point sources. we extracted a spectrum of the diffuse emission within the central 15"" (770 pe)."," After excising the emission from all detected point sources, we extracted a spectrum of the diffuse emission within the central $15^{\prime\prime}$ (770 pc)."
 A background spectrum was taken from a circular aperture 2.5’ from the center of the galaxy on the S3 chip., A background spectrum was taken from a circular aperture $2.5^{\prime}$ from the center of the galaxy on the S3 chip.
 The background subtracted spectrum was then fitted to an absorbed thermal (MEKAL) plus power-law model in the 5.0 keV band., The background subtracted spectrum was then fitted to an absorbed thermal (MEKAL) plus power-law model in the 0.3-5.0 keV band.
 The best-fit power-law component has [=1.7+0.3 and an unabsorbed 0.3-10.0 keV flux of 107 ergs em s7!. corresponding to L(0.3-10.0 keV)= 7.4«10°% eres s7!., The best-fit power-law component has $\Gamma = 1.7 \pm 0.3$ and an unabsorbed 0.3-10.0 keV flux of $5.5 \times 10^{-14}$ ergs $^{-2}$ $^{-1}$ corresponding to L(0.3-10.0 keV)= $7.4 \times 10^{38}$ ergs $^{-1}$.
 The best-fit thermal component has kT=0.6 (0.4-1.0 keV) and an unabsorbed 0.3-10.0 keV flux of 107 ergs em s7!. corresponding to L(0.3-10.0 keV)= 8.9&1077 eres s7!.," The best-fit thermal component has kT=0.6 (0.4-1.0 keV) and an unabsorbed 0.3-10.0 keV flux of $6.6 \times 10^{-15}$ ergs $^{-2}$ $^{-1}$, corresponding to L(0.3-10.0 keV)= $8.9 \times 10^{37}$ ergs $^{-1}$."
 ΑΙ errors are given at the confidence limit., All errors are given at the confidence limit.
 This analysis shows that of the remaining diffuse emission most likely arises from point sources with fluxes below the detection limit in the Chandra observation., This analysis shows that of the remaining diffuse emission most likely arises from point sources with fluxes below the detection limit in the Chandra observation.
 Assuming the luminosity function derived above is valid at lower luminosities. then the luminosity of the unresolved power-law component can be generated by approximately 40 sources below 3.0«10°7 ergs s7!.," Assuming the luminosity function derived above is valid at lower luminosities, then the luminosity of the unresolved power-law component can be generated by approximately 40 sources below $3.0\times 10^{37}$ ergs $^{-1}$."
" Using the best-fit emission measure in the thermal model and assuming a uniform gas distribution in the central 770 pe gives an number density of ης=0.01 cm. a gas mass of 5electron«10°M.. and a cooling time (7,=SKTMugs 24unyL) of 8«10° yr."," Using the best-fit emission measure in the thermal model and assuming a uniform gas distribution in the central 770 pc gives an electron number density of $n_e = 0.01$ $^{-3}$, a gas mass of $5 \times 10^5 \Mo$, and a cooling time $t_c = 5kTM_{gas}/2\mu m_p L$ ) of $8 \times 10^8$ yr."
 If the gas is able to cool unimpeded. this gives a mass cooling rate within the central 770 pe of M—6«107.yr7!.," If the gas is able to cool unimpeded, this gives a mass cooling rate within the central 770 pc of $\dot {\rm M} = 6 \times 10^{-4} \Mo yr^{-1}$."
 The estimated central gas density in NGC 3379 is a factor of 10 less than that found in X-ray luminous ellipticals whose emission is dominated by hot gas (e.g.. Lowenstein et al.," The estimated central gas density in NGC 3379 is a factor of 10 less than that found in X-ray luminous ellipticals whose emission is dominated by hot gas (e.g., Lowenstein et al."
 2001)., 2001).
 The optical luminosity density in. NGC3379 can. be determined. from the surface brightness profile and Abel's formula (Binney Tremaine 1987)., The optical luminosity density in NGC3379 can be determined from the surface brightness profile and Abel's formula (Binney Tremaine 1987).
 Using the de Vaucouleurs profile for NGC 3379 found by Capaccioli et al. (, Using the de Vaucouleurs profile for NGC 3379 found by Capaccioli et al. (
1990). gives Lp=2.3410°L.p Within the central 770 pe.,"1990), gives $L_B = 2.3 \times 10^{9} {\Lo}_B$ within the central 770 pc."
" Using the stellar mass loss rate from Faber Gallagher (1976) for elliptical galaxies of 0.15M./(10!°L.vr). gives a stellar mass loss rate of M,=0.035M. yr7!."," Using the stellar mass loss rate from Faber Gallagher (1976) for elliptical galaxies of $0.15\Mo/(10^{10} \Lo yr)$, gives a stellar mass loss rate of $\dot {\rm M}_s = 0.035 \Mo$ $^{-1}$ ."
 Thus. the presently observed gas massin NGC 3379 can be produced by stars in only 107 years. which is significantly less than the cooling time.and only 5 times the sound crossing time m this region.," Thus, the presently observed gas massin NGC 3379 can be produced by stars in only $10^7$ years, which is significantly less than the cooling time,and only 5 times the sound crossing time in this region."
 The continuous injection of gas shed by evolving stars with a temperature equal to the virial, The continuous injection of gas shed by evolving stars with a temperature equal to the virial
Simulations of clumpy flows have also indicated. that a Salpeter type clump mass spectrum can be reproduced. (?)..,Simulations of clumpy flows have also indicated that a Salpeter type clump mass spectrum can be reproduced \citep{Clark2006}.
 Spiral shocks have also been proposed. to explain the cdvnamies of molecular clouds (22)..," Spiral shocks have also been proposed to explain the dynamics of molecular clouds \citep{Bonnell2006,Zhang2001}."
 7. model giant molecular cloud formation as gas passes through a clumpwv spiral shock., \citet{Bonnell2006} model giant molecular cloud formation as gas passes through a clumpy spiral shock.
 The dynamics of the molecular clouds are determined on all scales simultaneously as the clouds form and the induced velocity dispersion. size. scale relation is consistent. with observations., The dynamics of the molecular clouds are determined on all scales simultaneously as the clouds form and the induced velocity dispersion size scale relation is consistent with observations.
 This can account for the observed. velocity dispersions that are found even in regions devoid of massive stars., This can account for the observed velocity dispersions that are found even in regions devoid of massive stars.
 Furthermore. there is no need for à continuous driving mechanism as the time for the decay of these velocities is proposed to be of similar magnitude to the cloud lifetime.," Furthermore, there is no need for a continuous driving mechanism as the time for the decay of these velocities is proposed to be of similar magnitude to the cloud lifetime."
 In this paper we investigate the velocity size relation in shock tests with uniform. clumpy ancl fractal distributions of isothermal eas.," In this paper we investigate the velocity size relation in shock tests with uniform, clumpy and fractal distributions of isothermal gas."
 The clumpy anc fractal distributions are chosen to reflect the highly. structured nature of the ISAL (????)..," The clumpy and fractal distributions are chosen to reflect the highly structured nature of the ISM \citep{Cox2005,Elmegreen2004,Dickey1990,
Perault1985}."
 We concentrate on modelling the passage of gas through a spiral arm by using a linear sinusoidal potential. although these results would apply generally to shocks between colliding Hows.," We concentrate on modelling the passage of gas through a spiral arm by using a linear sinusoidal potential, although these results would apply generally to shocks between colliding flows."
 We show that random velocities induced in non-uniform: shocks display a velocity size relation similar to those observed. and. provide simple analytical analysis alongside the results of our simulations., We show that random velocities induced in non-uniform shocks display a velocity size relation similar to those observed and provide simple analytical analysis alongside the results of our simulations.
 Thus the ‘turbulence’ in our results. describes. randonir motions of the gas and does not correspond to any theories of classical turbulence., Thus the 'turbulence' in our results describes random motions of the gas and does not correspond to any theories of classical turbulence.
 We use the 3D. smoothed particle hyerodynamics (SPILL) code based on the version by Benz (?).., We use the 3D smoothed particle hydrodynamics (SPH) code based on the version by Benz \citep{Benz1990}.
 The smoothing length: is allowed. to vary with space and time. with the constraint that the typical number of neighbours for each particle is kept near Ανν50.," The smoothing length is allowed to vary with space and time, with the constraint that the typical number of neighbours for each particle is kept near $N_{neigh} \thicksim50$."
 Artificial viscosity. is included: with either the standard: parameters a=1 and =? (7?) ora—2 and 3=4.," Artificial viscosity is included with either the standard parameters $\alpha=1$ and $\beta=2$ \citep{Monaghan1985,Monaghan1992} or $\alpha=2$ and $\beta=4$."
 We investigate uniform. clumpy and multi-scale shocks. starting with 3D shock tube tests first before considering gas subject to an external potential.," We investigate uniform, clumpy and multi-scale shocks, starting with 3D shock tube tests first before considering gas subject to an external potential."
 In. all calculations. the gas is isothermal. non sclleravitating and there are no magnetic fields.," In all calculations, the gas is isothermal, non self-gravitating and there are no magnetic fields."
 These calculations are dimensionless. and are characterized by the Mach number of the shock anc the initial clensity distribution of the gas.," These calculations are dimensionless, and are characterized by the Mach number of the shock and the initial density distribution of the gas."
 In all calculations. the particles are allocated velocities in the vw direction only and. except for the oblique shocks (Section 3.2). the gas shocks in the yz plane.," In all calculations, the particles are allocated velocities in the $x$ direction only and, except for the oblique shocks (Section 3.2), the gas shocks in the $yz$ plane."
 The parameters used in the dilferent. simulations are shown in ‘Table 1., The parameters used in the different simulations are shown in Table 1.
 We investigated. 4 distributions of sas - uniform. homogeneous spherical clumps in. pressure equilibrium. spherical clumps of different. racii/density and fractal cüstributions.," We investigated 4 distributions of gas - uniform, homogeneous spherical clumps in pressure equilibrium, spherical clumps of different radii/density and fractal distributions."
 The filling factors for cach clistribution are calculated by overlaving a 3D grid on each clistribution and determining the porosity., The filling factors for each distribution are calculated by overlaying a 3D grid on each distribution and determining the porosity.
 The filling factor is given bv P5NeaiNeos Where Nya ds the number of cells containing at least 1 particle anc Ανω the total number of cells., The filling factor is given by $F=N_{full}/N_{total}$ where $N_{full}$ is the number of cells containing at least 1 particle and $N_{total}$ the total number of cells.
 We take a 32° grid. so the mesh resolution is equivalent to a couple of smoothing lengths for the uniform case.," We take a $^3$ grid, so the mesh resolution is equivalent to a couple of smoothing lengths for the uniform case."
 Εις ensures that each cell contains at [east 1 particle or the uniform distribution so the filling factor isLOO%., This ensures that each cell contains at least 1 particle for the uniform distribution so the filling factor is.
. The maximum initial scale length corresponds to the range ofr values for the particles in the initial distribution., The maximum initial scale length corresponds to the range of $x$ values for the particles in the initial distribution.
 We [ist perform 3D. shock tube tests with initial distributions of uniform| and clumpy. gas to mocel colliding Hows., We first perform 3D shock tube tests with initial distributions of uniform and clumpy gas to model colliding flows.
 For both distributions. particles are. placed within a cuboid of dimensions 2<r<2. 1«y< and l«z<1.," For both distributions, particles are placed within a cuboid of dimensions $-2<x<2$, $-1<y<1$ and $-1<z<1$."
 Fo produce a clumpy shock we distribute the particles in uniform density spheres within these length scales., To produce a clumpy shock we distribute the particles in uniform density spheres within these length scales.
 “Phe clumps are a constant temperature. confined by either an external pressure field or a hotter cdilluse phase.," The clumps are a constant temperature, confined by either an external pressure field or a hotter diffuse phase."
 The clumps are initially allowed to settle into equilibrium xdore the simulation is carried out., The clumps are initially allowed to settle into equilibrium before the simulation is carried out.
 To produce clumps of different raclii. the external pressure (or pressure of the diffuse phase) can be increased or decreased.," To produce clumps of different radii, the external pressure (or pressure of the diffuse phase) can be increased or decreased."
 We then assign xwticles a velocity of ry=10 ὃν for με0 and ey=0 c. or eZ0., We then assign particles a velocity of $v_0=10$ $c_s$ for $x\le0$ and $v_0=-10$ $c_s$ for $x>0$.
 For the uniform distribution. this produces two approximately Mach. LO shocks. ancl multiple shocks of up ο Mach 20 for the clumpy distributions.," For the uniform distribution, this produces two approximately Mach 10 shocks, and multiple shocks of up to Mach 20 for the clumpy distributions."
 Phese caleulations were also repeated using twice the initial velocities and all ests use 2«.107 particles., These calculations were also repeated using twice the initial velocities and all tests use $2\times10^5$ particles.
 We mainly consider shocks in gas subject to an external potential., We mainly consider shocks in gas subject to an external potential.
 In this case the eas self shocks. similar to gas experiencing a stellar potential in a spiral galaxy.," In this case the gas self shocks, similar to gas experiencing a stellar potential in a spiral galaxy."
" A velocity dispersion relation similar to the observed mx£"" [aw has been shown to develop as gas passes through a elumpy spiral shock (?)..", A velocity dispersion relation similar to the observed $\sigma \propto r^{0.5}$ law has been shown to develop as gas passes through a clumpy spiral shock \citep{Bonnell2006}.
 Here we examine a simplified setup. where we can investigate the effect of the shock dynamics on the initial gas distribution.," Here we examine a simplified setup, where we can investigate the effect of the shock dynamics on the initial gas distribution."
 Instead of the spiral potential we use a 11) sinusoidal potential of the form where & is the wavenumber and 2 a length parameter to, Instead of the spiral potential we use a 1D sinusoidal potential of the form where $k$ is the wavenumber and $B$ a length parameter to
1998).,1998).
 Iu particular the deforiiatious produced iu a dark halo during an encounter could be one of the culprits for distorted morphologies like those observed iu lopsided aud warped ealactic disks (see e.g. Richter Sancisi 1991. Zaritsky Rix 1997. Ruduick Rix 1998. Swaters ct al.," In particular the deformations produced in a dark halo during an encounter could be one of the culprits for distorted morphologies like those observed in lopsided and warped galactic disks (see e.g. Richter Sancisi 1994, Zaritsky Rix 1997, Rudnick Rix 1998, Swaters et al."
 1995. Ilavnes et al.," 1998, Haynes et al."
 1998. IKorurcich. Tavues Lovelace 1998).," 1998, Kornreich, Haynes Lovelace 1998)."
 As recently shown by Weimbere (1995. 1998b). this could be the case for the Mill Way in which the deformations induced iu the dark halo by the Magellanic Clouds could be responsible for the observed warp iu the Galactic disk.," As recently shown by Weinberg (1995, 1998b), this could be the case for the Milky Way in which the deformations induced in the dark halo by the Magellanic Clouds could be responsible for the observed warp in the Galactic disk."
 Iu a recent investigation. Reshetuikov Combes (1998. 1999) have found that about of a sample of SLO salaxies have warped planes aud their analysis revealed that warped galaxies teud to be located im denser environments supporting a scenario m which iuteractious would play au miportaut role iu the formation of warps.," In a recent investigation, Reshetnikov Combes (1998, 1999) have found that about of a sample of 540 galaxies have warped planes and their analysis revealed that warped galaxies tend to be located in denser environments supporting a scenario in which interactions would play an important role in the formation of warps."
 Distortious of spheroidal galaxies are directly observable., Distortions of spheroidal galaxies are directly observable.
 Our results show that iu low velocity dispersion and dense cuvironmeuts. where low-velocity close eucounters are more likely to occur. even low-1ass perturbers can lead to significant asvnunetries iu the primary galaxy.," Our results show that in low velocity dispersion and dense environments, where low-velocity close encounters are more likely to occur, even low-mass perturbers can lead to significant asymmetries in the primary galaxy."
 Recent morphological studies of both local aud very distaut galaxies have vielded comparisous of the frequencies of syvste1s with distorted morphologics in different cuviroumenuts auc at different distances (see e.g. Zopt Whitmore 1993. Mendes de Oliveira Ticksou 1991. Abraham et al.," Recent morphological studies of both local and very distant galaxies have yielded comparisons of the frequencies of systems with distorted morphologies in different environments and at different distances (see e.g. Zepf Whitmore 1993, Mendes de Oliveira Hickson 1994, Abraham et al."
 1996a.1996b. van den Berel et al.," 1996a,1996b, van den Bergh et al."
 1996. Nain. Ratuatuuea Caifiths 1997. Brinchiuaun ct al.," 1996, Naim, Ratnatunga Griffiths 1997, Brinchmann et al."
 1998. Marleau Simard 1998. Cousclice Bershady 1998. Couselice Callagher 1999).," 1998, Marleau Simard 1998, Conselice Bershady 1998, Conselice Gallagher 1999)."
" Ou the theoretical side. the efforts have been larecly couceutrated on numerical N-body snaulatious aimed at iuvestigating the fate and the evolution of merging galaxies (see e.g. Darues Ieruquist 1992. Barnes 19098 aud references therein) aud the distortions and the morphological alterations produced in galaxies iu clusters and eroups bv the ensemble of ""tidal shocks” with other members (Moore et al."," On the theoretical side, the efforts have been largely concentrated on numerical $N$ -body simulations aimed at investigating the fate and the evolution of merging galaxies (see e.g. Barnes Hernquist 1992, Barnes 1998 and references therein) and the distortions and the morphological alterations produced in galaxies in clusters and groups by the ensemble of “tidal shocks” with other members (Moore et al."
 1996. 1998).," 1996, 1998)."
 This mechanisi cau dive the evolution of spiral galaxies iuto spheroidal svstems and it is a strong candidate explanation for the difference in the relative nuniber of spiral to elliptical ealaxics in clusters at ciffercut redshifts., This mechanism can drive the evolution of spiral galaxies into spheroidal systems and it is a strong candidate explanation for the difference in the relative number of spiral to elliptical galaxies in clusters at different redshifts.
 The plan of the paper is the following., The plan of the paper is the following.
 Iu $82 we outline the method adopted for this investigation while the mathematical details are described in the Appendix., In 2 we outline the method adopted for this investigation while the mathematical details are described in the Appendix.
 Tn 83 we show the results obtained for initial coucditious typical of ealactic dark halos and their dependence on the orbital paramcters of the perturber., In 3 we show the results obtained for initial conditions typical of galactic dark halos and their dependence on the orbital parameters of the perturber.
 Iu &1 initial conditions relevant for spleroidal galaxies aud a survev of orbital paraiucters of the perturber typical of differcut euvironnients are considered., In 4 initial conditions relevant for spheroidal galaxies and a survey of orbital parameters of the perturber typical of different environments are considered.
 Overall. our results quantify he relationship between the interaction rate for the environment and the probability of observing inregular uorphologies. aiding interpretation of any trend in the yaction of asviunietrie galaxies at different redshifts.," Overall, our results quantify the relationship between the interaction rate for the environment and the probability of observing irregular morphologies, aiding interpretation of any trend in the fraction of asymmetric galaxies at different redshifts."
 This latter aspect is of particular interest as data ou he amorpholoey of distaut galaxies are now available roni UST observations in the Tnibble Deep Field aud in the Medii: Deep Survey. projects (Abraham et al., This latter aspect is of particular interest as data on the morphology of distant galaxies are now available from HST observations in the Hubble Deep Field and in the Medium Deep Survey projects (Abraham et al.
 1996a. 1996b).," 1996a, 1996b)."
 Iu order to facilitate the comparison of our theoretical results with observational data. we will quantify the distortious produced by imeaus of the asviunetry parameter A (Abraham et al.," In order to facilitate the comparison of our theoretical results with observational data, we will quantify the distortions produced by means of the asymmetry parameter $A$ (Abraham et al."
 1996a) which has been determined for a ΙΙΙ of local aud distaut galaxies., 1996a) which has been determined for a number of local and distant galaxies.
 We will also introduce a generalized asvaiuunetry parameter Ade) which will provide an important information on the radial structure of the asvunuetry produced by the miechanisni we have considered and which will allow to test the theory presented here., We will also introduce a generalized asymmetry parameter $A(r)$ which will provide an important information on the radial structure of the asymmetry produced by the mechanism we have considered and which will allow to test the theory presented here.
 The main couchisious are sunuuarized and discussed in &5., The main conclusions are summarized and discussed in 5.
 This section outlines the perturbative solution of the collisionless Boltzinanu equation used to predict the dynamical response of galaxies to interactions., This section outlines the perturbative solution of the collisionless Boltzmann equation used to predict the dynamical response of galaxies to interactions.
 Additional mathematical detail is presented in the Appendix., Additional mathematical detail is presented in the Appendix.
 The linear perturbation theory used here has been applied in the past both to investigate the collisiouless stability of stellar disks (ναπια] 1977) aud of spherical stellar systems (Polvachenko Shukhinan 1981. Palmer Papaloizou 1957. Bertin Pegoraro 1989. Saha 1991. Weinbere 1991. Bertin ct al.," The linear perturbation theory used here has been applied in the past both to investigate the collisionless stability of stellar disks (Kalnajs 1977) and of spherical stellar systems (Polyachenko Shukhman 1981, Palmer Papaloizou 1987, Bertin Pegoraro 1989, Saha 1991, Weinberg 1991, Bertin et al."
 1991: see also Palmer 1994 aud references therein) aud to study the response of a galaxy to perturbations induced by satellite companions (Weinberg 1989. 1995. 1998b).," 1994; see also Palmer 1994 and references therein) and to study the response of a galaxy to perturbations induced by satellite companions (Weinberg 1989, 1995, 1998b)."
 This method is nuuercallv iuteusive but it avoids the problems of n-body simulatious due to the finite nuuber of particles which ean introduce spurious noise effects (see Woemberg 1998a) and it guarantees the resolution of resonances that lead to the excitation of patterus iu the primary svsteni., This method is numerically intensive but it avoids the problems of n-body simulations due to the finite number of particles which can introduce spurious noise effects (see Weinberg 1998a) and it guarantees the resolution of resonances that lead to the excitation of patterns in the primary system.
 With too few particles. for example. noise can cause orbital energies and angular momenta to dift on a time scale shorter than the pattern speed of the global respouse of interest.," With too few particles, for example, noise can cause orbital energies and angular momenta to drift on a time scale shorter than the pattern speed of the global response of interest."
 This clinumates the possibility of observing this structure in the simulation., This eliminates the possibility of observing this structure in the simulation.
 No doubt. real galaxies are uot sinooth aud inherent fluctuations may be iportaut to them evolution.," No doubt, real galaxies are not smooth and inherent fluctuations may be important to their evolution."
 However. many iillious of n-bodv particles are required to approach the estimates of natural auplitudes (Nelson Tremaine 1996. Weinberg Ἰθῆδα].," However, many millions of n-body particles are required to approach the estimates of natural amplitudes (Nelson Tremaine 1996, Weinberg 1998a)."
 We compute the perturbation induced ou a spherical stellar svstem by a point nass perturber lowdth mass m on a vectilinear trajectory with pericenter p., We compute the perturbation induced on a spherical stellar system by a point mass perturber with mass $m$ on a rectilinear trajectory with pericenter $p$.
 The response of the primary svsteni will be determined bx the sinultaneous solution of the linearized collisiouless Boltzmann and Poissou's equations. 2g vocas) where the subscript0 denotes the equilibrium quantities. subscript 1l denotes the first order perturbation of the distribution function f. the ILunitoniui Z7. and the potential $.," The response of the primary system will be determined by the simultaneous solution of the linearized collisionless Boltzmann and Poisson's equations, - =0, ^2 _1, where the subscript$0$ denotes the equilibrium quantities, subscript $1$ denotes the first order perturbation of the distribution function $f$, the Hamiltonian $H$, and the potential $\Phi$."
 The Boltzimannu equation has been written in terms of action-angle variables Low.," The Boltzmann equation has been written in terms of action-angle variables $\bf {I}, \bf{w}$."
" The perturbed potential. d. is the sum of the tidal potential due to the perturber. o ®,. aud the eravitational potential of the"," The perturbed potential, $\Phi_1$ , is the sum of the tidal potential due to the perturber, $\Phi_{p}$ , and the gravitational potential of the"
The red optical to NTR colors of these galaxies therefore appear to be a consequence of the SEDs rather than of dust reddening.,The red optical to NIR colors of these galaxies therefore appear to be a consequence of the SEDs rather than of dust reddening.
 The imuorphologies also show no color differentiation., The morphologies also show no color differentiation.
 This would not be expected if a reddened nucleus were beginniug to appear at the longer wavelengths., This would not be expected if a reddened nucleus were beginning to appear at the longer wavelengths.
 Photometric redshift fitting should be quite robust for this type of source., Photometric redshift fitting should be quite robust for this type of source.
 We used the photometric fitting routine of DDolzonella. Pelló.. Miralles (2001) anc our observed B. V. R. £I. J. H. and K imaguitudes to estimate redshifts for the N-rav sources.," We used the photometric fitting routine of \markcite{hyperz}B Bolzonella, Pelló,, Miralles (2001) and our observed $B$ , $V$ , $R$, $I$, $J$, $H$, and $K'$ magnitudes to estimate redshifts for the X-ray sources."
 The program compares the observed SED of the galaxy to a set of template spectra eoncrated with the Druzual Charlot evolutionary code (GISSEL9s. BBiuzual Charlot 1993).," The program compares the observed SED of the galaxy to a set of template spectra generated with the Bruzual Charlot evolutionary code (GISSEL98, \markcite{bc93}B Bruzual Charlot 1993)."
 We used the cight svuthetic star formation histories constructed to match tle observed. properties of local Ποιά galaxies from type E to hu. We asstuned no reddening., We used the eight synthetic star formation histories constructed to match the observed properties of local field galaxies from type E to Im. We assumed no reddening.
 Figure 3. shows the observed SED (solid squares) of Source 23 with its deep break between the J and JJ bauds., Figure \ref{figseds} shows the observed SED (solid squares) of Source 3 with its deep break between the $J$ and $H$ bands.
 The best-fit solution frouZgperz is overlaid on the observed. SED., The best-fit solution from is overlaid on the observed SED.
 The photometric redshifts for the three N-ray sources behind A2390 ire i44l.H(3d.3.—1.6) (Source 1). zppor=1.5(1.3.1.7) (Source 2). and. i55;=2.0(2.1.2.7) (Source 3). where the bracketed numbers are the confidence ranges.," The photometric redshifts for the three X-ray sources behind A2390 are $z_{phot}=1.4(1.3-1.6)$ (Source 1), $z_{phot}=1.5(1.3-1.7)$ (Source 2), and $z_{phot}=2.6(2.4-2.7)$ (Source 3), where the bracketed numbers are the confidence ranges."
 All of the SEDs are adequately fit by single burst models;, All of the SEDs are adequately fit by single burst models.
 The vouusest source (age 5«10? vit) ds also the highest redshift source (Source 3)., The youngest source (age $5\times 10^8$ yr) is also the highest redshift source (Source 3).
 While the photometric redshift estimates are not sensitive to the details of the models. the age estimate is affected by the choice of cosinology and the choice of models.," While the photometric redshift estimates are not sensitive to the details of the models, the age estimate is affected by the choice of cosmology and the choice of models."
 For example. a lower metallicity would require a larger age.," For example, a lower metallicity would require a larger age."
 The source behiud A3TO is found to have ρω=L0001.2). consistent with its spectroscopic redshift κ=1.060.," The source behind A370 is found to have $z_{phot}=1.0(0.8-1.2)$, consistent with its spectroscopic redshift $z_{spec}=1.060$."
" The inchision of reddening docs not change the redshift estimates,", The inclusion of reddening does not change the redshift estimates.
 We note that Τόπο et ((1998) and WWiluan. Fabian Gandhi (2000) put Source 3 at zjí4;220.5—0.6.," We note that \markcite{lemonon98}L Lémmonon et (1998) and \markcite{wilman2000}W Wilman, Fabian Gandhi (2000) put Source 3 at $z_{phot}\simeq 0.5-0.6$."
 The addition of the Z7 baud data breaks the degeneracy in the redshift tuformation and places the source at a mich lugher redshift in order to reproduce the strong JLL break., The addition of the $H$ band data breaks the degeneracy in the redshift information and places the source at a much higher redshift in order to reproduce the strong $J-H$ break.
 Subsequent to estimating the photometric redshifts we obtained spectra of Sources 1: aud 2 with CISCO., Subsequent to estimating the photometric redshifts we obtained spectra of Sources 1 and 2 with CISCO.
 Source Ll (Fig. D) , Source 1 (Fig. \ref{figspectra}) )
has an unambienous redshift of 1.167. based on the detection of 3727 aud 5007.," has an unambiguous redshift of 1.467, based on the detection of $\,3727$ and $\,5007$."
 The weaker L959 feature was too faint to be detected.," The weaker $\,4959$ feature was too faint to be detected."
 The 3727 line was coufiiied with a one hour spectrum (AAGOOO—10000.AL. 1.17 wide slit. ~12 rresolution) taken with LRIS on Keck Tou UT 1 October 2000.," The $\,3727$ line was confirmed with a one hour spectrum $\lambda\lambda6000-10000$, $''$ wide slit, $\sim12$ resolution) taken with LRIS on Keck I on UT 4 October 2000."
 This higher resolution spectrum yielded. a redshift of 1.166., This higher resolution spectrum yielded a redshift of 1.466.
 The apparent weakness of the emission lines in the spectrum is a consequence of the low resolution of the infrared spectroscopic observations. aud the equivalent widths of the lines are quite representative of other field ealaxies at these redshifts.," The apparent weakness of the emission lines in the spectrum is a consequence of the low resolution of the infrared spectroscopic observations, and the equivalent widths of the lines are quite representative of other field galaxies at these redshifts."
 Source 2 has only a single Lue which. if interpreted as 5007. would also place this source at 1.167. consistent with the photometric estimate.," Source 2 has only a single line which, if interpreted as $\,5007$ would also place this source at 1.467, consistent with the photometric estimate."
 2727 is not clearly preseut. however. so it is also possible that this line could be 5007. placing the galaxy at tape.=2.317.," $\,3727$ is not clearly present, however, so it is also possible that this line could be $\,5007$, placing the galaxy at $z_{spec}=2.317$."
 No LRIS spectruu is prescutly available for this source., No LRIS spectrum is presently available for this source.
 A one hour LRIS spectrum was also obtained for Source 3 covering the 6000 to 0000 region. but. as is reasonable given the faintucss of the source. no features could be clearly secu.," A one hour LRIS spectrum was also obtained for Source 3 covering the 6000 to 10000 region, but, as is reasonable given the faintness of the source, no features could be clearly seen."
 Based on ISOCAM. data (LLGuunonou et 11998. AAlticri et 11999) all three of the galaxies have sienificant excess flux in the mid-infrared waveleneths over and above the galaxy light (see Fie. 3)).," Based on ISOCAM data \markcite{lemonon98}L Lémmonon et 1998, \markcite{altieri99}A Altieri et 1999) all three of the galaxies have significant excess flux in the mid-infrared wavelengths over and above the galaxy light (see Fig. \ref{figseds}) ),"
 while none have strong 850 micron detectious (FFabian ct 22000: BBarecr νο 2001)., while none have strong 850 micron detections \markcite{fabian00}F Fabian et 2000; \markcite{barger01b}B Barger Kneib 2001).
 This longerwaveleneth excessin the infrared and the absence of strong sibuullimeter light is best interpreted with a model of hot dust surrounding the central AGN.as discussed in WWilnan. Fabian Gandhi (2000).," This longerwavelength excessin the infrared and the absence of strong submillimeter light is best interpreted with a model of hot dust surrounding the central AGN,as discussed in \markcite{wilman2000}W Wilman, Fabian Gandhi (2000)."
 We postpone a more detailed discussion, We postpone a more detailed discussion
lem In this appendix. we give the mathematical definitions of the three classic 2D. interpolation methods: Bilinear. Bicubic. Spline.,"1cm In this appendix, we give the mathematical definitions of the three classic 2D interpolation methods: Bilinear, Bicubic, Spline."
 For their logarithmic extensions.. Log-Dilinear. Log-Bicubic. Log-Spline). we only have one minor point to aclelress at the end of this section.," For their logarithmic extensions, Log-Bilinear, Log-Bicubic, Log-Spline), we only have one minor point to address at the end of this section."
 The Bilinear method is the simplest of the three., The Bilinear method is the simplest of the three.
 Let us write the coordinates of the grid points as Gr;.gj) Gop1.2.3...). and the signals as Gr.y;).," Let us write the coordinates of the grid points as $(x_i,y_j)$ $i,j=1,2,3...$ ), and the signals as $A(x_i,y_j)$."
 Suppose the point of our interest is Grey). which satisfies wr;xαπmoo;| and HjEogXgja. the Bilinear method defines Cr.jy) in the Following wav: where The Bicubie method includes higher order terms of { and i to achieve smoothness of the interpolated function.," Suppose the point of our interest is $(x,y)$, which satisfies $x_i\le x\le x_{i+1}$ and $y_j\le y\le y_{j+1}$, the Bilinear method defines $A(x,y)$ in the following way: where The Bicubic method includes higher order terms of $t$ and $u$ to achieve smoothness of the interpolated function."
 It requires the user to specify not only the signal lr;jy;). but also the spatial derivatives Q.l1/r. 0A/0g. and 07AOrdg al every grid point Gry.y;).," It requires the user to specify not only the signal $A(x_i,y_j)$, but also the spatial derivatives $\partial A/\partial x$, $\partial A/\partial y$, and $\partial^2A/\partial x\partial y$ at every grid point $(x_i,y_j)$."
 Since the spatial derivatives of the signal are usually not known a priori. we estimate then using the finite-dilference method: The interpolatecl function. inside cach grid square is written in the following polvnomial form: The values of the sixteen parameters Gy. are constrained. using eq.(15)) and the following three equations at the four corners of the grid square: where / and. η have been defined in eq.(13)).," Since the spatial derivatives of the signal are usually not known a priori, we estimate them using the finite-difference method: The interpolated function inside each grid square is written in the following polynomial form: The values of the sixteen parameters $c_{mn}$ are constrained using \ref{A_bicubic}) ) and the following three equations at the four corners of the grid square: where $t$ and $u$ have been defined in \ref{tu}) )."
 The 2D Spline method. simply refers to using the 1D Spline interpolation along each cimension., The 2D Spline method simply refers to using the 1D Spline interpolation along each dimension.
 The 1D (cubic) Spline interpolation method works as follows: Given the values of a function fre) at a set of points .r; ο the form of the function in the interval between wy andlor ds written as:," The 1D (cubic) Spline interpolation method works as follows: Given the values of a function $f(x)$ at a set of points $x_i$ $i=1...N$ ), the form of the function in the interval between $x_j$ and $x_{j+1}$ is written as:"
tvpes of interesting substructure.,types of interesting substructure.
 The most prominent of these is the inner region of NGC TOS. which exhibits an easily seen cust lane (see the illustrations of each one shown in Lauer et al.).," The most prominent of these is the inner region of NGC 708, which exhibits an easily seen dust lane (see the illustrations of each one shown in Lauer et al.)."
 Smaller dust lanes were also detected. within NGC 5193 and IC 4296. indicating likely accretion events in the recent past.," Smaller dust lanes were also detected within NGC 5193 and IC 4296, indicating likely accretion events in the recent past."
 Giovanellietal.(1982) found that A262. the host cluster for NGC! 708. contains galaxies deficient of II] gas. suggested as due to ranpressure stripping.," \citet{Gio82} found that A262, the host cluster for NGC 708, contains galaxies deficient of HI gas, suggested as due to ram-pressure stripping."
 This gas could be accreted by the central ο) galaxy. with its large dark halo.," This gas could be accreted by the central cD galaxy, with its large dark halo."
 A262 is a spiral-rich cluster. characterized by a X-ray source. οἱ 0152-36. which appears concentric with (he central galaxy.," A262 is a spiral-rich cluster, characterized by a X-ray source, 3U 015+36, which appears concentric with the central galaxy."
 The gas aud galaxy. virial temperatures for the cluster have recently been shown to be comparable by Neilletal.(2001)., The gas and galaxy virial temperatures for the cluster have recently been shown to be comparable by \citet{Nei01}.
. There seems to be a general pattern for IHE deficiency to be associated with X-ray emission. some other examples being A214 in ]lercules. or Coma.," There seems to be a general pattern for HI deficiency to be associated with X-ray emission, some other examples being A2147 in Hercules, or Coma."
 Our procedures for photometry of the four combined images follow the basic sequence outlined in more detail elsewhere (e.g.. Wooclworth&Harris(2000).. (2000))).," Our procedures for photometry of the four combined images follow the basic sequence outlined in more detail elsewhere (e.g., \citet{WoHa00}, \citet{Kav00}) )."
 OF the many detected [aint starlike objects on the frames. the vast majority are elobular clusters within (he DCGs.," Of the many detected faint starlike objects on the frames, the vast majority are globular clusters within the BCGs."
 At these distances (~ 55 Alpe) they appear as unresolved point sources. even in the PC frames.," At these distances $\sim$ 55 Mpc) they appear as unresolved point sources, even in the PC1 frames."
 Hence. it is easy {ο perform conventional point-spread [function (PSF) photometry on the frames.," Hence, it is easy to perform conventional point-spread function (PSF) photometry on the frames."
 Ixdependent empirical PSFs were constructed from several moderately. bright. uncrowded stars in each of the four WEPC? frames with DAOPIIOT.," Independent empirical PSFs were constructed from several moderately bright, uncrowded stars in each of the four WFPC2 frames with DAOPHOT."
 Then ALLSTAR (Stetson1994) was used to generate the final photometry., Then ALLSTAR \citep{Ste94} was used to generate the final photometry.
 A stanclard detection threshold of 3.5 times the rms scatter of the skv background. was adopted for the image detection., A standard detection threshold of 3.5 times the rms scatter of the sky background was adopted for the image detection.
 Crowding at all magnitudes was completely negligible in these high-latitude fields., Crowding at all magnitudes was completely negligible in these high-latitude fields.
 An image classification algorithm (C'LASSIFY: defined by Ixron.(1980) ancl (1991))) was then used to caleulate radial image moments of the candidate objects and thus to separate stellar Irom nonstellar objects in an objective way.," An image classification algorithm (CLASSIFY; defined by \citet{Kro80}
 and \citet{Har_et91}) ) was then used to calculate radial image moments of the candidate objects and thus to separate stellar from nonstellar objects in an objective way."
 By using artificial-star data passed through exactly the same measurement process. we established boundaries for the CLASSIFY parameters such that at least of the (rue stars were retrieved. αἱ all magnitudes (for verv similar examples with illustrative graphs. see Wooclworth(2000) ancl Navelaarsetal. (2000))).," By using artificial-star data passed through exactly the same measurement process, we established boundaries for the CLASSIFY parameters such that at least of the true stars were retrieved at all magnitudes (for very similar examples with illustrative graphs, see \citet{WoHa00} and \citet{Kav00}) )."
 As usual. for the faintest objects the image moments become verv uncertain. ancl it becomes difficult to distinguish between stellar ancl nonstellar objects.," As usual, for the faintest objects the image moments become very uncertain, and it becomes difficult to distinguish between stellar and nonstellar objects."
 Also. some nonstellar objects will be accidentally classified. as stellar.," Also, some nonstellar objects will be accidentally classified as stellar."
 These are statistically removed [rom the final GCLF bv subtracting a background luminosity [function (see below)., These are statistically removed from the final GCLF by subtracting a background luminosity function (see below).
 Finally. to define the photometric completeness function we constructed annular ringsaround the DCGs. and emploved extensive arUlicial-star tests (ο measure the detection completeness (mn. r). as a function of magnitude and radial distance.," Finally, to define the photometric completeness function we constructed annular ringsaround the BCGs, and employed extensive artificial-star tests to measure the detection completeness $f(m,r)$ , as a function of magnitude and radial distance."
 In practice we found, In practice we found
observed [act that there is star formation in regions that are below critical densities indicated a softer reduction of the star formation elliciency (han a simple truncation.,observed fact that there is star formation in regions that are below critical densities indicated a softer reduction of the star formation efficiency than a simple truncation.
 The value of the exponent (-1.5) produces a good match to the observed slope., The value of the exponent (-1.5) produces a good match to the observed slope.
 That such a reduction is required. (o match the observations is further evidence for a reduction of star formation efficiency below the rate predicted by the simple 4chimidt law at low surface densities., That such a reduction is required to match the observations is further evidence for a reduction of star formation efficiency below the rate predicted by the simple Schmidt law at low surface densities.
 The soft roll off of the Schmidt Law used in (his work does not imply that there is not a local critical densitv below which star formation does not occur., The soft roll off of the Schmidt Law used in this work does not imply that there is not a local critical density below which star formation does not occur.
 Even though the average density in the roughly 1 square kiloparsec area covered bv a single pixel may be below the critical density. there may be several areas where the local density is higher than the critical density.," Even though the average density in the roughly 1 square kiloparsec area covered by a single pixel may be below the critical density, there may be several areas where the local density is higher than the critical density."
 This is probably the best physical interpretation of the soft roll off on the kpe scale., This is probably the best physical interpretation of the soft roll off on the kpc scale.
" The main parameters that control the shape of the distribution Dunction are r, aud m.", The main parameters that control the shape of the distribution function are $r_e$ and $m^*$.
 The fit is relatively insensitive to a and the value of ó* mainly scales the distribution rather than change its shape., The fit is relatively insensitive to $\alpha$ and the value of $\phi^*$ mainly scales the distribution rather than change its shape.
" In hierarchical models of galaxy. formation and evolution both the value of m and r, will decrease as the redshilt increases.", In hierarchical models of galaxy formation and evolution both the value of $m^*$ and $r_e$ will decrease as the redshift increases.
" Lowering i"" decreases the value ol x where (he transition from power law to exponential occurs.", Lowering $m^*$ decreases the value of x where the transition from power law to exponential occurs.
" The effect of reducing r, is just the opposite. it increases the value of x where the transition occurs."," The effect of reducing $r_e$ is just the opposite, it increases the value of x where the transition occurs."
 Both of these effects make physical sense., Both of these effects make physical sense.
 The exact evolution of the intensity distribution will then depend on the evolution of (hese (wo parameters., The exact evolution of the intensity distribution will then depend on the evolution of these two parameters.
 In an attempt to determine (he effect of evolution the computation was performed for an epoch where the characteristic mass m* is 1/10 of the 3x10! AL. found from the matching at z= I., In an attempt to determine the effect of evolution the computation was performed for an epoch where the characteristic mass $m^*$ is 1/10 of the $3 \times 10^{10}$ $_{\odot}$ found from the matching at z = 1.
" The value of r, was then reduced to (1/10)! ol its value ad 2551.", The value of $r_e$ was then reduced to $(1/10)^{1/5}$ of its value at z=1.
 The result is the dashed line in fig. 1.., The result is the dashed line in fig. \ref{fig:hx}.
 The main elfect is a transition to the exponential function al a lower of x. If we reduce r. by a larger amount a distribution very close to the z = 1I distribution is obtained., The main effect is a transition to the exponential function at a lower of x. If we reduce $r_e$ by a larger amount a distribution very close to the z = 1 distribution is obtained.
 The HIST observations at redshifts greater than 1.5 do not adequately define the shape of L(x} to provide an observational test of the evolution of h(x)., The HST observations at redshifts greater than 1.5 do not adequately define the shape of h(x) to provide an observational test of the evolution of h(x).
 Galaxv evolution modelers predictions of the ου parameters ancl (he resulting L(x) can be compared at high redshift when NGST becomes operational., Galaxy evolution modeler's predictions of the free parameters and the resulting h(x) can be compared at high redshift when NGST becomes operational.
 We have shown that the star formation intensity distribution shape is a natural consequence ol a Schechter galaxy mass distribution. the Sehimnidt. law with a critical density. ancl star formation occurringe in exponential disks.," We have shown that the star formation intensity distribution shape is a natural consequence of a Schechter galaxy mass distribution, the Schmidt law with a critical density and star formation occurring in exponential disks."
 The primary parameters that control the shape, The primary parameters that control the shape
 tot! 1019 (1014. Lol <10° (Davidsonhuionov&Suuvaev1975) etal.1979).," $10^{14}$ $10^{15}$ $10^{11}$ $10^{13}$ $\lax 10^6$ \citep{davidson73,illarionov75} \citep{illarionov75, davies79}."
. 1086 10° («1% <P«10! (Ransomctal.2002:," $10^6$ $10^6$ $< 1$ $< P < 10^4$ \citep{ransom02, drake02, burwitz02}."
 108 (Braje&Romani2002):: 10 CP<10 P>10! XLot! >10% fis O yr ff:O/0 sd—pLM P Pare Tesh—P/2P. 60z3. ooLory c1. , $10^6$ \citep{braje02}; \citep{braje02}; $10^6$ $P<10$ $P>10^4$ $\lax 10^{14}$ $>10^6$ $I$ $\mu$ $\Omega$ $\mu$ $\vec{\mu}\cdot\vec{\Omega}/\Omega$ $\tau_{sd} \equiv \frac{P}{(n-1)\dot{P}}$ $P$ $\dot{P}$ $\tau_{csd} \equiv P/2\dot{P}$ $n\ne3$ $n < 1$$\tau_{sd}$ $n > 1$ 
considering pure expansion or expansion plus rotation xoduce similar results in this scale range.,considering pure expansion or expansion plus rotation produce similar results in this scale range.
 However. at separations Luger than 100 AU. the Vi function of 1o modeled spots decreases iore rapidly when rotation 1iofions are mceluded. which is qualitatively similar to 1ο V4 of the observed imasers at those separatious.," However, at separations larger than 100 AU, the $V_s$ function of the modeled spots decreases more rapidly when rotation motions are included, which is qualitatively similar to the $V_s$ of the observed masers at those separations."
 At separations «1 AU. a laveer dispersion is present i1 ιο velocity correlation fuuctious for the modeled spots shown by the erey curves). due to a lack of clustering. aud thus poor statistics.," At separations $<$ 1 AU, a larger dispersion is present in the velocity correlation functions for the modeled spots (shown by the grey curves), due to a lack of clustering, and thus poor statistics."
 This occurs because even though ιο total ummber of modeled spots is large. the ΠΟ’ Of pais with small separation is small because modeled sots do uot forme clusters (they are uniformly aud randomly distributed).," This occurs because even though the total number of modeled spots is large, the number of pairs with small separation is small because modeled spots do not form clusters (they are uniformly and randomly distributed)."
 From our results. we thus conclude that the trend describe by the ruis Doppler-velocity difference of maser spot pairs in W75 N VLA 2 as a function of separation can be explained bv the presence of organized motions. expansion. and rotation (both velocities with the same uaenitude ΓΕ... 13 within au arc of a circle observed with μα. inclination anele. 0 207.," From our results, we thus conclude that the trend described by the rms Doppler-velocity difference of maser spot pairs in W75 N VLA 2 as a function of separation can be explained by the presence of organized motions, expansion, and rotation (both velocities with the same magnitude $\sim$ 44 km $^{-1}$ ) within an arc of a circle observed with small inclination angle, $\theta\simeq$ ."
 Theher values of the angle 0 will produce a steeper eracdicu of the velocity difference toward larger separations., Higher values of the angle $\theta$ will produce a steeper gradient of the velocity difference toward larger separations.
 Moreover. he corresponding projected ellipse ou the plaue of the sky would lave its nünor axis much smaller han its uajor axis. Which is not observed.," Moreover, the corresponding projected ellipse on the plane of the sky would have its minor axis much smaller than its major axis, which is not observed."
 The expansion velocity required to fit the data is abou the same order of uaenitude as the measured proper motions (Torrellesal. 2003)., The expansion velocity required to fit the data is about the same order of magnitude as the measured proper motions \citep{Tor03}.
. The source W795 N VLA 1 appears to be different yon VLA 2 since proper-motion measurements of water nasers indicate a collimated outflow in this region (Torrellesetal.?2003)., The source W75 N VLA 1 appears to be different from VLA 2 since proper-motion measurements of water masers indicate a collimated outflow in this region \citep{Tor03}.
". In fact. the velocity correlation ""uctious of niaser spots present a areer dispersion frou. he fit."," In fact, the velocity correlation functions of maser spots present a larger dispersion from the fit."
 Despite the dispersion. the values estimated for he indices (similar to those of VLA 2) sueeestOO the xesence of an organized compouent of motion with a xeferential direction.," Despite the dispersion, the values estimated for the indices (similar to those of VLA 2) suggest the presence of an organized component of motion with a preferential direction."
 Iu this paper. we present a statistical analvsis of the spatial aud velocity distribution of water mascrs in the SFRs Ceplieus A and W75 N. using the data of the VLBA laser Observations carried out by Torrelles 2003).," In this paper, we present a statistical analysis of the spatial and velocity distribution of water masers in the SFRs Cepheus A and W75 N, using the data of the VLBA maser observations carried out by \citet{Tor01b,Tor03}."
. Our couclusious are as follows: L. U. is supported by Secretariaa de Estado de Universidaces ο Tuvestigacidun of MEC. (Spain)., Our conclusions are as follows: L. U. is supported by a de Estado de Universidades e Investigaciónn of MEC (Spain).
 G. À.. J. EF. ν J AL T. and L. OU. are partially supported by Ministerjo de Ciencia e Iunovaciónn (Spain). erat AYA 2008-)6189-COUS. Gucludiug FEDER funds). aud bv Cousejeríaa de Iunovacióuu. Ciencia v Empresa of Juuta de idalucíaa (Spiin).," G. A., J. F. G., J. M. T., and L. U. are partially supported by Ministerio de Ciencia e Innovaciónn (Spain), grant AYA 2008-06189-C03 (including FEDER funds), and by a de Innovaciónn, Ciencia y Empresa of Junta de a (Spain)."
 J.C. ALC. Re aud S. €. acknowledge the support of CONACVT (Mexico) erauts 615[7 aud G05SI.," J. C., A. C. R., and S. C. acknowledge the support of CONACyT (Mexico) grants 61547 and 60581."
" We are thankful to our referee for his/her use""lb comments that helpus to improve this paper.", We are thankful to our referee for his/her useful comments that helpus to improve this paper.
based on the observations of low-mass stars.,based on the observations of low-mass stars.
 In the current paradigm (Shu. Adams. Lizano 1987: Shu et al.," In the current paradigm (Shu, Adams, Lizano 1987; Shu et al."
 1993) the formation of low-mass stars is characterized bv an accretion phase. in which a central protostar and a cireumstellar disk form surrounded by an intalling envelope of dust aud gas. followed by a phase in which the protostar deposits linear ancl angular momentum. and mechanical energv into its surroundings (hrough jets and molecular outflows.," 1993) the formation of low-mass stars is characterized by an accretion phase, in which a central protostar and a circumstellar disk form surrounded by an infalling envelope of dust and gas, followed by a phase in which the protostar deposits linear and angular momentum, and mechanical energy into its surroundings through jets and molecular outflows."
 Although this paradigm has been very successful in explaining what is observationally known about the formation of low-mass stars (e.g. Lada 1991). its applicability to the formation of nmassive stars remains arguable.," Although this paradigm has been very successful in explaining what is observationally known about the formation of low-mass stars (e.g, Lada 1991), its applicability to the formation of massive stars remains arguable."
 la particular. alternative mechanisms such as (he merging of lower mass prolostars to form a massive protostar have received serious consideration lately (e. e. Bonnell. Date. Zinnecker 1998: Portegies Zwart et al.," In particular, alternative mechanisms such as the merging of lower mass protostars to form a massive protostar have received serious consideration lately (e. g. Bonnell, Bate, Zinnecker 1998; Portegies Zwart et al."
 1999). and there is evidence of strong dvnamieal interaction aud possibly merging deduced from the proper motions of sonie voung Massive stars in (he Orion IKL/DN reeion (Gomunez et al.," 1999), and there is evidence of strong dynamical interaction and possibly merging deduced from the proper motions of some young massive stars in the Orion KL/BN region (Gómmez et al."
 2005)., 2005).
 Whether (he massive protostars are Iormed by accretion or merging remains controversial., Whether the massive protostars are formed by accretion or merging remains controversial.
 If massive O stars are lormed by accretion. we expect (hat disks aud jets will be present in (heir earliest stages of evolution. as in the case of low mass stars.," If massive O stars are formed by accretion, we expect that disks and jets will be present in their earliest stages of evolution, as in the case of low mass stars."
 Due to their large Iuninosiües. hvpercompact LIT regions are also expected to be formed while the protostar is still undergoing accretion (Ixeto 2002. 2003).," Due to their large luminosities, hypercompact HII regions are also expected to be formed while the protostar is still undergoing accretion (Keto 2002, 2003)."
 On the other hand. if formed by random coalescence of lower-mass stars neither disks. jets or hvpercompact HII regions are expected.," On the other hand, if formed by random coalescence of lower-mass stars neither disks, jets or hypercompact HII regions are expected."
 Thus. the search for jets. disks and hvpercompact ILLI regions toward massive YSOs is crucial to understand their formation process.," Thus, the search for jets, disks and hypercompact HII regions toward massive YSOs is crucial to understand their formation process."
 The presence of jets ancl disks toward massive YSOs still lacks solid grounds., The presence of jets and disks toward massive YSOs still lacks solid grounds.
 Collimated outflows and/or disks have been found in a handful of D-tvpe protostars: IAS. 18162-2048 (£oLTx10!Liui. guez. Reipurth 1993): Cepheus A IIW2 (£—1x10!Los: guez οἱ al.," Collimated outflows and/or disks have been found in a handful of B-type protostars: IRAS 18162-2048 ${\cal L}\sim1.7\times10^4$;, guez, Reipurth 1993); Cepheus A HW2 ${\cal L} \sim1\times10^4$; guez et al."
 1994): IRAS 2012644104 (0~1.3x105 Cesaroni et al., 1994); IRAS 20126+4104 ${\cal L} \sim1.3\times10^4$; Cesaroni et al.
 1997): (£~3x103: Shepherd οἱ al., 1997); G192.16-3.82 ${\cal L}\sim3\times10^3$; Shepherd et al.
 1998: Devine et al., 1998; Devine et al.
 1999: Shepherd. Claussen. Kurtz POOL): AFGL 5142 (£~3.0x10*Los Zhang et al.," 1999; Shepherd, Claussen, Kurtz 2001); AFGL 5142 ${\cal L}\sim3.0\times10^3$; Zhang et al."
 2002)., 2002).
 None of these objects exceeds a bolometric luminositv of 2xLOAL. and are thus BO ZAMS or lower luminosity objects., None of these objects exceeds a bolometric luminosity of $2\times10^4$ and are thus B0 ZAMS or lower luminosity objects.
 The most luminous O-(vpe protostars associated. with a jet and a collimated molecular outflow is IRAS 16547-4247. a luminous infrared source with a bolometric luminosity of 6.2x10! (Garay οἱ al.," The most luminous O-type protostars associated with a jet and a collimated molecular outflow is IRAS $-$ 4247, a luminous infrared source with a bolometric luminosity of $6.2\times10^4$ (Garay et al."
 2002. 2007: guez οἱ al.," 2002, 2007; guez et al."
 2005)., 2005).
 There is. however. no known unquestionable case of a cireumstellar disk associated with a massive O-tvpe protostar.," There is, however, no known unquestionable case of a circumstellar disk associated with a massive O-type protostar."
 The search for disks around independent hieh-mass prolostars is inherently difficult., The search for disks around independent high-mass protostars is inherently difficult.
 The photoevaporation timescale of disks by the radiation from the central star (Yorke Welz 1996: Hollenbach 1997) ean be relatively short. and disks may be destroyed. before the dispersal of the molecular cloud core.," The photoevaporation timescale of disks by the radiation from the central star (Yorke Welz 1996; Hollenbach 1997) can be relatively short, and disks may be destroyed before the dispersal of the molecular cloud core."
 I other OB, If other OB
From. figure Y of⋅ MOI. we can infer⋠⋅ that 1|o.ON splits. galaxies. brighter. than the morphological. resolution. roughly into. cllipticals2. ancl spirals.,"From figure 7 of M01, we can infer that $B-V\sim0.8$ splits galaxies brighter than the morphological resolution roughly into ellipticals and spirals."
. Llowever. this. split 2.is not perfect.⋅ as noted above for⋅ the reference⋅ galaxies.," However, this split is not perfect, as noted above for the reference galaxies."
. Moreover. since. we want to highlightMM a progressive: change in. the spatial distribution of galaxies with a given property. we will proceed. as in the previous paragraph.," Moreover, since we want to highlight a progressive change in the spatial distribution of galaxies with a given property, we will proceed as in the previous paragraph."
 We split. the range covered by the D.V index of all galaxies above the uminosity. resolution limit of the simulation (0zD.V<1.42 ')) in a series of 7 bins called to (Galaxy Colour). from the bluest to the rededest.," We split the range covered by the $B-V$ index of all galaxies above the luminosity resolution limit of the simulation $0\la B-V \la 1.42$ ) in a series of 7 bins called to (Galaxy Colour), from the bluest to the reddest."
 The bluest bin is 0 15]. and the following bin thresholds are separated. by 0.1 in BV..," The bluest bin is [0 0.5], and the following bin thresholds are separated by 0.1 in $B-V$."
 The last' bin covers' the rangeD 1:15]., The last bin covers the range [1 1.5].
U ...Again.. the imits of the bins. and number of objects. are given in the sixth and seventh linesof Table ," Again, the limits of the bins, and number of objects, are given in the sixth and seventh lines of Table \ref{tab:BinLimits}."
We stress that we have binned up all galaxies above the resolution. limit. ofM our simulation., We stress that we have binned up all galaxies above the luminosity resolution limit of our simulation.
 This. upallows us the o go substantially fainter than if we restricted ourselves to galaxies for which the simulation gives reliable morphologies., This allows us to go substantially fainter than if we restricted ourselves to galaxies for which the simulation gives reliable morphologies.
 Our last sample selection makes use of the morphogical information in the simulation., Our last sample selection makes use of the morphogical information in the simulation.
 By definition. we need. here restrict. ourselves⋅⋅⊀ to⋅⇁ the galaxies brighter⋅⊀ than.We the ‘ morphology resolution. limitEN (Mg<.—— 18.46).," By definition, we need here to restrict ourselves to the galaxies brighter than the morphology resolution limit $M_{B}\leq
-18.46$ )."
 ⋅⊳∖The sample selection. is. done according. to AMisaauss—AlBaeua Gvhich. ranges. from» O to ~-7 among the simulated. galaxies2a which. possess a bulge).," The sample selection is done according to $M_{\rmn{B,bulge}}-M_{\rmn{B,total}}$ (which ranges from 0 to $\sim 7$ among the simulated galaxies which possess a bulge)."
 opThe bins. are called to (Galaxy Morphology) from the least bulge-dominated to the most. bulge-dominated., The bins are called to (Galaxy Morphology) from the least bulge-dominated to the most bulge-dominated.
o Note that for simplicity.1 we have put bulge-less galaxies. (above the morphology resolution)≜ together with. the objects. in. the first⋅ bin.," Note that for simplicity, we have put bulge-less galaxies (above the morphology resolution) together with the objects in the first bin."
oo Phe last two lines. of iTable 1. give. the limitsME of. these 7morphologv2l. bins. together with. the associated. number of⋅⊲ objects.," The last two lines of Table \ref{tab:BinLimits} give the limits of these “morphology” bins, together with the associated number of objects."
 ‘To compare the local environments of the various samples. and to look for a possible simple signature of a void population. we first compute the mass density (by smoothing the DAL field) ofB the region. surrounding.+ cach galaxy. and we compare the clistributions of this local densitv in our various samples.," To compare the local environments of the various samples, and to look for a possible simple signature of a void population, we first compute the mass density (by smoothing the DM field) of the region surrounding each galaxy, and we compare the distributions of this local density in our various samples."
 In addition. we estimate Iuminosity.n function1sity [i(per1 unitm mass)1 ofas objectsyects Ivingnpe] inuminosity environments of given density. focussing particularly on the variation of the overall shape.," In addition, we estimate the luminosity function (per unit mass) of objects lying in environments of given density, focussing particularly on the variation of the overall shape."
 characterize the environments of the galaxies by. mass to ↔ ‘ ∠⇂∢⋅⊔≱∖↓↿⊓⋅≱∖≱∖⊔↓∪∪↥↓↥∢⋅∠⇂∪∖⇁∢⋅↓⋅⊳↧⋅↱≻∣∣ smoothing. DEIscale. (aE smoothing.> of⋅ 10 is. Ealready too larget to bring.© out any trends between the samples)., We characterize the environments of the galaxies by mass densities smoothed over a 5 smoothing scale (a smoothing of 10 is already too large to bring out any trends between the samples).
" WeὉ assiassign the» DM distributiondistributi of the ⋅simulationall on ia fine regular grid with mesh spacing much smaller than one smoothing length (/2,). using a CIC scheme."," We assign the DM distribution of the simulation on a fine regular grid with mesh spacing much smaller than one smoothing length $R_{s}$ ), using a CIC scheme."
 The smoothing of the density. field is performed on this fine grid by means ofa Gaussian kernel which takes the form:, The smoothing of the density field is performed on this fine grid by means of a Gaussian kernel which takes the form:
212C4barders in their formation route is thus very useful.,barriers in their formation route is thus very useful.
" Lines 427C(C4pf eg. HCN and H:O are enhanced with respect to the fine-structure lines. especially for the highest ""hforementionedmechanical heating rate."," Lines of, e.g., HCN and $_2$ O are enhanced with respect to the aforementioned fine-structure lines, especially for the highest considered mechanical heating rate."
 It already pointed by consideredLoenenetal.(2008) that the HCwasN/HNC line ratios out when mechanical heating important. since at PUSH Hicreasetemperatures the HNC will be drivenis in to HCN (see glso Section 22)).," It was already pointed out by \citet{Loenen2008} that the HCN/HNC line ratios increase when mechanical heating is important, since at high temperatures the HNC will be driven in to HCN (see also Section \ref{Chemistry}) )."
 Moreover. the HCN/CO line ratios. also s4c3jncrease when mechanical heating effects are important. but 1¢-3the interpretation of an observed ratio is not straightforward. because the lines trace slightly different regions.," Moreover, the HCN/CO line ratios also increase when mechanical heating effects are important, but the interpretation of an observed ratio is not straightforward, because the lines trace slightly different regions."
 $609) [NII] fine-structure. lines. show a very strong response acefo LERenhanced cosmic ray rates. independent of mechanical 5.Q(-6feating effects.," The [NII] fine-structure lines show a very strong response to enhanced cosmic ray rates, independent of mechanical heating effects."
 The main concern with these lines. however. 4¢-Sys that one has to disentangle the contributions arising from 32C5fhese CR exposed clouds and HII regions. that also contribute 63658 significant amount of [NII] emission in starforming LiCABenvironments.," The main concern with these lines, however, is that one has to disentangle the contributions arising from these CR exposed clouds and HII regions, that also contribute a significant amount of [NII] emission in starforming environments."
NisseFit was run 10000. times for each galaxy. and the first 2000 runs were excluded to reach the fits where the 4? had converged to values near the minimum.,"NisseFit was run $10\,000$ times for each galaxy, and the first $2\,000$ runs were excluded to reach the fits where the $\chi^2$ had converged to values near the minimum."
 The resulting runs then reflect the actual best fit distribution of the parameter probability space fitted., The resulting runs then reflect the actual best fit distribution of the parameter probability space fitted.
 To study these probabilities. we collapse the runs into all possible 2-D combinations. and plot the significance. contours.," To study these probabilities, we collapse the runs into all possible 2-D combinations, and plot the significance contours."
 One example is shown in Fig. 3.., One example is shown in Fig. \ref{fig:params}.
 It is clear that metallicity (in all fits) and SER (in the constant SFR scenarios. not shown in Fig. 3))," It is clear that metallicity (in all fits) and SFR (in the constant SFR scenarios, not shown in Fig. \ref{fig:params}) )"
 cannot be determined accurately. and we ignore these results hereafter.," cannot be determined accurately, and we ignore these results hereafter."
 The parameters that are well determined are stellar mass. dust Ay and age. although the age fits in the 2 SSP fits suffer from a degeneracy at the older age when the young population age is very young.," The parameters that are well determined are stellar mass, dust $_V$ and age, although the age fits in the 2 SSP fits suffer from a degeneracy at the older age when the young population age is very young."
 In other words. when the young population age is very young. the light of these stars will outshine any older population. unless the mass fraction in the young population is very small (see Nilsson et al.," In other words, when the young population age is very young, the light of these stars will outshine any older population, unless the mass fraction in the young population is very small (see Nilsson et al."
 2010. for a longer discussion on this).," 2010, for a longer discussion on this)."
 The ages and mass fractions for all the 2 SSP fits (the only runs with more than one age fitted) indicate that we are not able to identify an older population in any of the objects. and only upper limits on the mass fraction in a potential old population may be inferred.," The ages and mass fractions for all the 2 SSP fits (the only runs with more than one age fitted) indicate that we are not able to identify an older population in any of the objects, and only upper limits on the mass fraction in a potential old population may be inferred."
 To determine the best fit parameters in. 1-D. a method similar to that of Nilsson et al. (," To determine the best fit parameters in 1-D, a method similar to that of Nilsson et al. ("
2009b) was used. where each parameter ts decided individually and independently from the other parameters by finding the median of the distribution of runs and integrating to of each wing to find the scatter.,"2009b) was used, where each parameter is decided individually and independently from the other parameters by finding the median of the distribution of runs and integrating to of each wing to find the scatter."
 The best fit results for the runs with two SSPs are found in Table 2.., The best fit results for the runs with two SSPs are found in Table \ref{tab:results}.
 Here we only give the results of the runs with two SSPs. as the results for the constant star formation rate and | SSP runs agree very well with those for two SSPs.," Here we only give the results of the runs with two SSPs, as the results for the constant star formation rate and 1 SSP runs agree very well with those for two SSPs."
 We see that most objects are fitted well. with reduced (P<3.," We see that most objects are fitted well, with reduced $\chi^2 \lesssim 3$."
 Objects with worse X? can be explained by features in the SEDs. from emission lines or too small error bars on the photometry due to systematic errors.," Objects with worse $\chi^2$ can be explained by features in the SEDs, from emission lines or too small error bars on the photometry due to systematic errors."
 From the U band photometry. probing restframe ~16002150 at +=1. we can calculate the UV-derived star formation rates of the galaxies.," From the U band photometry, probing restframe $\sim 1600-2150$ at $z = 1$, we can calculate the UV-derived star formation rates of the galaxies."
 Using the conversion of Kennicutt (1998). and extrapolating the flux in the U band to restframe 1500 assuming a flat continuum in η. the derived SFRs range from 2.21 M.. +. with a media value of 6.0 M... yr.1.," Using the conversion of Kennicutt (1998), and extrapolating the flux in the U band to restframe $1500$ assuming a flat continuum in $f_{\nu}$, the derived SFRs range from $2 - 21$ $_\odot$ $^{-1}$, with a median value of $6.0$ $_\odot$ $^{-1}$."
 These are the SFRs uncorrected for dust. which affects the UV light strongly.," These are the SFRs uncorrected for dust, which affects the UV light strongly."
 However. as the exact dust attenuation has been determined with the SED fitting. these car be corrected for the dust absorption using the same Calzetti et al. (," However, as the exact dust attenuation has been determined with the SED fitting, these can be corrected for the dust absorption using the same Calzetti et al. ("
2000) law as in the actual SED fitting.,2000) law as in the actual SED fitting.
 After correctior of the dust absorption. the dust corrected SERs range from 5. GEM. ο. with a median SFR of 25.4 M. |l," After correction of the dust absorption, the dust corrected SFRs range from $5-64$ $_\odot$ $^{-1}$, with a median SFR of $25.4$ $_\odot$ $^{-1}$."
 The results presented here come from fitting the SEDs of high redshift galaxies with. theoretical templates., The results presented here come from fitting the SEDs of high redshift galaxies with theoretical templates.
 This fitting is hampered by a number of systematic uncertainties., This fitting is hampered by a number of systematic uncertainties.
 Firstly. on the observational side. the inclusion of the grism spectrum in the photometric points helps constraining the fits very well. although for object LBG_449138 and LBG_553929. the strong [OIL] and [OH] emission lines had to be removed manually," Firstly, on the observational side, the inclusion of the grism spectrum in the photometric points helps constraining the fits very well, although for object 49138 and 53929, the strong [OIII] and [OII] emission lines had to be removed manually"
by a filament and also to the edges of a plasma sheet.,by a filament and also to the edges of a plasma sheet.
 The ellect. as seen on a screen. of refraction in a filament is an almost clark line with a bright region on either side. with an extended halo.," The effect, as seen on a screen, of refraction in a filament is an almost dark line with a bright region on either side, with an extended halo."
 Phe dark. patch is behind the region of the lens in which there is sullicient eracient to divert the rays out of the line of sight., The dark patch is behind the region of the lens in which there is sufficient gradient to divert the rays out of the line of sight.
 Lf the source is pulsed. the halo will show a delay which increases with angular distance.," If the source is pulsed, the halo will show a delay which increases with angular distance."
 The observed echoes conform well to this pattern., The observed echoes conform well to this pattern.
 As the lens approaches the line of sight to the pulsar. the observer is traversing the screen.," As the lens approaches the line of sight to the pulsar, the observer is traversing the screen."
 At the edge of the halo a faint echo is seen at maximum delay., At the edge of the halo a faint echo is seen at maximum delay.
 As the delay decreases to zero the observer crosses the cusp. with enhanced intensity.," As the delay decreases to zero the observer crosses the cusp, with enhanced intensity."
 The main pulse then disappears behind the lens (at day 740)., The main pulse then disappears behind the lens (at day 740).
 When the main pulse reappears LO days later it builds up to normal intensity. there is an increase in dispersion measure. and no clear echo appears.," When the main pulse reappears 10 days later it builds up to normal intensity, there is an increase in dispersion measure, and no clear echo appears."
 This indicates. that the lens effect is at the edge of a plasma sheet. in which there is a comparatively low gradient. of electron. content.," This indicates that the lens effect is at the edge of a plasma sheet, in which there is a comparatively low gradient of electron content."
 A departing echo is observed. starting LOO days later. at à time of enhanced intensity. ancl coinciding with a steep fall in dispersion measure.," A departing echo is observed starting 100 days later, at a time of enhanced intensity, and coinciding with a steep fall in dispersion measure."
 This is the other edge of the plasma sheet. although it is less well defined as can be seen in Figure 2bhb. The scale on the screen can be found rom the duration of the event. given only the transverse velocity of the edge across the line of sight.," This is the other edge of the plasma sheet, although it is less well defined as can be seen in Figure \ref{dispersion}b b. The scale on the screen can be found from the duration of the event, given only the transverse velocity of the edge across the line of sight."
 Several examples of echoes have been observed at various times. and all have similar durations and slopes.," Several examples of echoes have been observed at various times, and all have similar durations and slopes."
 We suggest that the transverse velocity is almost entirely due to the proper motion of the pulsar. which is observed. to be 120 km (kaplan et al2008).. with no significant contribution from the transverse velocity. of the sereen itself. and we assume initially that the edge is traversed orthogonally.," We suggest that the transverse velocity is almost entirely due to the proper motion of the pulsar, which is observed to be 120 km $^{-1}$ (Kaplan et al, with no significant contribution from the transverse velocity of the screen itself, and we assume initially that the edge is traversed orthogonally."
" Phe shadow region extended for 10 davs. Le. 9«10"" seconds. giving a scale of 1.010 nm. The significant parameters of the lens are its integrate electron. content IN. along a line of sight. and its gracdien CUNdae across the line of sieht."," The shadow region extended for 10 days, i.e. $9\times 10^5$ seconds, giving a scale of $1.0\times 10^{11}$ m. The significant parameters of the lens are its integrated electron content $N$ along a line of sight, and its gradient ${\rm d}N/{\rm
d}x$ across the line of sight."
 The overall increase of DAL is 0.10 em pe. giving à maximum electron conten Naas:=30JlOσας m02 7.," The overall increase of DM is 0.10 $^{-3}$ pc, giving a maximum electron content $N_{\rm max}=3\times
10^{21}$ $^{-2}$ ."
 Averaged over the shadow region.. the ateral gradient «Nd was 3101 7.," Averaged over the shadow region, the lateral gradient ${\rm d}N/{\rm d}x$ was $3\times 10^{10}$ $^{-3}$."
 Since the echo was observed. over a range of angles we assume that the clouc contains structure with larecr gradients which gave rise to he echoes at the maximum delay., Since the echo was observed over a range of angles we assume that the cloud contains structure with larger gradients which gave rise to the echoes at the maximum delay.
 We explore this by finding he angular deviation corresponding to this averaged overal eracient. and comparing with an independent estimate from he geometry of the observed. delays.," We explore this by finding the angular deviation corresponding to this averaged overall gradient, and comparing with an independent estimate from the geometry of the observed delays."
 The angular deviation in a gracient cVcr is fou rom the phase change in traversing a plasma with tota content IN., The angular deviation in a gradient ${\rm d}N/{\rm d}x$ is found from the phase change in traversing a plasma with total content $N$.
 The phase advance is Aru. where re ds the classical electron. radius οτιο}.=2.81015 m. qsThe angular ceviation 6=A7/(2z)reNiele. where A=0.5 m is the observing wavelength.," The phase advance is $\lambda r_{\rm e} N$, where $r_{\rm e}$ is the classical electron radius $e^2/(m_{\rm e}c^2) = 2.8 \times 10^{-15}$ m. The angular deviation $\theta =\lambda^2/(2\pi)r_{\rm e}{\rm d}N/{\rm
d}x$, where $\lambda=0.5$ m is the observing wavelength."
" For the gradient found above the average deviation 6,—3.2«10.""[E rad.", For the gradient found above the average deviation $\theta_{\rm av}=3.2\times 10^{-6}$ rad.
 The echo was first observed 50 days before the delay reduced to zero., The echo was first observed 50 days before the delay reduced to zero.
 A a velocity of 120 km + this is a lateral distance AOL52010H m. Using the average deviation angle the distance 2=A/G...1.610111T m= 5 pe. about .2.5 times. 1e radius of the nebula: a Larger local eraclient of electron content would give a larger deviation 6 and reasonably place 1e cloud. within the nebula.," At a velocity of 120 km $^{-1}$ this is a lateral distance $\Delta=5.2\times 10^{11}$ m. Using the average deviation angle the distance $R=\Delta/\theta_{\rm av}= 1.6\times 10^{17}$ m $\simeq 5$ pc, about 2.5 times the radius of the nebula; a larger local gradient of electron content would give a larger deviation $\theta$ and reasonably place the cloud within the nebula."
 The echo delay 9 is related to 2 and 6 by 6=SFR.," The echo delay $\delta$ is related to $R$ and $\theta$ by $\delta = {1
\over 2c}\theta^2R$."
" lor vw observed maximum cdelay. (5 milliseconds) Chis vields a maximum deviation angle 8,—43.10"". about 1.5 times 1f angle expected. from. the average eradient of. electron content."," For the observed maximum delay (5 milliseconds) this yields a maximum deviation angle $\theta_{\rm max} = 4.3 \times
10^{-6}$, about 1.5 times the angle expected from the average gradient of electron content."
 The two independent estimates of 6 are seen to be reasonably consistent., The two independent estimates of $\theta$ are seen to be reasonably consistent.
 The smaller angles of deviation involved. as the echo approaches zero delay are derived. from. parts of the edge with lower gradient., The smaller angles of deviation involved as the echo approaches zero delay are derived from parts of the edge with lower gradient.
 As expected. the echo delay. follows a parabolic curve. since it is proportional to 2X7," As expected, the echo delay follows a parabolic curve, since it is proportional to $\Delta^2$."
 The distance £ derived above places the electron cloud outside the nebula., The distance $R$ derived above places the electron cloud outside the nebula.
 As noted already. it would be reduced if the cloud contained gradients larger than the average: it would also be reduced. if the transverse velocity were reduced below 120 kms +. as would be the case if the edge were traversed at a non-orthogonal angle.," As noted already, it would be reduced if the cloud contained gradients larger than the average: it would also be reduced if the transverse velocity were reduced below 120 km $^{-1}$, as would be the case if the edge were traversed at a non-orthogonal angle."
 Our analysis is consistent with an electron cloud or filament located within the outer part of the nebula. and with a gradient ofelectron content given by the calculated average value.," Our analysis is consistent with an electron cloud or filament located within the outer part of the nebula, and with a gradient of electron content given by the calculated average value."
 The clement of delay. amounting to 1.2 ms. shown in Figure 2 to be independent of frequency. must. be related. to rav paths in the electron. cloud.," The element of delay, amounting to 1.2 ms, shown in Figure \ref{dispersion} to be independent of frequency, must be related to ray paths in the electron cloud."
 A lateral gradient. of electron content evidently exists in the region traversed by the rays for some davs after the event at day 747. and this can account for a geometric delay similar to that observed. in the echoes.," A lateral gradient of electron content evidently exists in the region traversed by the rays for some days after the event at day 747, and this can account for a geometric delay similar to that observed in the echoes."
 For this delay to be non-dispersive. however. the refracting region must be localised. so that rays at the three radio frequencies traverse nearly the same geometric path.," For this delay to be non-dispersive, however, the refracting region must be localised, so that rays at the three radio frequencies traverse nearly the same geometric path."
 We suggest that this path is close to the region where there gradient reduces to zero and the profile of electron content becomes Lat. as in the idealised model of Figure. 3..," We suggest that this path is close to the region where there gradient reduces to zero and the profile of electron content becomes flat, as in the idealised model of Figure \ref{sketch}."
 Almost 40 vears of recordings exist of the radio pulses from the Crab pulsar., Almost 40 years of recordings exist of the radio pulses from the Crab pulsar.
 Only three other events comparable to the 19907 event have been recorded. in 1974. 1992 and 1994: echoes were observed in all three. although in the 1974 event. which was the largest. there was insullicient. resolution of the echoes to allow a measurement of their delavs.," Only three other events comparable to the 1997 event have been recorded, in 1974, 1992 and 1994; echoes were observed in all three, although in the 1974 event, which was the largest, there was insufficient resolution of the echoes to allow a measurement of their delays."
 I0 was nevertheless remarked by Lyne and Thorne that the pulse profiles shown contained discrete components. which we now interpret as echoes: and a re-examination of their published: profiles shows that all three components of the pulse (precursor. main. and interpulse) were followed: by echoes.," It was nevertheless remarked by Lyne and Thorne that the pulse profiles shown contained discrete components, which we now interpret as echoes; and a re-examination of their published profiles shows that all three components of the pulse (precursor, main, and interpulse) were followed by echoes."
 The pulse intensity was observed to decrease to near zero in this event. as occurred in 1997.," The pulse intensity was observed to decrease to near zero in this event, as occurred in 1997."
 No other comparable event has been found in a close serutiny of more recent recordings (from 1997 to 2009)., No other comparable event has been found in a close scrutiny of more recent recordings (from 1997 to 2009).
 Pheecho delays. ancl their rates of change. were similar to those in the 1997 event.," Theecho delays, and their rates of change, were similar to those in the 1997 event."
 In our interpretation. this allows us to assume that the edge of," In our interpretation, this allows us to assume that the edge of"
el 22008).,et 2008).
 Note that. while Galloway et ((2008) identified 20 more candidate non-PRE bursts Chat satisfied their trigger criteria. spectral analyses of those bursts revealed. (hat the distinclive cooling associated with Type I bursts did not occur in these events. strongly sugeesting that (μον are Tvpe II instead (Lewin et 11993).," Note that, while Galloway et (2008) identified 20 more candidate non-PRE bursts that satisfied their trigger criteria, spectral analyses of those bursts revealed that the distinctive cooling associated with Type I bursts did not occur in these events, strongly suggesting that they are Type II instead (Lewin et 1993)."
 In order to analvze the PRE bursts. we extracted time resolved 2.5—25.0 keV N-rav spectra using the ftoolseerfref and included the data from all the RNTE/PCA lavers.," In order to analyze the PRE bursts, we extracted time resolved $-$ 25.0 keV X-ray spectra using the ftool and included the data from all the RXTE/PCA layers."
 We used the Science Event mode data with the 1125/5. 118 configuration. which has a nominal time resolution of 125 js in 64 spectral channels.," We used the Science Event mode data with the $\mu$ 1s configuration, which has a nominal time resolution of 125 $\mu$ s in 64 spectral channels."
 We binned the N-rav. spectra in 27 spectral channels and over 0.25 s (for count rates above 6000 cls 1) and over 0.5m 5 (for count rates between 3000 and 6000. ct Lj time intervals during each burst., We binned the X-ray spectra in 27 spectral channels and over 0.25 s (for count rates above 6000 ct $^{-1}$ ) and over 0.5 s (for count rates between 3000 and 6000 ct $^{-1}$ ) time intervals during each burst.
 Following Galloway el ((2008). we extracted a 16 5 spectrum prior to the burst ancl used it as background.," Following Galloway et (2008), we extracted a 16 s spectrum prior to the burst and used it as background."
5 We 5generated separate response matrix files for each burst using5 the PCARSP version 10.1 and took into account the offset pointing of the PCA curing the creation of the response matrix files., We generated separate response matrix files for each burst using the PCARSP version 10.1 and took into account the offset pointing of the PCA during the creation of the response matrix files.
" We fit (he extracted spectra with a blackbody function. using the hydrogen column density value of Ny,=L4x107 ? determined by Wijnands et ((2005) [rom Chandra observations."," We fit the extracted spectra with a blackbody function, using the hydrogen column density value of $_{\rm H} = 1.4 \times
10^{22}$ $^{-2}$ determined by Wijnands et (2005) from Chandra observations."
 We used NSPEC v12 (Arnaud 1996) lor our spectral analvsis., We used XSPEC v12 (Arnaud 1996) for our spectral analysis.
 For each spectrum. we calculated bolometric fIuxes using equation (3) of Galloway et ((2008).," For each spectrum, we calculated bolometric fluxes using equation (3) of Galloway et (2008)."
 Figure 1 shows an example countrate spectrum as well as the best [it blackbody model., Figure 1 shows an example countrate spectrum as well as the best fit blackbody model.
 There are no svslenmatic residuals in the fit. and the addition of anv other speteral components (e.g... a power-law model) is not statistically significant.," There are no systematic residuals in the fit, and the addition of any other spetcral components (e.g., a power-law model) is not statistically significant."
 In Figure 2. we show the distribution of the \?/eLo.f," In Figure 2, we show the distribution of the $\chi^2$ /d.o.f."
 values that we obtained by filling the N-rav. spectra of the source during the 2 PRE bursts ancl compare it to the expected distribution [or 25 degrees of freedom., values that we obtained by fitting the X-ray spectra of the source during the 2 PRE bursts and compare it to the expected distribution for 25 degrees of freedom.
 All fits with \7/d.o.f.<1.5 follow the expected distribution and are. therefore. statistically acceptable.," All fits with $\chi^2/{\rm d.o.f.} < 1.5$ follow the expected distribution and are, therefore, statistically acceptable."
 HLowever. the five spectra wilh 2/d.o.[.>1.5 are outliers. which are likely to be dominated by svstematic uncertainties.," However, the five spectra with $\chi^2/{\rm d.o.f.} >
1.5$ are outliers, which are likely to be dominated by systematic uncertainties."
 We rejected these fits [rom the subsequent analyses., We rejected these fits from the subsequent analyses.
 We show in Figure 3 the bolometric flux. the blackbody temperature. and (he normalization of the model spectra during the evolution of the two PRE bursts.," We show in Figure 3 the bolometric flux, the blackbody temperature, and the normalization of the model spectra during the evolution of the two PRE bursts."
 The characteristic decrease of the temperature and (he increase of the photospheric radius around the burst peak. as well as the cooling of the burst emission at a constant photospheric radius lor both bursts can be seen in both bursts.," The characteristic decrease of the temperature and the increase of the photospheric radius around the burst peak, as well as the cooling of the burst emission at a constant photospheric radius for both bursts can be seen in both bursts."
 In PRE bursts. the Eddington limit at the surface of the neutron star is thought to correspond to the point in each burst when the normalization of the blackbocly gets its lowest value while the temperature reaches its highest (Damen οἱ 11990).," In PRE bursts, the Eddington limit at the surface of the neutron star is thought to correspond to the point in each burst when the normalization of the blackbody gets its lowest value while the temperature reaches its highest (Damen et 1990)."
 The spectral, The spectral
"value of CI,",value of $CI_p$.
" An exponential surface briehtuess profile jas an index of 0.501 aud an rt ber profile results iu a CT, value of 0.371.", An exponential surface brightness profile has an index of 0.501 and an $r^{1/4}$ law profile results in a $CI_p$ value of 0.371.
 Again we find the ]xilge-dominated SU/a aud Sa galaxies to have light profiles that are close o the value expected for à pire kr! law. whereas the nreenlars aud disk-donminated spirals have couceutrations iat are closer to that predicted for à pure exponenutial xofile.," Again we find the bulge-dominated S0/a and Sa galaxies to have light profiles that are close to the value expected for a pure $r^{1/4}$ law, whereas the irregulars and disk-dominated spirals have concentrations that are closer to that predicted for a pure exponential profile."
" The relation|)etweecn AR-baud CT, gaindex aud Z- ype is found to give a Spearman p of 0.08. indicating strong correlation."," The relation between $R$ -band $CI_p$ index and $T$ -type is found to give a Spearman $\rho$ of 0.98, again indicating strong correlation."
 The significance of deviations from the vest-fit Tear trend is A similar analysis is presented Fie., The significance of deviations from the best-fit linear trend is A similar analysis is presented Fig.
 10 in ? which VArows the Petrosian couceutration nmidiees for 126 SDSS ealaxies against their T-type classification., 10 in \citet{shim01} which shows the Petrosian concentration indices for 426 SDSS galaxies against their T-type classification.
 The values for WaGs ealaxics appear to be systematically higher wn those for the SDSS galaxies., The values for the GS galaxies appear to be systematically higher than those for the SDSS galaxies.
 The imajoritv of SDSS Sa-lhn galaxies have iudices |votween 0.35 aud 0.50 whereas the equivalent range for CS galaxies is between 1E anc 0.57.," The majority of SDSS Sa-Im galaxies have indices between 0.35 and 0.50, whereas the equivalent range for GS galaxies is between 0.44 and 0.57."
 notedThe reasons for this offset are uot clear. but it should]© that the CS sample contains a large xoportion of low lhnuuimositv and low surface brightuess ealaxies. particularly amonest the late types. which may )o less fully represeuted in other samples.," The reasons for this offset are not clear, but it should be noted that the GS sample contains a large proportion of low luminosity and low surface brightness galaxies, particularly amongst the late types, which may be less fully represented in other samples."
 The PPetrvosian concentration iudex again las a less clean rend as a function of galaxy type than the R-baud index. but the Spearman rank test still indicates a strong correlation. p= 0.90.," The Petrosian concentration index again has a less clean trend as a function of galaxy type than the $R$ -band index, but the Spearman rank test still indicates a strong correlation, $\rho =$ 0.90."
 The significance of deviations frou he best-fit linear rend is, The significance of deviations from the best-fit linear trend is.
 For most types. the coutimmun cussion appears iore ceutrally concentrated han the eenissiou.," For most types, the continuum emission appears more centrally concentrated than the emission."
 Sab and Sh ealaxies are the exception aud appear to have statistically similar light distributions iu R-baud aud eenission., Sab and Sb galaxies are the exception and appear to have statistically similar light distributions in $R$ -band and emission.
 Overall we couclude that the C30 nudex seenis mareinally the best proxy for morphological type of the three investigated here., Overall we conclude that the C30 index seems marginally the best proxy for morphological type of the three investigated here.
 We also carried out an analysis of the effect of bars ou the mean concentration indices., We also carried out an analysis of the effect of bars on the mean concentration indices.
 Of the three iudices discussed above. the C30 iudex was preferred for this analysis as a result of the small scatter about mean values found above for this index.," Of the three indices discussed above, the C30 index was preferred for this analysis as a result of the small scatter about mean values found above for this index."
 Figure 2. shows the difference iu the mean value of this iudex between uubarred or weakly xured and strongly barred salaxies (ies A aud AB vs D classifications).," Figure \ref{fig:dc30_v_t} shows the difference in the mean value of this index between unbarred or weakly barred and strongly barred galaxies (i.e., A and AB vs B classifications)."
 The left-hand plot shows the effect of vars on the We-derived imdex. whereas the vielt-hand alot is for Π-ἱ wand cussion.," The left-hand plot shows the effect of bars on the -derived index, whereas the right-hand plot is for $R$ -band emission."
 The error bars show the standard error on the mean value of t1e difference for cach vpe., The error bars show the standard error on the mean value of the difference for each type.
 For points Wing above the line. the wnbarred/weakly xured galaxies have more centrally concentrated eimissiou hau do their strouglv barred counterparts.," For points lying above the line, the unbarred/weakly barred galaxies have more centrally concentrated emission than do their strongly barred counterparts."
 For the R- wand plot. this appears to be the case for most spiral vpes: all the poiuts for T.= 0 - 7 have positive values. indicating that stroug bars are associated with less ceutrally concentrated R-baud hWmuuimositv.," For the $R$ -band plot, this appears to be the case for most spiral types: all the points for $T=$ 0 - 7 have positive values, indicating that strong bars are associated with less centrally concentrated $R$ -band luminosity."
 For he later vpes (P= 8 - 10) anv effect of bars secuis o be in he opposite seuse. with the R- baud light !cing somewhat nore ceutrallv couceutrated than for those galaxies with ιο Or Weak bars," For the later types $T=$ 8 - 10) any effect of bars seems to be in the opposite sense, with the $R$ -band light being somewhat more centrally concentrated than for those galaxies with no or weak bars."
 The relation between bar effect aud ealaxy T-type was investigated using he two statistical ests introduced. above., The relation between bar effect and galaxy $T$ -type was investigated using the two statistical tests introduced above.
 The Spearman test indicates a significant correlation. with the p value being 0.82.," The Spearman test indicates a significant correlation, with the $\rho$ value being 0.82."
 The sjenificauce of deviations from the ]vost-fit linear trend isC, The significance of deviations from the best-fit linear trend is.
s Similar trends are seen in the left-hand plo of Fie., Similar trends are seen in the left-hand plot of Fig.
 2 for the distributions of lheht., \ref{fig:dc30_v_t} for the distributions of light.
 For al T-types. the scatter is larger. but again the earlier type spirals show less central coucentration in the barred galaxies. with the reverse being true for later types.," For all $T$ -types, the scatter is larger, but again the earlier type spirals show less central concentration in the barred galaxies, with the reverse being true for later types."
 Iu this case the transition takes place between types 5Γ aud (Scand Sed). although eiven the size of the error bars the differences for anv individual type are not siguificaut.," In this case the transition takes place between types 5 and 6 (Sc and Scd), although given the size of the error bars the differences for any individual type are not significant."
 The Spearman test indicates that this correlation is siguificaut (p =0.83)., The Spearman test indicates that this correlation is significant $\rho=$ 0.83).
 The significance of deviations from the best-fit linear trend is38%., The significance of deviations from the best-fit linear trend is.
.. The main result from this initial analvsis of barred ealaxies is the lower central concentration of both old aud forming stars m earlier-tvpe barred spiral galaxies than uon-barred., The main result from this initial analysis of barred galaxies is the lower central concentration of both old and forming stars in earlier-type barred spiral galaxies than non-barred.
 This effect will J)e studied in more detail iu Sect. L.., This effect will be studied in more detail in Sect. \ref{sec:bars}.
 The concentration indices considered iu the previous section convey onlv a simall fraction of the information on the spatial distribution of luminosity that is coutained iu our images., The concentration indices considered in the previous section convey only a small fraction of the information on the spatial distribution of luminosity that is contained in our images.
 Iu this section. we will present an analysis that attempts to extract more of that information. whilst still chabling mcanineful averages to be taken. thus minimising," In this section, we will present an analysis that attempts to extract more of that information, whilst still enabling meaningful averages to be taken, thus minimising"
helt curves aud spectra were extracted froii the cleaned photon list using CIAO V. 2.2.1 threads. which were also used for the generation of the relative response matrices.,"light curves and spectra were extracted from the cleaned photon list using CIAO V. 2.2.1 threads, which were also used for the generation of the relative response matrices."
 Spectral analysis was performed i in the same way as for the sspectra., Spectral analysis was performed in in the same way as for the spectra.
 Fie., Fig.
 l shows the images of the V892 Tau svsteii in the AACTS camera aud EEPIC-pn camera (before aud after the laree flare) as well as the Palomar Digital Sky Survey image of the field., \ref{fig:image} shows the images of the V892 Tau system in the ACIS camera and EPIC-pn camera (before and after the large flare) as well as the Palomar Digital Sky Survey image of the field.
 The L images are on the same sky coordinate scale., The 4 images are on the same sky coordinate scale.
 The source at the centre ofthe images is the V892 Tau svsten., The source at the centre of the images is the V892 Tau system.
 Vso2 Tan and V892 Tau NE are clearly resolved in the oobservatious., V892 Tau and V892 Tau NE are clearly resolved in the observations.
 The ppoiut spread fiction has a full width at half λα of 15 axesec iud therefore caunot resolve the system., The point spread function has a full width at half maximum of $15$ arcsec and therefore cannot resolve the system.
 The source at 38 arcsec to the SE of the V892 Tau system is [BUS9) MIIO 11. a T Tauri star (Bricenoctal. 1998)). first ideutified in ROSAT data by Strom&(1991).," The source at $38$ arcsec to the SE of the V892 Tau system is [BHS98] MHO 11, a T Tauri star \citealp{bhs+98}) ), first identified in ROSAT data by \cite{ss94}."
. The bright source iu the North-East coruer is ITubble I. a well known T Tawi.," The bright source in the North-East corner is Hubble 4, a well known T Tauri."
 In Table 1 owe report the coordinates of V892 Tau and Vs02 Tan NE as derived from the radio (VLA) observation of Skinneretal.(1993)... to be compared with the coordinates of the sources in the aud: ddata.," In Table \ref{tab:coord} we report the coordinates of V892 Tau and V892 Tau NE as derived from the radio (VLA) observation of \cite{sbs93}, to be compared with the coordinates of the sources in the and data."
 The source coordinates from the EPIC-pu data COTTOSDOLK to the peak of a gaussian distribution fitted to the source image. the source coordinates for the ddata simply correspond to the brightest pixel iu the ποιος nuage.," The source coordinates from the EPIC-pn data correspond to the peak of a gaussian distribution fitted to the source image, the source coordinates for the data simply correspond to the brightest pixel in the source image."
 The aerecinent between the SOULCCS positious in the jiinage and as determined from VLÀ observations is excellent (within 0.1. arcsec)., The agreement between the sources' positions in the image and as determined from VLA observations is excellent (within 0.4 arcsec).
 The source coordinates derived from the oobservatious before the flare have a 2.b arcsec offset from σος Tau (vell within the uncertainty expected for the determination of positions of EPIC X-ray sources) and a 6 aresee offset frou Ws92 Tau NE., The source coordinates derived from the observations before the flare have a 2.4 arcsec offset from V892 Tau (well within the uncertainty expected for the determination of positions of EPIC X-ray sources) and a 6 arcsec offset from V892 Tau NE.
 After the flare the position offsets become 3.1 arcsec and 6.8 arcsec frou V892 Tau and V892 Tau NE respectively., After the flare the position offsets become 3.1 arcsec and 6.8 arcsec from V892 Tau and V892 Tau NE respectively.
 The light curves of V802 Tau and V892 Tau NE durius the — I8 ks oobservation of March 2003 are shown in Fie. 2.., The light curves of V892 Tau and V892 Tau NE during the $\sim$ 18 ks observation of March 2003 are shown in Fig. \ref{fig:lc_chandra}.
 In the ddata. The Herbie Ac star Vao2 Tau is clearly resolved from the less huuinous companion. the WITS V892 Tau NE.," In the data, The Herbig Ae star V892 Tau is clearly resolved from the less luminous companion, the WTTS V892 Tau NE."
 During the oobservation V892 Tau shows significant variability. with its N-vav luminosity varving by a factor of 2 in less than l ks.," During the observation V892 Tau shows significant variability, with its X-ray luminosity varying by a factor of 2 in less than 1 ks."
 The rise in the source bhuniuositv is impulsive (the source doubles its ΕΠ m less than 1 ks) aud it is followed by a slow decay., The rise in the source luminosity is impulsive (the source doubles its luminosity in less than 1 ks) and it is followed by a slow decay.
 A proper time resolved spectral analysis of the rise phase is not possible with the signal to noise ratio of this data set., A proper time resolved spectral analysis of the rise phase is not possible with the signal to noise ratio of this data set.
 We have inspected the light curves of the source in a sof band (0.0 1.7 keV) and in a hard 1.7 8 keV) aud see evidence for a jardenimge of the spectriun during the rie phase., We have inspected the light curves of the source in a soft band $0.0$ $1.7$ keV) and in a hard $1.7$ $8$ keV) and see evidence for a hardening of the spectrum during the rise phase.
 The decay time το1 ber is typical of N-rav stellar flaring events., The decay time $\tau \simeq 1$ hr is typical of X-ray stellar flaring events.
 The bottom pancl of Fig., The bottom panel of Fig.
 2 shows the light curve yAon Vs92 Tan NE., \ref{fig:lc_chandra} shows the light curve of V892 Tau NE.
 The source is weak. on average 8 » 10 times less luminous than VWso2 Tau.," The source is weak, on average 8 to 10 times less luminous than V892 Tau."
 Within the A—wcertainty of the large errorbars the source docs not ypear senificautly variable., Within the uncertainty of the large error–bars the source does not appear significantly variable.
" The probability of coustaucy of V892 Tau NE according o the I&olmogorov-Siiiruov test (which measures the maxiumin deviation of the integral photon arrival times frou, a coustant source model) is .", The probability of constancy of V892 Tau NE according to the Kolmogorov-Smirnov test (which measures the maximum deviation of the integral photon arrival times from a constant source model) is .
. The ACIS spectra of V892 Tau and V892 Tau NE are shown in Fie. 3. , The ACIS spectra of V892 Tau and V892 Tau NE are shown in Fig. \ref{fig:ps_chandra}. .
As sununarised in Table 3 the spectrmn of V892 Tau is well fitted by an absorbed 1-T plasiuna model with ΑΠ=0.8340.08<1Peon2 and temperature AT—2.1040.19 keV. The metallicity of the plasina is not well coustrained by the data.," As summarised in Table \ref{tab:spectra}
 the spectrum of V892 Tau is well fitted by an absorbed 1-T plasma model with $N(H) = 0.83\pm0.08 \times 10^{22}~{\rm cm^{-2}}$ and temperature $kT = 2.10\pm0.19$ keV. The metallicity of the plasma is not well constrained by the data."
 The spectral data for V892 Tau NE have poor signa to nolse ratio. jievertheless a ft with an absorbed 1-T plasina model provide a useful estimate of the absorbing colin density aud the plasima teiiperature.," The spectral data for V892 Tau NE have poor signal to noise ratio, nevertheless a fit with an absorbed 1-T plasma model provide a useful estimate of the absorbing column density and the plasma temperature."
 We derive a value of NIT)7=ye1.2023:0.18reo«&+1077212σαD7. whichH isB similar.- to the value derived for V892 Tau and therefore consistent with the lbypothesis that V892 Tau NE is a plivsical companion of V892 Tau.," We derive a value of $N(H) =
1.20\pm0.18 \times 10^{22}~{\rm cm^{-2}}$, which is similar to the value derived for V892 Tau and therefore consistent with the hypothesis that V892 Tau NE is a physical companion of V892 Tau."
 The derived best-fit value of AT=1.0140.17 keV for the plasima temperature of V892 Tau NE is instead significautly different from the value derived for V892 Tau., The derived best-fit value of $kT = 1.04\pm0.17$ keV for the plasma temperature of V892 Tau NE is instead significantly different from the value derived for V892 Tau.
 The model dependent fluxes (as derived frou the best-fit models) are 7.1«10oper ein the baud 0.55 7.50 keV. and 5.7«10.Here1. oei the baud 0.67. 7.50 keV. respectively for V892 Tau aud V892 Tau NE.," The model dependent fluxes (as derived from the best-fit models) are $7.1 \times 10^{-13}$ in the band $0.55$ $7.50$ keV, and $5.7\times 10^{-14}$, in the band $0.67$ $7.50$ keV, respectively for V892 Tau and V892 Tau NE."
 The light curve of the V892 Tau system derived from the wo consecutive cexposures are shown in Fie. L., The light curve of the V892 Tau system derived from the two consecutive exposures are shown in Fig. \ref{fig:lc_xmm}.
 All the eaps iu the light curve excep the one around 61 ks are due to the filtering xocess. tha we have applied to the raw data in order o remove the effect of solar proton flares., All the gaps in the light curve except the one around 64 ks are due to the filtering process that we have applied to the raw data in order to remove the effect of solar proton flares.
 Iu particular roughlv 10 ks at the beginning of the first exposure have jen removed., In particular roughly 10 ks at the beginning of the first exposure have been removed.
 The gap around 61 ks is due to the time difference between the eud of the firstexposure aud the start of thesecond (Cone hour)., The gap around 64 ks is due to the time difference between the end of the firstexposure and the start of thesecond (one hour).
disk.,disk.
" The visibilities for FU Ori show hardly anv wavelength dependence ancl are very. flat from 3-13,am regardless of baseline.", The visibilities for FU Ori show hardly any wavelength dependence and are very flat from $\mu$ m regardless of baseline.
 Thus. its seems as if the flux distribution of FU Ori is smoother and the visibilities show no sign of a significant contribution from a hot inner rim.," Thus, its seems as if the flux distribution of FU Ori is smoother and the visibilities show no sign of a significant contribution from a hot inner rim."
 However. as FU Ori is surrounded by a heavily active accretion disk where the disk alone produces the majority of the observed fIux at almost all wavelengths (see section 7) differences in the intensity distribution aud thus in the visibilitv can be expected.," However, as FU Ori is surrounded by a heavily active accretion disk where the disk alone produces the majority of the observed flux at almost all wavelengths (see section 7) differences in the intensity distribution and thus in the visibility can be expected."
 Apart from differences in the shape of the visibility curves most ILXeDes are much better resolved. i.e. show lower visibilities than FU Ori (Leinertetal.2004)..," Apart from differences in the shape of the visibility curves most HAeBes are much better resolved, i.e. show lower visibilities than FU Ori \citep{leinert}."
. This. however. can at least partly be explained bv the distance to these objects which is in general less than 200 pc.," This, however, can at least partly be explained by the distance to these objects which is in general less than 200 pc."
 FU Ori on (the other hand has an assumed distance of 450 pe. and if it was closer to the Earth we would observe lower visibilities also for this object.," FU Ori on the other hand has an assumed distance of 450 pc, and if it was closer to the Earth we would observe lower visibilities also for this object."
 To derive a simple model for the geometry of the emitting regions we assume a simple Gaussian brightness distribution for each baseline., To derive a simple model for the geometry of the emitting regions we assume a simple Gaussian brightness distribution for each baseline.
 This is a reasonable first approximation for objects showing high visibilities., This is a reasonable first approximation for objects showing high visibilities.
 The FWIIM. and hence the physical size. of (his Gaussian in arcsec can be computed by where V(/) is the measured visibility for a certain spatial frequency /=2z (in 1) derived from the projected baseline 2 and the wavelength. A.," The FWHM, and hence the physical size, of this Gaussian in arcsec can be computed by where $V(f)$ is the measured visibility for a certain spatial frequency $f=\frac{B}{\lambda}$ (in $^{-1}$ ) derived from the projected baseline $B$ and the wavelength $\lambda$."
 Equation 2. results [rom a simple Fourier (transformation of the assumed. brightness distribution., Equation \ref{visi_eq} results from a simple Fourier transformation of the assumed brightness distribution.
 We computed the FWILIM for all three baselines al three different wavelengths (9.0. 11.0. and. 12.571).," We computed the FWHM for all three baselines at three different wavelengths (9.0, 11.0, and $\mu$ m)."
 For {his we averaged for each baseline !5 visibility points from Table 55. centered on the specified wavelengths., For this we averaged for each baseline 5 visibility points from Table \ref{visibilities_table} centered on the specified wavelengths.
 Table 6 gives the resulting sizes of the emitting regions in AU for an assiuned distance to FU Ori of 450 pe., Table \ref{table_sizes} gives the resulting sizes of the emitting regions in AU for an assumed distance to FU Ori of 450 pc.
 The results are also visualized in Fieure 5 where the EWIIM are shown in their orientation on (he sky., The results are also visualized in Figure \ref{figure_sizes} where the FWHM are shown in their orientation on the sky.
 As expected for thermal disk emission the FWIIM increases wilh wavelength for a given baseline., As expected for thermal disk emission the FWHM increases with wavelength for a given baseline.
 The only exception is the 12.5 size for the UT32-UT4 baseline which is surprisingly a little smaller (han that seen at 11.0jan. In addition to size estimations for each individual baseline our measurements based on different position angles allow to constrain the geometry of the disk., The only exception is the $\mu$ m size for the UT3-UT4 baseline which is surprisingly a little smaller than that seen at $\mu$ m. In addition to size estimations for each individual baseline our measurements based on different position angles allow to constrain the geometry of the disk.
 We fitted the derived FWIIM with an ellipse lor each of the considered wavelengths in order to derive a simple model for the spatial orientation of an assumed disk-like structure., We fitted the derived FWHM with an ellipse for each of the considered wavelengths in order to derive a simple model for the spatial orientation of an assumed disk-like structure.
 The resulting best fit ellipses are overplotted in Figure 5. and (heir parameters are summarized in the lower half of, The resulting best fit ellipses are overplotted in Figure \ref{figure_sizes} and their parameters are summarized in the lower half of
We shall see in the following that these differences do not play a significant role in the diagnosis of microturbulent magnetic fields by means of their Hanle effect in the line.,We shall see in the following that these differences do not play a significant role in the diagnosis of microturbulent magnetic fields by means of their Hanle effect in the line.
 The collisional cross-section obtained by Barklem O° Mara (BOM) ts 3 times larger than the classical Lindholm theory with the Van der Waals approximation: they estimate the relative accuracy of their results at about10%., The collisional cross-section obtained by Barklem O' Mara (BOM) is 3 times larger than the classical Lindholm theory with the Van der Waals approximation; they estimate the relative accuracy of their results at about.
. We have tested the effect on the line intensity and polarization of varying the elastic collision rate. namely we used the two values Γο=3.yvw (BOM standard result) and Γο2.2.yyar. where yyiy denotes the Van der Waals value.," We have tested the effect on the line intensity and polarization of varying the elastic collision rate, namely we used the two values $\Gamma_C= 3. \gamma_{VW}$ (BOM standard result) and $\Gamma_C= 2. \gamma_{VW}$, where $\gamma_{VW}$ denotes the Van der Waals value."
 In both cases the calculations were done with PFR and for the FALC model. the results are shown in Fig. (7)).," In both cases the calculations were done with PFR and for the FALC model, the results are shown in Fig. \ref{elasticrate}) )."
 We see that there are very few differences between the 2 cases. both in the line intensity and polarization profiles.," We see that there are very few differences between the 2 cases, both in the line intensity and polarization profiles."
 The line polarization is slightly smaller for the model with the largest collision rate. which is to be expected. but the effect is weak because. as explained previously. the line is formed in a low density medium where radiative broadening dominates over collisional broadening.," The line polarization is slightly smaller for the model with the largest collision rate, which is to be expected, but the effect is weak because, as explained previously, the line is formed in a low density medium where radiative broadening dominates over collisional broadening."
 We notice that in the line wings. which are not sensitive to the Hanle effect. the polarization obtained with the BOM value for the elastic collision rate is in slightly better agreement with the observed profile.," We notice that in the line wings, which are not sensitive to the Hanle effect, the polarization obtained with the BOM value for the elastic collision rate is in slightly better agreement with the observed profile."
 As far as depolarizing collisions are concerned. recently Derouich (2008) has shown that in the low chromosphere. where the D'' depolarizing collision rate is much smaller than the line radiative transition rate. the depolarization of the P;το level is mainly due to alignment transfer with the Dissτο metastable levels ofBat.," As far as depolarizing collisions are concerned, recently Derouich (2008) has shown that in the low chromosphere, where the $D^{(2)}$ depolarizing collision rate is much smaller than the line radiative transition rate, the depolarization of the $^2P_{3/2}$ level is mainly due to alignment transfer with the $D_{3/2,5/2}$ metastable levels of."
. The corresponding depolarization increases when the hydrogen density increases. because the alignments of the metastable Di;»5;; levels are vulnerable to collisions.," The corresponding depolarization increases when the hydrogen density increases, because the alignments of the metastable $^2D_{3/2,5/2}$ levels are vulnerable to collisions."
 In order to estimate this effect we solved the statistical equilibrium equations for the five Ball levels show1 in Fig., In order to estimate this effect we solved the statistical equilibrium equations for the five BaII levels shown in Fig.
 | including collisions. radiation and magnetic field effects as in Derouich (2008) Le. with a single scattering approximation. and we compared with the results of the equivalent two-level model. for 3 values of the hydrogel density.," 1 including collisions, radiation and magnetic field effects as in Derouich (2008) i.e. with a single scattering approximation, and we compared with the results of the equivalent two-level model, for 3 values of the hydrogen density."
 The results are given in the second column of Table 1. where Ap/Pings 1s the relative depolarization due to alignment transfer and ρω denotes the resonance polarization obtained with an equivalent two-level model.," The results are given in the second column of Table 1, where $\Delta p/p_{max}$ is the relative depolarization due to alignment transfer and $p_{max}$ denotes the resonance polarization obtained with an equivalent two-level model."
 We see that the relative depolarization varies between and40%., We see that the relative depolarization varies between and.
. A physical interpretation of the effect of multi-level coupling on the polarization of the D2 line is given here., A physical interpretation of the effect of multi-level coupling on the polarization of the D2 line is given here.
 By solving the statistical equilibrium equations for the five Ball levels shown in Fig., By solving the statistical equilibrium equations for the five BaII levels shown in Fig.
 | we can estimate the relative importance of the three absorption mechanisms responsible for the atomic polarization of the Ps. upper level., 1 we can estimate the relative importance of the three absorption mechanisms responsible for the atomic polarization of the $P_{3/2}$ upper level.
" We find that of absorptions take place from the fundamental S|;> level. from the D3,. metastable level and from the Ds;» one."," We find that of absorptions take place from the fundamental $S_{1/2}$ level, from the $D_{3/2}$ metastable level and from the $D_{5/2}$ one."
 The polarisability coefficient of those transitions are Ws=0.5. 0.32 and 0.02. respectively.," The polarisability coefficient of those transitions are $W_2=0.5$, 0.32 and 0.02, respectively."
 The Ds;—Pi transition has a low polarisability and it is responsible for an important fraction of the Pi» population., The $D_{5/2} \to P_{3/2}$ transition has a low polarisability and it is responsible for an important fraction of the $P_{3/2}$ population.
 This leads to a decrease of the polarization as compared to the results of the equivalent two-level atom model where the only possible polarizing mechanism ts the absorption of radiation from the fundamental level., This leads to a decrease of the polarization as compared to the results of the equivalent two-level atom model where the only possible polarizing mechanism is the absorption of radiation from the fundamental level.
 Let us now turn to the Hanle effect in the equivalent two-level atom model., Let us now turn to the Hanle effect in the equivalent two-level atom model.
 In Fig. (8)), In Fig. \ref{hanle}) )
" we compare the polarization profiles observed at three distances from the solar limb. namely at 4. 107 and 40"". to those derived from non-LTE radiative transfer modeling with partial frequency redistribution and the Hanle effect. for the atmospheric models FALC and FALX."," we compare the polarization profiles observed at three distances from the solar limb, namely at 4"", 10"" and 40"", to those derived from non-LTE radiative transfer modeling with partial frequency redistribution and the Hanle effect, for the atmospheric models FALC and FALX."
 We first remark that the observations in the line wings. which are not sensitive to the Hanle effect are well recovered in both cases.," We first remark that the observations in the line wings, which are not sensitive to the Hanle effect are well recovered in both cases."
 We recall that the central peak. due to the even Isotopes. Is sensitive to the Hanle effect of a microturbulent magnetic field. whereas the two secondary peaks. due to the odd isotopes. are not (Belluzzi 2007).," We recall that the central peak, due to the even isotopes, is sensitive to the Hanle effect of a microturbulent magnetic field, whereas the two secondary peaks, due to the odd isotopes, are not (Belluzzi 2007)."
 We can thus focus on the interpretation of the polarization observed in the central peak. which ts modeled here.," We can thus focus on the interpretation of the polarization observed in the central peak, which is modeled here."
 We see that it is quite well recovered when we introduce the depolarizing Hanle effect of a turbulent, We see that it is quite well recovered when we introduce the depolarizing Hanle effect of a turbulent
(Macchetto et al.,(Macchetto et al.
 1996. van Dokkum Franx 1995).," 1996, van Dokkum Franx 1995)."
" Statler (2001) argues that the 1.5"" (77 pc) dust ring is dynamically decoupled from the stars and likely has an external origin.", Statler (2001) argues that the $1.5^{\prime\prime}$ (77 pc) dust ring is dynamically decoupled from the stars and likely has an external origin.
 There is also some evidence for a central stellar disk based on the stellar rotation curves (Pastoriza et al., There is also some evidence for a central stellar disk based on the stellar rotation curves (Pastoriza et al.
" 2000) along with à central compact object with a mass of 1.03.9,10°M. (Magorrian et al.", 2000) along with a central compact object with a mass of $1.0-3.9 \times 10^8~ \Mo$ (Magorrian et al.
 1998. Haring Rix 2004).," 1998, Haring Rix 2004)."
 Thus. while the outer regions of NGC 3379 are dynamically relaxed. there are some indications of recent merger activity in the central parts of the galaxy.," Thus, while the outer regions of NGC 3379 are dynamically relaxed, there are some indications of recent merger activity in the central parts of the galaxy."
" In $2 we discuss the general properties of the X-ray point population detected in the Chandra observation of NGC 3379, including the luminosity function. spatial distribution and hardness ratios."," In $\S 2$ we discuss the general properties of the X-ray point population detected in the Chandra observation of NGC 3379, including the luminosity function, spatial distribution and hardness ratios."
 In $3. we discuss the spectral and temporal characteristics of the ULX., In $\S 3$ we discuss the spectral and temporal characteristics of the ULX.
 The properties of the diffuse emission and the central AGN are presented in $4., The properties of the diffuse emission and the central AGN are presented in $\S 4$.
 A discussion concerning the nature of the ULX and the dynamie state of the hot gas in NGC 3379 is given in $5. followed by à brief summary of our main results in $6.," A discussion concerning the nature of the ULX and the dynamic state of the hot gas in NGC 3379 is given in $\S 5$, followed by a brief summary of our main results in $\S 6$."
 NGC 3379 was observed by Chandra for 33.744 s on Feb. 13. 2001 with the ACIS-S detector in faint telemetry mode.," NGC 3379 was observed by Chandra for 33,744 s on Feb. 13, 2001 with the ACIS-S detector in faint telemetry mode."
 The data were reprocessed with the CIAO 3.2. version ofacis_process_events along with CALDB 3.0., The data were reprocessed with the CIAO 3.2 version of along with CALDB 3.0.
 Filtering the S3 data for background flares leaves 31.199 s of cleaned data.," Filtering the S3 data for background flares leaves 31,199 s of cleaned data."
" Since the archived data were processed several years ago. we also followed the thread on the CXC web pages to improve the astrometry. but this only produced a shift of0.12""."," Since the archived data were processed several years ago, we also followed the thread on the CXC web pages to improve the astrometry, but this only produced a shift of $0.12^{\prime\prime}$."
" We then checked for X-ray detections of stars in the USNO-B1.0 catalog and found one star with an optical position off-set by 0.2"" from its X-ray position. which is within the absolute astrometry uncertainties of Chandra and the USNO BI catalog."," We then checked for X-ray detections of stars in the USNO-B1.0 catalog and found one star with an optical position off-set by $0.2^{\prime\prime}$ from its X-ray position, which is within the absolute astrometry uncertainties of Chandra and the USNO B1 catalog."
 An adaptively smoothed 0.3-6.0 keV S3 image of the central 2’ by 2' region (6.2 kpe on a side) is shown in Fig., An adaptively smoothed 0.3-6.0 keV S3 image of the central $2^{\prime}$ by $2^{\prime}$ region (6.2 kpc on a side) is shown in Fig.
 | along with the optical isophotes of the galaxy., 1 along with the optical isophotes of the galaxy.
 This image shows that most of the X-ray emission from NGC 3379 is resolved by Chandra into point sources. many of which are aligned along the major optical axis of the galaxy.," This image shows that most of the X-ray emission from NGC 3379 is resolved by Chandra into point sources, many of which are aligned along the major optical axis of the galaxy."
 The point source labeled AGN in Fig., The point source labeled AGN in Fig.
 | is located within 1” of the optical centroid of the galaxy. and is probably associated with the central supermassive black hole with a dynamically measured mass of 1.03.9«108M. (Magorrian et al.," 1 is located within $1^{\prime\prime}$ of the optical centroid of the galaxy, and is probably associated with the central supermassive black hole with a dynamically measured mass of $1.0-3.9 \times 10^8 \Mo$ (Magorrian et al."
 1998. Haring Rix 2004).," 1998, Haring Rix 2004)."
 The brightest source in NGC 3379 (labeled ULX in Fig., The brightest source in NGC 3379 (labeled ULX in Fig.
 1) is located 7” (360 pe) to the NE of the AGN., 1) is located $7^{\prime\prime}$ (360 pc) to the NE of the AGN.
" The center of NGC 3379 has a high surface density of point sources. as can be seen in the full resolution. raw data image of the central 30"" by 30"" (1.5 kpe) region shown in Fig."," The center of NGC 3379 has a high surface density of point sources, as can be seen in the full resolution, raw data image of the central $30^{\prime\prime}$ by $30^{\prime\prime}$ (1.5 kpc) region shown in Fig."
 2., 2.
 Using a wavelet detection algorithm on the 0.3-6.0 keV image with a detection threshold of 107°. we detect 66 sources within the S3 field of view.," Using a wavelet detection algorithm on the 0.3-6.0 keV image with a detection threshold of $10^{-6}$, we detect 66 sources within the S3 field of view."
 The point source detection sensitivity varies across the field of view due to the presence of extended emission in the center of the galaxy with an average. exposure-corrected 36 threshold of 3.1«1077 et s7!.," The point source detection sensitivity varies across the field of view due to the presence of extended emission in the center of the galaxy with an average, exposure-corrected $3 \sigma$ threshold of $3.1 \times 10^{-4}$ ct $^{-1}$."
 Assuming an absorbed power-law model with galactic absorption (Nj;=2.79«1079 ema?) and index P=1.7. this count rate corresponds to an unabsorbed 0.3-10.0 keV flux of 1.6«.107? eres em s! and luminosity L(0.3-10.0 keV) = 2.2«IO Vergs sv!)," Assuming an absorbed power-law model with galactic absorption $N_H = 2.79 \times 10^{20}$ $^{-2}$ ) and index $\Gamma=1.7$, this count rate corresponds to an unabsorbed 0.3-10.0 keV flux of $1.6 \times 10^{-15}$ ergs $^{-2}$ $^{-1}$ and luminosity L(0.3-10.0 keV) = $2.2 \times 10^{37}$ ergs $^{-1}$."
 Forty sources above this threshold are detected in the central 1.6’ (5 kpc)., Forty sources above this threshold are detected in the central $1.6^{\prime}$ (5 kpc).
 Based on the Chandra deep field south (Giaccont et al., Based on the Chandra deep field south (Giacconi et al.
 2001). approximately one of these sources is likely a background object.," 2001), approximately one of these sources is likely a background object."
 Beyond 5 kpe. approximately 7 of the detected 26 sources are likely background objects.," Beyond 5 kpc, approximately 7 of the detected 26 sources are likely background objects."
 We therefore restrict most analysis to the central 5 kpe of the galaxy., We therefore restrict most analysis to the central 5 kpc of the galaxy.
 Based on our count rate conversion factor. the two brightest sources in NGC 3379 have luminosities of 2.4«10?? ergs s7! and 6.8«I0?* eres s7! The luminosity of the brightest source is consistent with that obtained by Swartz et al. (," Based on our count rate conversion factor, the two brightest sources in NGC 3379 have luminosities of $2.4 \times 10^{39}$ ergs $^{-1}$ and $6.8 \times 10^{38}$ ergs $^{-1}$ The luminosity of the brightest source is consistent with that obtained by Swartz et al. ("
2004). but the lummosity of the second brightest source is a factor of 1.8 less than that given by Swartz et al..,"2004), but the luminosity of the second brightest source is a factor of 1.8 less than that given by Swartz et al.,"
 after correcting for the slightly larger distance used by Swartz et al., after correcting for the slightly larger distance used by Swartz et al.
 of 11.1 Mpe from Ferrarese et al. (, of 11.1 Mpc from Ferrarese et al. (
2000).,2000).
 Swartz et al., Swartz et al.
 derived unabsorbed ray fluxes in the 0.5-8.0 keV band pass by fitting an absorbed power-law model to the spectra of each candidate ULX. treating the absorption and power-law index as free parameters.," derived unabsorbed X-ray fluxes in the 0.5-8.0 keV band pass by fitting an absorbed power-law model to the spectra of each candidate ULX, treating the absorption and power-law index as free parameters."
 They derived best-fit parameters of Nj;=2.2«107! em and P= for the second brightest source in NGC 3379. which accounts for," They derived best-fit parameters of $N_H=2.2 \times 10^{21}$ $^{-2}$ and $\Gamma=2.58$ for the second brightest source in NGC 3379, which accounts for"
The equation of motion. in terms of variables q aud i. becomes Similar to the CARI75 notation. I define the function P(q) as the sum of the independent terms in (he equation of motion (i.e.. the terms that are independent of w and εως). The equation of motion now becomes The problem of existence of a steady solution is thus reduced to determining whether a value of in and a normalized fiction (4) exists such Chat it satisfies (he boundary conditions and equation (6)).,"The equation of motion, in terms of variables $q$ and $\omega$, becomes Similar to the CAK75 notation, I define the function $h(q)$ as the sum of the independent terms in the equation of motion (i.e., the terms that are independent of $\omega$ and $d\omega/dq$ ), The equation of motion now becomes The problem of existence of a steady solution is thus reduced to determining whether a value of $\dot{m}$ and a normalized function $\omega(q)$ exists such that it satisfies the boundary conditions and equation \ref{equ_motion}) )."
 A steady solution for the hvdrodvnamie 1D model exists if and only if equation (6)) is integrable while simultaneously satisfving the boundary conditions., A steady solution for the hydrodynamic 1D model exists if and only if equation \ref{equ_motion}) ) is integrable while simultaneously satisfying the boundary conditions.
 Tvpically the boundary. condition is the position of the sonic point., Typically the boundary condition is the position of the sonic point.
 That is. the equation of motion must not only be inteerable. but. additionally there must exist a solution to (he equation of motion (hat presents a sonic [low speed al a predefined sonic position (boundary condition).," That is, the equation of motion must not only be integrable, but additionally there must exist a solution to the equation of motion that presents a sonic flow speed at a predefined sonic position (boundary condition)."
 To continue the discussion I recall the definition of the ;2 function given in Paper I.," To continue the discussion I recall the definition of the $\beta$ function given in Paper I,"
The currently most-accepted scheme for the structure of an active ealactic nuclei (AGN). the unified. model. consists of a supermassive black hole (SALBID). surrounded. by an accretion dise and. in the same plane at larger clistances. an axisvmametric distribution of gas and dust. usually referred to às à torus.,"The currently most-accepted scheme for the structure of an active galactic nuclei (AGN), the unified model, consists of a supermassive black hole (SMBH), surrounded by an accretion disc and, in the same plane at larger distances, an axisymmetric distribution of gas and dust, usually referred to as a torus."
 Within the inner region of the obscuring torus lies an ensemble of highly ionizecl and. dense gas clouds. where broad. emission. lines are. produced. the broad. line region (1211).," Within the inner region of the obscuring torus lies an ensemble of highly ionized and dense gas clouds, where broad emission lines are produced, the broad line region (BLR)."
 Outside the torus there is a region of low-ionizalion less dense gas clouds. the narrow line region (NLR) (e.g. Xntonucci 1993. Urry Padovani 1995).," Outside the torus there is a region of low-ionization less dense gas clouds, the narrow line region (NLR) (e.g. Antonucci 1993, Urry Padovani 1995)."
 The dust present in the putative obscuring torus blocks most radiation from the inner photoionizing source at optical. UV and soft X-ray wavelengths. but it is responsible for most of the AGN emission at mid-LIlt frequencies. (e.g. Antonucci&Miller 1O85.. Blancoetal.1990... Dopitaetal. 1998.. ‘Tristrametal. 20073).," The dust present in the putative obscuring torus blocks most radiation from the inner photoionizing source at optical, UV and soft X-ray wavelengths, but it is responsible for most of the AGN emission at mid-IR frequencies (e.g. \citealt{1985ApJ...297..621A}, , \citealt{1990MNRAS.242P...4B}, \citealt{1998ApJ...498..570D}, \citealt{2007A&A...474..837T}) )."
 The dust in the torus absorbs and re-radiates the emission from the core at these frequencies., The dust in the torus absorbs and re-radiates the emission from the core at these frequencies.
 Thus. mid-I1t observations. for example those probing rest-frame wavelengths: around yam. indirectly reveal radiation from regions close to the active nucleus.," Thus, mid-IR observations, for example those probing rest-frame wavelengths around $\mu$ m, indirectly reveal radiation from regions close to the active nucleus."
 This is supported. by observations that show the mid- to far-LR continuum emitted by powerful radio galaxies to be mainly due to AGN heating of cireumnuclear dust (eg. Dicken et al., This is supported by observations that show the mid- to far-IR continuum emitted by powerful radio galaxies to be mainly due to AGN heating of circumnuclear dust (e.g. Dicken et al.
 2009)., 2009).
 According to the most widely accepted model for the archetype structure of an AGN. in the aceretion disc. ancl potentially in regions connected. to the event. horizon of the black hole itself magnetized. inllowing material enables processes capable of driving two opposed collimated outllows. known as the jets (e.g. Longairetal.1973.. Blandford&Ixonigl. 1979... Meieretal. 20012).," According to the most widely accepted model for the archetype structure of an AGN, in the accretion disc, and potentially in regions connected to the event horizon of the black hole itself, magnetized inflowing material enables processes capable of driving two opposed collimated outflows, known as the jets (e.g. \citealt{1973MNRAS.164..243L}, \citealt{1979ApJ...232...34B}, \citealt{2001Sci...291...84M}) )."
 These jets push magnetised material out of the central region and feedlobes., These jets push magnetised material out of the central region and feedlobes.
 Lhe acceleration of ultrarelativistie charged. particles, The acceleration of ultrarelativistic charged particles
state.,state.
" As in the low-density case, we take initial conditions at t25 Myr from a low-resolution run, which follows gas cooling in a thermally unstable regime from an initial temperature of 2x109 K. We checked a variety of periods for the time-dependent heating rate: 10 Myr (5 Myr high state and 5 Myr low state), 2.5 Myr (0.2 2.3), 2 Myr (1 + 1), 1 Myr (0.2 + 0.8), and 1 Myr (0.5 + 0.5)."," As in the low-density case, we take initial conditions at $t=5$ Myr from a low-resolution run, which follows gas cooling in a thermally unstable regime from an initial temperature of $2\times10^6$ K. We checked a variety of periods for the time-dependent heating rate: 10 Myr (5 Myr high state and 5 Myr low state), 2.5 Myr (0.2 + 2.3), 2 Myr (1 + 1), 1 Myr (0.2 + 0.8), and 1 Myr (0.5 + 0.5)."
" In all the cases, we find results similar to the low-density case in terms of kinetic energy support, with Εμx0.05E7."," In all the cases, we find results similar to the low-density case in terms of kinetic energy support, with $E_k^{max}\approx0.05E_{th}^{max}$."
" The conclusions concerning Mach number, vorticity, and density variance evolution are essentially the same as in the low-density case."," The conclusions concerning Mach number, vorticity, and density variance evolution are essentially the same as in the low-density case."
" As far as the mass fractions are concerned, we find essentially more unstable gas on the average (e.g., fg29% in our experiment with a period of 1 Myr (0.5+0.5), averaged over five periods)."," As far as the mass fractions are concerned, we find essentially more unstable gas on the average (e.g., $f_G\approx29$ in our experiment with a period of 1 Myr (0.5+0.5), averaged over five periods)."
" The reason for this is twofold: first, the cold phase is always present and therefore so is the unstable interphase interface (fy saturates at ~55% in high state and at ο”88% in low state); second, the abrupt transition to the low state effectively populates the unstable regime at the expense of the warm stable phase, completely consuming it until this gas reexpands to join the warm stable phase of the low state that has a lower density than in the high state."," The reason for this is twofold: first, the cold phase is always present and therefore so is the unstable interphase interface $f_H$ saturates at $\sim55$ in high state and at $\sim88$ in low state); second, the abrupt transition to the low state effectively populates the unstable regime at the expense of the warm stable phase, completely consuming it until this gas reexpands to join the warm stable phase of the low state that has a lower density than in the high state."
 The bulk of this thermally unstable gas is then slowly (~0.3 Myr) redistributed between the cold (68%)) and warm (18%)) stable phases as the turbulent relaxation proceeds., The bulk of this thermally unstable gas is then slowly $\sim0.3$ Myr) redistributed between the cold ) and warm ) stable phases as the turbulent relaxation proceeds.
" In contrast to the low-density case, cores of some of the individual cold “clouds” can survive several heating events without being dissolved."," In contrast to the low-density case, cores of some of the individual cold “clouds” can survive several heating events without being dissolved."
" In addition, new cold clouds are forming by dynamic compressions in the disturbed warm phase (AP/P~ 10) following transitions to the high state."," In addition, new cold clouds are forming by dynamic compressions in the disturbed warm phase $\Delta P/P\sim10$ ) following transitions to the high state."
" This mechanism, described by Hennebelle&Pérault(1999),, is inefficient in our low-density simulations."," This mechanism, described by \citet{hennebelle.99}, is inefficient in our low-density simulations."
 Our model is limited in many respects., Our model is limited in many respects.
" First, it does not include effects of other driving mechanisms that can operate in parallel to thermal forcing."," First, it does not include effects of other driving mechanisms that can operate in parallel to thermal forcing."
" Second, by not including magnetic fields in the simulations, our description of ISM dynamics is incomplete."," Second, by not including magnetic fields in the simulations, our description of ISM dynamics is incomplete."
" Third, the model does not have self-regulation, linking production of the cold phase to the heating rate via star formation efficiency (Parravano1988)."," Third, the model does not have self-regulation, linking production of the cold phase to the heating rate via star formation efficiency \citep{parravano88}."
". Feedback could probably damp the oscillations occurring in response to the time variations of the heating rate introduced “by hand,” but this should be demonstrated by an explicit calculation."," Feedback could probably damp the oscillations occurring in response to the time variations of the heating rate introduced “by hand,” but this should be demonstrated by an explicit calculation."
" Fourth, the grain photoelectric heating rate is not simply proportional to gas density."," Fourth, the grain photoelectric heating rate is not simply proportional to gas density."
 This would slightly modify the shape of the thermal equilibrium curve as it shifts from a low to a high state., This would slightly modify the shape of the thermal equilibrium curve as it shifts from a low to a high state.
" Finally, the equation of state and the cooling function that we use do not take into account dynamics of ionization, recombination, formation of molecules, etc.,"," Finally, the equation of state and the cooling function that we use do not take into account dynamics of ionization, recombination, formation of molecules, etc.,"
 which are decoupled from hydrodynamics in our treatment., which are decoupled from hydrodynamics in our treatment.
 This could overestimate the minimum temperature of the cold phase and thus underestimate its density and the rms Mach number., This could overestimate the minimum temperature of the cold phase and thus underestimate its density and the rms Mach number.
" Nonetheless, our main conclusion that time-dependent heating supports turbulence in thermally unstable multiphase ISM remains valid in spite of the above-mentioned model simplifications."," Nonetheless, our main conclusion that time-dependent heating supports turbulence in thermally unstable multiphase ISM remains valid in spite of the above-mentioned model simplifications."
 Thermal forcing can potentially excite turbulent oscillations on a wide range of length scales because of the penetrative nature of FUV radiation., Thermal forcing can potentially excite turbulent oscillations on a wide range of length scales because of the penetrative nature of FUV radiation.
 The level of thermally sustained turbulence as a function of length scale at a given point within the Galactic disk is controlled by the local mean gas density and by the power spectrum of incident FUV flux time variations., The level of thermally sustained turbulence as a function of length scale at a given point within the Galactic disk is controlled by the local mean gas density and by the power spectrum of incident FUV flux time variations.
 It will share the pattern of spatial inhomogeneity of FUV energy density within the disk., It will share the pattern of spatial inhomogeneity of FUV energy density within the disk.
 The thermally unstable gas mass fractions in our low- and intermediate-density models (fg~1096 —30%)) are in reasonable agreement with measurements by Heiles&Troland as well as the fraction and morphology of the density distribution of the cold neutral phase (see, The thermally unstable gas mass fractions in our low- and intermediate-density models $f_G\sim 10\%-30$ ) are in reasonable agreement with measurements by \citet{heiles.02} as well as the fraction and morphology of the density distribution of the cold neutral phase (see
not allect the overall trend.,not affect the overall trend.
 Ehe increase in the satellite raction as one goes to lower redshifts can be explained by rw dynamical friction. of subhalos within their host. halos (Conroyetal.2006)., The increase in the satellite fraction as one goes to lower redshifts can be explained by the dynamical friction of subhalos within their host halos \citep{con06}.
. Subhalos are more likely to remain --ntact within massive halos. whereas in less massive halos rev are subject to more dynamical friction and can easily »e destroved.," Subhalos are more likely to remain intact within massive halos, whereas in less massive halos they are subject to more dynamical friction and can easily be destroyed."
 The clynanmical (riction becomes more/less ellicient as a function of the relative masses of subhatos o distinct halos., The dynamical friction becomes more/less efficient as a function of the relative masses of subhalos to distinct halos.
 This is to be compared to recent. results obtained by Zheng et al. (, This is to be compared to recent results obtained by Zheng et al. (
2007) who find that the evolution of the satellite fraction follows a trend. similar to what is seen here.,2007) who find that the evolution of the satellite fraction follows a trend similar to what is seen here.
 The following Fig., The following Fig.
 5. shows the evolution in the halo occupation. j/N4(m). for the extreme luminosity threshold samples obtained. from the best fit parameters for the two moclels.," \ref{Fig_hod} shows the evolution in the halo occupation, $N_g(m)$, for the extreme luminosity threshold samples obtained from the best fit parameters for the two models."
" Evidently. the minimum mass. M,,;, increases with the luminosity of the sample as is found locally in the SDSS (Zehavietal.2005).. again demonstrating that luminous galaxies occupy more massive halos."," Evidently, the minimum mass, $M_{min}$ increases with the luminosity of the sample as is found locally in the SDSS \citep{zeh05}, again demonstrating that luminous galaxies occupy more massive halos."
 In this section we will compare to results for the same LOD model (Z model) as used in Zehavi et al. (, In this section we will compare to results for the same HOD model (Z model) as used in Zehavi et al. (
2005).,2005).
" Figure 6 shows the comparison between the masses of halos that have at least one central galaxy. (M,,;,) and one satellite galaxy (Mj) on average as a function of Lin;unf Ls. where the ratio is in the D-band. ancl L;5,;. is the luminosity threshold given in Table Lo (Lin;s) and £, are at similar redshifts)."," Figure \ref{Fig_comparison} shows the comparison between the masses of halos that have at least one central galaxy $M_{min}$ ) and one satellite galaxy $M_1$ ) on average as a function of $L_{thresh}/L_*$ , where the ratio is in the B-band, and $L_{thresh}$ is the luminosity threshold given in Table \ref{table1}
 $L_{thresh}$ and $L_*$ are at similar redshifts)."
 Llere we will try to compare results obtained. at. clilferent redshifts., Here we will try to compare results obtained at different redshifts.
" For 40% of the VVDS samples. AM; is similar to the local SDSS results within the error bars. with the rest of the VVDS samples having higher values of A,,;,."," For $40 \%$ of the VVDS samples, $M_{min}$ is similar to the local SDSS results within the error bars, with the rest of the VVDS samples having higher values of $M_{min}$."
 Generally. the VVDS samples exhibit more massive halo masses. Aly. required to host satellite galaxies than what is seen locally.," Generally, the VVDS samples exhibit more massive halo masses, $M_1$, required to host satellite galaxies than what is seen locally."
 The value for the power law slope. à. is mostly similar to that for local galaxies. with the brightintermediate recshift ealaxies showing a higher slope.," The value for the power law slope, $\alpha$, is mostly similar to that for local galaxies, with the brightintermediate redshift galaxies showing a higher slope."
 We can see that. generally the samples with higher values for a. also have higher values of Aly and Adj; than present-clay galaxies.," We can see that generally the samples with higher values for $\alpha$, also have higher values of $M_1$ and $M_{min}$ than present-day galaxies."
 lt is interesting to note the ιν ratio. which is on average  45. rather high as compared to the value of ~ 23 for the SDSS galaxies.," It is interesting to note the $M_1/M_{min}$ ratio, which is on average $\sim$ 45, rather high as compared to the value of $\sim$ 23 for the SDSS galaxies."
 A direct. comparison. and interpretation of these results is complicated. as one ds looking at two different surveys taken in cillerent restframe bancs., A direct comparison and interpretation of these results is complicated as one is looking at two different surveys taken in different restframe bands.
 However. we can speculate that the high value of the ratio implies that the halo with one central galaxy needs to accrete roughly 45 times its mass in order to host a satellite ealaxv.," However, we can speculate that the high value of the ratio implies that the halo with one central galaxy needs to accrete roughly 45 times its mass in order to host a satellite galaxy."
 In other words. a halo of a given mass is likely to have fewer satellite galaxies at higher redshifts as opposed toa halo of the same mass observed locally.," In other words, a halo of a given mass is likely to have fewer satellite galaxies at higher redshifts as opposed to a halo of the same mass observed locally."
 The comparison of analytical models ancl data: provides useful information of how the distribution of galaxies depends on the underlving dark. matter., The comparison of analytical models and data provides useful information of how the distribution of galaxies depends on the underlying dark matter.
 Subsequently. the best-fitting parameters obtained as a result. of this comparison provide physical information regarding the dark matter halos ancl galaxies.," Subsequently, the best-fitting parameters obtained as a result of this comparison provide physical information regarding the dark matter halos and galaxies."
 The size of the VVDS dataset. allows one to stucly. with a unique sample. the global change in the uncerlying halo propertiesof an average galaxy down to z 1.," The size of the VVDS dataset allows one to study, with a unique sample, the global change in the underlying halo propertiesof an average galaxy down to $z \sim 1$ ."
 We attempt to follow the evolution in some properties of a, We attempt to follow the evolution in some properties of a
index La.,index 1.8.
 The ranee from 2 to 6 keV was used to avoid. ou the one hand the soft range. which is attributed to emission lines. aud the Fe EK-liue on the other.," The range from 2 to 6 keV was used to avoid, on the one hand the soft range, which is attributed to emission lines, and the Fe K-line on the other."
 The ealactic absorption plus anu intrinsic absorptiOl were included in the fitting process., The galactic absorption plus an intrinsic absorption were included in the fitting process.
 The iutrinsic absorption results in a value of δι=~|¢νpy?1077cin27. which is compatible with the foregroundτις extinction interred from the colour maps and the clumipy torus modeling.," The intrinsic absorption results in a value of $\mathrm{N_H} = 7_{-5}^{+6} \times 10^{22}~\mathrm{cm}^{-2}$, which is compatible with the foreground extinction inferred from the colour maps and the clumpy torus modeling."
 Thus. the absorption corrected Nay Duuünosifv is Lower=Loslo!ores1.," Thus, the absorption corrected X-ray luminosity is $\mathrm{L_{2-10~keV}} =1.7 \times 10^{41}~\mathrm{erg\, s^{-1}}$."
 This value is quite low when compared to the optical/imfrared Iuuinositv reprocessed by the blocking torus., This value is quite low when compared to the optical/infrared luminosity reprocessed by the blocking torus.
 However. it cau be reconciled by taking into account that Mrk 573 has beeu classified as a ACN (Cainazziotal. 2005).. ancl its N-rav luminosity has to be corrected hy a large factor to obtain the intrinsic hunünositv. since the column deusity is Ng>1.6«102!cm2 {μιch high value of the optical depth is also consistent with the value derived above frou the chuupy torus uxxdeling).," However, it can be reconciled by taking into account that Mrk 573 has been classified as a AGN \citep{Guainazzi05}, and its X-ray luminosity has to be corrected by a large factor to obtain the intrinsic luminosity, since the column density is $\rm{N_H > 1.6\times 10^{24}~cm^{-2}}$ (such high value of the optical depth is also consistent with the value derived above from the clumpy torus modeling)."
 Panessaetal.(2006). derived. a factor of 60 coupan18o a stuall suuple of Tvpe-2 Sevterts with a suuple of Type-l Sevterts. whereas Cappiotal.(2006) derived a factor about 100.," \citet{Panessa06} derived a factor of 60 comparing a small sample of Type-2 Seyferts with a sample of Type-1 Seyferts, whereas \citet{Cappi06} derived a factor about 100."
 These ναHes are in contrast with the work of Couzález-Martí1ct (2009).. who obtained a value of 12 using a sample of LINERS.," These values are in contrast with the work of \citet{Gonzalez-Martin09}, who obtained a value of 42 using a sample of LINERs."
 In addition. we lave to transform from N-ray to volometiic Inuinositv iiultiplviug by a factor 30 (Risaliti&Elvis 2001: see also Panessaetal. 2006)).," In addition, we have to transform from X-ray to bolometric luminosity multiplying by a factor 30 \citealt{Risaliti04}; see also \citealt{Panessa06}) )."
 Thus. we obtain a value for the volometric ACN huuinositv iu he range Ld.=3.15]10!eresto which is in ποσο agreement with the value derived frou the torus reprocessing. given the 1ucertaiuties mvolved.," Thus, we obtain a value for the bolometric AGN luminosity in the range $\mathrm{L^{bol}_{AGN}=3.1-5.1 \times 10^{44}~erg s^{-1}}$, which is in nice agreement with the value derived from the torus reprocessing, given the uncertainties involved."
 IKraeimeretal.(2009) derived a higrer value for the bolometric uuuinosity (3.2«107ergs ) based ou the |OTV| Iuniuositv. as nieasured by the Spitzer{IRS spectrum of Ak 573 (Moeléndezetal.20082.b).," \citet{Kraemer09} derived a higher value for the bolometric luminosity $3.2\times 10^{45}~erg s^{-1}$ ) based on the [OIV] luminosity, as measured by the Spitzer/IRS spectrum of Mrk 573 \citep{Melendez08a,Melendez08b}."
". In any case Ακ 573 seclus to be radiatiung near to the Eddinet«n limit assuniue the black hole mass to be around 2«OM, (Bian&Ga2007).", In any case Mrk 573 seems to be radiating near to the Eddington limit assuming the black hole mass to be around $2\times 10^7 \mathrm{M_{\odot}}$ \citep{BianGu07}.
. This result adds support t» the re-classification of Mik 573 as a hidden narrow-lc| Sevtert ] CRaniosAlmeidaetal.2008)., This result adds support to the re-classification of Mrk 573 as a hidden narrow-line Seyfert 1 \citep{Ramos08}.
. Απο is a uearby optically classifies Type-2 Sevtert. well-known for its extended. circnunuclear cluission-line reeious.," 573 is a nearby optically classified Type-2 Seyfert, well-known for its extended circumnuclear emission-line regions."
 It is this extension :ux the proximity of the source that convert 5573 as one of the ideal cases to study this cission coninioilv. found πι Tvpe-2 Sevtert galaxies., It is this extension and the proximity of the source that convert 573 as one of the ideal cases to study this emission commonly found in Type-2 Seyfert galaxies.
 We combine and to achieve high spectral aud spatial resohtion in order to diseutaugle the enüssion mechauisin of this extended eiission., We combine and to achieve high spectral and spatial resolution in order to disentangle the emission mechanism of this extended emission.
 We also used optical aud near-IR data in order to compare with the N-ray data., We also used optical and near-IR data in order to compare with the X-ray data.
 The main results are:manuscript: After this paper was submitted to the journal a paper was published by Bianchietal.(2000) with data in conumuon with OUL work., The main results are: After this paper was submitted to the journal a paper was published by \citet{Bianchi10} with data in common with our work.
" Most of their results are consistent with OUTS, although different approaches were used."," Most of their results are consistent with ours, although different approaches were used."
 They fitted twe ROS spectra using Cloudy phlotoiouzation imodols, They fitted the /RGS spectra using Cloudy photoionization models.
 Their best ft was obtained bv au hybrid inodel -photoionization | collisional excitatk1- where the collisional phase contributes 1/3 of the Hux in the band 0.5-0.8 keV. consistent with our work.," Their best fit was obtained by an hybrid model -photoionization + collisional excitation- where the collisional phase contributes 1/3 of the flux in the band 0.5-0.8 keV, consistent with our work."
 Moreover. they claiued the need of two pliotoioization phases (log U = 0.3 aud 1.8) to explain the spectimim in the range 0.1-7 keV. which is also in agreement with our results.," Moreover, they claimed the need of two photoionization phases (log U = 0.3 and 1.8) to explain the spectrum in the range 0.4-7 keV, which is also in agreement with our results."
 The results on the extended emission cannot be compared since they do not «listinguish between the two couc-like structures. extracting the spectrmu from a ciÀmcunmnnuclear anus.," The results on the extended emission cannot be compared since they do not distinguish between the two cone-like structures, extracting the spectrum from a circumnuclear annulus."
a position angle of 68+11 degs.,a position angle of $68\pm11$ degs.
 We note that the extended blue-shifted emission is consistent with the result on Na I D by Rupke et al. (, We note that the extended blue-shifted emission is consistent with the result on Na I D by Rupke et al. (
2005).,2005).
" The determination of the outflow mass rate depends on the wind geometry and on the conversion factor, a, between the CO luminosity and the molecular gas mass."," The determination of the outflow mass rate depends on the wind geometry and on the conversion factor, $\alpha$, between the CO luminosity and the molecular gas mass."
 We estimated the CO luminosity by fitting the observed line profile with a narrow plus a broad component (see Fig., We estimated the CO luminosity by fitting the observed line profile with a narrow plus a broad component (see Fig.
 1)., 1).
" The integrated CO luminosity of the broad component is L(CO)s=1.16x10? K km/s pc?, about 1/10 the luminosity of the narrow component."," The integrated CO luminosity of the broad component is $_B = 1.16\times 10^9$ K km/s $^2$, about 1/10 the luminosity of the narrow component."
" We converted the CO luminosity of the broad component into molecular gas mass M(H) by assuming a conservative conversion factor α= 0.5 Mo (K km s! pc2)-!, i.e. 1/10 the Galactic value."," We converted the CO luminosity of the broad component into molecular gas mass $_2$ ) by assuming a conservative conversion factor $\alpha =$ 0.5 $_\odot$ (K km $^{-1}$ $^2$ $^{-1}$, i.e. 1/10 the Galactic value."
" This is the lowest conversion factor found in different locations of M 82 (a typical starburst galaxy), including its molecular outflow (Weiss et al."," This is the lowest conversion factor found in different locations of M 82 (a typical starburst galaxy), including its molecular outflow (Weiss et al."
 2001)., 2001).
" We derive a mass of the outflowing molecular gas M(H5)=5.8x108 Mo, which is consistent with the lower limit of 7x10’ Mo inferred by Fischer et al. ("," We derive a mass of the outflowing molecular gas $(H_{2 }) = 5.8\times 10^ 8$ $_\odot$, which is consistent with the lower limit of $7\times10^7$ $_\odot$ inferred by Fischer et al. ("
2010) based on the absorption molecular lines detected byHerschel.,2010) based on the absorption molecular lines detected by.
". By assuming that this gas is uniformly distributed in a spherical volume of 0.6 kpc in radius, its inferred density is 25 ?."," By assuming that this gas is uniformly distributed in a spherical volume of 0.6 kpc in radius, its inferred density is $\sim 25$ $^{-3}$."
" Since the outflow velocity is at least 700 km s!, the inferred mass outflow rate is dM(H>)/dt= 2200 Μο ντ""."," Since the outflow velocity is at least 700 km $^{-1}$, the inferred mass outflow rate is $d$ $(H_2)/dt =$ 2200 $_\odot$ $^{-1}$."
" If we assume density profile scaling as r? for the gas distribution, the inferred outflow rate is 710 Mo yr-!."," If we assume density profile scaling as $r^{-2}$ for the gas distribution, the inferred outflow rate is 710 $_\odot$ $^{-1}$."
" An alternative, even more conservative estimate of the outflow rate can be derived by ignoring the Gaussian fit of the broad CO component and by using only the luminosity directly measured from the broad wings, i.e. by integrating their flux at velocities higher than +400 km s! and lower than -400 km s!."," An alternative, even more conservative estimate of the outflow rate can be derived by ignoring the Gaussian fit of the broad CO component and by using only the luminosity directly measured from the broad wings, i.e. by integrating their flux at velocities higher than +400 km $^{-1}$ and lower than -400 km $^{-1}$."
 The luminosity of the wings is L(CO)=3.2x10* K km s! pc? (~30% of the CO luminosity of the total broad component inferred through the gaussian fitting)., The luminosity of the wings is $ = 3.2 \times 10^8$ K km $^{-1}$ $^2$ $\sim$ of the CO luminosity of the total broad component inferred through the gaussian fitting).
" In this case we obtain a lower limit on the outflow rate of 600 Μο yr!, for the uniform gas distribution."," In this case we obtain a lower limit on the outflow rate of 600 $_\odot$ $^{-1}$, for the uniform gas distribution."
 By assuming a ος7? gas density profile we obtain the most conservative lower limit is ~260 Μο ντ]., By assuming a $\propto r^{-2}$ gas density profile we obtain the most conservative lower limit is $\sim$ 260 $_\odot$ $^{-1}$.
" We note that alternative geometries of the outflowing wind, such as shell-like or disk-like configurations, give higher outflow mass rate, both because the inferred outflowing gas density is higher and because the de-projected outflow velocity is higher (e.g. if the molecular outflow occurs on the galaxy disk plane)."," We note that alternative geometries of the outflowing wind, such as shell-like or disk-like configurations, give higher outflow mass rate, both because the inferred outflowing gas density is higher and because the de-projected outflow velocity is higher (e.g. if the molecular outflow occurs on the galaxy disk plane)."
" Similarly, it is easy to show that in the case of a bipolar outflow the inferred outflow rate is the same or may be even higher (the line of sight has to intercept the outflowing molecular gas, since it is seen in absorption withHerschel,, implying that the true, deprojected radius of the bipolar outflow may be equal or larger than the projected size)."," Similarly, it is easy to show that in the case of a bipolar outflow the inferred outflow rate is the same or may be even higher (the line of sight has to intercept the outflowing molecular gas, since it is seen in absorption with, implying that the true, deprojected radius of the bipolar outflow may be equal or larger than the projected size)."
 In any case the inferred outflow rate is much larger than the star-formation rate in the host galaxy of 200 Μο yr! (Taylor et al., In any case the inferred outflow rate is much larger than the star-formation rate in the host galaxy of 200 $_\odot$ $^{-1}$ (Taylor et al.
" 1999, Davies et al."," 1999, Davies et al."
 2004)., 2004).
 The mass loss rate being much larger than the rate at which gas is converted into stars implies a phase of rapid quenching of star formation in the regions reached by the outflow (~1 kpc scale)., The mass loss rate being much larger than the rate at which gas is converted into stars implies a phase of rapid quenching of star formation in the regions reached by the outflow $\sim$ 1 kpc scale).
" The total amount of molecular gas in the galaxy disk as inferred from the integrated emission of the narrow component, by using the CO-to-H» conversion factor oa appropriate for ULIRGs (Solomon Vanden Bout 2005), is M(H5)~1079 "," The total amount of molecular gas in the galaxy disk as inferred from the integrated emission of the narrow component, by using the $_2$ conversion factor $\alpha$ appropriate for ULIRGs (Solomon Vanden Bout 2005), is $(H_2)\sim 10^{10}$ "
We chose the ELAIS N2 region as having a wealth of information available from surveys at radio to X-ray wavelengths (see e.g. Rowan-Robinson et al.,We chose the ELAIS N2 region as having a wealth of information available from surveys at radio to X-ray wavelengths (see e.g. Rowan-Robinson et al.
 2004)., 2004).
 Moreover a bright calibrator 3345...1642]398) stands close by. essential for observations at this frequency.," Moreover a bright calibrator 345, 1642+398) stands close by, essential for observations at this frequency."
 To find the maximum number of sources. given source counts of slopes with which we are familiar in the racio regime. the worst survey strategy is a single deep exposure. which we estimate would have vielded x: 0.05 sources in the primary beam (1.00 aremin Full-Width HIalf-Maxinum. = FAWULIAL) area. assuming a 24-h integration and an rms of 250 μ.]ν in LO min.," To find the maximum number of sources, given source counts of slopes with which we are familiar in the radio regime, the worst survey strategy is a single deep exposure, which we estimate would have yielded $\le$ 0.05 sources in the primary beam (1.00 arcmin Full-Width Half-Maximum = FWHM) area, assuming a 24-h integration and an rms of 250 $\mu$ Jy in 10 min."
 At radio wavelengths. wide and shallow always beats narrow and deep: an analysis in Appendix Al quantifies this and shows that our adoption of the NVSS observing ¢vele (23 sec integration. 7 sec telescope settle) is near optimal for à series of independent. snapshots.," At radio wavelengths, wide and shallow always beats narrow and deep; an analysis in Appendix A1 quantifies this and shows that our adoption of the NVSS observing cycle (23 sec integration, 7 sec telescope settle) is near optimal for a series of independent snapshots."
 We estimated. from the Ile Telescope 15-Cllz source count available at the time that observing 2880 fields in 24h should vield about 5 sources at a d moy (7 40) survey limit., We estimated from the Ryle Telescope 15-GHz source count available at the time that observing 2880 fields in 24h should yield about 5 sources at a 4 mJy $\sim4\sigma$ ) survey limit.
 Phere is no significant. cillicultv in analysis of such a number of fieles: this is essentially a mapping process and the general emptiness of the sky at 43 Cllz implies that almost no deconvolution is required., There is no significant difficulty in analysis of such a number of fields; this is essentially a mapping process and the general emptiness of the sky at 43 GHz implies that almost no deconvolution is required.
 We used the technique of ‘referenced pointing” (seo www.vla.nrao.edu/memos/test/I89/) to position the primary beam with an accuracy of better than 5r aresec., We used the technique of `referenced pointing' (see www.vla.nrao.edu/memos/test/189/) to position the primary beam with an accuracy of better than 5 arcsec.
 The observations were made in late 2001. at which time all VLA antennas hack been newly equipped. with 43-CGllz receiving systems.," The observations were made in late 2001, at which time all VLA antennas had been newly equipped with 43-GHz receiving systems."
 Ln order to maximize sensitivity to extended emission and to minimize the ellects of the atmosphere on phase stability. we used. the the smallest VLA. Lo. the D configuration: and we observed in November. as the necessary phase stability at this frequency is generally available only in winter.," In order to maximize sensitivity to extended emission and to minimize the effects of the atmosphere on phase stability, we used the the smallest VLA, i.e. the D configuration; and we observed in November, as the necessary phase stability at this frequency is generally available only in winter."
 We observed in 4 6-hour runs to, We observed in 4 6-hour runs to
IXPNOTau9 on a 36e level. which might be due to the emission from the nearby source.,"KPNOTau9 on a $\sigma$ level, which might be due to the emission from the nearby source."
 All other objects have maps without contaminating sources., All other objects have maps without contaminating sources.
 As a complementary test. we searched the IRAS point source catalogue [or sources in a LOO” circle around. our targets.," As a complementary test, we searched the IRAS point source catalogue for sources in a 100"" circle around our targets."
 It turned out that only (wo objects — IKPNOTau9 and JO4I4I1--2811.— have an IRAS neighbour., It turned out that only two objects – KPNOTau9 and J041411+2811 – have an IRAS neighbour.
 In the case of IKPNOTau9. this neighbour has a LOOgam flux of JJv and is located at a distance of 917. and thus possibly could comlaminale the background measurement lor the brown dwarf.," In the case of KPNOTau9, this neighbour has a $\mu m$ flux of Jy and is located at a distance of 91"", and thus possibly could contaminate the background measurement for the brown dwarf."
 We therefore attribute the negative [lux level lor KPNOTan9 to improper background subtraction., We therefore attribute the negative flux level for KPNOTau9 to improper background subtraction.
 The ΗνΑΛ neighbour of JO41411-—2811 probably has no significant influence on our mmm flux measurement. as argued above.," The IRAS neighbour of J041411+2811 probably has no significant influence on our mm flux measurement, as argued above."
 We examined the sam RAS images for all our sources. and found that the flux level in the region which affects the background subtraction. is more or less constant.," We examined the $\mu m$ IRAS images for all our sources, and found that the flux level in the region which affects the background subtraction, is more or less constant."
 These results confirm that our objects. with the exception of IKPNO'Tau9. are fairly isolated and in regions without strong background inhomosgenities.," These results confirm that our objects, with the exception of KPNOTau9, are fairly isolated and in regions without strong background inhomogenities."
 An alternative assessment of the reliability of our mun fluxes can be made based on our measurenients itself., An alternative assessment of the reliability of our mm fluxes can be made based on our measurements itself.
 If the [Iuxes are pure noise (i.e. no significant emission from the target and not affected by improper background subtraction). we expect them (o scatter around zero. with Gaussian distribution.," If the fluxes are pure noise (i.e. no significant emission from the target and not affected by improper background subtraction), we expect them to scatter around zero, with Gaussian distribution."
 Ht is obvious that the complete sample is not consistent with Gaussian noise. because we have six objects will >30 fluxes. whereas we expect zero.," It is obvious that the complete sample is not consistent with Gaussian noise, because we have six objects with $>3\sigma$ fluxes, whereas we expect zero."
 After excluding all >30 detections (positive and negative). we expect 9.2 (66%)) to have flux levels within the lo and 13.3Js )) within the 26 uncertainties.," After excluding all $>3\sigma$ detections (positive and negative), we expect 9.2 ) to have flux levels within the $\sigma$ and 13.3 ) within the $\sigma$ uncertainties."
 The actual numbers for our sample are 8 and 13. respectively.," The actual numbers for our sample are 8 and 13, respectively."
 To verily if the average fIux after excluding >30 measurements (70.2 mmJ«v) is consistent wilh the expected. zero value. we carried out Monte Carlo sinmlations: Assuming pure Gaussian noise with e=0.78 (he average uncertaintv) and zero average. we generated 14 datapoints and computed the average.," To verify if the average flux after excluding $>3\sigma$ measurements $-0.2$ mJy) is consistent with the expected zero value, we carried out Monte Carlo simulations: Assuming pure Gaussian noise with $\sigma = 0.78$ (the average uncertainty) and zero average, we generated 14 datapoints and computed the average."
 From 10000 test runs. a substantial [raction of resulted in an average <—0.20 mmJv.," From 10000 test runs, a substantial fraction of resulted in an average $\le-0.20$ mJy."
 These (wo simple tests show that the scatter in Fie., These two simple tests show that the scatter in Fig.
 d is huly consistent will Gaussian noise plus an excessive number of 730 outliers. confirming again (hat the quoted fluxes ave (wilh the exception of IKRPNOTau9) most likely nol severely affected by imperlect background subtraction.," \ref{f1} is fully consistent with Gaussian noise plus an excessive number of $>3\sigma$ outliers, confirming again that the quoted fluxes are (with the exception of KPNOTau9) most likely not severely affected by imperfect background subtraction."
 We therefore conclude that the [Iuxes lor the positive 3e detections are related (0 our substellar targets., We therefore conclude that the fluxes for the positive $3\sigma$ detections are related to our substellar targets.
 Finally we note that two of our objects ΕΙΤΙ)Taul and 4 have already been observed with the same insirumentation (Ixleinetal.2003): their 1.2:zummn upper Lut (for CFIITDDTaul) and flux (Jor CFIITDDTaud) are completely consistent wilh our values., Finally we note that two of our objects – CFHTBDTau1 and 4 – have already been observed with the same instrumentation \citep{kap03}; their mm upper limit (for CFHTBDTau1) and flux (for CFHTBDTau4) are completely consistent with our values.
 We aimed to complement our nimm fluxes with near-infrared aud mid-infrared data to be able to constrain (he spectral energy. distribution (SED) [or the sources., We aimed to complement our mm fluxes with near-infrared and mid-infrared data to be able to constrain the spectral energy distribution (SED) for the sources.
 This is parücularlv interesting for the six detections. because il allows us to compare with disk," This is particularly interesting for the six detections, because it allows us to compare with disk"
displav the typical properties of Sevfert galaxies and are the sites of dust-enshrouded nuclear activitv.,display the typical properties of Seyfert galaxies and are the sites of dust-enshrouded nuclear activity.
 Also. objects of this class awe found much more frequently as the total luminosity creases.," Also, objects of this class are found much more frequently as the total luminosity increases."
 There are problems with optical diagnostics. (hough.," There are problems with optical diagnostics, though."
 The inner regions of ULIRGs are usually characterized by high obscurations at visible wavelengths. ancl extinction/reddening effects can reduce significantly Cie effectiveness of any emission-line criterion.," The inner regions of ULIRGs are usually characterized by high obscurations at visible wavelengths, and extinction/reddening effects can reduce significantly the effectiveness of any emission-line criterion."
 In particular. a number of objects end up with different classifications when considering different line ratios.," In particular, a number of objects end up with different classifications when considering different line ratios."
 This ambiguity affects mostly the objects classified either as ID regions or low-ionization nuclear emission-line regions (LINERs)., This ambiguity affects mostly the objects classified either as H regions or low-ionization nuclear emission-line regions (LINERs).
 The latter ean actually be a different manilestation ol the AGN family. possibly related to low accretion rates and/or low radiative efficiency (e.g. Maoz et al.," The latter can actually be a different manifestation of the AGN family, possibly related to low accretion rates and/or low radiative efficiency (e.g. Maoz et al."
 2005)., 2005).
 On the other hand. the same (low) degree of ionization can be due io SB-diriven thermal shocks and galactic winds (e.g. Sturm et al.," On the other hand, the same (low) degree of ionization can be due to SB-driven thermal shocks and galactic winds (e.g. Sturm et al."
 2006)., 2006).
 The nature of LINERs is then controversial. and the ambiguity can not be resolved without more elective diagnostics (e.e. in the hard X-rays: Gonzallez-\lartinn et al.," The nature of LINERs is then controversial, and the ambiguity can not be resolved without more effective diagnostics (e.g. in the hard X-rays; Gonzállez-Martínn et al."
 2009)., 2009).
 This forbids a complete census of AGN activity among ULIRGs at optical wavelengths. even if some improvement can be obtained through a revision of the classification boundaries (e.g. Yuan. IXewlev Sanders 2010).," This forbids a complete census of AGN activity among ULIRGs at optical wavelengths, even if some improvement can be obtained through a revision of the classification boundaries (e.g. Yuan, Kewley Sanders 2010)."
 A further limitation of optical diagnostics is their poorly quantitative nature: it is diflieult to correct the line ratios for extinction and to take into account the possible dilferential obscuration of the AGN and SB components., A further limitation of optical diagnostics is their poorly quantitative nature: it is difficult to correct the line ratios for extinction and to take into account the possible differential obscuration of the AGN and SB components.
 A more detailed look into this issue is provided in Fig. 5..," A more detailed look into this issue is provided in Fig. \ref{bo},"
 where the AGN bolometric contribution is plotted. against the optical classification (ie.. here we collapse the plot ol Fig.," where the AGN bolometric contribution is plotted against the optical classification (i.e., here we collapse the plot of Fig."
 4. along the luminosity axis)., \ref{sr} along the luminosity axis).
 We have defined. four different. regions in the oj space with respect to the AGNweight. that is: (region 1. as< 0.05): (region 2. 0.05«αι 0.25): (region 3. 0.25«Aggy< 0.60): (region 4. Apo2» 0.60).," We have defined four different regions in the $\alpha_\mathit{bol}$ space with respect to the AGN, that is: (region 1, $\alpha_\mathit{bol}<0.05$ ); (region 2, $0.05<\alpha_\mathit{bol}<0.25$ ); (region 3, $0.25<\alpha_\mathit{bol}<0.60$ ); (region 4, $\alpha_\mathit{bol}>0.60$ )."
 The separation values have been chosen as follows: αι=0.05 represents a sort of 3e. confidence limit. above which secure AGN detections are found.," The separation values have been chosen as follows: $\alpha_\mathit{bol}=0.05$ represents a sort of $\sigma$ confidence limit, above which secure AGN detections are found."
 Ilence all the SD-dominated sources and the more controversial cases fall into region 1 (beside many solid AGN detections)., Hence all the SB-dominated sources and the more controversial cases fall into region 1 (beside many solid AGN detections).
 In our previous works ανω=0.25 was selected [or the sake of a gross classification. corresponding to a luminosity ratio of 1:3 between the AGN and SB components (and. roughly. to the average AGN contribution).," In our previous works $\alpha_\mathit{bol}=0.25$ was selected for the sake of a gross classification, corresponding to a luminosity ratio of 1:3 between the AGN and SB components (and, roughly, to the average AGN contribution)."
 We keep this convention in the boundary between the region 2 and the region 3., We keep this convention in the boundary between the region 2 and the region 3.
 Finally. we have chosen αρ=0.60 to single," Finally, we have chosen $\alpha_\mathit{bol}=0.60$ to single"
16 pol sources and discrete fbuueutarv features iu ιο field.,the point sources and discrete filamentary features in the field.
 We chose the size of the backeround region o include =1100 total counts from the full set of observations., We chose the size of the background region to include $\approx 1400$ total counts from the full set of observations.
 Iu Figure 2.. we display the flux as a function of time caving the observations ou 2004 July 57.," In Figure \ref{fig:lc}, we display the flux as a function of time during the observations on 2004 July 5–7."
 Three minima are evident iu the light curve at about 0. 8. aud 31 h. As reported in Munoetal.(2005). we searched these observations for periodic variability using the Bavleigh statistic (Zi:Duccherietal.1983).," Three minima are evident in the light curve at about 0, 8, and 31 h. As reported in \citet{mun05}, we searched these observations for periodic variability using the Rayleigh statistic \citep[$Z_1^2$;][]{buc83}."
 The the power spectrin contains a strong signal with a period of 7.9 hi aud Zi=175. and its first Lavmonic with Zi=120.," The the power spectrum contains a strong signal with a period of 7.9 h and $Z_1^2 = 175$, and its first harmonic with $Z_1^2 = 120$."
 The exact significance of this signal is uncertain. because there may be red noise in the power spectrum.," The exact significance of this signal is uncertain, because there may be red noise in the power spectrum."
 However. oobservatious in 2001 August detected similaxy dips with the same 7.9 lL period (Bélangeretal.2005.D.Porquetetal.inprep.) making us confident that this signal represents a stable periodicity in the source. such as its orbital period.," However, observations in 2004 August detected similar dips with the same 7.9 h period \citep[][D. Porquet \etal, in prep.]{bel05}
 making us confident that this signal represents a stable periodicity in the source, such as its orbital period."
 Tn order to further explore the nature of the modulation. in the top panel of Figure 3. we display the 28 keV count rate folded about the 7.9 lL period.," In order to further explore the nature of the modulation, in the top panel of Figure \ref{fig:prof} we display the 2–8 keV count rate folded about the 7.9 h period."
 The most obvious feature iu the folded profile is the dip, The most obvious feature in the folded profile is the dip
 IXILII5D is a preanain sequence (PAIS) KY star located in the voung star cluster. 22264.,"}} 15D is a pre-main sequence (PMS) K7 star located in the young star cluster, 2264."
 Its large photometric variations are perioclic and possibly interpreted in terms of eclipses bv an opaque feature orbiling inside a circumstellar disk seen nearly edge-on (Herbstetal.2002)., Its large photometric variations are periodic and possibly interpreted in terms of eclipses by an opaque feature orbiting inside a circumstellar disk seen nearly edge-on \citep{Her02}.
. It is likely connected to the OOri star family. characterized by large photometric and polarimetric variabiliGes (eg.Nattaοἱal.1999).. but sets apart from them wilh (wo important differences: no inlrared excess in the disk ancl no significant color changes during the eclipses. both indicating particle sizes much larger than optical wavelengths.," It is likely connected to the Ori star family, characterized by large photometric and polarimetric variabilities \citep{Nal99}, but sets apart from them with two important differences: no infrared excess in the disk and no significant color changes during the eclipses, both indicating particle sizes much larger than optical wavelengths."
" As the VRI colours of (he star are consistent wilh ils spectral (wpe (Llamillonetal.2001).. the selective absorption is Ay«0.2. a value indeed in agreement with the loreground interstellar reddening. Ey,.x0.06—0.07. of stars in 22264 (Pérezetal.1987:Sung1997)."," As the $VRI$ colours of the star are consistent with its spectral type \citep{Ham01}, the selective absorption is $\Av < 0.2$, a value indeed in agreement with the foreground interstellar reddening, $\Ebv \simeq 0.06-0.07$, of stars in 2264 \citep{Perez87, Sung97}."
". On the other hand the grey absorption. Ao. during the out-ol-eclipse observations can be estimated indirectly using a well known relation (Barnesetal.1973). between the visual brightness parameter and the intrinsic colour (V—λα,"," On the other hand the grey absorption, $A_0$, during the out-of-eclipse observations can be estimated indirectly using a well known relation \citep{BEM78} between the visual brightness parameter and the intrinsic colour $(V-R)_0$."
 Using the absolute visual magnitude deduced Grom the distance modulus of (he cluster. this relation provides a value of the stellar radius consistent with that derived bv Hamiltonetal.(2001) and with the PAIS status only if the grev absorption is ~0.6—0.7 magnitude.," Using the absolute visual magnitude deduced from the distance modulus of the cluster, this relation provides a value of the stellar radius consistent with that derived by \citet{Ham01} and with the PMS status only if the grey absorption is $\sim 0.6 - 0.7$ magnitude."
 The five eclipses observed since the discovery of IKILII15D provide us with the main characteristics of the occultation: a relative depth ~0.95. an approximate duration of 18 davs and a period of 48.36 davs.," The five eclipses observed since the discovery of 15D \citep{Kear98} provide us with the main characteristics of the occultation: a relative depth $\sim 0.95$, an approximate duration of 18 days and a period of 48.36 days."
 The absorbing cloud is deduced to orbit at ~0.2 AU from the star and to extend over 1/3 of an orbit 2001)., The absorbing cloud is deduced to orbit at $\sim 0.2$ AU from the star and to extend over 1/3 of an orbit \citep{Ham01}.
. Such an occultation completely dilfers from a standard planetary. transit ancl also displavs other important peculiarities: (1) Ingress and egress are remarkably steep. indicating an occulting clump with a verv sharp edge: (ii) the minimum is not a flat land but contains a central bump wilh small scale features ancl a reverse peak close to mid-eclipse: (iii) when comparieg two successive occultations. (he reverse peak seems to shift from one side of the mid-eclipse to the other: (iv) since the first series of observations. (he occultation has clearly widened ancl deepened in secular fashion.," Such an occultation completely differs from a standard planetary transit and also displays other important peculiarities: (i) ingress and egress are remarkably steep, indicating an occulting clump with a very sharp edge; (ii) the minimum is not a flat land but contains a central bump with small scale features and a reverse peak close to mid-eclipse; (iii) when comparing two successive occultations, the reverse peak seems to shift from one side of the mid-eclipse to the other; (iv) since the first series of observations, the occultation has clearly widened and deepened in secular fashion."
"trend towards an increasing binary fraction with radius for the low-density (dynamically young) cluster, and no trend in the high-density (dynamically old) cluster.","trend towards an increasing binary fraction with radius for the low-density (dynamically young) cluster, and no trend in the high-density (dynamically old) cluster."
" ? do find that the ratio of wide-to-close binaries increases with radius within the inner pc, and is flat beyond this (see their fig."," \citet{Reipurth07} do find that the ratio of wide-to-close binaries increases with radius within the inner pc, and is flat beyond this (see their fig."
" 9; close binaries are 66 — 225 AU, and wide binaries are 225 — 670 AU)."," 9; close binaries are 66 – 225 AU, and wide binaries are 225 – 670 AU)."
" However, easily within the errors the ratios are flat beyond the inner 0.5 pc, and it is only within the inner 0.5 pc that there are significantly fewer wide binaries than close ones."," However, easily within the errors the ratios are flat beyond the inner 0.5 pc, and it is only within the inner 0.5 pc that there are significantly fewer wide binaries than close ones."
 This is not at all unexpected., This is not at all unexpected.
" While the Orion cluster has been expanding from its proposed denser initial state, it will still undergo dynamical processing as wider binaries are always more susceptible to destruction in the inner regions of a cluster than in the outer regions."," While the Orion cluster has been expanding from its proposed denser initial state, it will still undergo dynamical processing as wider binaries are always more susceptible to destruction in the inner regions of a cluster than in the outer regions."
" If Orion was always at its current size, then the wide binary population in the outer regions"," If Orion was always at its current size, then the wide binary population in the outer regions"
A population of new X-ray sources has been discovered by the observatory within a few tens of degrees of the direction to the Galactic center (Negueruela 2004:; Kuulkers 2005)).,A population of new X-ray sources has been discovered by the observatory within a few tens of degrees of the direction to the Galactic center (Negueruela \cite{Neg04}; Kuulkers \cite{Kuu05}) ).
 These sources display hard X-ray spectra. which are usually attributed to strong absorption by dense material lying close to the source.," These sources display hard X-ray spectra, which are usually attributed to strong absorption by dense material lying close to the source."
 Their continuum spectral parameters are typical of systems contaming neutron stars or black holes., Their continuum spectral parameters are typical of systems containing neutron stars or black holes.
 They are suspected to be high mass X-ray binaries (HMXBs) embedded in highly absorbing media. and for some of them massive companions were indeed identified (e.g.. Filliatre Chaty 2004:; Negueruela et al. 2005. 20061: ," They are suspected to be high mass X-ray binaries (HMXBs) embedded in highly absorbing media, and for some of them massive companions were indeed identified (e.g., Filliatre Chaty \cite{Fil04}; Negueruela et al. \cite{Neg05,Neg06}; ;"
Smith et al. 2006::, Smith et al. \cite{Smi06};
 Masetti et al. 2006:;, Masetti et al. \cite{Mas06};
 Pellizza et al. 2006::, Pellizza et al. \cite{Pel06};
 Chaty et al. 2008))., Chaty et al. \cite{Cha08}) ).
 Investigating the nature of these sources can provide important insight into the evolution of massive stars. the physics of compact objects. and the mechanisms driving the accretion process.," Investigating the nature of these sources can provide important insight into the evolution of massive stars, the physics of compact objects, and the mechanisms driving the accretion process."
 was discovered on April 7. 2005. by theINTEGRAL observatory (Soldi et al. 2005)).," was discovered on April 7, 2005, by the observatory (Soldi et al. \cite{Sol05}) ),"
 and displayec a significant increase in brightness over a timescale of a few days (Paizis et al. 2005))., and displayed a significant increase in brightness over a timescale of a few days (Paizis et al. \cite{Pai05}) ).
 The observatory performed follow-up observations on April 13 and 15. measuring a robust position of the source (Kennea et al. 2005::," The observatory performed follow-up observations on April 13 and 15, measuring a robust position of the source (Kennea et al. \cite{Ken05};"
 Beckmann et al. 2005).Swift , Beckmann et al. \cite{Bec05}) ).
XRT spectra are described well by both a absorbed power law and an absorbed black body., XRT spectra are described well by both an absorbed power law and an absorbed black body.
 The X-ray absorption is high (Ny= 0.4—1.7x107 οι) and variable. which suggests that it is intrinsic to the source rather tha interstellar.," The X-ray absorption is high $N_{\rm H} =
0.4$ $1.7 \times 10^{23}$ $^{-2}$ ) and variable, which suggests that it is intrinsic to the source rather than interstellar."
 The absorption by the Galactic interstellar medium in the source direction is almost an order of magnitude lower. hence supporting the intrinsic nature of the measured source absorption.," The absorption by the Galactic interstellar medium in the source direction is almost an order of magnitude lower, hence supporting the intrinsic nature of the measured source absorption."
 The precise position obtained by the XRT telescope onboard allowed Rodriguez Paizis (2005)) to search for optical and NIR counterparts. to assess the nature of the system.," The precise position obtained by the XRT telescope onboard allowed Rodriguez Paizis \cite{Rod05}) ) to search for optical and NIR counterparts, to assess the nature of the system."
" They found a NIR source (2MASS J16281083—4838560) only 1.7” from theSwift position. with estimated magnitudes 16.8. H>15.8. and K,=13.95+0.06."," They found a NIR source (2MASS $-$ 4838560) only $1.7\arcsec$ from the position, with estimated magnitudes $J > 16.8$ , $H > 15.8$, and $K_{\mathrm s} = 13.95 \pm
0.06$."
 A Magellan-Baade image in the K band (Steeghs et al. 2005)), A Magellan-Baade image in the $K$ band (Steeghs et al. \cite{Ste05}) )
 shows many point sources inside the Swiff-XRT error. circle. the position and magnitude (K~1.1) of the brightest one being fully consistent with those of the 2MASS source.," shows many point sources inside the -XRT error circle, the position and magnitude $K \sim 14.1$ ) of the brightest one being fully consistent with those of the 2MASS source."
 A mid-IR source found in the Galactic Legacy Midplane Survey Extraordinaire=) (GLIMPSE) was tentatively associated with the source by Beckmann et al. (2005))., A mid-IR source found in the Galactic Legacy Midplane Survey Extraordinaire (GLIMPSE) was tentatively associated with the source by Beckmann et al. \cite{Bec05}) ).
 No optical or UV counterpart of was found in Swéft-UVOT data (Beckmann et al. 2005)).," No optical or UV counterpart of was found in }-UVOT data (Beckmann et al. \cite{Bec05}) ),"
 nor in Digitized Sky Survey (DSS) images (Rodriguez Paizis 2005))., nor in Digitized Sky Survey (DSS) images (Rodriguez Paizis \cite{Rod05}) ).
 As pointed out by Beckmann et al. (2005)).," As pointed out by Beckmann et al. \cite{Bec05}) ),"
 if the association of the NIR and mid-IR sources with were correct. its spectral energy distribution would then resemble those of high-mass X-ray binaries.," if the association of the NIR and mid-IR sources with were correct, its spectral energy distribution would then resemble those of high-mass X-ray binaries."
 These arguments show the importance of additional multiwavelength investigations of to help identify its optical counterpart beyond any doubt., These arguments show the importance of additional multiwavelength investigations of to help identify its optical counterpart beyond any doubt.
 Unveiling its properties would shed light on the nature of the new population of X-ray sources discovered byINTEGRAL., Unveiling its properties would shed light on the nature of the new population of X-ray sources discovered by.
 With this aim. shortly after the discovery of we performed optical and NIR observations of this source using the ESO New Technology Telescope (NTT).," With this aim, shortly after the discovery of we performed optical and NIR observations of this source using the ESO New Technology Telescope (NTT)."
 In this paper. we present our observations (Sect. 2))," In this paper, we present our observations (Sect. \ref{obs}) )"
 and results (Sect. 3)).," and results (Sect. \ref{res}) ),"
 and discuss their implications for the nature of (Sect. 4))., and discuss their implications for the nature of (Sect. \ref{disc}) ).
 Our observations ofJ16283-4838.. which began only 11 days after the discovery. were carried out on the nights of 2005 April 18. 19. 21. and 28 (imaging). July 22 (imaging and polarimetry). and July 28 (spectroscopy). with the ESO 3.5-meter NTT at La Silla Observatory. Chile.," Our observations of, which began only 11 days after the discovery, were carried out on the nights of 2005 April 18, 19, 21, and 28 (imaging), July 22 (imaging and polarimetry), and July 28 (spectroscopy), with the ESO 3.5-meter NTT at La Silla Observatory, Chile."
 Optical and. NIR images of the field of the source were obtained with the ESO Superb-Seeing Imager 2 (SUSI2) and the Son of Isaac (SOFD) instruments respectively. as part of a target-of-opportunity program (ESO 075.D-0634. P. I. Chaty).," Optical and NIR images of the field of the source were obtained with the ESO Superb-Seeing Imager 2 (SUSI2) and the Son of Isaac (SOFI) instruments respectively, as part of a target-of-opportunity program (ESO 075.D-0634, P. I. Chaty)."
 NIR images were alsotaken with SOFI in polarimetric mode., NIR images were alsotaken with SOFI in polarimetric mode.
 Low-resolution NIR spectra of the brightest counterpart candidate proposed by Beckmann (2005)) were also obtained with SOFI., Low-resolution NIR spectra of the brightest counterpart candidate proposed by Beckmann \cite{Bec05}) ) were also obtained with SOFI.
The solar f-modes which are essentially surface gravity modes. provide a clagnostic of flows and magnetic fields present in the near surface region (Murawski Roberts 1993: Rosenthal Gough 1994: Rosenthal Christensen-Dalseaarcd 1995: Solia et al.,"The solar $f$ -modes which are essentially surface gravity modes, provide a diagnostic of flows and magnetic fields present in the near surface region (Murawski Roberts 1993; Rosenthal Gough 1994; Rosenthal Christensen-Dalsgaard 1995; Sofia et al."
 2005)., 2005).
 The f-mocle frequencies can also provide an accurate measure of solar radius (Schou οἱ al., The $f$ -mode frequencies can also provide an accurate measure of solar radius (Schou et al.
 1997: Antia 1993)., 1997; Antia 1998).
 The dependence of solar radius on the Ll-vear activity evele is still a matter of controversy., The dependence of solar radius on the 11-year activity cycle is still a matter of controversy.
 Different measurements of solar radius have given conflicting results about ils temporal variations (e.g.. Laclare et al.," Different measurements of solar radius have given conflicting results about its temporal variations (e.g., Laclare et al."
 1996: Noell. 2004: Ixuhn et al.," 1996; Noëll, 2004; Kuhn et al."
 2004: Chapman et al., 2004; Chapman et al.
 2008: Djafer et al., 2008; Djafer et al.
 2003)., 2008).
 Using the energv budget of the Sun. il is easy to see that the radius variation. if anv. would be localized only in the near surface regions which are well sampled by f-modes.," Using the energy budget of the Sun, it is easy to see that the radius variation, if any, would be localized only in the near surface regions which are well sampled by $f$ -modes."
 Hence. we can expect the /-node Lrequencies to reflect these variations in the solar radius.," Hence, we can expect the $f$ -mode frequencies to reflect these variations in the solar radius."
 Dziembowski et al. (, Dziembowski et al. (
2001) and Dziembowski Goode (2004) (henceforth DGOL) have obtained a relation between the f/-mocde [requeney variations and radius variations in the subsurface lavers.,2001) and Dziembowski Goode (2004) (henceforth DG04) have obtained a relation between the $f$ -mode frequency variations and radius variations in the subsurface layers.
 This relation was later used by Lefebvre ]xosovichev (2005. 2007) to show that helioseismic radius varies in ant-phase with solar activity in (he outer region of the Sun but there is a change in behavior in deeper lavers.," This relation was later used by Lefebvre Kosovichev (2005, 2007) to show that helioseismic radius varies in anti-phase with solar activity in the outer region of the Sun but there is a change in behavior in deeper layers."
These trends suggest (hat inlrared. colors of variable objects may support the cloud model [or variabilitv.,These trends suggest that infrared colors of variable objects may support the cloud model for variability.
" The J—A, color of those objects which have been surveved lor variabilitv is shown in Figure 2.", The $J-K_s$ color of those objects which have been surveyed for variability is shown in Figure 2.
 Although subject to small number statistics. suspected variables with spectral (vpes later than L2 tend to be bluer than the average L dwarl (IXirkpatricketal.2000) at the same spectral (wpe.," Although subject to small number statistics, suspected variables with spectral types later than L2 tend to be bluer than the average L dwarf \citep{kir00}
 at the same spectral type."
" Models. predict levetal.2001) that a hypothetical L dwarf with no clouds will be substantially bluer (c1.5mag:Marlev&Ackerman2001) at /—A, than a more realistic object with the sanie effective temperature and a cloudy atmosphere."," Models predict \citep{marl00,marl01} that a hypothetical L dwarf with no clouds will be substantially bluer \citep[$\sim1.5$ mag;][]{marl01b}
 at $J-K_s$ than a more realistic object with the same effective temperature and a cloudy atmosphere."
" Thus. if the average L cdwarf at a given spectral type is entirely covered with clouds. it would not likely be seen as a variable and it would have more (vpical J—A, color."," Thus, if the average L dwarf at a given spectral type is entirely covered with clouds, it would not likely be seen as a variable and it would have more typical $J-K_s$ color."
 In order for photometric variations to arise bv the cloud mechanism there must be non-uniformity in the cloud coverage. such as clearings in the clouds.," In order for photometric variations to arise by the cloud mechanism there must be non-uniformity in the cloud coverage, such as clearings in the clouds."
" So. not only would clear sections of the atmosphere (holes) provide a source for brightness variations. flix emerging through the hypothetical holes would cause the --fy, color to be somewhat bluer than the average object."," So, not only would clear sections of the atmosphere (holes) provide a source for brightness variations, flux emerging through the hypothetical holes would cause the $J-K_s$ color to be somewhat bluer than the average object."
 The lack of a similar trend for the early L and late M dwarls might indicate a different mechanism is at work in those atmosplieres., The lack of a similar trend for the early L and late M dwarfs might indicate a different mechanism is at work in those atmospheres.
 Clearly more data. including time resolved multi-color photometry. are required to determine if variabilitv is indeed connected to color.," Clearly more data, including time resolved multi-color photometry, are required to determine if variability is indeed connected to color."
 Schubert&Zhang(2000) show that the atmosphere of L cdwarls can exist in one of two states. chaotic or banded.," \citet{sch00} show that the atmosphere of L dwarfs can exist in one of two states, chaotic or banded."
 Since clouds respond to atmospheric motions. they. would presumably reflect one of these two morphologies.," Since clouds respond to atmospheric motions, they would presumably reflect one of these two morphologies."
 In general. higher mass objects will have more chaotic and three-dimensional internal dvnaimics than lower mass objects. meaning that the higher mass objects are less likely to have banded eloud features.," In general, higher mass objects will have more chaotic and three-dimensional internal dynamics than lower mass objects, meaning that the higher mass objects are less likely to have banded cloud features."
 It is not obvious which cloud morphology. would better produce photometric variations., It is not obvious which cloud morphology would better produce photometric variations.
 For example. objects with more chaotic atmospheres might be more likely to have uniformly distributed clouds.," For example, objects with more chaotic atmospheres might be more likely to have uniformly distributed clouds."
 When rotating such objects might show little photometric variation., When rotating such objects might show little photometric variation.
 If the chaotic atmospheres result in fairly complete cloud coverage. then we might expect that more massive L clwarls would be redder in J—A; and tend not to be variable.," If the chaotic atmospheres result in fairly complete cloud coverage, then we might expect that more massive L dwarfs would be redder in $J-K_s$ and tend not to be variable."
 Conversely. if chaotic atiospheres more often produce large holes in the clouds (μον may more easily produce photometric signatures (han banded atmospheres.," Conversely, if chaotic atmospheres more often produce large holes in the clouds they may more easily produce photometric signatures than banded atmospheres."
" In (is scenario. the more massive L cwarls would be bluer in J—A, and tend to be variable."," In this scenario, the more massive L dwarfs would be bluer in $J-K_s$ and tend to be variable."
 Furthermore. rapidly evolving. chaotic atiospheres mav be responsible lor the changes in photometric period observed [or some objects.," Furthermore, rapidly evolving, chaotic atmospheres may be responsible for the changes in photometric period observed for some objects."
 IC is easy {ο imagine similar scenarios for banded clouds., It is easy to imagine similar scenarios for banded clouds.
 The presence of clouds can also explain the different (pes of variability seen in b dwarts., The presence of clouds can also explain the different types of variability seen in L dwarfs.
 For example. periodic variations over (ime scales of days can be explained by a long-lived clearing or thickening in the clouds.," For example, periodic variations over time scales of days can be explained by a long-lived clearing or thickening in the clouds."
 Rotation modulation of the photometric signal will then produce a periodic signal., Rotation modulation of the photometric signal will then produce a periodic signal.
 The inhomogeneities could migrate latitudinallv or dissipate and relorm at a different latitude as does the Great Dark Spot in the atmosphere of Neptune, The inhomogeneities could migrate latitudinally or dissipate and reform at a different latitude as does the Great Dark Spot in the atmosphere of Neptune
the 1-D spectra. we ran each module through a three pixel 3-0 clipped mean.,"the 1-D spectra, we ran each module through a three pixel $\sigma$ clipped mean."
 We stitclied orders together using (the data at overlapping wavelengths in each lower wavelength order to attach onto its neighboring longer wavelength order., We stitched orders together using the data at overlapping wavelengths in each lower wavelength order to attach onto its neighboring longer wavelength order.
 At a few wavelengths in the overlap regions we interpolated values between neighboring orders to provide a smooth transition: the uncertainties assigned to these wavelengths were increased by50%... a conservative value based on examination of other overlap regions.," At a few wavelengths in the overlap regions we interpolated values between neighboring orders to provide a smooth transition; the uncertainties assigned to these wavelengths were increased by, a conservative value based on examination of other overlap regions."
 The same procedure was used (o attach the Long and Short modules. resulting in ten complete spectra ranging from 9.36 (o 37.2 jum.," The same procedure was used to attach the Long and Short modules, resulting in ten complete spectra ranging from 9.86 to 37.2 $\micron$."
 The final spectrum was produced by first normalizing each spectrum to have (he same flux al 15 yan. a region of the spectrum relatively free of residual noise features. and then taking a 3-7 clipped average of the ten sub-spectra.," The final spectrum was produced by first normalizing each spectrum to have the same flux at 15 $\mu$ m, a region of the spectrum relatively free of residual noise features, and then taking a $\sigma$ clipped average of the ten sub-spectra."
 Uncertainties were obtained by evaluating the standard deviation of the mean lor the spectra included in the average at each wavelength., Uncertainties were obtained by evaluating the standard deviation of the mean for the spectra included in the average at each wavelength.
 As a final step. the entire HiRes spectrum. (stard-disk) was adjusted by to. bring it into agreement with the LoRes spectrum. an amount which is within the quoted calibration uncertainties of the IHiltes spectrophotometry ane with our contemporaneous [liRes observations of Η05180.," As a final step, the entire HiRes spectrum (star+disk) was adjusted by to bring it into agreement with the LoRes spectrum, an amount which is within the quoted calibration uncertainties of the HiRes spectrophotometry and with our contemporaneous HiRes observations of HD68146."
 Observations of ILD69330 were made using the Michelle instrument on the Gemini-North telescope., Observations of HD69830 were made using the Michelle instrument on the Gemini-North telescope.
 Images were obtained on (wo nights (7-8 Mar 2007) at 11.2 and 13.5 jm. Careful flat-fielcling. registration ancl co-adcdition of individual frames resulted in images (hat proved to be unresolved compared to the reference star. IID61935.," Images were obtained on two nights (7-8 Mar 2007) at 11.2 and 18.5 $\mu$ m. Careful flat-fielding, registration and co-addition of individual frames resulted in images that proved to be unresolved compared to the reference star, HD61935."
 Some frames with significant trailing blur in (he direction of motion of the chopping secondary were rejected., Some frames with significant trailing blur in the direction of motion of the chopping secondary were rejected.
 The flux densities detected in cddiameter apertures (Table 2)) are consistent with the total startdisk flux density measured in the larger beams with Spitzer. confirming the compact nature of the disk.," The flux densities detected in diameter apertures (Table \ref{phottable}) ) are consistent with the total star+disk flux density measured in the larger beams with Spitzer, confirming the compact nature of the disk."
 The azimuthlially averaged I-dimensional scans of LD69330 and I1D61935 at 18.5 jan are shown in Figure 4 and are consistent with unresolved sources with a Full Width at Half Maximum (FWIIM) of0., The azimuthally averaged 1-dimensional scans of HD69830 and HD61935 at 18.5 $\mu$ m are shown in Figure \ref{Michelle} and are consistent with unresolved sources with a Full Width at Half Maximum (FWHM) of.
"3"".. The limit to the size of ILD69330's emitting region at 11.2 jam is + AU diameter. which is consistent with the inferred temperature and location for the dust discussed in the models below."," The limit to the size of HD69830's emitting region at 11.2 $\mu$ m is 4 AU diameter, which is consistent with the inferred temperature and location for the dust discussed in the models below."
 Similar results were reported by(2008)., Similar results were reported by.
. As discussed below. mid-IR observations with the MIDI Instrument on the VLT-Interferometer indicate that the," As discussed below, mid-IR observations with the MIDI Instrument on the VLT-Interferometer indicate that the"
of Vig were tabulated aud plotted. as iu Fie.,"of $_0$ were tabulated and plotted, as in Fig."
 l (upper). includiug estimates forLF color from the intriusic colors of Johnson(1966.1968) using lis tabulated valucs forJ andKW interpolated according to the effective waveleusthis of the filters.," \ref{fig1} (upper), including estimates for color from the intrinsic colors of \citet{jo66,jo68} using his tabulated values for and interpolated according to the effective wavelengths of the filters."
 Least squares fits to the data then established polvnonual relationships between the colors., Least squares fits to the data then established polynomial relationships between the colors.
 The derived relationships were then tested using 2\TASS data andDV observatious for 19 reasonably bright. nearbv and uureddened. photometric aud spectroscopic dwarf standards. stars with unsaturated observatious in the Ursa Major and .ILbvades clusters. which are both unreddeued. aud stars im the voung cluster NGC 2211 corrected for differcutial reddening within the Rosctte Nebula (seeTurner1976a).," The derived relationships were then tested using 2MASS data and observations for 19 reasonably bright, nearby and unreddened, photometric and spectroscopic dwarf standards, stars with unsaturated observations in the Ursa Major and Hyades clusters, which are both unreddened, and stars in the young cluster NGC 2244 corrected for differential reddening within the Rosette Nebula \citep[see][]{tu76a}."
. NCC 2211 stars were used in order to tie down the hot cud of the sequence: its member stars possess excelleut photometry ou the(BV system (Johnson1962) and lic in a region of well-established reddening law (πιο1976aj., NGC 2244 stars were used in order to tie down the hot end of the sequence; its member stars possess excellent photometry on the system \citep{jo62} and lie in a region of well-established reddening law \citep{tu76a}.
. Reddening corrections for carly-type stars were applied usingCDV colors and the relations = 0.295V). Ky) = 0.9ID]. and Ay=9.1τοIT) for R=AV/E(QB|V)=3.05. derived from van de IIulst's reddeuiug curve No.," Reddening corrections for early-type stars were applied using colors and the relations = 0.295, $_s$ ) = 0.49, and $A_V = 2.427 E(J-H)$ for $R = A_V/E(B-V) = 3.05$, derived from van de Hulst's reddening curve No."
 15 (seeJohnson 1966).., 15 \citep[see][]{jo66}. .
 Iu addition. the observations for Ivacdes stars were supplemented by data from C'arney(1982) in order to reduce the photometric scatter typical of 2M1ASS observatious.," In addition, the observations for Hyades stars were supplemented by data from \citet{ca82} in order to reduce the photometric scatter typical of 2MASS observations."
 The results are plotted in the lower portion of Fig. 1.., The results are plotted in the lower portion of Fig. \ref{fig1}.
 The observational treuds ofV... IL and I. ave fairly simular to the predicted trends sccu in the top portion of Fig. 1..," The observational trends of, , and $_s$ are fairly similar to the predicted trends seen in the top portion of Fig. \ref{fig1},"
 but with noticeable offsets. particularly inW," but with noticeable offsets, particularly in."
f The adopted intrinsic relations were therefore from establishedpolvuonmial fits to the observational data rather than from the predicted relations.with the results tabulated iu Table 1 for zero-age main sequence (ZAMS) stars.," The adopted intrinsic relations were therefore established from polynomial fits to the observational data rather than from the predicted relations,with the results tabulated in Table \ref{tab1} for zero-age main sequence (ZAMS) stars."
 The derived polvnomial relatiouships of Fig., The derived polynomial relationships of Fig.
 1 were used to derive intrinsic7 audLv colors for main-sequence stars as a function of iutriusic broad baud color iudex Τμ. aud the ZAXNIS calibration was that « Turner(1976b.1979).," \ref{fig1} were used to derive intrinsic and colors for main-sequence stars as a function of intrinsic broad band color index $_0$, and the ZAMS calibration was that of \citet{tu76b,tu79}."
. The intrinsic relationship for 2MASS colors in Table 1/ is similar to the intrinsic color-color relation for dwarf stars m the system. displaviug a noticeable “kink” near spectral type AQ.," The intrinsic relationship for 2MASS colors in Table \ref{tab1} is similar to the intrinsic color-color relation for dwarf stars in the system, displaying a noticeable “kink” near spectral type A0."
 2\TASSJ imaenitucdes sample the Brackett continua m hot stars; whileHand A; magnitudes sample the Pfund coutinmiun. so7 color should xovide a measure of the Pfuud ciscoutinuitv iu rot stars. niuch like color provides a nieasure of the Balmer discoutimuty.," 2MASS magnitudes sample the Brackett continuum in hot stars, while and $_s$ magnitudes sample the Pfund continuum, so color should provide a measure of the Pfund discontinuity in hot stars, much like color provides a measure of the Balmer discontinuity."
 ολΙαν. AY color should provides a measure of stellar temperature. uuch like color does iu the system.," Similarly, $_s$ color should provides a measure of stellar temperature, much like color does in the system."
 For cool stars all colors should closely track changes iu he slope of the black body contiunuai in the far infrared., For cool stars all colors should closely track changes in the slope of the black body continuum in the far infrared.
" Iuterstellar reddening affects JIN, colors uuch less thau it does colors. but the effects of circiuustellar emission are often more portant iu the far infrared than in the visible region."," Interstellar reddening affects $_s$ colors much less than it does colors, but the effects of circumstellar emission are often more important in the far infrared than in the visible region."
 A color-color diagram plotsD versusWV. for which a plot offf versus AL would be the closest approximation.," A color-color diagram plots versus, for which a plot of versus $_s$ would be the closest approximation."
 It appears. however. that1 colors may display a greater precision than AY colors in 2\TASS photometry. so tle diagnostic tool used here to establish reddening for cluster stars is a 2MASS color-color diagram in which AY is plottedversus 1.," It appears, however, that colors may display a greater precision than $_s$ colors in 2MASS photometry, so the diagnostic tool used here to establish reddening for cluster stars is a 2MASS color-color diagram in which $_s$ is plottedversus ."
 The intrinsic relatious were tested further using the three clusters used for thecalibration, The intrinsic relations were tested further using the three clusters used for thecalibration
a shredding of OPAL. lobular cluster near the vicinity of the ΓΗ.,a shredding of $10^5M_\odot$ globular cluster near the vicinity of the SBH.
 Subsequeuth. the cluster stars merge into a pre-existing disk of few 10*AZ... orbiting the ceutral black hole ou both prograde and retrograde. quasi-periodic. loop orbits.," Subsequently, the cluster stars merge into a pre-existing disk of few $10^7 M_\odot$, orbiting the central black hole on both prograde and retrograde, quasi-periodic, loop orbits."
 The s=1 mstabilitv which causes the large eccentricity is induced by resonant response to the counter-rotating orbits., The $m=1$ instability which causes the large eccentricity is induced by resonant response to the counter-rotating orbits.
 Iu this model. lopsideduess of the eccentric disk eeonetrv is created in response to the presence of retrograde orbits.," In this model, lopsidedness of the eccentric disk geometry is created in response to the presence of retrograde orbits."
 All Έπος models roughly reproduce tle double nucleus geometry and the dyes. such as the offset in the velocity dispersion profile from P2 by ~072. to varving degrees of accuracy.," All three models roughly reproduce the double nucleus geometry and the dynamics, such as the offset in the velocity dispersion profile from P2 by $\approx 0\farcs2$, to varying degrees of accuracy."
 However. one important difference between BEOQ0 aud BAOL | SSOL scenarios is that the BEOO model requires the stellar cluster be much less nassive than CZ 10) the black hole so that it can be disrupted.," However, one important difference between BE00 and BA01 + SS01 scenarios is that the BE00 model requires the stellar cluster be much less massive than $\lesssim 10\%$ ) the black hole so that it can be disrupted."
 By contrast. the BAOL aud SSOL models are constructed to have a mauch larger disk. about 20%LOW of the total central mass conceutration.," By contrast, the BA01 and SS01 models are constructed to have a much larger disk, about $20\%-40\%$ of the total central mass concentration."
 These two scenarios can be directly tested if a mass estimate of the disk can he robustly measured. and the uncertainty can be quantified.," These two scenarios can be directly tested if a mass estimate of the disk can be robustly measured, and the uncertainty can be quantified."
 The ambiguity in the bulee decomposition is reflected in the available plotometiry of Pl. which can siguificautlv differ in brightness frou. differeut studies;," The ambiguity in the bulge decomposition is reflected in the available photometry of P1, which can significantly differ in brightness from different studies."
 Furthermore. thus far it is not clear how much of P2 is part of the bulee or the eccentric disk.," Furthermore, thus far it is not clear how much of P2 is part of the bulge or the eccentric disk."
 Iu addition to the eccentric disk. the bulee structure of M31 and how it fits into the developing picture of ealaxv formation are interesting on thei own ierits.," In addition to the eccentric disk, the bulge structure of M31 and how it fits into the developing picture of galaxy formation are interesting on their own merits."
 However. their studies have been complicated by the double nucleus.," However, their studies have been complicated by the double nucleus."
 The steepness of galaxy ποιο aud correlations with other structural parameters (e.g. Faber et al., The steepness of galaxy nuclei and correlations with other structural parameters (e.g. Faber et al.
 1997) reflect the manner by which bulecs are formed., 1997) reflect the manner by which bulges are formed.
 Muuevical simulations show that black hole mergers cau flatten galaxy cores by ejecting stars from the ceuter (e.g. Ebisuzaki. Makino. Okwuura 1991: Nakano Makino 1999: and Milosavljevió Moerritt 2001).," Numerical simulations show that black hole mergers can flatten galaxy cores by ejecting stars from the center (e.g. Ebisuzaki, Makino, Okumura 1991; Nakano Makino 1999; and Milosavljević Merritt 2001)."
 This scenario appears promising for explaining the correlation fouud between huge bulees aud low central surface brightucss., This scenario appears promising for explaining the correlation found between large bulges and low central surface brightness.
 Iu this paper. I study the detailed propertics of the ADU» bulge by decomposing optical images to provide new structural parameters.," In this paper, I study the detailed properties of the M31 bulge by decomposing optical images to provide new structural parameters."
 With detailed decomposition one can address the following issues: Are there subtle structures in the bulee that are not obvious im full light?, With detailed decomposition one can address the following issues: Are there subtle structures in the bulge that are not obvious in full light?
 What are the relative contributions of the bulee. PI. aud P2 components?," What are the relative contributions of the bulge, P1, and P2 components?"
 What are their shapes?, What are their shapes?
 What is the bulee profile and how sharply is it peaked?, What is the bulge profile and how sharply is it peaked?
 Finally. I discuss what the new photometry of Pl and P2 reveal about the two competing scenarios that explain the formation ofthe eccentric disk.," Finally, I discuss what the new photometry of P1 and P2 reveal about the two competing scenarios that explain the formation of the eccentric disk."
" Although some of these questions can be addressed from other data in previous studies. this new analysis provides a unique look at the double nucleus based on a niore flexible rauge of assuniptious than foregoing «πο»,"," Although some of these questions can be addressed from other data in previous studies, this new analysis provides a unique look at the double nucleus based on a more flexible range of assumptions than foregoing studies."
 Iu the sections to follow. Section 2 discusses tle data. Section 3 briefly describes the analysis algorithun used. to debleud the bulge.," In the sections to follow, Section 2 discusses the data, Section 3 briefly describes the analysis algorithm used to deblend the bulge."
 Section. 1 discusses the decomposition. followed bv the cuviroument of the bulge iu Section 5.," Section 4 discusses the decomposition, followed by the environment of the bulge in Section 5."
 Section 6 compares the eccentric disk formation models., Section 6 compares the eccentric disk formation models.
 Conclusious follow in Section 7., Conclusions follow in Section 7.
 Throughout the discussion I assiune that the distance to M31 is D=770 kpc. following1000.," Throughout the discussion I assume that the distance to M31 is $D=770$ kpc, followingKB99."
 I also assume that the Calactic extinction is Ay0.21 (Burstein Teiles 1981). which is similar to Ay0.21 determined bv Schlegel ot al. (," I also assume that the Galactic extinction is $A_{\mbox{v}}=0.24$ (Burstein Heiles 1984), which is similar to $A_{\mbox{v}}=0.21$ determined by Schlegel et al. ("
1998).,1998).
 I&orineudy (1988) aud 100 dynamical models also suggest that AZ/Lz5.7 for the bulee stars., Kormendy (1988) and KB99 dynamical models also suggest that $M/L \approx 5.7$ for the bulge stars.
 This is similar to star formation mocels of Dell de Jong (2001). from which ouecau derive AZ/Lzz:6 based on the Wo£ color of the bulec.," This is similar to star formation models of Bell de Jong (2001), from which onecan derive $M/L\approx 6$ based on the $V-I$ color of the bulge."
 Tobtain V-band (=F£555W ) data (CO 5236. Westphal) ron the aarchive. as well as the decouvolved version roni T. Laucr.," I obtain $V$ -band $\approx F555W$ ) data (GO 5236, Westphal) from the archive, as well as the deconvolved version from T. Lauer."
 The two sets of nuages are used for different »irposes., The two sets of images are used for different purposes.
" The decouvolved version has been πο to a sinaller field of view (FOV) of 1178""«1178.", The deconvolved version has been trimmed to a smaller field of view (FOV) of $11\farcs8\arcsec\times11\farcs8$.
 While ideal or stucdving the nucleus. there is not enough realty for neasuring the bulee profile aud its shape.," While ideal for studying the nucleus, there is not enough realty for measuring the bulge profile and its shape."
" Therefore I also create a losaic image with a FOV of 150”«150"". with a missing echelon in the Planetary Camera (PC) quadrant."," Therefore I also create a mosaic image with a FOV of $150\arcsec\times150\arcsec$, with a missing echelon in the Planetary Camera (PC) quadrant."
 The net exposure time is 300s. which has sufficiently Heh signal-to-noise (S/N) to measure the bulge profile.," The net exposure time is 300s, which has sufficiently high signal-to-noise (S/N) to measure the bulge profile."
" Iu addition to the decouvolved aud mosaic images. I create a dithered image of the PC chip (FOV 2 35""ς35"") by combining four exposures of 300 secoud images."," In addition to the deconvolved and mosaic images, I create a dithered image of the PC chip (FOV = $35\arcsec\times35\arcsec$ ) by combining four exposures of 300 second images."
 This will be used later for comparisons with the decouvolved image., This will be used later for comparisons with the deconvolved image.
 I use a general ealaxy fitting program called CALFIT o do the 2-D decomposition., I use a general galaxy fitting program called GALFIT to do the 2-D decomposition.
" Detailed information on he software and how it is implemented are found in a colupanion paper by Peug. Ho. hupex. Rix (2002). but 1ο I describe it briefly,"," Detailed information on the software and how it is implemented are found in a companion paper by Peng, Ho, Impey, Rix (2002), but here I describe it briefly."
 One of the design capabilities of GALFIT is to accurately decompose nearby galaxies hat are highly resolved to study them closely aud uncover or extract galaxy sub-structures. such as nuclear disks. ws. aud double uuclei.," One of the design capabilities of GALFIT is to accurately decompose nearby galaxies that are highly resolved to study them closely and uncover or extract galaxy sub-structures, such as nuclear disks, bars, and double nuclei."
 To be highly flexible. GALFIT allows simultaneous fitting of staudard fiction types such as Sérrsic (1968). Cassian. expoucutial. and Nuker.," To be highly flexible, GALFIT allows simultaneous fitting of standard function types such as Sérrsic (1968), Gaussian, exponential, and Nuker."
 The xoeranm has the option to either couvolve the models with he PSF to simulate the seciug. or not convolve if the inage jas already been deconvolved.," The program has the option to either convolve the models with the PSF to simulate the seeing, or not convolve if the image has already been deconvolved."
 The απο: of compoucuts o fit is not limited a priori;, The number of components to fit is not limited a priori.
 GALETT iuininizesy. 4? residuals using a down-hill eradieut/parabolic expansion ucthod called Levenbere-Alarquardt (Press ot al., GALFIT minimizes $\chi^2$ residuals using a down-hill gradient/parabolic expansion method called Levenberg-Marquardt (Press et al.
 1997) w iteratively creating model tages. convolving them with the PSF. aud subtracting them from the data.," 1997) by iteratively creating model images, convolving them with the PSF, and subtracting them from the data."
 Even though the eradieut method is nof as “sinart™ as alternative Simulated Annealing or Metropolis aleoritlius. it has the virtue of being fast.," Even though the gradient method is not as “smart” as alternative Simulated Annealing or Metropolis algorithms, it has the virtue of being fast."
 As an example. fitting 6 components with Ll free parameters over the cutire 100.« pixel image. while doing couvolution. takes roughlv one müuute per iteration on a Peutinm £50. MITz conrputer. and converges m 30050 iterations.," As an example, fitting 6 components with 41 free parameters over the entire $400\times400$ pixel image, while doing convolution, takes roughly one minute per iteration on a Pentium 450 MHz computer, and converges in $30-50$ iterations."
 Speed is desirable because the 4? topolosv of the fit is complex: to adequately explore it aud fud an optimal fit thus requires testing various model combinations aud initial parameters., Speed is desirable because the $\chi^2$ topology of the fit is complex; to adequately explore it and find an optimal fit thus requires testing various model combinations and initial parameters.
 The iuerit fiction to minimize is the \?. or in normalized form. AZ= \7/Naut. defined as: where σεν Is the uncertainty. or weight. at cach pixel: Mace ," The merit function to minimize is the $\chi^2$ , or in normalized form, $\chi^2_\nu = \chi^2/N_{\rm dof}$ , defined as: where $\sigma_{x,y}$ is the uncertainty, or weight, at each pixel; $N_{\rm dof}$ "
and therefore presume that thev represent real effects in (he NUYV emission lines emitted bv the star.,and therefore presume that they represent real effects in the NUV emission lines emitted by the star.
 We discuss the velocities compared to those expected from our flare models in Section 3 below., We discuss the velocities compared to those expected from our flare models in Section 3 below.
 Figure 8 shows the time evolution of the Mg ID k emission for the largest flares. F2 and F9.," Figure 8 shows the time evolution of the Mg II k emission for the largest flares, F2 and F9."
 Flare F2 exhibits rather simple behavior. with the narrow central profile strongly enhanced on the blue side in the impulsive phase (E2a). and decreasing monotonically during the flare decay (F2b. F2c).," Flare F2 exhibits rather simple behavior, with the narrow central profile strongly enhanced on the blue side in the impulsive phase (F2a), and decreasing monotonically during the flare decay (F2b, F2c)."
 In contrast. Hare F9 is quite unusual. initially (F9a) showing a strong red enhancement in the central profile. but (hen developing very. broad. svimetric wines in F9b.," In contrast, flare F9 is quite unusual, initially (F9a) showing a strong red enhancement in the central profile, but then developing very broad, symmetric wings in F9b."
 Upon closer inspection of Figure 4. it appears that there is evidence for similarly broad wings in FT and F9a.," Upon closer inspection of Figure 4, it appears that there is evidence for similarly broad wings in F7 and F9a."
 These data are reminiscent of those in Dovle and Byrne (1987) where extensive. symmetric line broadening was seen in the hydrogen Balmer lines during; a flare on YZ CAL.," These data are reminiscent of those in Doyle and Byrne (1987) where extensive, symmetric line broadening was seen in the hydrogen Balmer lines during a flare on YZ CMi."
 Droad hydrogen lines during flares are usually interpreted as being due to Stark broadening (Donati-Falchi οἱ al., Broad hydrogen lines during flares are usually interpreted as being due to Stark broadening (Donati-Falchi et al.
 1935. Hawley ancl Pettersen 1991. Johns-IXrull et al.," 1985, Hawley and Pettersen 1991, Johns-Krull et al."
 1997. Jevremovic et al.," 1997, Jevremovic et al."
 1993) but Dovle and Byrne were unable to find a satisfactory fit to the data with Stark profiles., 1998) but Doyle and Byrne were unable to find a satisfactory fit to the data with Stark profiles.
 Instead they inferred svimetric red ancl blue shifted emitting regions. or turbulent motions. of order a few hundred km |.," Instead they inferred symmetric red and blue shifted emitting regions, or turbulent motions, of order a few hundred km $^{-1}$."
 This is typical of the broad component often seen in transition region and coronal emission lines which is often ascribed (o overlapping emission from multiple explosive microfLIaring events. leading to high velocity mass motions (Antonneci οἱ al.," This is typical of the broad component often seen in transition region and coronal emission lines which is often ascribed to overlapping emission from multiple explosive microflaring events, leading to high velocity mass motions (Antonucci et al."
 1993. Wood et al.," 1993, Wood et al."
 1996. Wood. Linsky Avres 1997. Villu οἱ al.," 1996, Wood, Linsky Ayres 1997, Vilhu et al."
 1993)., 1998).
 ILowever. those broad components are not associated with large Lares. but appear even in (he quiescent spectra of active (earlier-tvpe) stars. and (he Sun.," However, those broad components are not associated with large flares, but appear even in the quiescent spectra of active (earlier-type) stars, and the Sun."
 The phenomenon we have observed appears to be directly related to the laree flare heating event E9. and does nol appear in quiescence or in (he other flaring events (except possibly flare ET).," The phenomenon we have observed appears to be directly related to the large flare heating event F9, and does not appear in quiescence or in the other flaring events (except possibly flare F7)."
 In particular it does not occur in the other large flare we observed. {lave F2.," In particular it does not occur in the other large flare we observed, flare F2."
 In the present case. we do not expect that Stark broadening will be important for the Ae IH resonance lines.," In the present case, we do not expect that Stark broadening will be important for the Mg II resonance lines."
 Fleurier. Salial-Drechot Chapelle (1977) determined experimentally that the Stark broadening of these lines varied. with core width between0.05-0.10A.. and wines extending to ~0.1-0.2A.. for temperatures of 1—3x 101 and an electron density of NS10 7 much higher than expected in the M dwarf atmosphere.," Fleurier, Sahal-Brechot Chapelle (1977) determined experimentally that the Stark broadening of these lines varied with core width between, and wings extending to $\sim$, for temperatures of $1-3 \times 10^{4}$ K and an electron density of $N_e = 10^{17}$ $^{-3}$, much higher than expected in the M dwarf atmosphere."
 Figure 9 illustrates a simple Gaussian fit to the underlying broad component in F9b vielding a redshift of + ((53 £8 km !) and FWIIM of + (050 + 60 km !|)., Figure 9 illustrates a simple Gaussian fit to the underlying broad component in F9b yielding a redshift of $\pm$ (53 $\pm$ 8 km$^{-1}$ ) and FWHM of $\pm$ (250 $\pm$ 60 km $^{-1}$ ).
 IE interpreted as velocities. the EWIIM would suggest. turbulent emitiing material with tvpical speeds > LOO km | within a bulk condensation moving downward at roughly 50 km |.," If interpreted as velocities, the FWHM would suggest turbulent emitting material with typical speeds $>$ 100 km $^{-1}$ within a bulk condensation moving downward at roughly 50 km $^{-1}$."
 Both the bulk and turbulent. velocities exceed the sound speed in the chromosphere. which is roughly 10-20 km 1," Both the bulk and turbulent velocities exceed the sound speed in the chromosphere, which is roughly 10-20 km $^{-1}$ ."
 Adopting a chromospheric mass, Adopting a chromospheric mass
systematic errors (Drell et al.,systematic errors (Drell et al.
 2000)., 2000).
 Phe method described in this paper provides an independent determination of uA that relies only on the assumption of isotropy. and relatively well understood redshift cistortion models.," The method described in this paper provides an independent determination of $\Lambda$ that relies only on the assumption of isotropy, and relatively well understood redshift distortion models."
" The value of 3,(2~L4) obtained is almost identical to that. of local galaxies: e.g. af.~0)=0.43+0.07 determined using the 2dEP Galaxy Redshift Survey (Peacock et al.", The value of $\beta_q(z\sim1.4)$ obtained is almost identical to that of local galaxies; e.g. $\beta_{\rm g}(z\sim0) = 0.43\pm0.07$ determined using the 2dF Galaxy Redshift Survey (Peacock et al.
 2001)., 2001).
 If QSOs trace average galaxy environments. as sugeested by deep imaging around. QSOs (e.g. Croom Shanks 1999). measurements of QSO clustering (Croom οἱ al.," If QSOs trace average galaxy environments, as suggested by deep imaging around QSOs (e.g. Croom Shanks 1999), measurements of QSO clustering (Croom et al."
 2001a. Llovle et al.," 2001a, Hoyle et al."
 2001a). and the observation that most. if not all. major galaxies appear to host a supermassive black hole (Alagorrian 1998). then : appears to vary very little as a [function of redshift.," 2001a), and the observation that most, if not all, major galaxies appear to host a supermassive black hole (Magorrian 1998), then $\beta$ appears to vary very little as a function of redshift."
 With a larger QSO sample. we should be able to investigate the evolution of 3 by measuring it in narrower redshift bins.," With a larger QSO sample, we should be able to investigate the evolution of $\beta$ by measuring it in narrower redshift bins."
 “Phe main caveat would be whether we really are comparing the same population of objects at different recishifts., The main caveat would be whether we really are comparing the same population of objects at different redshifts.
 Using the simulation. we have demonstrated that a significant constraint on the A and 3? parameters will » available once the QSO survey is complete.," Using the simulation, we have demonstrated that a significant constraint on the $\Lambda$ and $\beta$ parameters will be available once the QSO survey is complete."
 The current survey. however. has fewer QSOs and a complicated window unction.," The current survey, however, has fewer QSOs and a complicated window function."
 To investigate if this window function has any systematic elfect on the LOk results we can again turn to heVolianc., To investigate if this window function has any systematic effect on the 10k results we can again turn to the.
 We have repeated the power spectral analysis of the simulation. using the current 2QZ window unction. when constructing a mock catalogue of 10000 QSOs.," We have repeated the power spectral analysis of the simulation, using the current 2QZ window function when constructing a mock catalogue of 10000 QSOs."
 The results are shown in figure 13.., The results are shown in figure \ref{fig10000hv}.
 The constraint on the A and; parameters obtained. [rom the LOk simulation is very similar to that. obtained from the Lok 20Z catalogue. although with slightly larger uncertainty.," The constraint on the $\Lambda$ and $\beta$ parameters obtained from the 10k simulation is very similar to that obtained from the 10k 2QZ catalogue, although with slightly larger uncertainty."
" The joint best fit values obtained. taking into account the mass clustering evolution method described in section 77.. are 3=O42 and O,,=0.20."," The joint best fit values obtained, taking into account the mass clustering evolution method described in section \ref{massclu}, are $\beta = 0.42$ and $\Omega_{\rm m} = 0.20$."
" Comparing the results derived from the LOk and 25k mock catalogues. there is a shift in the best fitting parameters to a slightly higher 3 for a given Oy, when the incomplete window function is considered. however the two constraints are entirely consistent within the errors. indicating that possible systematic elfects due to the Lok window function are fairly small."," Comparing the results derived from the 10k and 25k mock catalogues, there is a shift in the best fitting parameters to a slightly higher $\beta$ for a given $\Omega_{\rm m}$ when the incomplete window function is considered, however the two constraints are entirely consistent within the errors, indicating that possible systematic effects due to the 10k window function are fairly small."
 Lf this shift is systematic then taking it into account would slightly weaken the significance of our A constraint from the 10k 2QZ catalogue., If this shift is systematic then taking it into account would slightly weaken the significance of our $\Lambda$ constraint from the 10k 2QZ catalogue.
"Onemay - xf MAUR ΤΕΝ.αλ) = oh, eOUPE(QN. A) =GNI) = Ζ= ΕΙΝ.","where This can be obtained either from the full second order equations of motion, or from the first order $N=1$ supersymmetry equations \cite{r2stabeq} similarly to \cite{SJR}."
 A)− lim Fu. Y)-qY valid dueto, At the horizon $r=0$ ) we have: The horizon solution agrees with previous results \cite{r2entropy}.
 (he Vo= 2 offlat χοπ/y2fs (X τα PX Zis deter, The entropy is given by the Wald formula for the supersymmetric case with $R^2$ -terms \cite{r2entropy}: where $A$ is the area of the event horizon.
minedby (he charges andthe asvanptotce modulivalues at infinity. and , Plugging in the solution gives the expected result: In the supersymmetric case the ADM mass saturates the BPS bound: This result is exact and is identical to that of the $R$ -level.
does not equality (33)) imposed Our caleu," Note that in the special case $k^3=k_0$, the $\xi$ -functions vanish in all space."
lalions doneusing Maple with GRTensor. Black Holes, This means that the $R^2$ -level solution is simply the $R$ -level ansatz with $F(\epsilon=0)$replaced by $F(\epsilon)$.
 3.1 Supersymmetric In the supersvi, One may ask whether this behavior continues to higher orders in $\epsilon$.
metric (qu 0) withA! A? &Á*. the solutionreads," Assuming that the prepotential \ref{prep}) ) itself does not contain higher orders in $\epsilon$ , the solution to order $\epsilon^2$ would be"
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"R (eq. [2]]),","$R$ (eq. \ref{eq:R}] ]),"
 then the distribution of the R values will be givenby where C is an appropriate normalization constant., then the distribution of the R values will be givenby where $C$ is an appropriate normalization constant.
 The distribution P(E;c) peaks at R=0 for all values of σ and drops quickly to zero such that the median value of R for this distribution is Rs9%= c., The distribution $P(R;\sigma)$ peaks at $R=0$ for all values of $\sigma$ and drops quickly to zero such that the median value of $R$ for this distribution is $R_{\rm 50\%}=\sigma$ .
 Given that half of our 6, Given that half of our 6
Lorentz [actor in the initial data being ο8.,Lorentz factor in the initial data being $\sim 8$.
 It was concluded that the Dlandford-MceIxee blastwave constitutes a problem Chat is too demanding to be a useful validation test unless very significant computer resources are emploved., It was concluded that the Blandford-McKee blastwave constitutes a problem that is too demanding to be a useful validation test unless very significant computer resources are employed.
 We have (hus opted (o simulate the 3D-equivalent of the shock tube problem: the initial condition for the 3D relativistie blastwave problem comprises (wo piecewise constant states separated bv a discontinuity al an arbitrary radius r=1.0., We have thus opted to simulate the 3D-equivalent of the shock tube problem: the initial condition for the 3D relativistic blastwave problem comprises two piecewise constant states separated by a discontinuity at an arbitrary radius $r=1.0$.
 A munerical solution with spherical sviumetry. and initial conditions that produced mareinally relativistic [low speed. was presented by vanPutten(1994): here we adopt an initial state (he leads to higher Lorentz factors of order those encountered in (he applications discussed in (he next sections.," A numerical solution with spherical symmetry, and initial conditions that produced marginally relativistic flow speed, was presented by \citet{vp3}; here we adopt an initial state the leads to higher Lorentz factors – of order those encountered in the applications discussed in the next sections."
 The inner and outer states have pressure and rest density: p;=104. po=10.0. n;=1.0. and no—0.1.," The inner and outer states have pressure and rest density: $p_I=10^4$ , $p_O=10.0$, $n_I=1.0$, and $n_O=0.1$."
 The fluid is initially at rest ancl (he acliabatic index is taken to be Y=4/3., The fluid is initially at rest and the adiabatic index is taken to be $\Gamma=4/3$.
 For time |o0 a spherical shock wave propagates to larger radius., For time $t>0$ a spherical shock wave propagates to larger radius.
 Substantial Lorentz factors (~5) are generated. and by the end of the simulation the propagating structure encompasses e35 times (he volume of the initial high pressure sphere. which is ample development for isolating asvmmnietries (hal might arise.," Substantial Lorentz factors $\sim 5$ ) are generated, and by the end of the simulation the propagating structure encompasses $\sim 35$ times the volume of the initial high pressure sphere, which is ample development for isolating asymmetries that might arise."
 Some asvmmeltry is (o be expected. because (he criteria used for flageine cells to be refined is applied automatically. ancl a combination ol differences due to rounding. coupled with the clustering algorithm. leads to a non-uniform distribution of refined patches.," Some asymmetry is to be expected, because the criteria used for flagging cells to be refined is applied automatically, and a combination of differences due to rounding, coupled with the clustering algorithm, leads to a non-uniform distribution of refined patches."
 Taking the computational domain to be c—[—10.0.10.0]. y=[—10.0.10.0). and 10.0). Figure 3. shows the final evolution (at /—2.3967) of a run with 3 levels of refinement. vielding ~180 finest level cells across the blastwave diameter at the end of the computation.," Taking the computational domain to be $x = [-10.0,10.0]$, $y = [-10.0,10.0]$, and $z = [-10.0,10.0]$ , Figure \ref{blast} shows the final evolution (at $t=2.3967$ ) of a run with $3$ levels of refinement, yielding $\sim 180$ finest level cells across the blastwave diameter at the end of the computation."
 The (vo panels show laboratory frame density (42) and Lorentz factor (5) for cuts along the (wo orthogonal coordinate directions in the plane wr=0. plus additional cuts that bisect these two axes.," The two panels show laboratory frame density $R$ ) and Lorentz factor $\gamma$ ) for cuts along the two orthogonal coordinate directions in the plane $x=0$, plus additional cuts that bisect these two axes."
 The peak laboratory [rame density. along the diagonal cuts is 13% higher than along the coordinate directions., The peak laboratory frame density along the diagonal cuts is $13\%$ higher than along the coordinate directions.
 This is associated with an oscillation in the Lorentz lactor near lo ils peak value ~5. in the cuts that bisect the coordinate directions.," This is associated with an oscillation in the Lorentz factor near to its peak value $\sim 5$, in the cuts that bisect the coordinate directions."
" The oscillation in the Lorentz factor cuts is a result of the fact that even with an effective resolution of ~5.8x10"" cells encompassing the blastwave. the scale of the leading edge of the evolving siructure is so narrow that the Lorentz factor varies [rom ~5 to ~4 over a scale of a couple of the finest cells."," The oscillation in the Lorentz factor cuts is a result of the fact that even with an effective resolution of $\sim 5.8 \times
10^6$ cells encompassing the blastwave, the scale of the leading edge of the evolving structure is so narrow that the Lorentz factor varies from $\sim 5$ to $\sim 4$ over a scale of a couple of the finest cells."
 OL axis. the chance. and changing. location of cells with respect to the thin shell of high Lorentz factor leads to a small sunplitude irregularity in the peak value. evident as f[Iuctuation along the radial eut.," Off axis, the chance, and changing, location of cells with respect to the thin shell of high Lorentz factor leads to a small amplitude irregularity in the peak value, evident as fluctuation along the radial cut."
 Notice however that an inability to fully capture (he extreme values encountered on the finest scale. does not in general influencethe global location of major flow. structures or their values away [rom these extremes.," Notice however that an inability to fully capture the extreme values encountered on the finest scale, does not in general influencethe global location of major flow structures or their values away from these extremes."
 In particular. (here is no global asvimetry.," In particular, there is no global asymmetry."
planet detections.,planet detections.
 Even when zero planets are detected (a null result) such a calculation places upper limits on the planet fraction., Even when zero planets are detected (a null result) such a calculation places upper limits on the planet fraction.
 In this paper we describe a Monte Carlo method. for calculating detection probabilities (and false alarm. rates) of transiting planets based on photometric data. as functions of various parameters. taking into account the following factors: We then apply the method to the transit survey deseribed in Bramichetal.2005. there on referred to as BRAOS) to determine the expected number transiting planet detections and place limits on the planet fraction ofas a function of star and planet type.," In this paper we describe a Monte Carlo method for calculating detection probabilities (and false alarm rates) of transiting planets based on photometric data, as functions of various parameters, taking into account the following factors: We then apply the method to the transit survey described in \citealt{bra05} (here on referred to as BRA05) to determine the expected number of transiting planet detections and place limits on the planet fraction as a function of star and planet type."
 In Section 2 we describe the lighteurve data used in the analysis and in Section 3 we define the detection probabilities and false alarm probabilities for an extra-solar planet based on photometric data., In Section 2 we describe the lightcurve data used in the analysis and in Section 3 we define the detection probabilities and false alarm probabilities for an extra-solar planet based on photometric data.
 In Section + we present the Monte Carlo method that we used to calculate these probabilities and derive limits on the hot Jupiter fraction in the field of NGC 7789 as a function of star and planet type., In Section 4 we present the Monte Carlo method that we used to calculate these probabilities and derive limits on the hot Jupiter fraction in the field of NGC 7789 as a function of star and planet type.
 In Section 5 we discuss the results and in Section 6 we present our conclusions., In Section 5 we discuss the results and in Section 6 we present our conclusions.
 A transit survey of the field of NGC 7789 was presented in BRAOS in which ~33000 stars were photometrically monitored in the Sloan r band over three separate runs with dates 1999 June 22-30. 1999 July 22-31 and 2000 September 10-20.," A transit survey of the field of NGC 7789 was presented in BRA05 in which $\sim$ 33000 stars were photometrically monitored in the Sloan $r^{\prime}$ band over three separate runs with dates 1999 June 22-30, 1999 July 22-31 and 2000 September 10-20."
 For brevity these runs shall be refered to from now on as 1999-06. 1999-07 and 2000-09 respectively.," For brevity these runs shall be refered to from now on as 1999-06, 1999-07 and 2000-09 respectively."
 To summarise. in BRAOS. Sloan +’i’ colour indices were used to construct a colour-magnitude diagram (CMD) and thereby identify the cluster main sequence.," To summarise, in BRA05, Sloan $r^{\prime}-i^{\prime}$ colour indices were used to construct a colour-magnitude diagram (CMD) and thereby identify the cluster main sequence."
 Fig., Fig.
 | shows the CMD for the stars from chip 4 which was centred on the cluster., \ref{fig:CMD} shows the CMD for the stars from chip 4 which was centred on the cluster.
 Although the cluster main sequence is visible. it is clear that most stars in the sample are field stars and not cluster stars.," Although the cluster main sequence is visible, it is clear that most stars in the sample are field stars and not cluster stars."
" A theoretical main sequence model for the stellar mass range 0.05Al,—M.<=LAOAL. was adopted and fitted to the cluster main sequence via magnitude offsets.", A theoretical main sequence model for the stellar mass range $0.08~M_{\sun}~\leq~M_{*}~\leq~1.40~M_{\sun}$ was adopted and fitted to the cluster main sequence via magnitude offsets.
" Using the known cluster distance d,=2337pe and reddening (D.V)=0.217. and adopting an Einasto law for the distribution of the interstellar medium in the Milky Way (Robinetal.20033). a distance d; was determined for each star such that the theoretical main sequence passes through the star's position on the CMD."," Using the known cluster distance $d_{\mbox{\small c}} = 2337$ pc and reddening $E(B~-~V)~=~0.217$, and adopting an Einasto law for the distribution of the interstellar medium in the Milky Way \citealt{rob03}) ), a distance $d_{*}$ was determined for each star such that the theoretical main sequence passes through the star's position on the CMD."
 It was argued that giant stars lie beyond the edge of the galaxy in order to be in the image data., It was argued that giant stars lie beyond the edge of the galaxy in order to be non-saturated in the image data.
" Hence. it was assumed that each star is on the main sequence. and after determining the star's distance d. the star's mass AZ, and radius £2, could be read off from its position on the theoretical main sequence."," Hence, it was assumed that each star is on the main sequence, and after determining the star's distance $d_{*}$, the star's mass $M_{*}$ and radius $R_{*}$ could be read off from its position on the theoretical main sequence."
 In this paper we consider the 32027 stars from this data set that have a lighteurve from the 2000-09 run. and an assigned distance. mass and radius.," In this paper we consider the 32027 stars from this data set that have a lightcurve from the 2000-09 run, and an assigned distance, mass and radius."
 The remaining stars with lightcurves lack a colour measurement or were too blue to be assigned a mass and radius using the adopted theoretical main. sequence., The remaining stars with lightcurves lack a colour measurement or were too blue to be assigned a mass and radius using the adopted theoretical main sequence.
 We also consider the lightcurve data from the 1999-07 run where it exists., We also consider the lightcurve data from the 1999-07 run where it exists.
 BRAOS searched for transits in the LO-night 1999-07 run and the |1-night 2000-09 run. I4 months later.," BRA05 searched for transits in the 10-night 1999-07 run and the 11-night 2000-09 run, 14 months later."
 The 1999-06 run was too sparsely sampled in time to support transit hunting by the adopted search technique., The 1999-06 run was too sparsely sampled in time to support transit hunting by the adopted search technique.
 We ure interested in the expected number of transiting planet detections (and false alarms) for stars of different masses or. equivalently. spectral types.," We are interested in the expected number of transiting planet detections (and false alarms) for stars of different masses or, equivalently, spectral types."
 To facilitate this analysis we consider 4+ mutually exclusive subsets of stars which in union make up the set of all 32027 stars., To facilitate this analysis we consider 4 mutually exclusive subsets of stars which in union make up the set of all 32027 stars.
 These sets are the late F stars. G stars. K stars and M stars respectively.," These sets are the late F stars, G stars, K stars and M stars respectively."
 Table | shows the number of stars in each set and the spectral type/mass ranges to which they correspond., Table \ref{tab:spectypes} shows the number of stars in each set and the spectral type/mass ranges to which they correspond.
 The table also includes the number of stars for which 1999-07 lightcurve data exists., The table also includes the number of stars for which 1999-07 lightcurve data exists.
 The mass ranges for the various spectral types are taken from Lang(1992)., The mass ranges for the various spectral types are taken from \citet{lan92}.
" In BRAOS. a matched filter algorithm was used to search for transits in the lighteurves by adopting a square ""boxear"" shape for the transit lighteurve of total width S.A/ (where Af is the transit duration searched for)."," In BRA05, a matched filter algorithm was used to search for transits in the lightcurves by adopting a square “boxcar” shape for the transit lightcurve of total width $5 \Delta t$ (where $\Delta t$ is the transit duration searched for)."
" This search was based on the transit detection statistic: where xj is the chi squared of the boxcar transit fit. \onst is the chi squared of the constant fit. vou, is the chi squared of the boxcar transit fit for the Noy, out-of-transit data points."," This search was based on the transit detection statistic: where $\chi^{2}_{\mbox{\small tra}}$ is the chi squared of the boxcar transit fit, $\chi^{2}_{\mbox{\small const}}$ is the chi squared of the constant fit, $\chi^{2}_{\mbox{\small out}}$ is the chi squared of the boxcar transit fit for the $N_{\mbox{\small out}}$ out-of-transit data points."
 Transit candidates were chosen using a threshold of Si;-—ὃmin= 10., Transit candidates were chosen using a threshold of $S_{\mbox{\small tra}}~\geq~S_{\mbox{\small min}}~=~10$ .
Stellar evolutioi theories. which potentially provide tlie best tools lor computing z. give results that critically deperd on the adopted initial mass fuuction (IME) aud the low masses cutolL,"Stellar evolution theories, which potentially provide the best tools for computing $\tau$, give results that critically depend on the adopted initial mass function (IMF) and the low masses cutoff."
 These Parameters are rat]er uncertain aud unply large differences on the derived 7., These parameters are rather uncertain and imply large differences on the derived $\tau$.
 For iustauce. for a single coeval stellar population. assumiug a sinele-slope Salpeter INF with slope 2.35. 7 chiauges by a factor 1.5 varyiue[n] the low masses cutoll [rom 0.15 to 0.05 MN...," For instance, for a single coeval stellar population, assuming a single-slope Salpeter IMF with slope 2.35, $\tau$ changes by a factor 4.5 varying the low masses cutoff from 0.15 to 0.05 $_{\odot}$."
 Note that this problem is ouly uuildly alleviatecd usiug a 1uore realistic multi-slope IMF (e.g..Scalo1986). at low masses ," Note that this problem is only mildly alleviated using a more realistic multi-slope IMF \citep[e.g.,][]{scalo86} at low masses \citep{renzini93}."
Au uubiased way to estimate 7 relies on the analysis of the stars in the solar neighborhood. where sars of all spectral type in all possible evolutionary phases can be stucied in great detalL," An unbiased way to estimate $\tau$ relies on the analysis of the stars in the solar neighborhood, where stars of all spectral type in all possible evolutionary phases can be studied in great detail."
 Fro; BinneyaudMerrifield (1998).. the total mass of stars in LO! pe? is 356M.. + 30M. of whie dwa‘fs.," From \citet{binney98}, , the total mass of stars in $10^4$ $^3$ is $_{\odot}$ + $_{\odot}$ of white dwarfs."
 For this sample it is found 7Q.T., For this sample it is found $\tau = 0.7$.
 [unis sunple of stars can be used to estimate 7 for any stellar population simpy integratiu the light over al the relevant spectral types. while maiutainiug the total mass fixed.," This sample of stars can be used to estimate $\tau$ for any stellar population simply integrating the light over all the relevant spectral types, while maintaining the total mass fixed."
 This will result ina =yper limit to the true 7. because more than one geueratiou of stars contributed to it ai therefore the numjer of low lass stars is overrepreseηος.," This will result in an upper limit to the true $\tau$, because more than one generation of stars contributed to it and therefore the number of low mass stars is overrepresented."
 h ellipical exàilixles the stellar population is composed by stars iu: 5tb-ejaut branch. red-giant branch. horizonliil |anch. aud main sequen'e up to specral type In. I1.104 pe? there are | red eiauts of Ix type axl 0.25 of M type poptlatine ie plant branch. which 1ave average luminosity of 5T and 120 solar hiuninuosities. respectivey.," In elliptical galaxies the stellar population is composed by stars in: sub-giant branch, red-giant branch, horizontal branch, and main sequence up to spectral type K. In $10^4$ $^3$ there are 4 red giants of K type and 0.25 of M type populating the giant branch, which have average luminosity of 57 and 120 solar luminosities, respectively."
 For ach star in the giant brauch there are ~ 21 stars in he sub-giai| bratch aud 0.75 in the horiZOII€c ‘auch (Rewink1998)... wuich contribute for 5 and 22 L.. respectively.," For each star in the giant branch there are $\sim$ 21 stars in the sub-giant branch and 0.75 in the horizontal branch \citep{renzini98}, which contribute for 5 and 52 $L_{\odot}$, respectively."
 Stars still in the maill secce add alother 18 L., Stars still in the main sequence add another 18 $L_{\odot}$.
.F om these Consideratious it follows tlal 7=.386/120~0.9., From these considerations it follows that $\tau=386/420 \sim 0.9$.
 This valie is cousistent Wwith that derived for globular clusters (Maucdushev.StauevaandSpasova1991:RetziuianclCiotti1993) that lave a similar old stellar »opulation.," This value is consistent with that derived for globular clusters \citep{mandushev91,renzini93}
 that have a similar old stellar population."
 These objects. however. are slighty different form eliptical galaxies because {μον cout contalu gas and at least a fraction of the less massive stars. which contributes to the mass but not o the luminosity. have probably left the clusters long ago (e.g.Aguilar.HutancSuithaudBurkert 2002).," These objects, however, are slightly different form elliptical galaxies because they don't contain gas and at least a fraction of the less massive stars, which contributes to the mass but not to the luminosity, have probably left the clusters long ago \citep[e.g.][]{aguilar88, smith02}."
. So. lor globular cluster 7=1 should be fairly correct. while rere T shall sel 7=2 [or all elliptical galaxies.," So, for globular cluster $\tau=1$ should be fairly correct, while here I shall set $\tau=2$ for all elliptical galaxies."
 A fairly conservative assumptio1 that allows for the presence of large amount of gas. particularly relevant in LSB galaxies.," A fairly conservative assumption that allows for the presence of large amount of gas, particularly relevant in LSB galaxies."
 For spiral galaxies. setting 7=1.2. that is about twice as bij as that observed in the solar nelelborboocd. allows for plenty of gas aud the bulge old stellar population.," For spiral galaxies, setting $\tau=1.2$, that is about twice as big as that observed in the solar neighborhood, allows for plenty of gas and the bulge old stellar population."
 Finally. for cluster of ealaxies I assuie a uix of spiral and ellipticals. with 71.6.," Finally, for cluster of galaxies I assume a mix of spiral and ellipticals, with $\tau \sim 1.6$."
 EL stressthat the resultsreported iu this work do not siguificantly cepeud on the assumed value of 7 ithe rauge 1-2., I stressthat the resultsreported in this work do not significantly depend on the assumed value of $\tau$ inthe range 1-2.
 I compute here the quantity g/gx aud investigate its relation with o. r.. aud £ for a sample ol elliptical galaxies," I compute here the quantity $g/g_N$ and investigate its relation with $\sigma$ , $r_e$ , and $L$ for a sample of elliptical galaxies."
The authors acknowledge the work of Claudio Linares in early stages of the project and Alexander IXnebe for his help in setting up his AUF code to be used in our simulations.,The authors acknowledge the work of Claudio Llinares in early stages of the project and Alexander Knebe for his help in setting up his AHF code to be used in our simulations.
 The simulation has been done at (he Barcelona Supercomputer Center and analyzed at NIC Juelich., The simulation has been done at the Barcelona Supercomputer Center and analyzed at NIC Juelich.
 GY acknowledges support of MEC. (Spain) through research grants FPA2009-08958 and AYA2009-13815-C03-02., GY acknowledges support of MEC (Spain) through research grants FPA2009-08958 and AYA2009-13875-C03-02.
 We also acknowledge the CONSOLIDER-INGENIO projects MULTIDARIE (CSD2009-00064) and SvEC (CSD2007-0050) for supporting our collaboration., We also acknowledge the CONSOLIDER-INGENIO projects MULTIDARK (CSD2009-00064) and SyEC (CSD2007-0050) for supporting our collaboration.
 We would also like to thank the anonvinous referee for her/his valuable comments., We would also like to thank the anonymous referee for her/his valuable comments.
than 40 AU escape. although several of the objects that did escape evolved to perihelia as large as 50toGO AU before escaping.,"than 40 AU escape, although several of the objects that did escape evolved to perihelia as large as 50–to–60 AU before escaping."
 A simple property we demand of our debiased model is that it be nearly stationary in lime. i.e.. (hat the model distribution not Favor a special time: this provides basic confidence in our debiasing procedure.," A simple property we demand of our debiased model is that it be nearly stationary in time, i.e., that the model distribution not favor a special time; this provides basic confidence in our debiasing procedure."
 To determine if our model distribution satisfies (his. we perform a Ixolnogorov-Snmirnov test (adapted [rom ?)) to compare the distributions of the orbital elements (semi-major axis. eccentricitv. perihelion distance and inclination). e. e.g. aad 7. al various times in the simulation.," To determine if our model distribution satisfies this, we perform a Kolmogorov-Smirnov test (adapted from \citet{press92}) ) to compare the distributions of the orbital elements (semi-major axis, eccentricity, perihelion distance and inclination), $a$ , $e$ , $q$, and $i$, at various times in the simulation."
 The test measures the absolute difference in the cumulative distributions of (wo samples of any given orbital element and (ranslates (lis into a probability. which we will call the INS probability’. (hat the (wo samples are drawn Irom (he same parent distribution.," The test measures the absolute difference in the cumulative distributions of two samples of any given orbital element and translates this into a probability, which we will call the `KS probability', that the two samples are drawn from the same parent distribution."
 This is accomplished by comparing the absolute clillerence between (he samples cumulative distributions to the distribution of differences expected lor (wo samples that are drawn from one distribution., This is accomplished by comparing the absolute difference between the samples' cumulative distributions to the distribution of differences expected for two samples that are drawn from one distribution.
 If the debiasing procedure does create a model distribution that is representative of the scattered disk. we expect the IX-9 test to vield a high NS probability for the distributions of each orbital element compared at dillerent (mes in the simulation.," If the debiasing procedure does create a model distribution that is representative of the scattered disk, we expect the K-S test to yield a high KS probability for the distributions of each orbital element compared at different times in the simulation."
 We find that the INS probability is nearly [or the inclination distributions compared al the beginning and the end of the simulation. indicating that the inclination distribution obtained in the debiasing procedure is stationary on 4 Qwr timescales.," We find that the KS probability is nearly for the inclination distributions compared at the beginning and the end of the simulation, indicating that the inclination distribution obtained in the debiasing procedure is stationary on 4 Gyr timescales."
 The eccentricity distributions at 0 Gyr and al 4 Gvr give a IS probability of81%... whereas the distributions al 1 Gyr and at 4 Gyr vield nearlyL00%.," The eccentricity distributions at 0 Gyr and at 4 Gyr give a KS probability of, whereas the distributions at 1 Gyr and at 4 Gyr yield nearly."
. The INS probabilities for the semi-major axis and perihelion distance distributions at 0 Gyr and at 4 Gyr are only and556... respectively.," The KS probabilities for the semi-major axis and perihelion distance distributions at 0 Gyr and at 4 Gyr are only and, respectively."
 The probability for the semimajor axis reaches nearly when we consider the distributions ab 1 Gyr and 4 Gyr., The probability for the semimajor axis reaches nearly when we consider the distributions at 1 Gyr and 4 Gyr.
 For the perihelion distance. the INS probability increases to when we consider the distributions at 1 Gyr and 4 Gyr. and it is nearly for the comparison between 2 Gvr and 4 Gyr.," For the perihelion distance, the KS probability increases to when we consider the distributions at 1 Gyr and 4 Gyr, and it is nearly for the comparison between 2 Gyr and 4 Gyr."
 The improvement in the INS probability when (he first gigavear of (he simulation is excluded indicates that the initial population has a transient subset. but the distribution stabilizes alter 1 Gvr and remains stationary for the remainder of the simulation.," The improvement in the KS probability when the first gigayear of the simulation is excluded indicates that the initial population has a transient subset, but the distribution stabilizes after 1 Gyr and remains stationary for the remainder of the simulation."
 This can be seen most clearly in the evolution of the perihelion distribution. which is shown in Figure 4..," This can be seen most clearly in the evolution of the perihelion distribution, which is shown in Figure \ref{f:qdist_time}."
 The objects with a verv low probability of discovery. that. initially dominate parts of the distributions of q aud α due to the debiasing procedure are either lost from the scattered disk or evolve to more stable regions of phase space during the first eigavear., The objects with a very low probability of discovery that initially dominate parts of the distributions of $q$ and $a$ due to the debiasing procedure are either lost from the scattered disk or evolve to more stable regions of phase space during the first gigayear.
 This supports our hypothesis that the imperfectionsofthe debiasing procedure. specifically," This supports our hypothesis that the imperfectionsofthe debiasing procedure, specifically"
Discovered in. 1963. (Schmidt: 1963). quasars are the most uminous continuously emitting objects in the Universe and represent the veh luminosity enc of the class. οἱ objects known as Active Galactic Nuclei (AGN).,"Discovered in 1963 (Schmidt 1963), quasars are the most luminous continuously emitting objects in the Universe and represent the high luminosity end of the class of objects known as Active Galactic Nuclei (AGN)."
 Like their ower Luminosity cousins - Sevlert 1 galaxies - the bull of he energy produced. in quasars is thought to arise [rom accretion onto a compact object (the putative super-massive Mack hole)., Like their lower luminosity cousins - Seyfert 1 galaxies - the bulk of the energy produced in quasars is thought to arise from accretion onto a compact object (the putative super-massive black hole).
 Vhis central engine is also thought to be where he N-ravs. that are observed from both quasars and Sevfert ls. originate [rom.," This central engine is also thought to be where the X-rays, that are observed from both quasars and Seyfert 1s, originate from."
 In one model. the UV. photons produced. by. viscous dissipation in an accretion disk are Comptonised. to. N-rav energies by a hot corona above the surface of this cisk (Llaardt Maraschi 1993).," In one model, the UV photons produced by viscous dissipation in an accretion disk are Comptonised to X-ray energies by a hot corona above the surface of this disk (Haardt Maraschi 1993)."
" These hare X-rays in turn illuminate the accretion disk. being either ""rellected back towards the observer or thermalised. into optical or UV photons."," These hard X-rays in turn illuminate the accretion disk, being either `reflected' back towards the observer or thermalised into optical or UV photons."
" Evidence for these ""rellection! features (in the form of an iron Ίνα emission line. Fe Ix. absorption edge and Compton scattering. hump) is commonly. observed. in. the X-ray spectral band in Sevfert 1: galaxies (e.g. Pounds 1990. Nandra Pounds 1994)."," Evidence for these `reflection' features (in the form of an iron $\alpha$ emission line, Fe K absorption edge and Compton scattering hump) is commonly observed in the X-ray spectral band in Seyfert 1 galaxies (e.g. Pounds 1990, Nandra Pounds 1994)."
 The detection of these reflection features however. in quasars. remains more ellusive (Reeves 1997. Lawson Turner 1997).," The detection of these reflection features however, in quasars, remains more ellusive (Reeves 1997, Lawson Turner 1997)."
 In thefoud quasars the situation is somewhat further complicated by the presence of a powerful relativistic radio-jet., In the quasars the situation is somewhat further complicated by the presence of a powerful relativistic radio-jet.
 In the X-ray band. these racio-loucl quasars have Blatter. X-ray spectral emission (e.g. Wilkes," In the X-ray band, these radio-loud quasars have flatter X-ray spectral emission (e.g. Wilkes"
The same single-mode analysis shows that the limit πι--o vields an uncertainty in the power spectrum amplitude aac71V2Lea sample variance limit equal to that for Infer in the MS method.,"The same single-mode analysis shows that the limit $nP\rightarrow\infty$ yields an uncertainty in the power spectrum amplitude $\sigma_{\ln G}\rightarrow 1/\sqrt{2}$, a sample variance limit equal to that for $\ln fG$ in the MS method."
" Figures 2. plots the uncertainty in the growth measures vs the minimum: mass My, of halo included in a fiducial redshift survey.", Figures \ref{sigtheta} plots the uncertainty in the growth measures vs the minimum mass $M_{\rm min}$ of halo included in a fiducial redshift survey.
 We make the following assumptions for the [icbucial survey: lt ds desirable to measure the growth function to percent-level accuracy dini several redshift bins of. this width at z«1 in order to test general relativity (in combination with percent-level constraints on a(/) from distance-measurement methods)., We make the following assumptions for the fiducial survey: It is desirable to measure the growth function to percent-level accuracy in several redshift bins of this width at $z<1$ in order to test general relativity (in combination with percent-level constraints on $a(t)$ from distance-measurement methods).
" We note that the total number of independent Fourier mocdes in the fiducial survey volume is only N,,,=1150 (48.000) for Ais=0.035Alpet (0.1)."," We note that the total number of independent Fourier modes in the fiducial survey volume is only $N_m=1150$ (43,000) for $k_{\rm max}=0.03h\,{\rm Mpc}^{-1}$ (0.1)."
" We recall that the standard RSD method. has a sample-variance floor of ej,re;=VerAS and therefore a measure of fC is not attainable for fu:0.1hAlpe"," We recall that the standard RSD method has a sample-variance floor of $\sigma_{\ln fG}=\sqrt{21/N_m}$ and therefore a measure of $fG$ is not attainable for $k_{\rm max}\le0.1h\,{\rm Mpc}^{-1}$."
" The sample variance limit of the MS. method. is 9,je;=velfyFNAL2N,, and the meaurement is. potentially. achievable. in the optimistic case.", The sample variance limit of the MS method is $\sigma_{\ln fG}=\sqrt{1/2N_m}$ and the meaurement is potentially achievable in the optimistic case.
" Figures 2. plot the uncertainty in InfC€ for the Standard RSD and MS methods. and then plot ay, and oye: independentlv for the Fixed. Bias and the Single Bias Aleasurement cases."," Figures \ref{sigtheta} plot the uncertainty in $\ln fG$ for the Standard RSD and MS methods, and then plot $\sigma_{\ln f}$ and $\sigma_{\ln G}$ independently for the Fixed Bias and the Single Bias Measurement cases."
 In this last case.we plot results from two levels of uncertainty on the weighted bias 6.," In this last case,we plot results from two levels of uncertainty on the weighted bias $\bar b$."
" “Pwo horizontal axes plot the Ai, for the halo survey. ancl the total number of halos in the volume. the number of redshift measurements recuired for the survey."," Two horizontal axes plot the $M_{\rm min}$ for the halo survey, and the total number of halos in the volume, the number of redshift measurements required for the survey."
" The Standard RSD analysis (in blue) approaches a sample-variance-limited. plateau as expected. with very little improvement once Ny10 10"" halo redshifts are obtained."," The Standard RSD analysis (in blue) approaches a sample-variance-limited plateau as expected, with very little improvement once $N_H>10^5$ $10^6$ halo redshifts are obtained."
 The AMIS analysis (purple) shows significant improvement over the Standard. RSD analysis when c10° redshifts are obtained. or halos «10175.ΣΑΙ.," The MS analysis (purple) shows significant improvement over the Standard RSD analysis when $>10^6$ redshifts are obtained, or halos $<10^{13}h^{-1} M_\odot$."
" ""Ehe cosmic variance limit for the AIS analysis is oufe:l/v2N, (dotted. purple line) but. this is not attained even with surveys of =»10° redshifts.", The cosmic variance limit for the MS analysis is $\sigma_{\ln fG} > 1/\sqrt{2N_m}$ (dotted purple line) but this is not attained even with surveys of $>10^9$ redshifts.
 This is attributable to the limited number of halos with 6>1. as explained below.," This is attributable to the limited number of halos with $b>1$, as explained below."
 La the conservative case of fuss=0.03Alpe που(αλα!Les of «10% in ου in this bin for the ALS method.," In the conservative case of $k_{\rm max}=0.03h\,{\rm Mpc}^{-1}$, uncertainties of $<10\%$ in $fG$ require $N_h>10^7$ redshifts to be measured in this bin for the MS method."
" If the theory. supports baas=0.15Alpe then only Ny,<10"" reclshifts are needed to reach accuracy. and 107 redshifts viele a measure of fC."," If the theory supports $k_{\rm max}=0.1h\,{\rm Mpc}^{-1}$, then only $N_h<10^5$ redshifts are needed to reach accuracy, and $10^7$ redshifts yield a measure of $fG$."
" Uncertainties on fC scale roughly as NV,L7 when using the standard or AIS methods at IN;«107.", Uncertainties on $fG$ scale roughly as $N_h^{-1/2}$ when using the standard or MS methods at $N_h<10^5$.
 In practice a survey has a trade between depth ancl sky coverage. aud one can ask whether better cosmological constraints result [rom a larec. sparse redshift survey or a deeper. smaller-area survey given a fixed survey cost (cluration).," In practice a survey has a trade between depth and sky coverage, and one can ask whether better cosmological constraints result from a large, sparse redshift survey or a deeper, smaller-area survey given a fixed survey cost (duration)."
" HE every recishift can be attained at equal cost. then this scaling implies a Hat trade. i.c. constraint on {Co are independent. of survey area when Ny,«107."," If every redshift can be attained at equal cost, then this scaling implies a flat trade, i.e. constraint on $fG$ are independent of survey area when $N_h<10^5$."
" However the o-vs-Ny, curve Lattens at higher Ny. whieh indicates that one should favor area over depth when Ny,7107."," However the $\sigma$ $N_h$ curve flattens at higher $N_h$, which indicates that one should favor area over depth when $N_h>10^5$."
 Phis would be even more true if redshifts are more expensive to obtain in lower-mass halos., This would be even more true if redshifts are more expensive to obtain in lower-mass halos.
 The red lines show the uncertainties on In€ (solid) and Inf£ (dashed) in the Fixed Bias case of perfect lensing calibration of all galaxy biases., The red lines show the uncertainties on $\ln G$ (solid) and $\ln f$ (dashed) in the Fixed Bias case of perfect lensing calibration of all galaxy biases.
 We note first that constraints on Inf (dashed red) are obtained at accuracy similar to the InfC constraints with the MS method at equal Ny., We note first that constraints on $\ln f$ (dashed red) are obtained at accuracy similar to the $\ln fG$ constraints with the MS method at equal $N_h$.
 Next we note that the In€ constraint (solid. red) approaches a sample-variance limited. plateau much more rapidly than does the AIS method., Next we note that the $\ln G$ constraint (solid red) approaches a sample-variance limited plateau much more rapidly than does the MS method.
 For A=0.035Alpe one can obtain errors in InC with a survey of only ~LO’ halos at Al>10thTAL... 104 times fever redshifts (han are needed [or similar {Co constraints using the AIS method alone. and completely. impossible with Standard. RSD methods.," For $k_{\rm max}=0.03h\,{\rm Mpc}^{-1}$, one can obtain errors in $\ln G$ with a survey of only $\sim10^5$ halos at $M>10^{14}h^{-1}M_\odot$, $10^4$ times fewer redshifts than are needed for similar $fG$ constraints using the MS method alone, and completely impossible with Standard RSD methods."
 This is attributable to the lensing information breaking the degencracy between bias and matter power spectrum. amplitude., This is attributable to the lensing information breaking the degeneracy between bias and matter power spectrum amplitude.
" The dashed red line shows that the constraint on f does indeed continue to decrease without apparent bound as Ny, inereases and shot noise is reduced.", The dashed red line shows that the constraint on $f$ does indeed continue to decrease without apparent bound as $N_h$ increases and shot noise is reduced.
" Lt is disappointing. however. that ej,y islarger than σε; until the survey exceeds 10 redshifts in this bin."," It is disappointing, however, that $\sigma_{\ln f}$ islarger than $\sigma_{\ln G}$ until the survey exceeds $10^9$ redshifts in this bin."
 Using Equation (17)) as a &uide to the behavior of ayyr. we can see that the survey can attain very large n? but still have the f constraint degrade by an unfavorable lever arm factor (6|fgvf Var(b).," Using Equation \ref{sigmaf}) ) as a guide to the behavior of $\sigma_{\ln f}$, we can see that the survey can attain very large $nP$ but still have the $f$ constraint degraded by an unfavorable lever arm factor $\langle b+f\mu^2\rangle / \sqrt{{\rm Var}(b)}$ ."
 While massive halos have b>1 and a significant range of bias. al of the halos with masses LOM 10775.1AZ. have b within a few percent of unity.," While massive halos have $b>1$ and a significant range of bias, all of the halos with masses $10^{10}$ $10^{12}h^{-1}M_\odot$ have $b$ within a few percent of unity."
 As low-mass halos dominate the ful sample. Var(b) drops. degrading the gains from decreasec shot noise.," As low-mass halos dominate the full sample, ${\rm Var}(b)$ drops, degrading the gains from decreased shot noise."
". rpsThe scaling. is. roughly opxIN,01/23 77.", The scaling is roughly $\sigma_{\ln f}\propto N_h^{-1/3}$ .
 The green lines show the effect. of degrading perfec, The green lines show the effect of degrading perfect
we repeat this calculation for the five different resolution levels of the Aq-A halo (Fig. 2)).,we repeat this calculation for the five different resolution levels of the Aq-A halo (Fig. \ref{MOA532}) ).
" We see that AqAl together with Aq-A3, Aq-A4, and Aq-A5 has an angular momentum distribution broadly of the same form as Aq-A2, with increasing noise as the resolution decreases because of the smaller number of subhaloes."," We see that Aq-A1 together with Aq-A3, Aq-A4, and Aq-A5 has an angular momentum distribution broadly of the same form as Aq-A2, with increasing noise as the resolution decreases because of the smaller number of subhaloes."
" Each resolution level is dominated by a different subhalo mass; the minimum subhalo mass in Aq-A5 is ~10’Mo, while in Aq-AI1 it is three orders of magnitude smaller."," Each resolution level is dominated by a different subhalo mass; the minimum subhalo mass in Aq-A5 is $\sim10^{7}\mathrm{M}_{\odot}$, while in Aq-A1 it is three orders of magnitude smaller."
 We find a similar degree of convergence with numerical resolution for haloes Aq-B through to Aq-F. In Fig., We find a similar degree of convergence with numerical resolution for haloes Aq-B through to Aq-F. In Fig.
 3 we probe the orientation of the angular momentum vector of different populations., \ref{MassBins} we probe the orientation of the angular momentum vector of different populations.
" In the top panel, we compare the distribution for the 1000 largest subhaloes at the final redshift (particle number > 1222, equivalent to subhalo mass of 1.7x 10M) with that the 100 subhaloes present at z=0 that had the most massive progenitors and that of the entire halo population."," In the top panel, we compare the distribution for the 1000 largest subhaloes at the final redshift (particle number $>1222$ equivalent to subhalo mass of $1.7\times10^{7}\mathrm{M}_{\odot}$ ) with that the 100 subhaloes present at $z=0$ that had the most massive progenitors and that of the entire halo population."
 The most massive progenitor is defined as the halo in the merger tree that contained the largest number of particles over the entire history of the simulation., The most massive progenitor is defined as the halo in the merger tree that contained the largest number of particles over the entire history of the simulation.
 'This mass is very close to the mass that the subhalo had at the time it fell into the main halo., This mass is very close to the mass that the subhalo had at the time it fell into the main halo.
" It is these subhaloes that are most likely to host satellite galaxies, according to Libeskindetal.(2009)."," It is these subhaloes that are most likely to host satellite galaxies, according to \citet{li09}."
". Of the subhaloes that had the 100 largest progenitors, all bar 6 are among the top 1000 most massive subhaloes at redshift zero."," Of the subhaloes that had the 100 largest progenitors, all bar 6 are among the top 1000 most massive subhaloes at redshift zero."
 The distributions of cos θµ.ς for all three populations of subhaloes are consistent within the errors., The distributions of cos $\theta_{\mathrm{H\cdot S}}$ for all three populations of subhaloes are consistent within the errors.
" To establish whether the angular momentum orientation of the subhalo population is special, in the lower panel of Fig."," To establish whether the angular momentum orientation of the subhalo population is special, in the lower panel of Fig."
 3 we compare subhaloes in Aq-A2 with particles from the main halo., \ref{MassBins} we compare subhaloes in Aq-A2 with particles from the main halo.
 We create a special sample of halo particles with the same radial distribution as the subhaloes., We create a special sample of halo particles with the same radial distribution as the subhaloes.
 This is made by first defining a set of about 30 radial bins between the halo centre and the virial radius., This is made by first defining a set of about 30 radial bins between the halo centre and the virial radius.
" The halo subsample is produced by first noting how many subhaloes lie in a particular bin, and then randomly selecting the same number of halo particles from the that same bin."," The halo subsample is produced by first noting how many subhaloes lie in a particular bin, and then randomly selecting the same number of halo particles from the that same bin."
 This is always possible as the number of halo particles in any bin exceeds the corresponding number of subhaloes., This is always possible as the number of halo particles in any bin exceeds the corresponding number of subhaloes.
 We compare this particle sample's distribution of cos θµΗ.ς with that for the Aq-A2 subhaloes and for the entire set of main halo particles., We compare this particle sample's distribution of cos $\theta_{\mathrm{H\cdot S}}$ with that for the Aq-A2 subhaloes and for the entire set of main halo particles.
 The three distributions are statistically inconsistent with each other., The three distributions are statistically inconsistent with each other.
" The subhalo population has a larger fraction of aligned and antialigned members, with the radially selected subsample being intermediate between the subhaloes and the halo particles as a whole."," The subhalo population has a larger fraction of aligned and antialigned members, with the radially selected subsample being intermediate between the subhaloes and the halo particles as a whole."
" Although even the latter has a non-uniform distribution of angular momenta cosines, it is significantly flatter than that of other two populations."," Although even the latter has a non-uniform distribution of angular momenta cosines, it is significantly flatter than that of other two populations."
" This suggests that the accretion mechanism that supplies subhaloes (of all masses) is somewhat different from the mechanism by which halo particles are accreted, or that the evolution of subhaloes differs from that of halo particles."," This suggests that the accretion mechanism that supplies subhaloes (of all masses) is somewhat different from the mechanism by which halo particles are accreted, or that the evolution of subhaloes differs from that of halo particles."
" To investigate the orientation of the orbital spins in more detail, we plot the angular momentum vectors of each subhalo on an all-sky Mollweide projection, one for each halo at resolution L2."," To investigate the orientation of the orbital spins in more detail, we plot the angular momentum vectors of each subhalo on an all-sky Mollweide projection, one for each halo at resolution L2."
" Each map displayed here was divided into ~45000 pixels, with angular width ~1°, and smoothed with a Gaussian beam of FWHM 10? using Healpix routines (Górskietal.2005)."," Each map displayed here was divided into $\sim45000$ pixels, with angular width $\sim1^{\circ}$, and smoothed with a Gaussian beam of FWHM $10^{\circ}$ using Healpix routines \citep{Go04}."
". We identify the pixel with the highest density after smoothing, and call this the 'densest point vector’."," We identify the pixel with the highest density after smoothing, and call this the `densest point vector'."
 The pre-smoothing maps for all six L2 haloes are displayed in Fig 4.., The pre-smoothing maps for all six L2 haloes are displayed in Fig \ref{MapOrbits}.
" The main halo spin vector is marked in red, its antipole in blue, and the densest point vector in green."," The main halo spin vector is marked in red, its antipole in blue, and the densest point vector in green."
" Aq-A2 exhibits the cleanest structure of all the haloes, with strong clustering around the pole and antipole, joined by two strands."," Aq-A2 exhibits the cleanest structure of all the haloes, with strong clustering around the pole and antipole, joined by two strands."
" Aq-B2 is, in contrast, characterised by irregular structures concentrated around regions distant from the main halo spin poles."," Aq-B2 is, in contrast, characterised by irregular structures concentrated around regions distant from the main halo spin poles."
" All of the other haloes exhibit clustering around the main halo spin, with other, local, features apparent."," All of the other haloes exhibit clustering around the main halo spin, with other, local, features apparent."
 The densest point vector position is always closer to the main halo spin than to its antipole., The densest point vector position is always closer to the main halo spin than to its antipole.
 One may think of Figs., One may think of Figs.
 1 to 3 as an integration around lines of equal angle from the red and blue circles., \ref{MOT1} to \ref{MassBins} as an integration around lines of equal angle from the red and blue circles.
" As noted above, we are particularly interested in those subhaloes that are most likely to host satellites, and so we repeat this plot for the 100 subhaloes with largest progenitors in Fig. 5.."," As noted above, we are particularly interested in those subhaloes that are most likely to host satellites, and so we repeat this plot for the 100 subhaloes with largest progenitors in Fig. \ref{MapOrbits2}."
" As expected from Fig. 3,"," As expected from Fig. \ref{MassBins},"
 the 100 subhaloes with the largest progenitors trace the underlying structure of subhaloes in the map traced in Fig. 4.., the 100 subhaloes with the largest progenitors trace the underlying structure of subhaloes in the map traced in Fig. \ref{MapOrbits}.
" A few of them lie in regions where there are few subhaloes of any mass,and so we might expect to find satellite galaxies spatially removed from the disc-of-satellites for at least some portions"," A few of them lie in regions where there are few subhaloes of any mass,and so we might expect to find satellite galaxies spatially removed from the disc-of-satellites for at least some portions"
"electrons then spiral in the strong ambient magnetic fields, and generate gyrosynchrotron radiation that can be detected at radio wavelengths.","electrons then spiral in the strong ambient magnetic fields, and generate gyrosynchrotron radiation that can be detected at radio wavelengths."
 The energy released during reconnection also serves to heat the stellar corona to temperatures sufficient to generate detectable thermal brehmsstrahlung X-ray emission., The energy released during reconnection also serves to heat the stellar corona to temperatures sufficient to generate detectable thermal brehmsstrahlung X-ray emission.
" Stars more massive than about 2—3 M. are expected to move towards the main sequence on radiative tracks, and are not anticipated to have strong superficial fields."," Stars more massive than about 2–3 $M_{\odot}$ are expected to move towards the main sequence on radiative tracks, and are not anticipated to have strong superficial fields."
" Such stars are, therefore, not expected to be magnetically active."," Such stars are, therefore, not expected to be magnetically active."
 X-ray emission as well as non-thermal radio emission have been detected [rom massive O and WR stars Berghóller et al., X-ray emission as well as non-thermal radio emission have been detected from massive O and WR stars Berghöffer et al.
" 1997, Pittard Dougherty 2006), but are believed to be the result of shocks in winds and wind interactions."," 1997, Pittard Dougherty 2006), but are believed to be the result of shocks in winds and wind interactions."
" For intermediate mass stars (spectral type A and late B), however, the situation appears somewhat more uncertain."," For intermediate mass stars (spectral type A and late B), however, the situation appears somewhat more uncertain."
" Those stars should not be magnetically active, and have no strong winds."," Those stars should not be magnetically active, and have no strong winds."
" Yet a small fraction of them (perhaps about5%:,; Montmerle et 22005) do show evidence of strong magnetic acüvity (Stelzer et 22005. 2006; Wade et 22009; Hubrig et 22009)."," Yet a small fraction of them (perhaps about; Montmerle et 2005) do show evidence of strong magnetic activity (Stelzer et 2005, 2006; Wade et 2009; Hubrig et 2009)."
" Several explanations have been put forward, but arguably the most plausible one is that the magnetic activity is in fact associated with a low-mass companion rather than with the intermediate mass primary itself SStelzer et 22005. 2006)."," Several explanations have been put forward, but arguably the most plausible one is that the magnetic activity is in fact associated with a low-mass companion rather than with the intermediate mass primary itself Stelzer et 2005, 2006)."
" Several recent X-ray observations, however, might be more easily interpreted if the Herbig AeBe stars themselves were magnetically active TTelleschi et 22007, Günnther Schmitt 2009. Huenemoerder et 22009)."," Several recent X-ray observations, however, might be more easily interpreted if the Herbig AeBe stars themselves were magnetically active Telleschi et 2007, Günnther Schmitt 2009, Huenemoerder et 2009)."
" Unul the present observations, EC 95 was believed to be a single intermediate-mass young star, and was, therefore. not expected to show magnetic acüvity."," Until the present observations, EC 95 was believed to be a single intermediate-mass young star, and was, therefore, not expected to show magnetic activity."
" Indeed, Smith et ((1999) and Giardino et ((2007) did proposed that the non-thermal radio emission as well as the X-ray emission [rom that source might be provided by a low-mass companion rather than by the AeBe protostar."," Indeed, Smith et (1999) and Giardino et (2007) did proposed that the non-thermal radio emission as well as the X-ray emission from that source might be provided by a low-mass companion rather than by the AeBe protostar."
" Preibisch (1999) also mentüoned that a low-mass companion could be at the origin of the X-ray emission, though he initially did not consider this possibility particularly likely since the X-ray luminosity of EC 95 based on ROSAT observations appeared to be nearly three orders of magnitude larger than that of typical T Tauri stars."," Preibisch (1999) also mentioned that a low-mass companion could be at the origin of the X-ray emission, though he initially did not consider this possibility particularly likely since the X-ray luminosity of EC 95 based on ROSAT observations appeared to be nearly three orders of magnitude larger than that of typical T Tauri stars."
" However, more recent XMM-Newton observations have shown that the X-ray luminosity of EC 95 is significantly smaller than originally thought, and is in fact within the range of normal T Tauri stars (albeit near the top of the X-ray luminosity function)."," However, more recent XMM-Newton observations have shown that the X-ray luminosity of EC 95 is significantly smaller than originally thought, and is in fact within the range of normal T Tauri stars (albeit near the top of the X-ray luminosity function)."
" Thus, it appears plausible again that the X-ray emission might come [rom a low-mass companion."," Thus, it appears plausible again that the X-ray emission might come from a low-mass companion."
 Our observations have revealed that EC 95 is a binary system most likely composed of a ~ 4-5 M. primary. and a low-mass T Tauri companion.," Our observations have revealed that EC 95 is a binary system most likely composed of a $\sim$ 4-5 $M_\odot$ primary, and a low-mass T Tauri companion."
" Moreover, the low-mass companion (EC 95b) appears to be usually brighter at radio wavelengths than the more massive primary 11; 33)."," Moreover, the low-mass companion (EC 95b) appears to be usually brighter at radio wavelengths than the more massive primary 1; 3)."
" Given the exisung correlation between the X-ray and the radio emission of T Tauri stars, one would expect EC 95b to also be the brightest member of the system at X-ray wavelengths."," Given the existing correlation between the X-ray and the radio emission of T Tauri stars, one would expect EC 95b to also be the brightest member of the system at X-ray wavelengths."
" Thus, the magnetic activity [rom EC 95 is likely to be largely dominated by a Jow-mass companion, rather than by the intermediate mass primary. as proposed by Smith et (1999)."," Thus, the magnetic activity from EC 95 is likely to be largely dominated by a low-mass companion, rather than by the intermediate mass primary, as proposed by Smith et (1999)."
" However, our observations show that the intermediate mass star EC 95a, although on average somewhat weaker than EC 95b, is a non-thermal radio Thus, the existence of a"," However, our observations show that the intermediate mass star EC 95a, although on average somewhat weaker than EC 95b, is a non-thermal radio Thus, the existence of a"
"thing is worth noticing in equation 3:: because of the hypothesis of pressure equilibrium and constant temperature of the cold phase. 1, is proportional to pressure Z: at the same time. the fraction of gas mass in the hot phase is always very low. so f, is very similar to the fraction of gas in molecular form.","thing is worth noticing in equation \ref{eq:sfr}: because of the hypothesis of pressure equilibrium and constant temperature of the cold phase, $n_c$ is proportional to pressure $P$; at the same time, the fraction of gas mass in the hot phase is always very low, so $f_{\rm mol}$ is very similar to the fraction of gas in molecular form."
" As a consequence. the particle star formation rate is primarily regulated by gas pressure (with the complication that fava is computed at the beginning of a star formation cycle and then kept frozen. while fj,,; is computed at each time-step)."," As a consequence, the particle star formation rate is primarily regulated by gas pressure (with the complication that $t_{\rm dyn}$ is computed at the beginning of a star formation cycle and then kept frozen, while $f_{\rm mol}$ is computed at each time-step)."
 The SFR term of equation 3. deposits a fraction (1 of the transformed mass into the stellar component. while a fraction fy... restored from massive stars in an Instantaneous Recycling Approximation ΙΚΑ). is given back to the hot phase.," The SFR term of equation \ref{eq:sfr} deposits a fraction $(1-f_{\rm
 re})$ of the transformed mass into the stellar component, while a fraction $f_{\rm re}$, restored from massive stars in an Instantaneous Recycling Approximation (IRA), is given back to the hot phase."
 The formed stellar com»onent is accumuated within the partice. and contributes to its inertia but not to its gas mass in all SPH computations.," The formed stellar component is accumulated within the particle, and contributes to its inertia but not to its gas mass in all SPH computations."
 The evaporation rate is assumed to be due to the destruction of molecular clouds and amounts to a fraction 0.1 of the SFR.," The evaporation rate is assumed to be due to the destruction of molecular clouds and amounts to a fraction $f_{\rm
 ev}=0.1$ of the SFR."
 Production of star particles is done according to the stochastic star formation algorithm of ?. (see paper I for details)., Production of star particles is done according to the stochastic star formation algorithm of \cite{Springel03} (see paper I for details).
 We allow for 4 generations of star particles to be spawned by each parent gas particle., We allow for 4 generations of star particles to be spawned by each parent gas particle.
 Each new star particle is produced at the expense of the stellar component and. if needed. of the cold phase.," Each new star particle is produced at the expense of the stellar component and, if needed, of the cold phase."
 The three components of a multi-phase particle are of course assumed to be subject to the same hydrodynamical forces., The three components of a multi-phase particle are of course assumed to be subject to the same hydrodynamical forces.
 To alleviate the effect of this unphysical assumption. and to mimic the destruction of a star-forming cloud after a few dynamical times (2).. the code forces each particle to leave the multi-phase regime after two dynamical times fu. computed as specified above.," To alleviate the effect of this unphysical assumption, and to mimic the destruction of a star-forming cloud after a few dynamical times \citep{Monaco04b}, the code forces each particle to leave the multi-phase regime after two dynamical times $t_{\rm dyn}$, computed as specified above."
" One SN is generated each AZ,«v=120M. of stars formed. and each SN generates LO erg of energy."," One SN is generated each $M_{\star,SN}=120\ {\rm M}_\odot$ of stars formed, and each SN generates $10^{51}$ erg of energy."
 Of the energy generated in the IRA. a small fraction fj;=0.02 is given to the local hot phase to sustain its high temperature. while a fraction fii...=0.3 is disributed to the hot phases of neighbour particles in a 60-degree wide cone anti-aligned with the gas density gradient.," Of the energy generated in the IRA, a small fraction $f_{\rm fb,i}=0.02$ is given to the local hot phase to sustain its high temperature, while a fraction $f_{\rm fb,o}=0.3$ is distributed to the hot phases of neighbour particles in a 60-degree wide cone anti-aligned with the gas density gradient."
 Energy contributions to particles are weighted by their distance from the cone axis. to mimic the expansion of SN-driven blasts along the least resistance path (2)..," Energy contributions to particles are weighted by their distance from the cone axis, to mimic the expansion of SN-driven blasts along the least resistance path \citep{McKee77}."
 The present version of the code distributes only thermal energy., The present version of the code distributes only thermal energy.
 The assumption of a molecular fraction regulated by pressure (equation 2)) is very important because it makes the evolution of the system intrinsically runaway: star formation generates SNe. energy feedback from SNe pressurizes the hot phase. the increase in pressure leads to an increase in molecular fraction and. thus to an increase in SER.," The assumption of a molecular fraction regulated by pressure (equation \ref{eq:fmol}) ) is very important because it makes the evolution of the system intrinsically runaway: star formation generates SNe, energy feedback from SNe pressurizes the hot phase, the increase in pressure leads to an increase in molecular fraction and thus to an increase in SFR."
 The runaway halts when the molecular fraction saturates to unity., The runaway halts when the molecular fraction saturates to unity.
 However. the dynamical response of the pressurized particle is able to limit this runaway through the expansion work done on neighbours.," However, the dynamical response of the pressurized particle is able to limit this runaway through the expansion work done on neighbours."
 This intrinsic runaway behaviour. together with the long cooling times. are the main reasons for our efficient thermal feedback.," This intrinsic runaway behaviour, together with the long cooling times, are the main reasons for our efficient thermal feedback."
 The first three simulations. the MW (at. standard and. high resolution) and DW test cases of paper I. start from already formed dise galaxies with 10 and 20 per cent gas fractions. and are shown at 0.5 Gyr. after the initial transients due to the switching on of stellar feedback have (almost) died out and while gas consumption is still negligible.," The first three simulations, the MW (at standard and high resolution) and DW test cases of paper I, start from already formed disc galaxies with 10 and 20 per cent gas fractions, and are shown at 0.5 Gyr, after the initial transients due to the switching on of stellar feedback have (almost) died out and while gas consumption is still negligible."
 The fourth simulation. the SH halo. forms in an isolated. static halo filled with rotating gas initially in virial equilibrium: we analyze the simulation at 0.7 Gyr. i.e. at the peak of its SFR. when the dise is still largely gas-dominaed.," The fourth simulation, the SH halo, forms in an isolated, static halo filled with rotating gas initially in virial equilibrium; we analyze the simulation at 0.7 Gyr, i.e. at the peak of its SFR, when the disc is still largely gas-dominated."
 In all cases. conclusions are unchanged when simulations are considered at other times.," In all cases, conclusions are unchanged when simulations are considered at other times."
 We show in figure |. face-on and edge-on mays of gas surface densiies and temperatures for the four simulations., We show in figure \ref{fig:maps} face-on and edge-on maps of gas surface densities and temperatures for the four simulations.
 It is interesting to notice that in the regions interested by star formation the dises (most of the MW and MW.HHR dises and the inner few kpe of SH) are relatively hot and surrounded by thick coronae of gas heated by feedback and eireulating above or below the dise in a galactic fountain., It is interesting to notice that in the regions interested by star formation the discs (most of the MW and HR discs and the inner few kpc of SH) are relatively hot and surrounded by thick coronae of gas heated by feedback and circulating above or below the disc in a galactic fountain.
 The effect of star formation on the DW galaxy is much less evident., The effect of star formation on the DW galaxy is much less evident.
 Figure 2. shows the standard. HI and molecular SK relations of he four simulations. compared with the data of ?. for normal spiral galaxies.," Figure \ref{fig:skone} shows the standard, HI and molecular SK relations of the four simulations, compared with the data of \cite{Bigiel08} for normal spiral galaxies."
 Gas surface densities are always meant to include contribution from helium., Gas surface densities are always meant to include contribution from helium.
 In the upper panel the thin line represents he fit proposed by ?.., In the upper panel the thin line represents the fit proposed by \cite{Kennicutt98}.
 Simulations have been processed as follows., Simulations have been processed as follows.
 Our analyses are restricted to cold gas: in this paper by cold gas we mean the cold phase of multi-phase particles plus all single-phase yarticles colder han LO’ K. The molecular gas surface density is computed using the molecular fraction of equation 2.. applied only o multi-phase particles: HI gas is just cold minus molecular gas.," Our analyses are restricted to cold gas; in this paper by cold gas we mean the cold phase of multi-phase particles plus all single-phase particles colder than $10^5$ K. The molecular gas surface density is computed using the molecular fraction of equation \ref{eq:fmol}, applied only to multi-phase particles; HI gas is just cold minus molecular gas."
 A galaxy frame is detined by the inertia tensor of stars and cold gus. the z-axis corresponding to the largest eigenvalue.," A galaxy frame is defined by the inertia tensor of stars and cold gas, the z-axis corresponding to the largest eigenvalue."
 The angular momentum of the same particles is always found to be at most a few degrees off the z-axis., The angular momentum of the same particles is always found to be at most a few degrees off the z-axis.
 Then. radial surface densities (in cylindrical coordinates) of (cold. HI. molecular) gas and SFR are computed: these are reported in the figure as colored thick lines.," Then, radial surface densities (in cylindrical coordinates) of (cold, HI, molecular) gas and SFR are computed; these are reported in the figure as colored thick lines."
 The same quantities are computed on a square grid in the x-y plane. with bin," The same quantities are computed on a square grid in the x-y plane, with bin"
"=a, fux πο... The function C,CX) is GiX) 2 qvi - Vin quay",= X. The function $G_{1}(X)$ is (X) = ( - V ) 	 	 - X'.
"Here V'=VCX""y.",Here $V' \equiv V(X')$.
" Note that q, and C,CX) vanish automatically if one uses the step function velocity profile of an infinitely thinshock. with VOX)=V, for X«0 and VCX)=V. for X>0."," Note that $\qbar_{1}$ and $G_{1}(X)$ vanish automatically if one uses the step function velocity profile of an infinitely thinshock, with $V(X) = V_{1}$ for $X < 0$ and $V(X) = V_{2}$ for $X \ge 0$."
" At order & one finds the following equation for ονΧο: Wie Hee, [i p. subject to the boundary condition G»(—o)=Gaf+e)0."," At order $\varepsilon$ one finds the following equation for $G_{2}(X)$: + = (1 + ) , subject to the boundary condition $G_{2}( - \infty) = G_{2}(+ \infty) = 0$."
" Integrating (57)) from X=—o to X=+o yields an equation for q»: d.— —- dX.. GX), ία TS p Ξ avc, y [", Integrating \ref{secondG}) ) from $X = - \infty$ to $X = +\infty$ yields an equation for $\qbar_{2}$: = - X (X) (1 + ) 	 	 = X (X) (.
 This procedure can be extended to higher order. but little is gained at the expense of increasingly complex mathematies.," This procedure can be extended to higher order, but little is gained at the expense of increasingly complex mathematics."
 We use two examples of immediate importance for a test of the numerical scheme advocated here., We use two examples of immediate importance for a test of the numerical scheme advocated here.
 Table | gives the parameters as used in the numerical simulations presented below., Table 1 gives the parameters as used in the numerical simulations presented below.
 The first example is the case of a uniform diffusivity Dov=Dj. which formally corresponds to c=| and £4=oo.," The first example is the case of a uniform diffusivity $D(x) = D_{1}$, which formally corresponds to $\sigma = 1$ and $L_{\rm d} = \infty$."
 This is the case where the CES is known to yield good results., This is the case where the CES is known to yield good results.
 This case has been treated before by Axford. Drury Summers (1982). who show that for the hyperbolic tangent velocity profile (29)) adopted here the cosmic ray transport equation can be solved analytically.," This case has been treated before by Axford, Drury Summers (1982), who show that for the hyperbolic tangent velocity profile \ref{Vprofile}) ) adopted here the cosmic ray transport equation can be solved analytically."
" They find an asymptotic slope of the momentum distribution equal to a=""E (ac xp.", They find an asymptotic slope of the momentum distribution equal to = ( 1 + ).
 Here @=3r/tr—1)., Here $\qbar_{0} = 3r/(r - 1)$.
 The perturbation expansion used here (and in a slightly different form by Drury (1983)) reproduces this result., The perturbation expansion used here (and in a slightly different form by Drury (1983)) reproduces this result.
 The hyperbolic tangent velocity law (29)) impliesWV , The hyperbolic tangent velocity law \ref{Vprofile}) ) implies = - .
Substituting this into the generally valid expression (559). together with dX=Vd/Dj. one finds: ου.," Substituting this into the generally valid expression \ref{qone}) ), together with ${\rm d}X = V \: {\rm d}x/D_{1}$, one finds: = - x ( ) =."
 This agrees with result (59)) of Drury et al (1982)., This agrees with result \ref{DASq}) ) of Drury et al (1982).
" Using this in relation (56)) one finds that C,CX)=0. which implies C,2q,=0 for n>2."," Using this in relation \ref{G1}) ) one finds that $G_{1}(X) = 0$, which implies $G_{n} = q_{n} = 0$ for $n \ge 2$."
 Here the perturbation expansion breaks off at order e and yields the asymptotic result. as noted before by Drury (1983).," Here the perturbation expansion breaks off at order $\varepsilon$ and yields the asymptotic result, as noted before by Drury (1983)."
 In terms of the compression ratio r=Vi/V». the diffusion length far upstream Lap7Di/Vi and g=4—3 one has: q=3 ," In terms of the compression ratio $r = V_{1}/V_{2}$, the diffusion length far upstream $L_{\rm diff}^{-\infty} = D_{1}/V_{1}$ and $q = \qbar - 3$ one has: 	q = ( 1 + ) = (1 + )."
As a second example we consider the case of a constant diffusion length: ea., As a second example we consider the case of a constant diffusion length: = =.
 This example is important as a test case of the predictor-corrector algorithm used here as it has a strong gradient in the ditfusivity., This example is important as a test case of the predictor-corrector algorithm used here as it has a strong gradient in the diffusivity.
" Formally it corresponds to 7= rand £4=L, so that in the shock el. which becomes large if &«| for thin shocks."," Formally it corresponds to $\sigma = r$ and $L_{\rm d} = L_{\rm s}$ so that in the shock | |, which becomes large if $\varepsilon \ll 1$ for thin shocks."
 If one adopts the hyperbolic tangent profile (29))/(60)) and uses the fact that dXX- =—. relation (559) yields: (1 ji Inf: jr.," If one adopts the hyperbolic tangent profile \ref{Vprofile}) \ref{Vder}) ) and uses the fact that X =, relation \ref{qone}) ) yields: = - ( ) x ( ) 	 	 = ( ) ( ) ."
" The function C, canbecalculated. but the integral over C, that determinesthe next order correction g to the slope can not be expressed in elementary functions."," The function $G_{1}$ canbecalculated, but the integral over $G_{1}$ that determinesthe next order correction $\qbar_{2}$ to the slope can not be expressed in elementary functions."
 Limiting ourselves to the correction one has in terms of g=qg- 3and r=V/V:, Limiting ourselves to the first-order correction one has in terms of $q = \qbar - 3$ and $r = V_{1}/V_{2}$:
and therefore the cloud continues to transfer its momentunir to the mixed gas and thus to the corona.,and therefore the cloud continues to transfer its momentum to the mixed gas and thus to the corona.
 In order to quantify the deceleration of the cold cloud. in the simulations we monitor the evolution of its centroid velocity. defined as the total momentum of the cold gas in the clouds direction of motion c divided the total mass of the cold gas.," In order to quantify the deceleration of the cold cloud, in the simulations we monitor the evolution of its centroid velocity, defined as the total momentum of the cold gas in the cloud's direction of motion $x$ divided the total mass of the cold gas."
 We compare the centroid velocity with the analytic estimate (2) where the characteristic drag time is given in terms of the cloud's initial mass Adj). geometrical cross section e. and coronal density. py. by In all our dissipative standard simulations. independent of the initial speed. ey. after Myr the centroid. velocity. is 0.75ry. significantly lower than the values predicted. by equation (1)). which treats the cloudas a rigidbody.," We compare the centroid velocity with the analytic estimate \citetalias{FraternaliB08}) ) where the characteristic drag time is given in terms of the cloud's initial mass $M_{\rm cl}$, geometrical cross section $\sigma$ , and coronal density $\rho_{\rm h}$, by In all our dissipative standard simulations, independent of the initial speed $v_0$ , after $\Myr$ the centroid velocity is $\sim
0.75~v_0$, significantly lower than the values predicted by equation \ref{eq:drag}) ), which treats the cloudas a rigidbody."
 “Phis, This
Schwartz. aud Paulo Holvorcem. in determining coordinates for the novae and in preparing the LAU cireulars announcing their discovery.,"Schwartz, and Paulo Holvorcem, in determining coordinates for the novae and in preparing the IAU circulars announcing their discovery."
 We also thank the anonymous referee for valuable comments on the manuscript., We also thank the anonymous referee for valuable comments on the manuscript.
 This research. has made use of the NASA/ IPAC Infrared Science Archive. which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Aclministration.," This research has made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
 This publication makes use of data products from the Two Micron. All Sky Survey. which is à 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.," 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."
 This research draws upon data provided by Dr. Philip Alassev as distributed by the NOAO Seience Archive., This research draws upon data provided by Dr. Philip Massey as distributed by the NOAO Science Archive.
 NOAO is operated bv the Association of Universities for Research in Astronomy (AURA). Ine. under a cooperative agreement with the National Science Foundation.," NOAO is operated by the Association of Universities for Research in Astronomy (AURA), Inc. under a cooperative agreement with the National Science Foundation."
provide a fruitful route to study the first stars.,provide a fruitful route to study the first stars.
 Iu thisLetter we assess the contribution to hydrogen relonization from ionidng photous produced in fast radiative accretion shocks., In this we assess the contribution to hydrogen reionization from ionising photons produced in fast radiative accretion shocks.
 We note that virialisation (Miniatictal.2001). and fast accretion shocks (Dopitaetal.2001) should lead to similar ionizing huninosities siuce the cooling radiation following the virialization shock contains cucrey comparable to the eravitational potential energv available to drive the fast accretion shock., We note that virialisation \citep[][]{Miniati2004} and fast accretion shocks \citep[][]{Dopita2011} should lead to similar ionizing luminosities since the cooling radiation following the virialization shock contains energy comparable to the gravitational potential energy available to drive the fast accretion shock.
" In our nuuerical examples. we adopt the standard set of cosmological parameters (Ixonatsuetal. 2011).. with values of O1,20.01. O4,=0.21 and O4=0.76 for the matter. barvon. aud dark cucreyv fractional deusity respectively, 5h=0.73. for the dimensionless Iubble constant. and os=0.82,"," In our numerical examples, we adopt the standard set of cosmological parameters \citep{Komatsu2011}, , with values of $\Omega_{\rm b}=0.04$, $\Omega_{\rm m}=0.24$ and $\Omega_\Lambda=0.76$ for the matter, baryon, and dark energy fractional density respectively, $h=0.73$, for the dimensionless Hubble constant, and $\sigma_8=0.82$."
 We beein by computing the uunboer of ionizations per barvon processed through shocks., We begin by computing the number of ionizations per baryon processed through shocks.
 We also discuss the ionisation historv due to stars for comparison., We also discuss the ionisation history due to stars for comparison.
 Dopitaetal.(2011) presented a fitting Γωνία for the nunuber of ionising photons that euter the ICAI per barvon processed through a shock. which iu our termunoloey is where eds the velocity of the shocked eas.," \citet[][]{Dopita2011} presented a fitting formula for the number of ionising photons that enter the IGM per baryon processed through a shock, which in our terminology is where $v$ is the velocity of the shocked gas."
" Dopitaetal.(20011) argue that ¢=V2eq, where ey. is the virial velocity of the halo. and we utilise this throughout the current work."," \citet[][]{Dopita2011} argue that $v=\sqrt{2}v_{\rm vir}$ where $v_{\rm vir}$ is the virial velocity of the halo, and we utilise this throughout the current work."
" The virial velocity is estimated from the halo mass using where ος) ds close to unity and defined as à=ο μι12)Ff. A.[(Q,,=/05ids?WA.|δρα39d? and d=OFTanl (seo equatious 2225 in Barkana Loch 2001 for more details)."," The virial velocity is estimated from the halo mass using where $\zeta(z)$ is close to unity and defined as $\zeta\equiv [(\Omega_m/\Omega_m^z)(\Delta_c/18\pi^2)]$ , $\Omega_m^z \equiv [1+(\Omega_\Lambda/\Omega_m)(1+z)^{-3}]^{-1}$, $\Delta_c=18\pi^2+82d-39d^2$, and $d=\Omega_m^z-1$ (see equations 22–25 in Barkana Loeb 2001 for more details)."
" Under the assuuptioun that eas shocks once when a halo forms. the number of photous produced per hydrogen atom in the Universe can be estimated frou the Pross&Schechter(1971). amass function da/fdAL (withmodificationsduetoSheth&Tormen1999) where AM,=(ΟιΟι) aud py=(0,/0,)0,. are the barvonic mass Iuside a dark matter halo aud barvonic massdensity in the ICAL respectively."," Under the assumption that gas shocks once when a halo forms, the number of photons produced per hydrogen atom in the Universe can be estimated from the \citet[][]{press1974} mass function $dn/dM$ \citep[with modifications due to][]{Sheth1999}
 where $M_{\rm b}=(\Omega_{\rm b}/\Omega_{\rm m})M$ and $\rho_{\rm
b}=(\Omega_{\rm b}/\Omega_{\rm m})\rho_{\rm m}$, are the baryonic mass inside a dark matter halo and baryonic massdensity in the IGM respectively."
 Results are shown as a function of redshift in the upper left panel of Figure 1.. assiuniue ninm halo masses corresponding to raj=101a8.| (the cooling threshold for livdrogen]. and ea230uns|! (theJeausthresholdiuauionizedIGALDijkstraetal.2001).," Results are shown as a function of redshift in the upper left panel of Figure \ref{plot1}, assuming minimum halo masses corresponding to $v_{\rm vir}=10~{\rm km~s^{-1}}$ (the cooling threshold for hydrogen), and $v_{\rm vir}=30~{\rm
km~s^{-1}}$ \citep[the Jeans threshold in an ionized
IGM, ][]{Dijkstra2004}."
. The results are independent of this choice owing to the dominance of massive halos in oxoducimg radiation from fast accretion shocks (Dopitaetal.2011)., The results are independent of this choice owing to the dominance of massive halos in producing radiation from fast accretion shocks \citep[][]{Dopita2011}.
. We find that ouly a few percent of the ICAL is reionized by :~6. aud <1% at i~δν indicating iat there are 1-2 orders of magnitude too few ionising photons produced im fast accretion shocks to reionize the Universe.," We find that only a few percent of the IGM is reionized by $z\sim6$, and $<1\%$ at $z\sim8$, indicating that there are 1-2 orders of magnitude too few ionising photons produced in fast accretion shocks to reionize the Universe."
 We note that virialisation shocks do not exist in low uass halos due to the presence of cold flows., We note that virialisation shocks do not exist in low mass halos due to the presence of cold flows.
 Caven he halo mass dependent fraction of cold flow accretion feota Where no shock is produced (seresetal.2009:Faucher-Cüguereetal.2010).. one should exclude. cold node accretion material from the calculation of iouisine ununositv associated with virialisation shocks (Miniatietal.2001).," Given the halo mass dependent fraction of cold flow accretion $f_{\rm cold}$ where no shock is produced \cite[][]{Keres2009,Faucher2010}, one should exclude cold mode accretion material from the calculation of ionising luminosity associated with virialisation shocks \citep[][]{Miniati2004}."
. This would reduce the predictions in Miuiatietal.(2001). by a factor of (1.foa). where feoiq Is weighted over the mass aud number of halos.," This would reduce the predictions in \citet[][]{Miniati2004} by a factor of $(1-f_{\rm
cold})$, where $f_{\rm cold}$ is weighted over the mass and number of halos."
 However Dopitaetal.(2011) arene that with sufficieut resolution the cold flow material is fouud to shock at the imtersection with the nascent ealactie disk., However \citet[][]{Dopita2011} argue that with sufficient resolution the cold flow material is found to shock at the intersection with the nascent galactic disk.
 As a result no correction for cold flow accretion should be applied to estimates of fast accretion shock produced ionising photons iu equation (3)) or later iu thisL, As a result no correction for cold flow accretion should be applied to estimates of fast accretion shock produced ionising photons in equation \ref{fcol}) ) or later in this.
"etter, We compare the above result to the nmuuber of lonizations obtained for stars (lower right panel of Figure 1)).", We compare the above result to the number of ionizations obtained for stars (lower right panel of Figure \ref{plot1}) ).
 Tere we utilise equation (3)) but assume AN.=JO00(ff.)15. appropriate for a Salpeter IME with a fiducial value of fifico~0.001. where f. and fixe are the star-formation efficiency aud escape fraction of iouising photons respectively.," Here we utilise equation \ref{fcol}) ) but assume $N_\gamma=4000(f_\star f_{\rm
esc})=15$, appropriate for a Salpeter IMF with a fiducial value of $f_\star f_{\rm esc}\sim0.004$, where $f_\star$ and $f_{\rm esc}$ are the star-formation efficiency and escape fraction of ionising photons respectively."
 We again show minim halo masses corresponding to e=lOkus! and ocu=30kinsd. leading to significant variation in the totalnunmiber of ionising plotous produced.," We again show minimum halo masses corresponding to $v_{\rm vir}=10~{\rm km~s^{-1}}$ and $v_{\rm
vir}=30~{\rm km~s^{-1}}$, leading to significant variation in the totalnumber of ionising photons produced."
 Iu difference to shock produced photous. these fiducial stellar populations are casily able to reionize the Universe by 26 (as is well known).," In difference to shock produced photons, these fiducial stellar populations are easily able to reionize the Universe by $z\sim6$ (as is well known)."
 This result is plausible since the nuclear efficiency. of stars is larger by many orders of maguitude than the efficicucy of converting rest lass to radiation by a shock. ~(cey?=10° for the shock speeds of interest (6<300kins 1).," This result is plausible since the nuclear efficiency of stars is larger by many orders of magnitude than the efficiency of converting rest mass to radiation by a shock, $\sim (v/c)^2 \la 10^{-6}$, for the shock speeds of interest $v\la 300~{\rm km~s^{-1}}$ )."
 The calculation iu 2.1. utilises each barvou oulv once. Whereas. unlike the case for stars. a barvon may be processed through shocks several times duriug the hierarchical formation of a galaxy.," The calculation in \ref{ioncol} utilises each baryon only once, whereas, unlike the case for stars, a baryon may be processed through shocks several times during the hierarchical formation of a galaxy."
 We therefore recaleulate the contribution to reionization frou fast aceretion shocks based ou the merger rate of halos., We therefore recalculate the contribution to reionization from fast accretion shocks based on the merger rate of halos.
"Specifically, when a halo of mass Mo<AA merges with a halo of mass A4. we assmuue that the birvons contained within halo 1 are shocked at e2νοAL,| Alo).","Specifically, when a halo of mass $M_2<M_1$ merges with a halo of mass $M_1$ , we assume that the baryons contained within halo 1 are shocked at $v=\sqrt{2}v_{vir}(M_1+M_2)$ ."
 We assume that the star-formation officlenev is ueeleible so that all barvous are available to shock during each merecr., We assume that the star-formation efficiency is negligible so that all baryons are available to shock during each merger.
 The resulting expression for the rate of production of all baryous in galaxies is, The resulting expression for the rate of production of all baryons in galaxies is
Kaiser predicted from the spherical infall model. the infall pattern surrounding the cluster appears as (wo trumpet horns stuck together.,"Kaiser predicted from the spherical infall model, the infall pattern surrounding the cluster appears as two trumpet horns stuck together."
 Were. however. detailed analyses demonstrate that the boundary of the eluster is not a caustic.," Here, however, detailed analyses demonstrate that the boundary of the cluster is not a caustic."
 In fact the causties are. in effect. smeared out by the peculiar motions within the infalling süructures.," In fact the caustics are, in effect, smeared out by the peculiar motions within the infalling structures."
 Even so. (he underlving pattern remains.," Even so, the underlying pattern remains."
 Dialerio and Geller (1997) ancl Dialerio (1999) used earlier simulations to explore approaches io the determination of the mass distribution in the infall region where. obviously. the equilibrium assumptions used to analvze cluster cores do not apply.," Diaferio and Geller (1997) and Diaferio (1999) used earlier simulations to explore approaches to the determination of the mass distribution in the infall region where, obviously, the equilibrium assumptions used to analyze cluster cores do not apply."
 They showed that the amplitude of the pattern in Figure 2 is a measure of (he escape velocity from the cluster., They showed that the amplitude of the pattern in Figure \ref{fig:simulation} is a measure of the escape velocity from the cluster.
 They developed a method of identifving the steep change in phase density Chat defines the infall pattern and called (heir mass estimation technique the “caustic” technique in recognition of Kaiser's early work., They developed a method of identifying the steep change in phase density that defines the infall pattern and called their mass estimation technique the “caustic” technique in recognition of Kaiser's early work.
 Their approach enables measurement of the cluster mass prolile to large raclius., Their approach enables measurement of the cluster mass profile to large radius.
 Fieure 2. provides a guide to the translation from (he simulations to analvsis of the data., Figure \ref{fig:simulation} provides a guide to the translation from the simulations to analysis of the data.
 The third. column shows snapshots of the evolution of the cluster in redshift space. but in contrast with the snapshots in the second column. the cluster center is on the left edee of the plot.," The third column shows snapshots of the evolution of the cluster in redshift space, but in contrast with the snapshots in the second column, the cluster center is on the left edge of the plot."
 The abscissa is the projected comoving distance [rom (he cluster center and (the ordinate is (he same as in (he second column. the line-oF-sight rest lrame peculiar velocily relative to the cluster center.," The abscissa is the projected comoving distance from the cluster center and the ordinate is the same as in the second column, the line-of-sight rest frame peculiar velocity relative to the cluster center."
 In. essence this plot shows an azimutlal sunm over (he representation in (he second column., In essence this plot shows an azimuthal sum over the representation in the second column.
 Of course. we can never observe the dark matter distribution: we sample it bv measuring redshilts of individual galaxies in the cluster.," Of course, we can never observe the dark matter distribution; we sample it by measuring redshifts of individual galaxies in the cluster."
" Figure 2. (left column) also shows what happens when we sample velocity distribution with 500 ""galaxies. about the number of objects we could reasonably observe in a single svslem."," Figure \ref{fig:simulation} (left column) also shows what happens when we sample velocity distribution with 500 “galaxies”, about the number of objects we could reasonably observe in a single system."
" The ""galaxies in Figure 2. are a ranclomly chosen subset of the dark. matter parücles: ie we assume (he galaxies are unbiased (racers of the dark matter distribution (e.g. Diaferio 1999: Dialerio et al.", The “galaxies” in Figure \ref{fig:simulation} are a randomly chosen subset of the dark matter particles; i.e we assume the galaxies are unbiased tracers of the dark matter distribution (e.g. Diaferio 1999; Diaferio et al.
 1999)., 1999).
 This sampling makes it painfully obvious that we are observationallv blind to most of the complex structure in the infall region: there just are nol enough galaxies to sample the dark matter distribution densely enough to reveal all of the subsystems that come together to form the cluster., This sampling makes it painfully obvious that we are observationally blind to most of the complex structure in the infall region; there just are not enough galaxies to sample the dark matter distribution densely enough to reveal all of the subsystems that come together to form the cluster.
 In spite of these fundamental limitations. application of the caustic technique to data on many clusters has provided some interesting. new insights.," In spite of these fundamental limitations, application of the caustic technique to data on many clusters has provided some interesting, new insights."
 Figure 3. shows data lor 4 from the HeCS (Ilectospec Cluster Survey) sample of Rines et al. (, Figure \ref{fig:hecs} shows data for 4 from the HeCS (Hectospec Cluster Survey) sample of Rines et al. (
2011. in preparation).,"2011, in preparation)."
" The clusters range in Moo from 9.5x 10M, to 14x 10M..", The clusters range in $_{200}$ from $9.5 \times 10^{13}$ $_\odot$ to $1.4 \times 10^{15}$ $_\odot$.
 For the most massive svstem in Figure 3.. ~ 200 ealaxies define (he infall pattern: lor the least massive (here are  100 ealaxies.," For the most massive system in Figure \ref{fig:hecs}, $\sim$ 200 galaxies define the infall pattern; for the least massive there are $\sim$ 100 galaxies."
 The observations match (he qualitative appearance of the infall patterns in the simulations., The observations match the qualitative appearance of the infall patterns in the simulations.
Light scattered by the wings of the PSF profile was shown to contribute noticeably to the color profiles of the outskirts of spiral galaxies.,Light scattered by the wings of the PSF profile was shown to contribute noticeably to the color profiles of the outskirts of spiral galaxies.
" In his paper, ? argues that the choice of PSF model used for image deconvolution is crucial for analyzing data of spiral galaxies."," In his paper, \cite{de_jong_point_2008} argues that the choice of PSF model used for image deconvolution is crucial for analyzing data of spiral galaxies."
 The author further claims that the red stellar halos found by previous studies were artificially produced by PSF models that were poorly sampled at large radii., The author further claims that the red stellar halos found by previous studies were artificially produced by PSF models that were poorly sampled at large radii.
 In this appendix we confirm this observation and show that the widely used synthetic PSF images may lead to inaccurate color gradients even in the peaky light profiles of massive red galaxies., In this appendix we confirm this observation and show that the widely used synthetic PSF images may lead to inaccurate color gradients even in the peaky light profiles of massive red galaxies.
 In order to test the effects of PSF profile selection on the properties of the derived Sérrsic model we followed the procedure described in subsection 3.2 to deconvolve the stacks using two different PSF images in each observed band., In order to test the effects of PSF profile selection on the properties of the derived Sérrsic model we followed the procedure described in subsection \ref{sur_bri} to deconvolve the stacks using two different PSF images in each observed band.
" The first profile is identical to the one used throughout this study and it combines a synthetically produced PSF model with a bright star stack (subsection 3.1)), while the second utilizes only the synthetic PSF image produced by the Read Atlas Image code."," The first profile is identical to the one used throughout this study and it combines a synthetically produced PSF model with a bright star stack (subsection \ref{psf_stacks}) ), while the second utilizes only the synthetic PSF image produced by the Read Atlas Image code."
" While the u,g,r and z models are less than larger when using only the synthetic PSF image, the i-band stack exhibits an increase of more than in this case."," While the u,g,r and z models are less than larger when using only the synthetic PSF image, the i-band stack exhibits an increase of more than in this case."
 The increased growth in model size in the i-band compared to the other stacks implies that the halo of these galaxies may erroneously appear to be red., The increased growth in model size in the i-band compared to the other stacks implies that the halo of these galaxies may erroneously appear to be red.
" This is most clearly evident in figure A,, where we plot the r-i color profiles of the two PSF deconvolved models and show that in the synthetic only case an artificial red halo appears outside of roughly 40 kpc."," This is most clearly evident in figure \ref{fig:psf_comp}, where we plot the r-i color profiles of the two PSF deconvolved models and show that in the synthetic only case an artificial red halo appears outside of roughly 40 kpc."
" We conclude, in agreement with ?,, that a proper choice of well sampled PSF model is critical for studying the faint outskirts of galaxies."," We conclude, in agreement with \cite{de_jong_point_2008}, that a proper choice of well sampled PSF model is critical for studying the faint outskirts of galaxies."
with approximaely the same luinosity.,with approximately the same luminosity.
 The ouly observed feature which has uot a theoretical couuterpart is t]je almost horizoial tail extencdiug redwards of 22 abonrp«] just fainter“il an 21., The only observed feature which has not a theoretical counterpart is the almost horizontal tail extending redwards of 2 at just fainter than 24.
 CiveLits position at the edge of the ACB and TRGB. we tend to ascribe it to thermally plsine ACB stars. a poorly understood evolutiona' phase whose models are not «lisplaved i the adopted tracXS.," Given its position at the edge of the AGB and TRGB, we tend to ascribe it to thermally pulsing AGB stars, a poorly understood evolutionary phase whose models are not displayed in the adopted tracks."
 Whatever t1e assumed netaicity. [rom he tracks overlap on the empirical ChD of the whole observed region. it turns out that the blue plune of NGC 1705 is populated by bot1 intermecliate- aud high-auass stars ou the MS axl by inassive stars at the blue edge of the core ieliuiu. burning phases. the red plume is populated by RSCs aud ACB stars. aud the faint red chmp with 1 ~uw AE 2 and 215 ~a ~a 20.5 is populated by RGB stas.," Whatever the assumed metallicity, from the tracks overlap on the empirical CMD of the whole observed region, it turns out that the blue plume of NGC 1705 is populated by both intermediate- and high-mass stars on the MS and by massive stars at the blue edge of the core helium burning phases, the red plume is populated by RSGs and AGB stars, and the faint red clump with 1 $\lsim$ $\lsim $ 2 and 24.5 $\lsim$ $\lsim $ 26.5 is populated by RGB stars."
 dt is also appareut llat very [ew stars more massive tlan 30 are likely to be prese in theresolvec Ποια of the galaxy., It is also apparent that very few stars more massive than 30 are likely to be present in the resolved field of the galaxy.
 The more massive objects present in M.iguificant numbers in the empirica CALD appear to be stars of 15-30AL., The more massive objects present in significant numbers in the empirical CMD appear to be stars of 15–30.
.. I1 [act the theoretical ‘racks Correspoucing to tliis mass reulge Calaccount lor several objects spauniug in the CMD from due to 'ed colors and. counectiug he blue aud the red plume with a sort o“bright bridge., In fact the theoretical tracks corresponding to this mass range can account for several objects spanning in the CMD from blue to red colors and connecting the blue and the red plume with a sort of bright bridge.
 This Sweests Ilat an enliauceilelt of the SF acivity has probady occurred at the epoch correspoucdiug o the lieime of a 20AL... j.e. aromd 15 Myr ago.," This suggests that an enhancement of the SF activity has probably occurred at the epoch corresponding to the lifetime of a 20, i.e. around 15 Myr ago."
 Fig., Fig.
 26 shows the alalogous superposition of tje same three sets of evollolary tracks ou he most tightly seleced CMD cerived from the NIC2 daa aud coutaiuiug 829 stars., \ref{cmdtracknic} shows the analogous superposition of the same three sets of evolutionary tracks on the most tightly selected CMD derived from the NIC2 data and containing 829 stars.
 The stroug cillerentia| effect due to the severe incompleteness alectii£& blue faint objects 1 ilese bands (see also ALOL results in the extreme paucity of blue plume st:us already inentioned ist Section {., The strong differential effect due to the severe incompleteness affecting blue faint objects in these bands (see also Al01) results in the extreme paucity of blue plume stars already mentioned in Section 4.
 The red vertical ris pepulated by RSGs aud AGB sta SOWxile the RGB is too faiul o be reachable with these stars so restrictively selected: we do reach it if we relax the 6540 criterion (see Fig. 18))," The red vertical is populated by RSGs and AGB stars, while the RGB is too faint to be reachable with these stars so restrictively selected; we do reach it if we relax the $\sigma_{DAO}$ criterion (see Fig. \ref{cm_vjh}) )"
 lo opa40<0.2., to $\sigma_{DAO}\le 0.2$.
 Iu agreement with what we find lor the optical CMD. stellar inodels with 15-30 fit quite well ie. bridge of brightest points connecting the blue and the red plume.," In agreement with what we find for the optical CMD, stellar models with 15–30 fit quite well the bridge of brightest points connecting the blue and the red plume."
 An accurate e«timate of the oxygen abuncdauce iu (GC 1705 was given by Storclii-Bergmana et al. (, An accurate estimate of the oxygen abundance in NGC 1705 was given by Storchi-Bergmann et al. (
1991). who derive 88.36 D‘Ol1 HII regious.,"1994), who derive 8.36 from HII regions."
 This implies hei the voungest objects have a metallicity of al)out. Z==00.00[., This implies that the youngest objects have a metallicity of about 0.004.
 No cli‘ect informati¢ is available [or older objects., No direct information is available for older objects.
 The stellar metallicities (Z sightly higher or sliguly lower that iat of the LN1C) quoted by MEDC and Sterubev (1998 are ouly iudicative., The stellar metallicities (Z slightly higher or slightly lower than that of the LMC) quoted by MFDC and Sternberg (1998) are only indicative.
 Wide band p10ometry is cer‘tainly not the best way to iner netal abudauces. but we can »ovide some his.," Wide band photometry is certainly not the best way to infer metal abundances, but we can provide some hints."
 By simply compariug the color distributiolis ο theoretical tracks aud observed objects iu t CALDs of bot1 Fig., By simply comparing the color distributions of theoretical tracks and observed objects in the CMDs of both Fig.
 25 and Fie. 26..," \ref{cmdtrackwfpc} and Fig. \ref{cmdtracknic},"
 it is appareu that tje vast majority of tLe stars resolve in NGC 1705 a'e more than 00.0001. inelucling the old stars ou the RGB.," it is apparent that the vast majority of the stars resolved in NGC 1705 are more metal-rich than 0.0004, including the old stars on the RGB."
 Iu (act. even accountit& lor the large uncertaluties ii the photomeVic conversiols from the theoretical to the observatyal plane ancl for the photometric errors of the data. this set of stellar models appears to be 00 blue to be recoucilable wih the observed optical aud NIR colors in all the red evolutionary jxhases.," In fact, even accounting for the large uncertainties in the photometric conversions from the theoretical to the observational plane and for the photometric errors of the data, this set of stellar models appears to be too blue to be reconcilable with the observed optical and NIR colors in all the red evolutionary phases."
 Ou the other hand. the 00.008 models seem sliehtly too red. specially iu the red CMD regiou.," On the other hand, the 0.008 models seem slightly too red, specially in the red CMD region."
 Iu this, In this
A large fraction⋅. of psClusters of galaxies. shows a very complex structure. either. in. the spatial. anc velocity. distribution⊲⊀⊀ of the galaxies. themselves (e.g. West. Bothun; 1990). or in. N-rayv maps. (e.g. Sarazin. LOSS).. or in both.,"A large fraction of Clusters of galaxies shows a very complex structure, either in the spatial and velocity distribution of the galaxies themselves (e.g. West Bothun 1990), or in the X-ray maps (e.g. Sarazin 1988), or in both."
have Although the substructures are often clearly visible in the optical (c.g. Alellier et al. (1988))), Although the substructures are often clearly visible in the optical (e.g. Mellier et al. \shortcite{Mell}) )
 and on the X-ray maps (Crebeney et al., and on the X-ray maps (Grebenev et al.
 1995) it is not always easy to make a quantitative analysis ofthe amount of substructure present., 1995) it is not always easy to make a quantitative analysis of the amount of substructure present.
 During the vast vears. many methods borrowed from statistical theory nave been applied to attempt determining. the number and properties of the subgroups.," During the past years, many methods borrowed from statistical theory have been applied to attempt determining the number and properties of the subgroups."
 Most. of these methods rave been applied to 2D angular position galaxy maps or 1D line-of-sight velocity distributions., Most of these methods have been applied to 2D angular position galaxy maps or 1D line-of-sight velocity distributions.
 ln a recent paper (Gambera et al., In a recent paper (Gambera et al.
 1997.. hereafter. GPAB) we have. tried o make use of the full 3D information available from a complete sample of galaxy projected: positions ancl line-of-sight velocities for Abell 1656 (the Coma cluster) in order to recover and determine some morphological propertics of the substructure.," 1997, hereafter GPAB) we have tried to make use of the full 3D information available from a complete sample of galaxy projected positions and line-of-sight velocities for Abell 1656 (the Coma cluster) in order to recover and determine some morphological properties of the substructure."
 In this paper we will explain and test the . ⊔↓⋖⋅⇂↓↕⋯⇂↓⊔⇂↓⋅⋯⇂⊔⊓⋅∠⇂↓⊔≺∣↓↴⇀∖∐∪⊔⊳∖↓⊔⋯↓⋜⋯⋅∠⇂≼⇍↓⊔⊳∖↿⋖⋅↓⋅≼∼⋜⊔⋜↧↓∪⋏∙≟⊳∖ mM . and we will. determine. the minimalD. requirements.. which. to be satisfied in ⊲⋅⊀order to apply it with. confidence.: Our method is based on the use of the Discrete Wavelet Transform (hereafter DWT). a technique originally introduced in turbulence studies (see c.g. Paree 1992).," In this paper we will explain and test the method introduced in GPAB on simulated cluster catalogs and we will determine the minimal requirements which have to be satisfied in order to apply it with Our method is based on the use of the Discrete Wavelet Transform (hereafter DWT), a technique originally introduced in turbulence studies (see e.g. Farge 1992)."
 Sometimes the Continuos Wavelet Transform (hereafter CW'D) has been applied. to characterize. substructure in galaxy clusters (e.g. Escalera Mazure (1988). hereafter∙ EM).. but it. has been convincinely shown that the two methods are unrelated to each other (Le.. CWE unis not the continuous. limitD of DWT).: and that CW'E vsis unsuitable. to characterize.the structure of clusters (Fang&Pando1997:Fang1996).," Sometimes the Continuos Wavelet Transform (hereafter CWT) has been applied to characterize substructure in galaxy clusters (e.g. Escalera Mazure \shortcite{Esc1}, hereafter EM), but it has been convincingly shown that the two methods are unrelated to each other (i.e. CWT is not the continuous limit of DWT), and that CWT is unsuitable to characterizethe structure of clusters \cite{fp,pf}."
. The DWT admits à complete. compact supported. orthogonal bases (Daubechies1988).. and. for this reason it represents a suitable tool for the analvsis of samples of objects.," The DWT admits a complete, compact supported orthogonal bases \cite{da}, and for this reason it represents a suitable tool for the analysis of samples of objects,"
crust of the neutrou star comes to resemble a “layer cake.” with alternating layers of lattice aud licuic.,"crust of the neutron star comes to resemble a “layer cake,” with alternating layers of lattice and liquid."
 This paper is relevant for the brightest low-mass x-ray binaries., This paper is relevant for the brightest low-mass x-ray binaries.
 These weakly magnetized ueutron stars are cousilered possible progenitors of millisecond pulsarstacharya 1995).. aud there has been much theoretical interest in the evolution of the crustmagneticfield(RomaniLO90::Geppert&Urpin199[: Urpin&Geppert1995:: Ixonar&BhattacharyaLOOT:: Brown&Bildsten 1998:: Urpin.Geppert.&Ilxoneukov 1998)).," These weakly magnetized neutron stars are considered possible progenitors of millisecond pulsars\citep[for a review, see][]{bhattacharya95}, and there has been much theoretical interest in the evolution of the crustmagneticfield\citealt{romani90};;\citealt{geppert94};; \citealt{urpin95}; \citealt{konar97}; \citealt{brown98a}; \citealt*{urpin98a}) )."
 Many of these neutrou stars rotatewithin au apparently uarrow range of spin [requeucies 300Hz (vanderWlis1995.e.g.)..," Many of these neutron stars rotatewithin an apparently narrow range of spin frequencies $\sim 300\rm\,Hz$ \citep[e.g.]{vdKlis98}."
 One possibility for this convergence of spin frequeucies is thatgravitational radiation from the neutroi star balauces the accretion torque (Bildsten1008:Aucerssonetal.1999).," One possibility for this convergence of spin frequencies is thatgravitational radiation from the neutron star balances the accretion torque \citep{bildsten98:gravity-wave,andersson99}."
. The source for the eravitatioual radiation could be amass quadrupole formect by misaligued electron capture layers i the crust(Bildsten1998)oracurrentquadrupole[rom au r-mocle instability in the core 1998:: Eriediman&Morsink 1998:: Lindblom.Owen.&Morsink 1998: Owenοἱal.1998: Auclersson.]xokkotas.&Schutz 1999:: Andersson.Wokkotas.&Stergioulas 1999))., The source for the gravitational radiation could be amass quadrupole formed by misaligned electron capture layers in the crust\citep{bildsten98:gravity-wave} oracurrentquadrupolefrom an r-mode instability in the core \citealt{andersson98:_new_class}; \citealt{friedman98:_axial_instab}; \citealt*{lindblom98:_gravit_radiat_instab}; \citealt{owen98:_gravit}; \citealt*{andersson99:_gravit_radiat}; \citealt*{andersson99}) ).
 All of these problems depen« ou the thermodynamics of theneutron stars crust aid core and motivate tliis paper., All of these problems depend on the thermodynamics of theneutron star's crust and core and motivate this paper.
 A neutronstar hasseveral cistinctive regions., A neutronstar hasseveral distinctive regions.
 The consists of uniform ape (inthe least dense parts)., The consists of uniform $npe^-$ (inthe least dense parts).
" At a baryon deusity less than »20.60. where,=0.164ahu.pE"" isH the saturationηv.ouclearmalter."," At a baryon density less than $n \approx 0.6 n_s$, where $n_s=0.16\fermi^{-3}$ is the saturation \citep*{pethick95}."
"iudividualiimcteiappear(Pelhick. íHIaveuhall.KL Pethick&Bavenhall1998)). z-equilibrium. zz0.00230, (1%10.ffm3) O.OLAL. A/À/ LOSvr (Webbink.Rappaport.&Savonije1983): MALvr.1. ", $\approx 30\MeV$ \citealt{pethick98}) $\beta$ $\approx 0.0023 n_s$ $4\ee{-4}\fermi^{-3}$ $0.01 \Msun$ $M/\dot{M}$ $10^8\yr$ \citep*{webbink83}; $\dot{M}=10^{-10}\Msun\yr^{-1}$ 
"Voor and Ving, to the cool and hot components respectively we obtain densities. pressures and. thermal energy. content given in Table 7..","$V_{\rm corr}$ and $V_{\rm bub}$ to the cool and hot components respectively we obtain densities, pressures and thermal energy content given in Table \ref{tab:inferred}."
 From the simulation we know the true properties of all the material in the bubble., From the simulation we know the true properties of all the material in the bubble.
 The density in the asma contributing the majority of the O.1-24keV [lux (all the material with volume emissivities within an order of magnitude of the maximum emissivity) varies between 2.5Son.&2.2m averaging r.Scm.," The density in the plasma contributing the majority of the $0.1$ $2.4 \keV$ flux (all the material with volume emissivities within an order of magnitude of the maximum emissivity) varies between $2.5 \ltsimm n_{\rm e} 
\ltsimm 28.2 \pcc$, averaging $7.8 \pcc$."
" In the hotter rubble interior the electron. density is D,&0.09en5", In the hotter bubble interior the electron density is $n_{\rm e} \approx 0.09 \pcc$.
" The shocked wind and cold shell are practically isobaric at P=13.lo"" vn cem7. and the majority of the total hermal energv in the bubble of py=1.17107""ere is in the hot shocked wind material that occupies most of the rubble volume."," The shocked wind and cold shell are practically isobaric at $P = 1.3 \times 10^{-9}$ dyn $\pcm2$, and the majority of the total thermal energy in the bubble of $E_{\rm TH} = 1.17 \times 10^{49} \erg$ is in the hot shocked wind material that occupies most of the bubble volume."
 Comparing the inferred properties of the bubble from. Table 7 with the true properties (see for example Fig. 4)), Comparing the inferred properties of the bubble from Table \ref{tab:inferred} with the true properties (see for example Fig. \ref{fig:weaver}) )
 we find: ]t is interesting to note that. despite. the dillerence between the true bubble properties and those for a bubble with thermal conduction (compare the radial profiles in Fig. 4)).," we find: It is interesting to note that despite the difference between the true bubble properties and those for a bubble with thermal conduction (compare the radial profiles in Fig. \ref{fig:weaver}) ),"
 the properties inferred from theAT PSPC two temperature spectral fitting are generally consistent with a bubble with conduction: The true properties of the bubble in the high mass loss rate simulation are practically identical to those of the low mass loss rate bubble., the properties inferred from the PSPC two temperature spectral fitting are generally consistent with a bubble with conduction: The true properties of the bubble in the high mass loss rate simulation are practically identical to those of the low mass loss rate bubble.
 The only dillerences are that the density and temperature in the shocked. wind are a factor of two cdilferent., The only differences are that the density and temperature in the shocked wind are a factor of two different.
 The intrinsic X-ray luminosity in the 0.1-2.4keV band is higher. but most importantly theROSAT count rate is double that of the low mass rate simulation.," The intrinsic X-ray luminosity in the $0.1$ $2.4 \keV$ band is higher, but most importantly the count rate is double that of the low mass rate simulation."
 ‘The intrinsic luminosity is again dominated by the cool. dense material near that shocked wind/cold shell interface. whose properties are almost identical between the two simulations.," The intrinsic luminosity is again dominated by the cool, dense material near that shocked wind/cold shell interface, whose properties are almost identical between the two simulations."
The bottom two plots of Figure 2 show another example of a simulated pulse profile.,The bottom two plots of Figure \ref{figure1} show another example of a simulated pulse profile.
 The pulse period is 200 ms and tsamp 7:20.39 ms.," The pulse period is 200 ms and $t_{\rm
 samp}\approx$ 0.39 ms."
" The “intrinsic” simulated profile has 3 orthogonal PA jumps, which occur in quick succession between phases -0.05 and 0."," The “intrinsic” simulated profile has 3 orthogonal PA jumps, which occur in quick succession between phases -0.05 and 0."
" The 7; used is the same as the previous example, ranging from 2.4 to 1.25 ms in the band."," The $\tau_s$ used is the same as the previous example, ranging from 2.4 to 1.25 ms in the band."
" Given the shorter pulse period and the smaller value of tsamp, the ratio R is larger ranging from R& for the lowest frequency channel to Rαγ3 for the highest 6frequency channel for the same configuration of filterbank (i.e. 16 16-MHz channels centred on 1.4 GHz)."," Given the shorter pulse period and the smaller value of $t_{\rm samp}$, the ratio $R$ is larger ranging from $R\approx 6$ for the lowest frequency channel to $R\approx 3$ for the highest frequency channel for the same configuration of filterbank (i.e. 16 16--MHz channels centred on 1.4 GHz)."
 The total power profile does not show evidence of scattering., The total power profile does not show evidence of scattering.
" However, the PA swing diverges significantly from the simple geometric model."," However, the PA swing diverges significantly from the simple geometric model."
" The kinks and wiggles in the PA swing on the right are common in real data, and cause severe difficulties both in fitting the PA data to the rotating vector model, and in theoretical treatment of the magnetospheric processes."," The kinks and wiggles in the PA swing on the right are common in real data, and cause severe difficulties both in fitting the PA data to the rotating vector model, and in theoretical treatment of the magnetospheric processes."
 Figure 3 illustrates the way a polarization profile with orthogonal PA jumps is affected by different values of τε., Figure \ref{strip} illustrates the way a polarization profile with orthogonal PA jumps is affected by different values of $\tau_s$ .
" Again, the profile has a resolution of tsamp ¥0.39 ms."," Again, the profile has a resolution of $t_{\rm samp}\approx$ 0.39 ms."
" Large 7. (24 ms) values result in obvious smearing of the total power profile, however the polarization PA is affected significantly from low 7, values."," Large $\tau_s$ $\geq$ 4 ms) values result in obvious smearing of the total power profile, however the polarization PA is affected significantly from low $\tau_s$ values."
" Note how the steep, middle part of the PA swing, becomes flat at 7,—2 ms."," Note how the steep, middle part of the PA swing, becomes flat at $\tau_s=2$ ms."
" The fact that small amounts of scattering significantly distort the pulsar PA swing, has the direct consequence that observed PA curves are not independent of frequency."," The fact that small amounts of scattering significantly distort the pulsar PA swing, has the direct consequence that observed PA curves are not independent of frequency."
" This complicates the interpretation of PA swings, especially in comparing data from multiple, widely spaced frequencies."," This complicates the interpretation of PA swings, especially in comparing data from multiple, widely spaced frequencies."
 Another consequence of scattering manifests itself within the relatively narrow band of an observation at a given frequency., Another consequence of scattering manifests itself within the relatively narrow band of an observation at a given frequency.
" Figure 4 shows a simulated pulse profile with an orthogonal PA jump and the PA swings of two frequency channels, 256 MHz apart and centred at 1.4 GHz."," Figure \ref{figure3} shows a simulated pulse profile with an orthogonal PA jump and the PA swings of two frequency channels, 256 MHz apart and centred at 1.4 GHz."
 The circles correspond to the highest and the crosses to the lowest frequency., The circles correspond to the highest and the crosses to the lowest frequency.
" The PA swings are obviously different, in particular after the pulse phase of the orthogonal jump."," The PA swings are obviously different, in particular after the pulse phase of the orthogonal jump."
 This difference will result in slightly different polarization profilesfor observations with different bandwidth: for two, This difference will result in slightly different polarization profilesfor observations with different bandwidth: for two
onlv has the effects of the planetary Doppler shift imprinted on it.,only has the effects of the planetary Doppler shift imprinted on it.
 For the spectra that we present in the following section. we assume that the stellar Doppler shift is known and has been appropriately accounted for.," For the spectra that we present in the following section, we assume that the stellar Doppler shift is known and has been appropriately accounted for."
 The entire 2291-2350 min section of (he hot Jupiter (iransmission spectrum for (he a(mosphere model is shown in Figure 5 with no velocity shifts., The entire 2291-2350 nm section of the hot Jupiter transmission spectrum for the drag-free atmosphere model is shown in Figure \ref{transmission_full} with no velocity shifts.
 In (he unshifted spectrum. the first overtone band of CO is clearly visible at the model spectral resolution of 105.," In the unshifted spectrum, the first overtone band of CO is clearly visible at the model spectral resolution of $10^6$."
 Many additional spectral features from water are also present in (his region of the spectrum along with weaker CII; features that originate from cooler regions in (he atmosphere., Many additional spectral features from water are also present in this region of the spectrum along with weaker $_4$ features that originate from cooler regions in the atmosphere.
 The comb of strong and well-separated CO lines around 2.3 microns is particularly useful for measuring Doppler shifts as long as the spectral resolution exceeds —10/., The comb of strong and well-separated CO lines around 2.3 microns is particularly useful for measuring Doppler shifts as long as the spectral resolution exceeds $\sim$ $^4$.
 At the much lower spectral resolution of most transmission spectrum measurements to date. the 2.3 jam CO features are unresolved in a single broad. absorption band. and RY measurements al a precision of ~1 km |! are not possible.," At the much lower spectral resolution of most transmission spectrum measurements to date, the 2.3 $\mu$ m CO features are unresolved in a single broad absorption band, and RV measurements at a precision of $\sim 1$ km $^{-1}$ are not possible."
 Figure G6. shows a representative 2.3 nm segment of the fully Doppler shifted spectra from 2308.0 to 2310.3 nm. (, Figure \ref{spectra_snapshot} shows a representative 2.3 nm segment of the fully Doppler shifted spectra from 2308.0 to 2310.3 nm. (
The rest of the spectrum shows similar RV shifts fom winds. but 1 is not feasible to show the entire spectral range of our model in a single plot and to still see (he Doppler shilts bv eve.),"The rest of the spectrum shows similar RV shifts from winds, but it is not feasible to show the entire spectral range of our model in a single plot and to still see the Doppler shifts by eye.)"
 A net blueshift in all four models is clearly apparent in the top panel of Figure 6..., A net blueshift in all four models is clearly apparent in the top panel of Figure \ref{spectra_snapshot}.
 Additionally. a significant amount of line broadening occurs in the (ransiission spectra [rom the joint effects of winds and rotation. which effectively. weakens the peak strength of each of the individual spectral lines ancl also causes many of the adjacent lines to blend together.," Additionally, a significant amount of line broadening occurs in the transmission spectra from the joint effects of winds and rotation, which effectively weakens the peak strength of each of the individual spectral lines and also causes many of the adjacent lines to blend together."
 The effect of broadening due to rotation alone is shown in the bottom panel of Figure 6.. clearly indicating the extent to which the Doppler broadening is due to rolation as opposed to winds.," The effect of broadening due to rotation alone is shown in the bottom panel of Figure \ref{spectra_snapshot}, clearly indicating the extent to which the Doppler broadening is due to rotation as opposed to winds."
 Rotation is the dominant source of broadening (quite similar to the elfect of rotational broadening on stellar spectra)., Rotation is the dominant source of broadening (quite similar to the effect of rotational broadening on stellar spectra).
 ILowever. the winds themselves also cause some hurther broadening of the (transmission spectra. since the wind patterns are nol entirely coherent ancl instead some considerable variation in the line-ol-sight wind speeds occurs around the terminator as shown in Figure 4..," However, the winds themselves also cause some further broadening of the transmission spectra, since the wind patterns are not entirely coherent and instead some considerable variation in the line-of-sight wind speeds occurs around the terminator as shown in Figure \ref{fig:dshift}."
 Some of (he models clearly show a higher level of incoherence in the wind pattern than others (e.g. the standard drag-Iree model as opposed to the model with enhanced hivperdissipation). which has a small effect on the width of the broadened spectral lines from model to mocel.," Some of the models clearly show a higher level of incoherence in the wind pattern than others (e.g. the standard drag-free model as opposed to the model with enhanced hyperdissipation), which has a small effect on the width of the broadened spectral lines from model to model."
 By cross correlating the Doppler-shifted. spectra against the unshifted. transmission, By cross correlating the Doppler-shifted spectra against the unshifted transmission
consisting of a primary that is massive enough to have evolved to the RGB in less than a Hubble time and an M dwarf companion.,consisting of a primary that is massive enough to have evolved to the RGB in less than a Hubble time and an M dwarf companion.
 We illustrate this formation pathway with an example computed using the Binary Star Evolution code (BSE)) described in (2002)., We illustrate this formation pathway with an example computed using the Binary Star Evolution code ) described in \citet{Hurley:2002}.
". Initially, the binary consisted of a 1.5 primaryand a 0.25 companion with an orbital periodMo of 75 days."," Initially, the binary consisted of a $1.5 ~M_{\sun}$ primaryand a $0.25~M_{\sun}$ companion with an orbital period of 75 days."
" After Me2.9 Gyr, the primary began moving up the giant branch and expanded to fill its Roche lobe, resulting in unstable mass transfer."," After 2.9 Gyr, the primary began moving up the giant branch and expanded to fill its Roche lobe, resulting in unstable mass transfer."
 The system went through a CE phase during which the primary’s H envelope was ejected., The system went through a CE phase during which the primary's H envelope was ejected.
 The system emerged as a 0.33Mo He WD and a 0.25Mo companion with an orbital period of 4 hours., The system emerged as a $0.33~M_{\sun}$ He WD and a $0.25~M_{\sun}$ companion with an orbital period of 4 hours.
" The system then underwent a period of tidal readjustment that further reduced the orbital separation and caused the low-mass, hydrogen burning star to fill its Roche lobe."," The system then underwent a period of tidal readjustment that further reduced the orbital separation and caused the low-mass, hydrogen burning star to fill its Roche lobe."
" Low mass main sequence stars are deeply convective, so the mass transfer was unstable and the stars coalesced into a 0.57Mo star with a 0.33M; He core."," Low mass main sequence stars are deeply convective, so the mass transfer was unstable and the stars coalesced into a $0.57~M_{\sun}$ star with a $0.33~M_{\sun}$ He core."
" Using andSSE,, we calculated how long it would take a Mzays= star to evolve to a similar structure and found times Moof 79 and 82 Gyr, respectively."," Using and, we calculated how long it would take a $M_{ZAMS} =
0.6~M_{\sun}$ star to evolve to a similar structure and found times of 79 and 82 Gyr, respectively."
" However, through accelerated stellar evolution the coalesced binary reaches this evolutionary state in only 5.1 Gyr."," However, through accelerated stellar evolution the coalesced binary reaches this evolutionary state in only 5.1 Gyr."
" While was adequate for modeling the binary evolution, we needed to use and to investigate the evolution of the merger product."," While was adequate for modeling the binary evolution, we needed to use and to investigate the evolution of the merger product."
" After the binary coalesced, continued to evolve the merger product as a single star and, in the illustrative case described above, the merger product eventually reached the sdB stage."," After the binary coalesced, continued to evolve the merger product as a single star and, in the illustrative case described above, the merger product eventually reached the sdB stage."
" However, based on the merger product's core mass, assumed that it was a Mzams=2.5Mg star at the base of the RGB and evolved the star accordingly."," However, based on the merger product's core mass, assumed that it was a $M_{ZAMS} = 2.5~M_{\sun}$ star at the base of the RGB and evolved the star accordingly."
" This assumption is clearly not appropriate to model the evolution of a 0.57Mo star, so we used and to find stars with core-envelope structures similar to the merged star from themodel."," This assumption is clearly not appropriate to model the evolution of a $0.57~M_{\sun}$ star, so we used and to find stars with core-envelope structures similar to the merged star from themodel."
" As described above, the stellar evolution models showed that a star with Mzams=0.6Mo eventually formed a 0.33Mo He core surrounded by a ~0.25Mo H rich envelope."," As described above, the stellar evolution models showed that a star with $M_{ZAMS} = 0.6~M_{\sun}$ eventually formed a $0.33~M_{\sun}$ He core surrounded by a $\sim 0.25~M_{\sun}$ H rich envelope."
" When we continued the evolution of this star, the mass of the He core grew through H shell burning and ignited at the tip of the giant branch while Reimers mass loss removed the remaining H envelope."," When we continued the evolution of this star, the mass of the He core grew through H shell burning and ignited at the tip of the giant branch while Reimers mass loss removed the remaining H envelope."
 This additional evolution from the RGB to the sdB stage took 140 Myr., This additional evolution from the RGB to the sdB stage took 140 Myr.
 The evolutionary tracks computed by and are shown in[ure21., The evolutionary tracks computed by and are shown in.
" If the merged He WD-M dwarf maintains the distinct core-envelope structure of an RGB star, then the binary evolution scenario described above forms a 0.48Mo sdB in 5.2 Gyr."," If the merged He WD–M dwarf maintains the distinct core-envelope structure of an RGB star, then the binary evolution scenario described above forms a $0.48~M_{\sun}$ sdB in 5.2 Gyr."
 It is not clear whether the merged star will maintain the core-envelope structure of an RGB star., It is not clear whether the merged star will maintain the core-envelope structure of an RGB star.
" Instead, some mixing might occur that would alter the chemical profile assumed above."," Instead, some mixing might occur that would alter the chemical profile assumed above."
" To bracket the range of possible outcomes, we explored models in which the He WD mixed completely with the M dwarf and formed a homogeneous, He-rich star."," To bracket the range of possible outcomes, we explored models in which the He WD mixed completely with the M dwarf and formed a homogeneous, He-rich star."
" In the mergers considered here, we are mixing ~0.2Mo of material with standard abundances with 0.Mg H depleted material, resulting in a star with Y~0.6."," In the mergers considered here, we are mixing $\sim 0.2~M_{\sun}$ of material with standard abundances with $\sim 0.3~M_{\sun}$ H depleted material, resulting in a star with $Y \sim 0.6$."
" Usingstar,, we have modeled the evolution of these completely mixed stars and found that they too will form sdBs."," Using, we have modeled the evolution of these completely mixed stars and found that they too will form sdBs."
" shows evolutionary tracks in the g) plane for 0.6Mo and 0.7M mixed stars(log withT;r;,log η=0.5 and 0.7, respectively, and initial Y—0.6."," shows evolutionary tracks in the $\log~T_{eff},
\log~g$ ) plane for $0.6~M_{\sun}$ and $0.7~M_{\sun}$ mixed stars with $\eta = 0.5$ and 0.7, respectively, and initial $Y=0.6$."
" The less massive star takes 5.8 Gyr to become a sdB and remains in the sdB “box” for 210 Myr, while the moremassive star evolves to the sdB stage in only 3.3 Gyr and remains in this stage for 110 Myr."," The less massive star takes 5.8 Gyr to become a sdB and remains in the sdB “box” for 210 Myr, while the moremassive star evolves to the sdB stage in only 3.3 Gyr and remains in this stage for 110 Myr."
" The sdB stars have masses of 0.44Mo and 0.50Mo, respectively."," The sdB stars have masses of $0.44~M_{\sun}$ and $0.50~M_{\sun}$, respectively."
" For comparison, the evolutionary track of the 0.6Mc; mixed star is also shown in2]."," For comparison, the evolutionary track of the $0.6~M_{\sun}$ mixed star is also shown in."
" 'These models demonstrate that even if the preformed[Figure He core dissolves during the merger, these systems can still evolve into sdBs within a Hubble time."," These models demonstrate that even if the preformed He core dissolves during the merger, these systems can still evolve into sdBs within a Hubble time."
 This formation channel depends on several assumptions made in modeling binary stellar evolution., This formation channel depends on several assumptions made in modeling binary stellar evolution.
 The initial configuration of the binary must be such that the primary evolves to the giant branch and triggers an episode of unstable mass transfer that completely removes its envelope before the merger., The initial configuration of the binary must be such that the primary evolves to the giant branch and triggers an episode of unstable mass transfer that completely removes its envelope before the merger.
" Furthermore, the total mass of the merged star must be in the range 0.55—0.84Mo for the merged object to evolveinto a sdB. We have computed a grid of models with η=0.3 to make a preliminary exploration of the parameter space of progenitor systems with primary masses in the range 1—3 Mo, secondary masses in the range 0.1—0.8 Mo, periods in the range 1— 350d, and CE ejection efficiencies, «cg, in the range 0.5— 1.5."," Furthermore, the total mass of the merged star must be in the range $0.53 - 0.84~M_{\sun}$ for the merged object to evolveinto a sdB. We have computed a grid of models with $\eta = 0.3$ to make a preliminary exploration of the parameter space of progenitor systems with primary masses in the range $1 - 3~M_{\sun}$ , secondary masses in the range $0.1
- 0.8~M_{\sun}$ , periods in the range $1 - 350$ d, and CE ejection efficiencies, $\alpha_{CE}$ , in the range $0.5 - 1.5$ ."
" Of the 5x104 models in our grid, produce merger products"," Of the $5\times10^{4}$ models in our grid, produce merger products"
Analysis of the dy line profile can help us determine some of the intrinsic properties of the central black hole in AGN and its accretion disc (e.g. Fabian ct al.,Analysis of the K line profile can help us determine some of the intrinsic properties of the central black hole in AGN and its accretion disc (e.g. Fabian et al.
 1989: Laor 1991)., 1989; Laor 1991).
 Foremost in the current climate is the determination of black hole (BLD) spin. for example recent observations by Miniutti et al. (," Foremost in the current climate is the determination of black hole (BH) spin, for example recent observations by Miniutti et al. ("
2007 2009) and Schmoll et al. (,2007 2009) and Schmoll et al. (
2009) have analysed the broad. Wha region in order to constrain the spin of the central BLL.,2009) have analysed the broad $\alpha$ region in order to constrain the spin of the central BH.
 A black hole can be characterised simply by ils mass and its spin., A black hole can be characterised simply by its mass and its spin.
 Many objects have now been classified in terms of their mass into three categories: Galactic DII. Intermediate Mass. BIL and Supermassive DII.," Many objects have now been classified in terms of their mass into three categories: Galactic BH, Intermediate Mass BH and Supermassive BH."
 The spin of the black hole determines he nature of the space-time metric in the regions close to it., The spin of the black hole determines the nature of the space-time metric in the regions close to it.
" The spin parameter e=ο)fOAl-"" (where J=aneular momentum and 0<a« 0.998) is used to describe the spin oft1e BLL where e is limited toa maximum value of e=0.998 at the Thorne limit."," The spin parameter $a=cJ/GM^2$ (where J=angular momentum and $0<a<0.998$ ) is used to describe the spin of the BH, where $a$ is limited to a maximum value of $a=0.998$ at the Thorne limit."
 This is due to photon capture in which photons travelling on negative angular momentum orbits are »eferentially captured by the black hole therefore producing an upper bound to the spin parameter e and hence limiting the innermost stable circular orbit to a minimum of risco.=1.235r4 (Phorne 1974).," This is due to photon capture in which photons travelling on 'negative' angular momentum orbits are preferentially captured by the black hole therefore producing an upper bound to the spin parameter $a$ and hence limiting the innermost stable circular orbit to a minimum of $r_{\rm ISCO}=1.235\,r_{\rm g}$ (Thorne 1974)."
 Constraining the spin of SMDlIs in AGN and studying the distribution of black hole spin can aid our understanding of the evolution of AGN and the black holes themselves e.g. mergers. relativistic jets and variability (Blandford Znajek 1971: Volonteri et al.," Constraining the spin of SMBHs in AGN and studying the distribution of black hole spin can aid our understanding of the evolution of AGN and the black holes themselves e.g. mergers, relativistic jets and variability (Blandford Znajek 1977; Volonteri et al."
 2007: Ixing. Prinele Hofmann 2008).," 2007; King, Pringle Hofmann 2008)."
 Line emission from the inner regions of the accretion disc can become broadened: due to relativistic effects. and Doppler motions (Fabian et al., Line emission from the inner regions of the accretion disc can become broadened due to relativistic effects and Doppler motions (Fabian et al.
 1989) resulting in an asvmmetrie profile., 1989) resulting in an asymmetric profile.
 Evidence for such broadening was found using the A-rav CCD detectors onboardUA (Tanaka. Inoue Πο 1994) by Alushotzky et al. (," Evidence for such broadening was found using the X-ray CCD detectors onboard (Tanaka, Inoue Holt 1994) by Mushotzky et al. ("
1995). ‘Tanaka et al. (,"1995), Tanaka et al. ("
1995) and Nanelra ct al. (,1995) and Nandra et al. (
1997) typically over the kkeV. range.,1997) typically over the keV range.
 Now with N-rav spectra of increasing quality and over wider energy ranges such as that obtained with the NIS. anc high. energy. UND detectors. onboard GO.OkkeV. Alitsucda et al.," Now with X-ray spectra of increasing quality and over wider energy ranges such as that obtained with the XIS and high energy HXD detectors onboard keV, Mitsuda et al."
 2007). the Wkly region of AGN can be examined in detail.," 2007), the K region of AGN can be examined in detail."
 High cnerey X-ray, High energy X-ray
current observations.,current observations.
 In addition. one can estimate the dust content of individual DLA from the ratio of depleted onto uncepleted elements it contains.," In addition, one can estimate the dust content of individual DLA from the ratio of depleted onto undepleted elements it contains."
 H is assumed that a fraction of the elements is not detected in the gas phase because it ds locked. into dust erains (e.g. Pettinietal.(2000)., It is assumed that a fraction of the elements is not detected in the gas phase because it is locked into dust grains (e.g. \citet[]{pettini2000}.
. Such studies have shown that the dust. depletion. levels observed in DLAs are small. lower than those measured in the Galactic interstcllar-meditun clouds with similar ΑΗ) column clensities.," Such studies have shown that the dust depletion levels observed in DLAs are small, lower than those measured in the Galactic interstellar-medium clouds with similar N(HI) column densities."
 In fact. the observations resemble the dust depletion of the Galactic halo warm gas (Ixulkarniet1997:Savaglioetal. 9000).," In fact, the observations resemble the dust depletion of the Galactic halo warm gas \citep[]{kulkarni1997, savaglio2000}."
".. Another more direct way to probe the dust content of DLAs is to look for a 2175 feature (the so-called ""dust bump"") in the quasar spectrum.", Another more direct way to probe the dust content of DLAs is to look for a 2175 feature (the so-called “dust bump”) in the quasar spectrum.
 This feature is known to be more predominant in the Milky Way than in the Small and Large Magellanie Clouds., This feature is known to be more predominant in the Milky Way than in the Small and Large Magellanic Clouds.
 Many searches Lead to no-detections (Jura1977:Smithetal.1979:1989:Peietal. 1991).," Many searches lead to no-detections \citep[]{jura1977, smith1979, boisse1988, lanzetta1989, fall1989, pei1991}."
. Using Sloan Digital Sky Survey (SDSS) spectra. Wangetal.(2004). claimed the detections of bumps in three individual Alell absorbers at. z=1.5. while Junkkarinenetal.(2004) have detected a clear bump ina DLA at z=0.524 toward a BL Lac at z-0.94.," Using Sloan Digital Sky Survey (SDSS) spectra, \citet[]{wang2004} claimed the detections of bumps in three individual MgII absorbers at z=1.5, while \citet[]{junkkarinen2004}{ have detected a clear bump in a DLA at z=0.524 toward a BL Lac at z=0.94."
 Other searches remained unsuccessful. indicating that most DLAs do not show a preclominant 2175 feature.," Other searches remained unsuccessful, indicating that most DLAs do not show a predominant 2175 feature."
 A promising tool for finding the signatures of dust are searches for the strong silicate absorption features in the infrared. Ixulkarnictal.(2007).," A promising tool for finding the signatures of dust are searches for the strong silicate absorption features in the infrared. \citet[]{kulkarni2007},"
.. using Spitzer data. report on the first indication of 9.7 p/m silicate absorption in a DL at ως=0.52. and find it to be stronger than expected rom extrapolation of the relation from Galactic dilluse ISM.," using Spitzer data, report on the first indication of 9.7 $\mu$ m silicate absorption in a DLA at $_{abs}=0.52$, and find it to be stronger than expected from extrapolation of the relation from Galactic diffuse ISM."
 A consequence of DLAs being dusty. is that it might event the inclusion of their. background. quasar in colour-colour selected samples., A consequence of DLAs being dusty is that it might prevent the inclusion of their background quasar in colour-colour selected samples.
 In. order to overcome such ases. raclio-selected quasar samples have been built recently: while 2001Qéllison) did not discover a missing »opulation. a more recent survey finds 4 fields (out of ~60) where the radio-Dux selected quasars have no optical counterparts down to very faint limits (Jorgensonetal. 2006).," In order to overcome such biases, radio-selected quasar samples have been built recently: while \citet[]{ellison2001}{ did not discover a missing population, a more recent survey finds 4 fields (out of $\sim$ 60) where the radio-flux selected quasars have no optical counterparts down to very faint limits \citep[]{jorgenson2006}."
. “These objects are the potential candidates of cust. DLAs obscuring background. In this paper. we investigate νο another possible signature of the dust. content of DLAs.," These objects are the potential candidates of dusty DLAs obscuring background In this paper, we investigate yet another possible signature of the dust content of DLAs."
 Indeed. dust in a quasar absorber will redden the colour of its background quasar. Peietal.(1991)," Indeed, dust in a quasar absorber will redden the colour of its background quasar. \citet[]{pei1991}{"
 were the first to look for reddening of the quasar continua in a sample containing DLAs and detect reddening at level., were the first to look for reddening of the quasar continua in a sample containing DLAs and detect reddening at level.
 More recently. Murphy&Liske(2004) have used 79 DLAs from SDSS Data ltelease 2 to compare the slope of power laws fitting quasar continua with and withoutDLAs.," More recently, \citet[]{murphy2004}{ have used 79 DLAs from SDSS Data Release 2 to compare the slope of power laws fitting quasar continua with and withoutDLAs."
 These authors found a colour excess due to Small Magellanie Cloud type of dust of reddening. <LE(B-V)>> «0.01 mae., These authors found a colour excess due to Small Magellanic Cloud type of dust of reddening $<$ $>$ $<$ 0.01 mag.
 Using colour excess measurements lor SDSS siehtlines containing DLAs and matching sightlines without them. Vladiloetal.(2008)report a detection of reddening in a sample of DRS SDSS QSOs towards sightlines containing DLAs at. the mean level of <LE(B-V)> <6.3x10 7.," Using colour excess measurements for SDSS sightlines containing DLAs and matching sightlines without them, \citet[]{vladilo2008} report a detection of reddening in a sample of DR5 SDSS QSOs towards sightlines containing DLAs at the mean level of $<$ $>$ $\sim$ $^{-3}$."
 They derive their estimate by comparing colours of QSOs with intervening DLAs to the mean colour. of a set of QSO of similar emission redshift and. brightness. ancl bootstrapping the resulting broad distribution of colour excesses to ascertain the slightly redder colours of the DLA QSOs.," They derive their estimate by comparing colours of QSOs with intervening DLAs to the mean colour of a set of QSO of similar emission redshift and brightness, and bootstrapping the resulting broad distribution of colour excesses to ascertain the slightly redder colours of the DLA QSOs."
 X byproduct of the Pontzen&Pettini(2009) work. who performed a Bayesian analysis on a variety. of optical and. radio datasets. to determine the loss fraction of QSOs due to dust obsceuration. is an estimate of the probability density distribution for dust reddening in optical spectra. that also peaks around the same value for «I(D-V)7.," A byproduct of the \citet[]{pontzen2009}{ work, who performed a Bayesian analysis on a variety of optical and radio datasets to determine the loss fraction of QSOs due to dust obscuration, is an estimate of the probability density distribution for dust reddening in optical spectra, that also peaks around the same value for $<$ $>$."
 Here. we build composite spectra of quasar spectra with DLAs from SDSS Data Releases 5 and 7. and compare them with reference composite spectra built to match the original spectra in magnitudes ane redshifts.," Here, we build composite spectra of quasar spectra with DLAs from SDSS Data Releases 5 and 7, and compare them with reference composite spectra built to match the original spectra in magnitudes and redshifts."
 The organisation of the paper is as follows: after the introduction. we describe the selection of the DLA samples in section 2. followed by the construction of the matching non-absorber samples. and the procedure for obtaining composite spectra in section 3. where we also presents the method for deriving extinction as well as the results for the full samples.," The organisation of the paper is as follows: after the introduction, we describe the selection of the DLA samples in section 2, followed by the construction of the matching non-absorber samples, and the procedure for obtaining composite spectra in section 3, where we also presents the method for deriving extinction as well as the results for the full samples."
 After discussing subsample analyses (section 4). we summarise and conclude in section We start the selection of our saniples from two lists of SDSS surveys lor DLAs.," After discussing subsample analyses (section 4), we summarise and conclude in section We start the selection of our samples from two lists of SDSS surveys for DLAs."
 Prochaska&Wolfe(2009) searched in Data Release (Di) 5. following a partially automated: procedure. tested on earlier data releases (Prochaska&Lerbert-Fort2004:rochaskactal. 2005).," \citet[]{prochaska2009} searched in Data Release (DR) 5, following a partially automated procedure tested on earlier data releases \citep[]{prochaska2004, prochaska2005}."
. Their cleanest Cstatistical’) sample contains 738 DLAs., Their cleanest ('statistical') sample contains 738 DLAs.
 We have chosen not to include objects hat are listed. in their non-statistical sample. as these can be towards QSOs with strong intrinsic absorption. or WE Zab. Zoso.," We have chosen not to include objects that are listed in their non-statistical sample, as these can be towards QSOs with strong intrinsic absorption, or have $_{abs} \sim$ $_{QSO}$."
 Broad absorption line QSO (DALs) are known to be more strongly. redcdened in the rest-frame UV than non-BAL QSOs. and hence this list of DLAs is hen cross-correlatecd with the BAL catalogue of Gibsonetal. (2009).. in order to avoid such BALs entering our analysis.," Broad absorption line QSO (BALs) are known to be more strongly reddened in the rest-frame UV than non-BAL QSOs, and hence this list of DLAs is then cross-correlated with the BAL catalogue of \citet[]{gibson2009}, in order to avoid such BALs entering our analysis."
 Furthermore. we have restricted ourselves to those sightlines towards QSOs in the Prochaska&Wolfe(2009) sample that only exhibit one single DLA.," Furthermore, we have restricted ourselves to those sightlines towards QSOs in the \citet[]{prochaska2009}{ sample that only exhibit one single DLA."
 Phus. we exclude the 59 respectively 5 cases. where two or three DLAs are founcl along the same line of sight.," Thus, we exclude the 59 respectively 5 cases, where two or three DLAs are found along the same line of sight."
 Vhis leaves us with a sample of 526 SDSS DRS OSO sightlines with securely identified DLA well suited. for the analysis of their. dust Noterdaemeetal.(2009) have surveyed DRT. and present a list of 1426 strong LLL absorbers at recdshilts 2.415x ο x 5.2. of which 937 svstems have log N(IIL 20.3.," This leaves us with a sample of 526 SDSS DR5 QSO sightlines with securely identified DLA well suited for the analysis of their dust \citet[]{noterdaeme2009}{ have surveyed DR7, and present a list of 1426 strong HI absorbers at redshifts $\leq$ z $\leq$ 5.2, of which 937 systems have log N(HI $\geq$ 20.3."
 Applying the same criteria for filtering out potentially problematic sightlines as above. we arrive at 731 SDSS DRT OSO sightlines for this DLA," Applying the same criteria for filtering out potentially problematic sightlines as above, we arrive at 731 SDSS DR7 QSO sightlines for this DLA"
2 are (he name of the filler set and the name of the filter (either broad- or narrow-band).,2 are the name of the filter set and the name of the filter (either broad- or narrow-band).
 In the next three columns we show Ay (column 3). \/7F (column 4) and A (column 5) computed directly Irom the transmittance.," In the next three columns we show $\lambda_0$ (column 3), $\sqrt{\mu^2}$ (column 4) and $\Delta$ (column 5) computed directly from the transmittance."
 Column 6 is the estimated mean value of À. computed with Eq. 12..," Column 6 is the estimated mean value of $\lambda_z$ computed with Eq. \ref{eq:mean2},"
 concordance cosmology and the line. not including the corrections for the lines (see bellow).," concordance cosmology and the line, not including the corrections for the lines (see bellow)."
 With A. we compute the elective width of the filter in column 7., With $\bar\lambda_z$ we compute the effective width of the filter in column 7.
 The difference between commonly used cosmologies and different emission lines is less than in all cases., The difference between commonly used cosmologies and different emission lines is less than in all cases.
 Note that the filter #1194. when used together wilh #2208 has an effective width roughly three times its width.," Note that the filter 194, when used together with 208 has an effective width roughly three times its width."
 The wavelength where the emission line is located has low transmittance., The wavelength where the emission line is located has low transmittance.
 In this case. using A instead of A’ would underestimate the continuum this.," In this case, using $\Delta$ instead of $\Delta'$ would underestimate the continuum flux."
 In Table 2.. we stummarize the mean values of the combined width and ó parameter (Eq 32 lor and [Oru]--computed. quantities with concordance cosmology.," In Table \ref{tab:width:mline}, we summarize the mean values of the combined width and $\phi$ parameter (Eq \ref{eq:combinecover} for and -computed quantities with concordance cosmology."
 Columns 1 and 2 have the same meaning in Table 1.., Columns 1 and 2 have the same meaning in Table \ref{tab:properties}. .
 Columns 3 to 5 correspond to Io--computed quantities with r—I([NuA6584) )—0.32 (Sect. 4.4.1)), Columns 3 to 5 correspond to -computed quantities with $r=$ =0.32 (Sect. \ref{sec:hanii}) )
 and columns 6 to 8 to with r—JHOE) /E(O0n1A5007))—1.05 (Sect.4.4.2))., and columns 6 to 8 to with $r=$ )=1.05 \ref{sec:hboiii}) ).
 Columns 3 and 6 are the mean wavelengths computed according to Eq 12.. using the value of A”.," Columns 3 and 6 are the mean wavelengths computed according to Eq \ref{eq:mean2}, using the value of $\Delta''$."
" Columns 4 and 7 the value of the combined width at the computed mean wavelength value and columns 5 and 8 (he parameter ©, that denotes the the fraction of the total flux that comes from the main line."," Columns 4 and 7 the value of the combined width at the computed mean wavelength value and columns 5 and 8 the parameter $\phi$, that denotes the the fraction of the total flux that comes from the main line."
 ForHo.. the narrow-band filter in the CAFOSS200 filter set has the same © value than the broad-band fillers.," For, the narrow-band filter in the CAFOS8200 filter set has the same $\phi$ value than the broad-band filters."
 This means that is wide enough to contain simultaneously the two nitrogen lines andIla., This means that is wide enough to contain simultaneously the two nitrogen lines and.
. The narrow-band fillers of WECS200 and WFC9200 have a greater oO. meaning (hat the lines only enter partially inside the narrow line filter.," The narrow-band filters of WFC8200 and WFC9200 have a greater $\phi$, meaning that the lines only enter partially inside the narrow line filter."
 In this (wo last cases. and of the fIux in the narrow-band filter come fromΠα... instead of the nominal71%.," In this two last cases, and of the flux in the narrow-band filter come from, instead of the nominal."
.. Using the standard correction with very narrow filters such as the narrow-band filters of WFC. produce a systematic under estimation of the line flux of about15%.," Using the standard correction with very narrow filters such as the narrow-band filters of WFC, produce a systematic under estimation of the line flux of about."
.. In the case of[Ort]... © is very close to unity for the narrow-band filters.," In the case of, $\phi$ is very close to unity for the narrow-band filters."
 This means that will enter alone in the filter., This means that will enter alone in the filter.
 In the case of broad-band. the ratio is close io the limit value (0.42 with the assumed value of 7). meaning that the three lines enter the broad-band filler.," In the case of broad-band, the ratio is close to the limit value (0.42 with the assumed value of $r$ ), meaning that the three lines enter the broad-band filter."
 As a summary. we can consider that enters alone in the narrow-band filter. but we have to include the three lines in the broad-band filter.," As a summary, we can consider that enters alone in the narrow-band filter, but we have to include the three lines in the broad-band filter."
 We use galaxv templates to caleulate the color evolution of different.ealaxies with, We use galaxy templates to calculate the color evolution of differentgalaxies with
2008:: Flovd.Bate.&Webster 2009)).,; \citealt{fbw09}) ).
 Iu those analyses. we mareinalised over the smooth matter percentage m the leus as a nuisance parameter.," In those analyses, we marginalised over the smooth matter percentage in the lens as a nuisance parameter."
 Tere. we turu the problem around aud iustead iiarginalise over the quasar paralucters to obtain coustraiuts on the smooth matter percentage m the Ileus at the image positions.," Here, we turn the problem around and instead marginalise over the quasar parameters to obtain constraints on the smooth matter percentage in the lens at the image positions."
 Rough microleusiug measurements of smooth matter percentages have been repored previously, Rough microlensing measurements of smooth matter percentages have been reported previously.
 Spectroscopy of SDSS JO921|0219 uudertaken by I&eetonetal.-(2006) sugeested a sinootl matter percentage of SO to per ceut in that leus at the locaion of the D and A images., Spectroscopy of SDSS J0924+0219 undertaken by \citet{keeton+06} suggested a smooth matter percentage of 80 to 85 per cent in that lens at the location of the $D$ and $A$ images.
 Using X-ray mouitoring of TE 110111805. Chartaset reported a smooth matter percentage of SNO per cent is favoured.," Using X-ray monitoring of HE 1104+1805, \citet{chartas+09} reported a smooth matter percentage of $\sim80$ per cent is favoured."
 Pooleyetal.(2009) measured the snooth matter percentage in PC L115|080 to be ~90 per cent. using X-ray observations.," \citet{pooley+09} measured the smooth matter percentage in PG 1115+080 to be $\sim90$ per cent, using X-ray observations."
 Moreauctal.(2008) found a weal trend supporting this result., \citet{metal08} found a weak trend supporting this result.
 Most receutly. Daietal.(2009). favoured a uooth matter fraction of ~το per cent using X-ray and optical monitoriug of RNJ 1131-1231.," Most recently, \citet{dai+09} favoured a smooth matter fraction of $\sim70$ per cent using X-ray and optical monitoring of RXJ 1131-1231."
 Microleusimg analyses conusisteutlv xediet a significant smooth matter percentage in the eusing galaxy at the position of anomalous imiages., Microlensing analyses consistently predict a significant smooth matter percentage in the lensing galaxy at the position of anomalous images.
" Iu this paper. we present constraints on the dark uatter percentages in three lensing ealaxies: MC MILL05.1. SDSS J092110219 and Q2237|0305,"," In this paper, we present constraints on the dark matter percentages in three lensing galaxies: MG 0414+0534, SDSS J0924+0219 and Q2237+0305."
 MG MLLO53L and SDSS J092110219 are both leused by earlv-tvpe galaxies. and consist of close image pairs displaving a flux ratio anomaly.," MG 0414+0534 and SDSS J0924+0219 are both lensed by early-type galaxies, and consist of close image pairs displaying a flux ratio anomaly."
 MC 04111053115 uoderatelv anomalous. whereas SDSS J0921]|0219 is he nost anomalous lensed quasar currentlv known.," MG 0414+0534 is moderately anomalous, whereas SDSS J0924+0219 is the most anomalous lensed quasar currently known."
 (92237|0305 differs frou the previous sources in two key ware: it is lensed by a barred spiral ealaxy. an it does rot contain a close dmaee pair.," Q2237+0305 differs from the previous sources in two key ways: it is lensed by a barred spiral galaxy, and it does not contain a close image pair."
 Nevertheless. is known o be affected by microleusiug(e.g. Irwinetal.it 1989)).," Nevertheless, it is known to be affected by microlensing (e.g. \citealt{irwin+89}) )."
 This paper is laid out as follows: iu Section 2 we discuss he observational data ou the three svstems of interest., This paper is laid out as follows: in Section \ref{sec:obs} we discuss the observational data on the three systems of interest.
 The simulation technique is briciiy described in Section 3.., The simulation technique is briefly described in Section \ref{sec:sims}.
 We present our results and discussion iu Section L.. and conclude in Section 5..," We present our results and discussion in Section \ref{sec:results}, and conclude in Section \ref{sec:conclusions}."
" Throughout this paper we use a cosmology with My=vüknisMpe+. Q,,=0.3 aud O4= 0.7. A"," Throughout this paper we use a cosmology with $H_0=70\rm{kms^{-1}Mpc^{-1}}$, $\Omega_m=0.3$ and $\Omega_{\Lambda}=0.7$ ."
ICCO LL11)053 Lwas ciscovered by Hewittetal.(1992)., MG 0414+0534 was discovered by \citet{hewitt+92}.
 It consists of a background quasar at τν=2.61 (Lawrenceetal.1995) aud a foreground carly-type lensing galaxy at uo0.96 (Tonry&Iochauek1999)., It consists of a background quasar at $z_s=2.64$ \citep{lejt95} and a foreground early-type lensing galaxy at $z_l=0.96$ \citep{tk99}.
. Four images of the quasar are observed. with the close image pair (mages ly and ly) displawing a flux ratio anomaly.," Four images of the quasar are observed, with the close image pair (images $A_1$ and $A_2$ ) displaying a flux ratio anomaly."
 This anomaly is weak in both the mid-infrared (ο=0.90+0.01 on 2005 October 10. Minezakietal. 2009)) and the radio οι=0.90c0.02 on 1990 April 2. Katz 1993)). but somewhat stronger in the optical +£0.06 on 1991 November 2-1. Schechter&Moore 1993)).," This anomaly is weak in both the mid-infrared $A_2/A_1 = 0.90 \pm 0.04$ on 2005 October 10, \citealt{minezaki+09}) ) and the radio $A_2/A_1 = 0.90 \pm 0.02$ on 1990 April 2, \citealt{kh93}) ), but somewhat stronger in the optical $A_2/A_1 = 0.45\pm 0.06$ on 1991 November 2-4, \citealt{sm93}) )."
 Tn our analysis. we used three epochs of ultiwavelength ALG ΟΠΟΡΟΙ observations. preseuted in Table 1..," In our analysis, we used three epochs of multi-wavelength MG 0414+0534 observations, presented in Table \ref{0414obs}."
 The first two epochs were archival IIST data. obtained frou the CASTLES Survey (Falco. 1997).," The first two epochs were archival HST data, obtained from the CASTLES Survey \citep{fls97}."
. The third epoch was obtained by us using the Maeclau 6.5-netre Baade telescope., The third epoch was obtained by us using the Magellan 6.5-metre Baade telescope.
 These data were first prescuted in Batectal.(2008)., These data were first presented in \citet{bfww08}.
. SDSS J092L|0219 is the most anomalous lensed quasar currently known., SDSS J0924+0219 is the most anomalous lensed quasar currently known.
 The τὰinmun image A has been observed to be a factor of ~20 brighter than the saddle point inage D in the optical (IXeetonetal.2006)., The minimum image $A$ has been observed to be a factor of $\sim20$ brighter than the saddle point image $D$ in the optical \citep{keeton+06}.
. The quasar was discovered by Ivadaetal.(2003) in Sloan Digital Sky Survey. (SDSS) Huaging. and consists of an carly-type lensing galaxy at uoc0.391 (Eigeubrod2006) and a backerouned qiasar at zy=1.521 (Inadaetal. 2003).," The quasar was discovered by \citet{inada03} in Sloan Digital Sky Survey (SDSS) imaging, and consists of an early-type lensing galaxy at $z_l = 0.394$ \citep{eigenbrod06a} and a background quasar at $z_s=1.524$ \citep{inada03}."
. Again. we use three οιοι» of observational data.," Again, we use three epochs of observational data."
 These are preseuted in Tate 2., These are presented in Table \ref{0924obs}.
 The 2008 Much 21 data were obtained bv ous using the Magellan: 6.5-metre Baade telescope (Floyd.Bate.&Webster2009)., The 2008 March 21 data were obtained by us using the Magellan 6.5-metre Baade telescope \citep{fbw09}.
. The 20053 November 18-25 data were taken using the TST/NICALIOS aud WEPC2 instruments as part of the CASTLES Survey (Ixeetonetal.2006)., The 2003 November 18-23 data were taken using the HST/NICMOS and WFPC2 instruments as part of the CASTLES Survey \citep{keeton+06}.
. The 2001 December 15 were obtained by Inadaetal.(2003). using the MagIC instrmucut on fhe Baacle telescope. aud reduced by us (details can be found in Flovd.Bate.&Webster2009) ).," The 2001 December 15 were obtained by \citet{inada03} using the MagIC instrument on the Baade telescope, and re-reduced by us (details can be found in \citealt{fbw09}) )."
 exavitationallyQ2237|]0305. is perhaps the most welbstudied lensed) quasar., Q2237+0305 is perhaps the most well-studied gravitationally lensed quasar.
 It was discovered by IIuchraetal.(1985)... aud consists of a lensing ealaxy at uo=0.0391 aud a background quasar at ty=1.6095.," It was discovered by \citet{huchra+85}, and consists of a lensing galaxy at $z_l=0.0394$ and a background quasar at $z_s=1.695$."
 The two previous sources had early type leusing galaxies: the lens in Q2237|0305 is 1| barred spiral., The two previous sources had early type lensing galaxies; the lens in Q2237+0305 is a barred spiral.
 Near-perfect aliguinent between observev. leus and quasar results in four virtually sviunietric nuages of the backeround source. located in the bulex| of the lensing galaxy.," Near-perfect alignment between observer, lens and quasar results in four virtually symmetric images of the background source, located in the bulge of the lensing galaxy."
 The optical depth to stars is therefore quite high. maline the svstem an excellent tare’et for microlensing analyses.," The optical depth to stars is therefore quite high, making the system an excellent target for microlensing analyses."
 Typically. this is taken to mean the smeoth matterpercentage in microlensing simulations cau be set to zero.," Typically, this is taken to mean the smooth matterpercentage in microlensing simulations can be set to zero."
 We test this assumption1 here., We will test this assumption here.
 (Q2237willit|0305 also cliffers froni the two previous sources in that does not contain a closeimage pair displaviug, Q2237+0305 also differs from the two previous sources in that it does not contain a close image pair displaying
emission at the location of the weaker OVI line (Figure 5).,emission at the location of the weaker OVI line (Figure 5).
" NGC 1549: There is very little Galactic absorption or extinction LO?"" 7. and there is little evidence of Hs absorption (Figure 6)."," NGC 1549: There is very little Galactic absorption or extinction $\times$ $^{20}$ $^{-2}$, and there is little evidence of $_{2}$ absorption (Figure 6)."
 Absorption by Galactic atomic gas is present and the CII AI036 line (0.6 wwide) lies in the middle of the redshifted OVI A1032 line. although emission near this region is seen in the Lif2b channel as well.," Absorption by Galactic atomic gas is present and the CII $\lambda$ 1036 line (0.6 wide) lies in the middle of the redshifted OVI $\lambda$ 1032 line, although emission near this region is seen in the Lif2b channel as well."
 The weaker line. OVI AI033. is in an uneontaminatec region and it is detected with à FWIIM of about 1.2A.," The weaker line, OVI $\lambda$ 1038, is in an uncontaminated region and it is detected with a FWHM of about 1.2."
 The flux is elevated just to the blue side of the CII A1036 line. whichmax be the unabsorbed part of the OVI A1032 Line.," The flux is elevated just to the blue side of the CII $\lambda$ 1036 line, whichmay be the unabsorbed part of the OVI $\lambda$ 1032 line."
 The line width is about 0.7 of the FWIIAI of the stellar velocity dispersion., The line width is about 0.7 of the FWHM of the stellar velocity dispersion.
 There is no emission from CIII A977., There is no emission from CIII $\lambda$ 977.
 There is a narrow emission line at the location of the Galactic OVI line., There is a narrow emission line at the location of the Galactic OVI line.
 NGC 3115: Galactic molecular absorption lines are strong. but the redshifted OVI lines lie in uncontaminated parts of the spectrum.," NGC 3115: Galactic molecular absorption lines are strong, but the redshifted OVI lines lie in uncontaminated parts of the spectrum."
 There is no emission from OVI (Figure 7)., There is no emission from OVI (Figure 7).
 NGC 3319: The redshift of this galaxy places the OVI lines in a part of the spectrum uncontaminated bv Galactic absorption., NGC 3379: The redshift of this galaxy places the OVI lines in a part of the spectrum uncontaminated by Galactic absorption.
 No OVI emission is found (Figure 8)., No OVI emission is found (Figure 8).
 NGC 3585: This fairly isolated. E'7/S0 galaxy lies in a loose group and has moderate Galactic reddening and an HI column of 5.6x 10720 7., NGC 3585: This fairly isolated E7/S0 galaxy lies in a loose group and has moderate Galactic reddening and an HI column of $\times$ $^{20}$ $^{-2}$.
 There were problems that occurred when (his spectrum was being obtained. aud of the 16 ksee exposure. only 3 ksec was acceptable. nearly all during the day (Figure 9).," There were problems that occurred when this spectrum was being obtained, and of the 16 ksec exposure, only 3 ksec was acceptable, nearly all during the day (Figure 9)."
 Consequently. the airglow lines are strong.," Consequently, the airglow lines are strong."
 Absorption by Galactic atomic and molecular gas is nearly absent. which is surprising given the column clensity of the gas.," Absorption by Galactic atomic and molecular gas is nearly absent, which is surprising given the column density of the gas."
 Due to the absence of the strong Galactic C II. Ar I ancl Fe II (1063.13 A) lines. we do not believe this to be a reliable spectrum.," Due to the absence of the strong Galactic C II, Ar I and Fe II (1063.18 ) lines, we do not believe this to be a reliable spectrum."
 This object may be reobserved αἱ a future date but Chis data only provides a poorly constrained upper limit., This object may be reobserved at a future date but this data only provides a poorly constrained upper limit.
 NGC 3607: The Galactic HE column and extinction are low toward this svstem x 107 7) and there are no strong Is absorption lines that can be identified., NGC 3607: The Galactic HI column and extinction are low toward this system $\times$ $^{20}$ $^{-2}$ ) and there are no strong $_{2}$ absorption lines that can be identified.
 The only clear Galactic atomic absorptlion line is (he CL) A1036 feature. aud the strong OVI line would be redshifted to 1035.2A. where there appear to be an emission feature (Figure 10).," The only clear Galactic atomic absorption line is the CII $\lambda$ 1036 feature, and the strong OVI line would be redshifted to 1035.2, where there appear to be an emission feature (Figure 10)."
 This feature also occurs in the lower S/N Lif2b channel. which does not add much to the S/N of (his result. but. provides some consistency.," This feature also occurs in the lower S/N Lif2b channel, which does not add much to the S/N of this result, but provides some consistency."
 However. (he emission from the weaker OVI line is nol present. although since (his line is only half the strength of the other OVI line. its absence does not lead to an inconsistency wilh the presence of the other line.," However, the emission from the weaker OVI line is not present, although since this line is only half the strength of the other OVI line, it's absence does not lead to an inconsistency with the presence of the other line."
" The center of the OVI line is al 1035.3A. a redshift of 979 km ! (and à EWIIM of 0.523:0.1.A. or 150 kin "")1). which is similar to the redshift of the galaxy. 935 km |."," The center of the OVI line is at 1035.3, a redshift of 979 km $^{-1}$ (and a FWHM of $\pm$ 0.1, or 150 km $^{-1}$ ), which is similar to the redshift of the galaxy, 935 km $^{-1}$."
 In addition to the weak OVI line. the CIII A977 line is also detected at à somewhat higher level of significance and at exactly the redshift of the galaxy.," In addition to the weak OVI line, the CIII $\lambda$ 977 line is also detected at a somewhat higher level of significance and at exactly the redshift of the galaxy."
 It has a width of 0.722:0.15 A. which is consistent with the width of the OVI line.," It has a width of $\pm$ 0.15 , which is consistent with the width of the OVI line."
 A line width of 0.6, A line width of 0.6
loose constraints on the source spectrum (see below).,loose constraints on the source spectrum (see below).
" Alerlonietal.(2003) ancl Falckeetal.(2004) have quantified an empirical relation (or ""Dundamental plane”) among X-ray and 5 GlIIz radio Iuminositv and DIL mass: we use the Merlonietal.(2003). relation LyxLSA. Mer", \citet{mer03} and \citet{fal04} have quantified an empirical relation (or “fundamental plane”) among X-ray and 5 GHz radio luminosity and BH mass; we use the \citet{mer03} relation $L_R\propto L_{\rm X}^{0.6}M_{\rm BH}^{0.78}$.
lonietal.(2003) analyzed this relation in (he context of accretion flows and jets associated with massive DIIs., \citet{mer03} analyzed this relation in the context of accretion flows and jets associated with massive BHs.
 One might expect some general relationship among these three quantities. if an. N-rav-enitting accretion flow onto a massive DII leads to creation of a svnchrotron-emitting radio jet. with the detailed correlation providing some insight into the nature of that flow.," One might expect some general relationship among these three quantities, if an X-ray-emitting accretion flow onto a massive BH leads to creation of a synchrotron-emitting radio jet, with the detailed correlation providing some insight into the nature of that flow."
 By comparing the empirically determined relation wil expectations from theoretical models. Merlonietal.(2003). deduced that the data for BIIs emitting at only a lew percent of (he Eddineton rate are consistent with radiativelv inefficient. accretion flows and a svnchrotron jet. but inconsistent with standard disk accretion models.," By comparing the empirically determined relation with expectations from theoretical models, \citet{mer03} deduced that the data for BHs emitting at only a few percent of the Eddington rate are consistent with radiatively inefficient accretion flows and a synchrotron jet, but inconsistent with standard disk accretion models."
 Maccarone(2004) scaled the fundaimental-plaue relation (to values appropriate for an IMDIIELin a Galactic globular cluster: we rescale their equation here to find a predicted radio flux density of Using the previously cited. X-ray luminosity and DIL mass for Gl and our adopted distance modulus. this predicts a 5 GIIz flix density of 77 jiJx for GI.," \citet{mac04} scaled the fundamental-plane relation to values appropriate for an IMBH in a Galactic globular cluster; we rescale their equation here to find a predicted radio flux density of Using the previously cited X-ray luminosity and BH mass for G1 and our adopted distance modulus, this predicts a 5 GHz flux density of 77 $\mu$ Jy for G1."
 However. taking into account the uncertainty in the IMDII mass. the unknown spectral index of the radio enission. and (he dispersion of 0.55 in log£j (Merlonietal.2003).. the predicted 3.4 GllIz flux densitv for GI is in the range of tens to a few hundred microjansky.," However, taking into account the uncertainty in the IMBH mass, the unknown spectral index of the radio emission, and the dispersion of 0.88 in $\log L_R$ \citep{mer03}, the predicted 8.4 GHz flux density for G1 is in the range of tens to a few hundred microjansky."
 Thus. our radio detection of 28 yJyv at 8.4 GIHz is consistent with the predictions for a 1.3x10!AZ. ΙΔΗΣ. but strongly inconsistent with a 10AZ. BIL," Thus, our radio detection of 28 $\mu$ Jy at 8.4 GHz is consistent with the predictions for a $1.8\times 10^4~M_\odot$ IMBH, but strongly inconsistent with a $10~M_\odot$ BH."
 Since neutron star X-ray binaries in a variety of states have radio/X-rav ratios much lower (han BIL X-ray binaries 2006).. and thus another 2 orders of magnitude below the observed value. stellaa-1nass X-ray binaries of any (vpe are ruled out as the possible origin of the radio emission in Gl.," Since neutron star X-ray binaries in a variety of states have radio/X-ray ratios much lower than BH X-ray binaries \citep{mig06}, and thus another 2 orders of magnitude below the observed value, stellar-mass X-ray binaries of any type are ruled out as the possible origin of the radio emission in G1."
 We can use (he radio/X-rav ratio to assess other possible origins for the radio emission., We can use the radio/X-ray ratio to assess other possible origins for the radio emission.
" lere. we use the ratio Ry=rpL,(8.4Gllz)/L4(2—10keV) as a fiducial marker."," Here, we use the ratio $R_{\rm X} = \nu L_\nu({\rm 8.4\ GHz})/L_{\rm X}({\rm 2-10\ keV)}$ as a fiducial marker."
 For GI. Rx5x107. which is considerably lower than Asz107 that is common to the Galactic supernova remnant Cas A. low-Iuminositv active galactic nuclei (supposing G1 might be a," For G1, $R_{\rm X}\approx 5\times 10^{-5}$, which is considerably lower than $R_{\rm X} \approx 10^{-2}$ that is common to the Galactic supernova remnant Cas A, low-luminosity active galactic nuclei (supposing G1 might be a"
"Several studies have proposed that dust forms in Eta Car during periastron (Falceta-Gongalvesetal.2005;Kashi&Soker2008a),, although Smith(2010) showed that the dust formation is cycle-dependent, occurring preferentially in the earlier documented spectroscopic events of 1981.4 and 1992.5.","Several studies have proposed that dust forms in Eta Car during periastron \citep{diego05,kashi08b}, , although \citet{smith10} showed that the dust formation is cycle-dependent, occurring preferentially in the earlier documented spectroscopic events of 1981.4 and 1992.5."
" Significant dust formation is uncertain during the 2003.5 event Smith(2010),, and no near-infrared photometry has been reported for the 2009.0 event."," Significant dust formation is uncertain during the 2003.5 event \citet{smith10}, and no near-infrared photometry has been reported for the 2009.0 event."
 The J-band flux increased by ca., The $J$ -band flux increased by ca.
" just before the 2003.5 event (Whitelocketal.2004),, which has been interpreted as due to free-free (Whitelocketal.2004) or hot dust emission (Kashi&Soker2008a)."," just before the 2003.5 event \citep{whitelock04}, which has been interpreted as due to free-free \citep{whitelock04} or hot dust emission \citep{kashi08b}."
". More importantly, interferometric observations in the K-band during the 2009.0 spectroscopic event, obtained simultaneously to our VLT/CRIRES measurements, do not show a significant change in the size of the K-band emitting region (Weigelt et al."," More importantly, interferometric observations in the K-band during the 2009.0 spectroscopic event, obtained simultaneously to our VLT/CRIRES measurements, do not show a significant change in the size of the K-band emitting region (Weigelt et al."
" 2010, in preparation), arguing against significant emission from hot dust in the inner 70 milli-arcseconds of Eta Car."," 2010, in preparation), arguing against significant emission from hot dust in the inner 70 milli-arcseconds of Eta Car."
" Therefore, a photospheric radius of Eta Car A at 1.08 um of 2.2 AU is hereafter assumed as the size of the continuum emission, based on the direct interferometric measurements in the K-band (vanBoekeletal.2003;Weigelt 2007)) scaled to 1.08 jum and on the value that we computed using the CMFGEN radiative transfer model of Eta Car A (Hillieret 2001)."," Therefore, a photospheric radius of Eta Car A at 1.08 $\mu$ m of 2.2 AU is hereafter assumed as the size of the continuum emission, based on the direct interferometric measurements in the $K$ -band \citealt{vb03,weigelt07}) ) scaled to 1.08 $\mu$ m and on the value that we computed using the CMFGEN radiative transfer model of Eta Car A \citep{hillier01}."
". If a binary companion is evoked, the periodicity might be explained as due to brief ejections of high-velocity material by Eta Car A triggered during each periastron passage."," If a binary companion is evoked, the periodicity might be explained as due to brief ejections of high-velocity material by Eta Car A triggered during each periastron passage."
" However, this scenario presents several difficulties, given that previous spectroscopic observations suggested that the wind of Eta Car A becomes roughly spherical during periastron (Smithetal. 2003a)."," However, this scenario presents several difficulties, given that previous spectroscopic observations suggested that the wind of Eta Car A becomes roughly spherical during periastron \citep{smith03}."
". It would also imply that, during periastron, material from Eta Car A at ~2000kms""! (instead of the usual 500-600 kms~!)) collides with the shock front."," It would also imply that, during periastron, material from Eta Car A at $\sim2000~\kms$ (instead of the usual 500–600 ) collides with the shock front."
 This increased velocity from Eta Car A would produce a much higher X-ray luminositythan what is currently observed., This increased velocity from Eta Car A would produce a much higher X-ray luminositythan what is currently observed.
" Both issues could be circumvented if the density and volume-filling factor of the ~2000kms""! transient wind are sufficiently low so as not to affect the X-ray hardness luminosity and the Ha absorption profiles measured by Smithetal.(2003a).", Both issues could be circumvented if the density and volume-filling factor of the $\sim2000~\kms$ transient wind are sufficiently low so as not to affect the X-ray hardness luminosity and the $\alpha$ absorption profiles measured by \citet{smith03}.
". However, it is unlikely that such a thin wind would produce detectable absorption in 410833."," However, it is unlikely that such a thin wind would produce detectable absorption in $\lambda$ 10833."
" The existence of a brief high-velocity wind from Eta Car A would be very unlikely in a single star scenario, although we cannot rule out that possibility based on our present data."," The existence of a brief high-velocity wind from Eta Car A would be very unlikely in a single star scenario, although we cannot rule out that possibility based on our present data."
" In particular, asingle-star scenario would have to invoke a yet unknown mechanism that would produce a periodic episode of high-velocity wind like clockwork every 2022.7+1.3 days, as measured by Daminelietal.(2008b)."," In particular, asingle-star scenario would have to invoke a yet unknown mechanism that would produce a periodic episode of high-velocity wind like clockwork every $2022.7 \pm 1.3$ days, as measured by \citet{damineli08_period}."
". The edge velocity of the high-velocity absorption component seen in 410833 and 441394, 1403 appears to approach the velocity expected of the wind of Eta Car B, 3000kms!, based upon X-ray spectroscopic modeling by Pittard&Corcoran(2002)."," The edge velocity of the high-velocity absorption component seen in $\lambda$ 10833 and $\lambda$$\lambda$ 1394, 1403 appears to approach the velocity expected of the wind of Eta Car B, $3000~\kms$, based upon X-ray spectroscopic modeling by \citet{pc02}."
" To date, Eta Car B has not been observed directly."," To date, Eta Car B has not been observed directly."
 Could the high-velocity absorption component form directly in the wind of Eta Car B?, Could the high-velocity absorption component form directly in the wind of Eta Car B?
" In the next two subsections, we investigate that possibility."," In the next two subsections, we investigate that possibility."
" Such a hypothesis would correspond to the classical detection of a companion in normal massive binary systems, such as in WR+OB binaries."," Such a hypothesis would correspond to the classical detection of a companion in normal massive binary systems, such as in WR+OB binaries."
" However, the flux of Eta Car B is several orders of magnitude lower than that of Eta Car A in the near-infrared continuum around theHer 410833 2006)."," However, the flux of Eta Car B is several orders of magnitude lower than that of Eta Car A in the near-infrared continuum around the $\lambda$ 10833 \citep{hillier06}."
". Therefore, even if the wind of Eta Car B could produce a saturatedΗει 410833 absorption profile when observed in isolation, an undetectable amount of absorption (0.5—1%) would be seen in the combined spectrum of Eta Car A and B. One possible way to observe the wind of Eta Car B, should it contain significant amounts of neutral He, would be if its 410833 absorption zone is extended and dense enough to absorb continuum radiation from Eta Car A, in a “wind-eclipse” scenario."," Therefore, even if the wind of Eta Car B could produce a saturated $\lambda$ 10833 absorption profile when observed in isolation, an undetectable amount of absorption $\sim0.5-1\%$ ) would be seen in the combined spectrum of Eta Car A and B. One possible way to observe the wind of Eta Car B, should it contain significant amounts of neutral He, would be if its $\lambda$ 10833 absorption zone is extended and dense enough to absorb continuum radiation from Eta Car A, in a ``wind-eclipse'' scenario."
" In order to detect the wind of Eta Car B only during a brief period, at phases 0.976<91.023 (i.e., dd), the binary system would again have to be oriented in the sky with a longitude of periastron of w~90°, since Eta Car B would need to be located between the observer and the continuum source for only a brief period around periastron passage."," In order to detect the wind of Eta Car B only during a brief period, at phases $0.976 \leq \phi \leq 1.023$ (i.e., d), the binary system would again have to be oriented in the sky with a longitude of periastron of $\omega \sim 90\degr$, since Eta Car B would need to be located between the observer and the continuum source for only a brief period around periastron passage."
" In addition, the material in the wind of Eta Car B must have a sufficiently high column density of neutral He to absorb enoughHer 210833 photons, but this is not predicted by the Hillieretal.(2006) radiative transfer model either."," In addition, the material in the wind of Eta Car B must have a sufficiently high column density of neutral He to absorb enough $\lambda$ 10833 photons, but this is not predicted by the \citet{hillier06} radiative transfer model either."
 A much larger M and/or lower Te would again be needed to produce enough optical depth in the wind of Eta Car B for 210833., A much larger $\mdot$ and/or lower $\teff$ would again be needed to produce enough optical depth in the wind of Eta Car B for $\lambda$ 10833.
" In principle, this would argue for the presence of a Wolf-Rayet (WR) instead of an O-type star companion, since WRs have a higher wind density than O-type stars."," In principle, this would argue for the presence of a Wolf-Rayet (WR) instead of an O-type star companion, since WRs have a higher wind density than O-type stars."
" Even in the unlikely possibility thatthe wind of Eta Car B could significantly absorb 110833 photons, the “wind-eclipse” scenario also fails to reproduce the observed duration of the high-velocity absorption component for the assumed orbital parameters, even allowing for significant uncertainties in these parameters."," Even in the unlikely possibility thatthe wind of Eta Car B could significantly absorb $\lambda$ 10833 photons, the ``wind-eclipse'' scenario also fails to reproduce the observed duration of the high-velocity absorption component for the assumed orbital parameters, even allowing for significant uncertainties in these parameters."
" We show in Figure 10aa a pole-on view of the geometry of the orbit for the “wind-eclipse” scenario, assuming masses of 90 and 30Μο for Eta Car A and B, respectively, orbital period of P=2022.7 d, semi-major axis of a=15.4 AU, eccentricity of e= 0.9, and w= 90°."," We show in Figure \ref{wind_eclipse}a a a pole-on view of the geometry of the orbit for the “wind-eclipse” scenario, assuming masses of 90 and $30~\msun$ for Eta Car A and B, respectively, orbital period of $P=2022.7$ d, semi-major axis of $a=15.4$ AU, eccentricity of $e=0.9$ , and $\omega=90\degr$ ."
" A photospheric radius at 1.08 wm of 2.2 AU is assumed for Eta Car A, as discussed in Section 5.1.."," A photospheric radius at 1.08 $\mu$ m of 2.2 AU is assumed for Eta Car A, as discussed in Section \ref{cont}. ."
" In thissubsection, we assume an inclination angle of i=90° to derive an upper limit for the timescale of the “wind-eclipse”."," In thissubsection, we assume an inclination angle of $i=90\degr$ to derive an upper limit for the timescale of the “wind-eclipse”."
 For (more realistic), For (more realistic)
The problem of Monte-Carlo radiation transfer in. very optically thick regions - such as in. the midplane of circumstellar disks - 1$ challenging.,The problem of Monte-Carlo radiation transfer in very optically thick regions – such as in the midplane of circumstellar disks – is challenging.
 Without any approximations. photon packets can get trapped for millions of interactions. increasing the required computational time by several orders of magnitude.," Without any approximations, photon packets can get trapped for millions of interactions, increasing the required computational time by several orders of magnitude."
 Minetal.(2009.hereafterMO9) presented two complementary methods to greatly improve the efficiency of Monte-Carlo radiation transfer codes in very optically thick regions: a modified random walk (MRW) and a partial diffusion approximation (PDA)., \citet[][hereafter M09]{Min:09:155} presented two complementary methods to greatly improve the efficiency of Monte-Carlo radiation transfer codes in very optically thick regions: a modified random walk (MRW) and a partial diffusion approximation (PDA).
 The MRW prevents photons from getting stuck in very optically thick regions. and the PDA allows temperatures to be calculated in regions that see few or no photons.," The MRW prevents photons from getting stuck in very optically thick regions, and the PDA allows temperatures to be calculated in regions that see few or no photons."
 The essence of the MRW method is that instead of computing thousands to millions of individual absorption or scattering events for a single photon in these optically thick regions. one can make use of the solution to the diffusion approximation inside small regions to propagate the photon efficiently.," The essence of the MRW method is that instead of computing thousands to millions of individual absorption or scattering events for a single photon in these optically thick regions, one can make use of the solution to the diffusion approximation inside small regions to propagate the photon efficiently."
 Monte-Carlo radiation transfer codes propagate photons in a grid made up of cells of constant density and temperature., Monte-Carlo radiation transfer codes propagate photons in a grid made up of cells of constant density and temperature.
" Therefore if the mean optical depth to the edge of a cell is much larger than unity. one can set up a sphere whose radius is smaller than the distance to the closest wall. inside which the density will be constant. and travel to the edge of a sphere in a single step using the diffusion approximation,"," Therefore if the mean optical depth to the edge of a cell is much larger than unity, one can set up a sphere whose radius is smaller than the distance to the closest wall, inside which the density will be constant, and travel to the edge of a sphere in a single step using the diffusion approximation."
 A probability distribution function is used to sample the true distance traveled to exit this sphere (since the photon would follow a random walk inside the sphere. rather than moving in a straight line).," A probability distribution function is used to sample the true distance traveled to exit this sphere (since the photon would follow a random walk inside the sphere, rather than moving in a straight line)."
 This true distance. which depends on the radius of the sphere and the local diffusion coefficient D. can then be used along with the mass absorption coefficient k to compute the total amount of energy deposited in the dust during the diffusion.," This true distance, which depends on the radius of the sphere and the local diffusion coefficient $D$, can then be used along with the mass absorption coefficient $\bar{\kappa}$ to compute the total amount of energy deposited in the dust during the diffusion."
 This is required in order to compute the temperature in the cell accurately., This is required in order to compute the temperature in the cell accurately.
 Finally. the photon is emitted from a random position on the surface of the sphere with a frequency sampled from the Planek function.," Finally, the photon is emitted from a random position on the surface of the sphere with a frequency sampled from the Planck function."
 M09 provide equations for D. &. and the dust emission coefficient 7. taking into account that photons can be both scattered and absorbed and re-emitted.," M09 provide equations for $D$, $\bar{\kappa}$ and the dust emission coefficient $\eta_\nu$, taking into account that photons can be both scattered and absorbed and re-emitted."
 In M09. the suggested algorithm is to first calculate ην iteratively. and to then use 7. to compute D and x.," In M09, the suggested algorithm is to first calculate $\eta_\nu$ iteratively, and to then use $\eta_\nu$ to compute $D$ and $\bar{\kappa}$."
" In this note. | show that 7j, does not need to be solved iteratively. but can be solved directly. and | use this solution to show that D and & can in fact very easily be computed. resulting in both a simpler implementation of the MRW. and in some cases performance gains."," In this note, I show that $\eta_\nu$ does not need to be solved iteratively, but can be solved directly, and I use this solution to show that $D$ and $\bar{\kappa}$ can in fact very easily be computed, resulting in both a simpler implementation of the MRW, and in some cases performance gains."
" In the presence of isotropic scattering. the emissivity of dustin local thermodynamie equilibrium (LTE) is given by where v is the frequency of the radiation. &, 1s the Nass absorption. coefficient. B,(7) is the Planck function at the temperature 7 of the dust. c, is the mass scattering coefficient. and J, is the mean intensity of the radiation field."," In the presence of isotropic scattering, the emissivity of dust in local thermodynamic equilibrium (LTE) is given by where $\nu$ is the frequency of the radiation, $\kappa_\nu$ is the mass absorption coefficient, $B_\nu(T)$ is the Planck function at the temperature $T$ of the dust, $\sigma_\nu$ is the mass scattering coefficient, and $J_\nu$ is the mean intensity of the radiation field."
" Assuming that the radiation ts isotropic. J,=7,. where 7, is the intensity of the radiation."," Assuming that the radiation is isotropic, $J_\nu = I_\nu$, where $I_\nu$ is the intensity of the radiation."
" In the optically thick regime. 7.=S,.. where ὃν Is the source function."," In the optically thick regime, $I_\nu=S_\nu$, where $S_\nu$ is the source function."
" Therefore. The source function is defined as the ratio of the total emissivity to the total extinction. which in this case ts where y, is the mass extinction coefficient (y,=a,+ c0)."," Therefore, The source function is defined as the ratio of the total emissivity to the total extinction, which in this case is where $\chi_\nu$ is the mass extinction coefficient $\chi_\nu = \kappa_\nu + \sigma_\nu$ )."
 Therefore. Equation (2)) can be rewritten as Re-arranging this equation. one obtains which can be simplified. since: Therefore.The source function for this emissivity Is which is expected for thermal emission from dust in LTE in the optically thick regime.," Therefore, Equation \ref{eq:initials}) ) can be rewritten as Re-arranging this equation, one obtains which can be simplified, since: Therefore,The source function for this emissivity is which is expected for thermal emission from dust in LTE in the optically thick regime."
2c,2cm
order not to overpredict the nuniber of quasars in the ΠΟΕ (IIainan. Madau Loeb 1999).,"order not to over–predict the number of quasars in the HDF (Haiman, Madau Loeb 1999)."
 The data points are from recent ROSAT measurements. and the dashed curve in this figure. shows a fitting formula from Alivaji ct al. (," The data points are from recent ROSAT measurements, and the dashed curve in this figure shows a fitting formula from Miyaji et al. ("
1998).,1998).
 Note that the faintest quasar actually observed has Ly<10ttevesft. and the fitting formmla below this huuinositv is liehly uncertain.," Note that the faintest quasar actually observed has $L_X\lsim 10^{44}~{\rm erg~s^{-1}}$, and the fitting formula below this luminosity is highly uncertain."
 Nevertheless. it is iuportaunt to remember that our model was calibrated so as to fit the observed. optical Iuuunositv. function of quasars: the existence of a population of obscured quasars which are faint iu the optical baud but bright im Xravs could increase the number counts bevoud our model predictious.," Nevertheless, it is important to remember that our model was calibrated so as to fit the observed optical luminosity function of quasars; the existence of a population of obscured quasars which are faint in the optical band but bright in X--rays could increase the number counts beyond our model predictions."
 Existing estimates of the Nrav background (XRB) provide another useful check on our quasar model., Existing estimates of the X–ray background (XRB) provide another useful check on our quasar model.
 The resolved backgrouud flux at a photon enerev £ is giveu by (Poebles 1993) τε 1G where E.=E(l|2): aud ντ). is the comoving cluissivity at a localphoton energv £.. in units of ΚονcmnPsbthkeWft. from quasars shining at a redshift +.," The unresolved background flux at a photon energy $E$ is given by (Peebles 1993) F(E)= c dz , where $E_z=E(1+z)$; and $j(E_z,z)$ is the comoving emissivity at a localphoton energy $E_z$, in units of ${\rm keV~cm^{-3}~s^{-1}~sr^{-1}~keV^{-1}}$, from quasars shining at a redshift $z$."
 This cuiissivity is a sum over all quasars whose individual observed fux at 2=0 is below the ROSAT PSPC detection limit for discrete sources of 10Pergcn7s+ (Hasueer Zamora 1997)., This emissivity is a sum over all quasars whose individual observed flux at $z=0$ is below the ROSAT PSPC detection limit for discrete sources of $2\times10^{-15}~{\rm erg~cm^{-2}~s^{-1}}$ (Hasinger Zamorani 1997).
 Fieve 3 shows the predicted spectrmm of the NRD in our model at +=0 (solid lines)., Figure \ref{fig:xrb} shows the predicted spectrum of the XRB in our model at $z=0$ (solid lines).
 Tn computing the backeround spectrum. we ignored the UT absorption iu the ICM. since it is negligible at energiesabove 100 eV. We also carried out the integral in equation (3)) only for 2o>2. the redshift range where our model is valid (IHaimiau Loch 1998).," In computing the background spectrum, we ignored the HI absorption in the IGM, since it is negligible at energiesabove 100 eV. We also carried out the integral in equation \ref{eq:xrb}) ) only for $z>2$, the redshift range where our model is valid (Haiman Loeb 1998)."
 The short dashed lines show the predicted fluxes assuming a steeper spectral slope bevoud 10 keV (a=0.5. or a photon iudex of -1.5).," The short dashed lines show the predicted fluxes assuming a steeper spectral slope beyond 10 keV $\alpha=-0.5$, or a photon index of -1.5)."
 The lone dashed line shows the unresolved fraction of the soft NRB observed with ROSAT (Mivaji ct al., The long dashed line shows the unresolved fraction of the soft XRB observed with ROSAT (Miyaji et al.
 1998b: Fabian Darcous 1992)., 1998b; Fabian Barcons 1992).
 This fraction represeuts the observational upper limit ou the component of the soft NRB that could iu principle arise from hieliredshift quasars., This fraction represents the observational upper limit on the component of the soft XRB that could in principle arise from high-redshift quasars.
 As the figure shows. our quasar model predicts an unresolved flix just below this lnüt in the 0.5-3 keV. range.," As the figure shows, our quasar model predicts an unresolved flux just below this limit in the 0.5-3 keV range."
 The model also predicts that most 290%) of this vet unresolved fraction arises from quasars bevoud 2=5., The model also predicts that most $\gsim 90\%$ ) of this yet unresolved fraction arises from quasars beyond $z=5$.
" By sununing the buninosity function over redshifts. we obtain the uuniber counts of quasars per solid anele expected to be detectable in a flux interval ον arouncl Py=lyLid;(D) for all sources above a redshift 1 λοςΕνxc) led}ls), where (JVdid) is the comoving volume clement per nuit redshift aud solid angle. and dp(2) is the Iuninositv distance at a redshift :."," By summing the luminosity function over redshifts, we obtain the number counts of quasars per solid angle expected to be detectable in a flux interval $dF_X$ around $F_X=L_X/4\pi d_L^2(z)$, for all sources above a redshift $z$: = dz ) (L_X,z) d_L^2(z), where $(dV/dzd\Omega)$ is the comoving volume element per unit redshift and solid angle, and $d_L(z)$ is the luminosity distance at a redshift$z$ ."
 In Figure L. we show the predicted counts frou equation (1)) iu the O.LGkeV. enerev baud of the CCD," In Figure \ref{fig:counts}, , we show the predicted counts from equation \ref{eq:counts}) ) in the 0.4–6keV energy band of the CCD"
longer.,longer.
 Let's then consider a 1017. GC pulsar with an initial period of 60 ms (the pulsar produced in SN 386 is such a source)., Let's then consider a $10^{12}$ G pulsar with an initial period of 60 ms (the pulsar produced in SN 386 is such a source).
 We then have Za~3107 eres. so that (ignoring expansion losses). the energy deposited in the PWN over 20 years would be ~2«107 ores.," We then have $\dot{E}_{\rm rot} \sim 3\times 10^{36}$ erg/s, so that (ignoring expansion losses), the energy deposited in the PWN over 20 years would be $\sim 2\times 10^{45}$ ergs."
 For a volume similar to that for SN 1986J. above. the equipartition magnetic field would be 10 mC. corresponding to a lifetime at 2 keV of about 10 days.," For a volume similar to that for SN 1986J above, the equipartition magnetic field would be $\sim10$ mG, corresponding to a lifetime at 2 keV of about 10 days."
 This is still a short enough lifetime for our purposes., This is still a short enough lifetime for our purposes.
 Alternatively. for 2)~5 ms and D~107 C. we have Lua~οLol! orefs. and over 20 vears. this vields Aaον+107 eres.," Alternatively, for $P_0\sim 5$ ms and $B \sim 10^{12}$ G, we have $\dot{E}_{\rm rot} \sim 6\times 10^{40}$ erg/s, and over 20 years, this yields $E_{\rm tot} \sim 4\times 10^{49}$ ergs."
 This implies D21 GC. so that laa15 minutes.," This implies $B
> 1$ G, so that $t_{\rm synch} \sim 15$ minutes."
" Therefore. we conclude that. overall. the magnetic fields in voung PWNe are likely strong enough to justify the use of the £L,fon relation even for the voungest objects in our sample."," Therefore, we conclude that, overall, the magnetic fields in young PWNe are likely strong enough to justify the use of the $L_x-\dot{E}_{\rm rot}$ relation even for the youngest objects in our sample."
 One further point to notice with respect to the νο£a correlation is the fact that it is based. on a cliversity of objects.," One further point to notice with respect to the $L_x-\dot{E}_{\rm
rot}$ correlation is the fact that it is based on a diversity of objects."
 Ehe low end range of the relation. in particular. is populated with Millisecond Pulsars (AISPs). which are spun up neutron stars.," The low end range of the relation, in particular, is populated with Millisecond Pulsars (MSPs), which are spun up neutron stars."
 It is possible that this class of objects might bias the correlation of the voungest. isolated: pulsars in the sample.," It is possible that this class of objects might bias the correlation of the youngest, isolated pulsars in the sample."
 Generally speaking. once they are spun up. he AISPs form PWNe again (c.g. Cheng et al.," Generally speaking, once they are spun up, the MSPs form PWNe again (e.g. Cheng et al."
" 2006). and he conversion of Ey. into L,. which is practically an instantaneous relationship (as compared to the ages under consideration). should not be dependent on the history of he svstem."," 2006), and the conversion of $\dot{E}_{\rm rot}$ into $L_x$, which is practically an instantaneous relationship (as compared to the ages under consideration), should not be dependent on the history of the system."
 The magnetic field of the objects (Lower for the AISPs than for the voung. isolated pulsars). however. might influence the conversion ellicieney. (Ciuseinov et al.," The magnetic field of the objects (lower for the MSPs than for the young, isolated pulsars), however, might influence the conversion efficiency (Guseinov et al."
 2004). renee biasing the overall slope of the correlation.," 2004), hence biasing the overall slope of the correlation."
 Overall. in our analysis a steeper slope would lead to tighter limits on he NS spin birth distribution. and viceversa for a shallower slope.," Overall, in our analysis a steeper slope would lead to tighter limits on the NS spin birth distribution, and viceversa for a shallower slope."
 What would be alfected the most by a slope change is he high [Esa tail of the population., What would be affected the most by a slope change is the high $\dot{E}_{\rm rot}$ tail of the population.
" Hence. in £44. besides deriving results using the £L,[Esa for all the pulsars. we will also examine the ellects of a change of slope for the astest pulsars."," Hence, in 4, besides deriving results using the $L_x-\dot{E}_{\rm rot}$ for all the pulsars, we will also examine the effects of a change of slope for the fastest pulsars."
 The rotational energv loss of the star. under the assumption that it is dominated by magnetic dipole losses. is given by where A2 is the NS radius. which we take to be 10 km. £2 the NS magnetic field. f)—2x/P the star angular velocity. and 6 the angle between the magnetic and. spin axes.," The rotational energy loss of the star, under the assumption that it is dominated by magnetic dipole losses, is given by where $R$ is the NS radius, which we take to be 10 km, $B$ the NS magnetic field, $\Omega=2\pi/P$ the star angular velocity, and $\theta$ the angle between the magnetic and spin axes."
 We take sin&=1 for consistency with what generally assumed in pulsar radio studies., We take $\sin\theta=1$ for consistency with what generally assumed in pulsar radio studies.
 With sin@=1I and a constant D field. the spin evolution of the pulsars is simply given by bPu-Jue. 15 2o. ∖∖⇁⇂⊔⋅↓⋅⋖⋅∫≈↓∪⋏∙≟≼∙⊔↓↓⊳∖∣↓↥≺⋅⊔↓∪⊔↓≺⊾⊔↥∪⇂↓⊔≺⊾↓⋅↿⋯∪⇂⋅⋠. the star. and £% is its initial spin period.," With $\sin\theta=1$ and a constant $B$ field, the spin evolution of the pulsars is simply given by P(t) =, where $I\approx 10^{45} {\rm g}\, {\rm cm}^2$ is the moment of inertia of the star, and $P_0$ is its initial spin period."
" The X-ray Iuminositv of the pulsar at. time / (whieh traces Eau) correspondingly declines as L,=νο|ffl)= where fy=BLOοBP(2x)6500vrReLBDuo. having defined 4,4=I1/(107 g en?). Rig=R/Okm). Poo=Lu/C10 ms)."," The X-ray luminosity of the pulsar at time $t$ (which traces $\dot{E}_{\rm
rot}$ ) correspondingly declines as $ L_x=L_{x,0}(1+t/t_0)^{-2}$, where $t_0\equiv 3Ic^3P_0^2/B^2R^6(2\pi)^2 \sim 6500\, {\rm
yr}\;I_{45}B^{-2}_{12}R_{10}^{-6}P^2_{0,10}$, having defined $I_{45}\equiv I/(10^{45}$ g $^2$ ), $R_{10}\equiv R/(10\; {\rm km})$, $P_{0,10}\equiv P_0/(10\;{\rm ms})$ ."
 For Εστω the Dux does not vary significantly., For $t\la t_0$ the flux does not vary significantly.
 Since the ages fsx of the SNe in our sample are all vr. we deduce that. for typical pulsar fields. fsxlu.," Since the ages $t_{\rm SN}$ of the SNe in our sample are all $\la 77$ yr, we deduce that, for typical pulsar fields, $t_{\rm SN}\ll
t_0$."
 The luminosities of the pulsars associated: with the SNe in our sample are therefore expected to be still in the plateau region. and thus they directly probe the initial birth parameters. before evolution alfects the periods appreciably.," The luminosities of the pulsars associated with the SNe in our sample are therefore expected to be still in the plateau region, and thus they directly probe the initial birth parameters, before evolution affects the periods appreciably."
 In order to compute the X-ray luminosity distribution of a population of voung pulsars. the magnetic fields and the initial periods of the pulsars need to be known.," In order to compute the X-ray luminosity distribution of a population of young pulsars, the magnetic fields and the initial periods of the pulsars need to be known."
 As discussed in $11. a number of investigations have been mace over the last few decades in order to infer the birth parameters of Ns. and in particular the distribution of initial periods and magnetic fields.," As discussed in 1, a number of investigations have been made over the last few decades in order to infer the birth parameters of NSs, and in particular the distribution of initial periods and magnetic fields."
 Llere. we begin our study by comparing the SN data with the results of a pulsar population calculation that assumes one of such distributions. and specifically one that makes predictions for birth periods in the millisecond range.," Here, we begin our study by comparing the SN data with the results of a pulsar population calculation that assumes one of such distributions, and specifically one that makes predictions for birth periods in the millisecond range."
 After establishing that the SN. cata are. highly inconsistent with such short initial spins. we then generalize our analysis by inverting the problem. and. performing a parametric study. aimed at finding the minimum values of the birth periods that result in predicted X-ray bIuminosities consistent with the SN X-ray data.," After establishing that the SN data are highly inconsistent with such short initial spins, we then generalize our analysis by inverting the problem, and performing a parametric study aimed at finding the minimum values of the birth periods that result in predicted X-ray luminosities consistent with the SN X-ray data."
 As a starting point to constrain pulsar birth parameters. we consider the results of one of the most. recent and comprehensive radio studies. based on large-scale raclio pulsar surveys. that is the one carried out by Arzoumanian et al. (," As a starting point to constrain pulsar birth parameters, we consider the results of one of the most recent and comprehensive radio studies, based on large-scale radio pulsar surveys, that is the one carried out by Arzoumanian et al. ("
2002; ACC in the following).,2002; ACC in the following).
" They find that. if spin clown is dominated. by dipole radiation losses (ice. braking index equal to 3). and the magnetic field. does not appreciably decay. the magnetic field strength. (taken as Gaussian in log) has a mean dogBuC]=12.35 and a standard. deviation of 0.4. while the initial birth period distribution. (also taken as a log-CGaussian). is found. to have à mean flog{κο=2.3 with a standard deviation ap,I0.2 (within the searched range of 0.10.7)."," They find that, if spin down is dominated by dipole radiation losses (i.e. braking index equal to 3), and the magnetic field does not appreciably decay, the magnetic field strength (taken as Gaussian in log) has a mean $\langle
\log B_0[G] \rangle =12.35$ and a standard deviation of $0.4$, while the initial birth period distribution (also taken as a log-Gaussian), is found to have a mean $\langle \log P_0(s) \rangle =-2.3$ with a standard deviation $\sigma_{P_0}> 0.2$ (within the searched range of $0.1-0.7$ )."
 In the first part of our. paper. as a specific example. of a distribution that predicts a large fraction of pulsars to be born with millisecond periods. we use their. inferred parameters described above.," In the first part of our paper, as a specific example of a distribution that predicts a large fraction of pulsars to be born with millisecond periods, we use their inferred parameters described above."
" Since in their model the standard. deviation for the initial period distribution is constrained only by the lower limit ap,20.2. here we adopt ap,—0.3."," Since in their model the standard deviation for the initial period distribution is constrained only by the lower limit $\sigma_{P_0}>0.2$, here we adopt $\sigma_{P_0}=0.3$."
 As the width of the velocity dispersion increases. the predicted. X-ray. luminosity distribution becomes more and more heavily weighed towards higher Iuminosities (sec Figure 1 in Perna Stella 2004).," As the width of the velocity dispersion increases, the predicted X-ray luminosity distribution becomes more and more heavily weighed towards higher luminosities (see Figure 1 in Perna Stella 2004)."
" Therefore. the limits that we derive in this section would be even stronger if ap, were larger than what we assume."," Therefore, the limits that we derive in this section would be even stronger if $\sigma_{P_0}$ were larger than what we assume."
 We then assume that the Lx[Esa correlation is described bv the PO2 best fit with the corresponding scatter., We then assume that the $L_X-\dot{E}_{\rm rot}$ correlation is described by the P02 best fit with the corresponding scatter.
" In order to test the resulting theoretical predictions for the pulsar distribution of X-ray Iuminosities against the limits of the SNe. we perform 10"" Monte Carlo realizations of the compact object remnant. population."," In order to test the resulting theoretical predictions for the pulsar distribution of X-ray luminosities against the limits of the SNe, we perform $10^6$ Monte Carlo realizations of the compact object remnant population."
 Each realization is made up of Nowy=Nex100. with ages equal to the ages of the SNe in our sample.," Each realization is made up of $N_{\rm obj}=N_{\rm
SN}=100$, with ages equal to the ages of the SNe in our sample."
 Phe fraction of massive stars that leave behind a black hole (BLT) has been theoretically estimated in the simulations by Heger et al. (, The fraction of massive stars that leave behind a black hole (BH) has been theoretically estimated in the simulations by Heger et al. (
2003).,2003).
 For a solar metallicity and a Salpeter IAL. they find that this fraction is about of the total.," For a solar metallicity and a Salpeter IMF, they find that this fraction is about of the total."
 However we need to, However we need to
 , 
"56Ni) heating, adiabatic expansion continues until the remnant is as old as the effective diffusion time tg~ where is the opacity, after which the entropy(&Mej/v,c)!/?, is lost.","$^{56}$ Ni) heating, adiabatic expansion continues until the remnant is as old as the effective diffusion time $t_d\sim (\kappa \Mej/v_t c)^{1/2}$ , where $\kappa$ is the opacity, after which the entropy is lost."
 Such thermally powered light-curves (e.g. Type IIp's) have a luminosity LincEsnte/t?., Such thermally powered light-curves (e.g. Type IIp's) have a luminosity $L_{\rm th}\sim \Esn \te/ \td^2$.
 The large amount of adiabatic expansion that has occurred by the time t~tg leads to low luminosities., The large amount of adiabatic expansion that has occurred by the time $t\sim t_d$ leads to low luminosities.
" Now consider the impact of late time (t>>t.) energy injection from a young magnetar with radius Rp,=10km and initial spin Q;=2-/P;.", Now consider the impact of late time $t \gg \te$ ) energy injection from a young magnetar with radius $\Rns=10 \ {\rm km}$ and initial spin $\Omega_i=2\pi/P_i$.
 The magnetar rotational energy is where Pio=P;/10 ms and we set the NS moment of inertia to be [ης=1045gcm?., The magnetar rotational energy is where $\Pt = P_i/10$ ms and we set the NS moment of inertia to be $I_{\rm ns} = 10^{45}\ {\rm g \ cm^{2}} $.
" This magnetar loses rotational energy at the rate set by magnetic dipole radiation (with the angle, o, between rotation and magnetic dipole fixed at sin?a= 1/2), injecting most of the energy into the expanding remnant on the down timescale where By,=G."," This magnetar loses rotational energy at the rate set by magnetic dipole radiation (with the angle, $\alpha$, between rotation and magnetic dipole fixed at $\sin^2\alpha=1/2$ ), injecting most of the energy into the expanding remnant on the spin-down timescale where $\Bf=B/10^{14} \ {\rm G}$."
" To input this energy at a time tpS;ta requires a Β/1014minimum P field of where Κος=&/0.2cm?g~!, Ms=Μαι/5δMo and Εσι=Egn/10°! ergs~+."," To input this energy at a time $\tp \la \td$ requires a minimum $B$ field of where $\kapes = \kappa/0.2 \ {\rm cm^2 \ g^{-1}}$, $M_5 = \Mej/5~\Msun$ and $E_{51} = \Esn/10^{51}~{\rm ergs}^{-1}$ ."
 The required fields are in the magnetar range., The required fields are in the magnetar range.
 This late time entropy injection resets the interior energy scale to Ejc and overwhelms the initial thermal energy when EyEp>Esn(te/tp)., This late time entropy injection resets the interior energy scale to $\Eint \sim \Ep$ and overwhelms the initial thermal energy when $\Ep>\Esn (t_e/t_p)$.
" Thus even low magnetar energies E,<Esp play an important role.", Thus even low magnetar energies $E_p<\Esn $ play an important role.
 The resulting peak luminosity is which is primarily a function of the magnetic field., The resulting peak luminosity is which is primarily a function of the magnetic field.
 This shows that Lypeak~1099—1095ergs~! SNe can be achieved from magnetars with [514=1—10 and initial spins in the P;—220 ms range., This shows that $L_{\rm peak}\sim 10^{43}-10^{45} \ {\rm erg \ s^{-1}}$ SNe can be achieved from magnetars with $\Bf=1-10$ and initial spins in the $P_i= 2-20$ ms range.
 Our simple estimate ignores the details of how the deposited energy is distributed throughout the interior of the expanding SNe remnant., Our simple estimate ignores the details of how the deposited energy is distributed throughout the interior of the expanding SNe remnant.
" Since the dissipation mechanism for the pulsar wind in this medium is poorly understood, we assume the injected magnetar energy is thermalized spherically at the base of the supernova ejecta."," Since the dissipation mechanism for the pulsar wind in this medium is poorly understood, we assume the injected magnetar energy is thermalized spherically at the base of the supernova ejecta."
" The remnant is assumed to be in homologous expansion with a shallow power law density structure in the interior where vp=(2Egn/? is the characteristic ejecta velocity, and the density Ma)falls off sharply abovet."," The remnant is assumed to be in homologous expansion with a shallow power law density structure in the interior where $\vt = (2 \Esn/ \Mej)^{1/2}$ is the characteristic ejecta velocity, and the density falls off sharply above."
". The central overpressure caused by the energy deposition blows a bubble in the SN remnant, similar to the dynamics studied in the context of pulsar wind nebulae Chevalier1977;&Frans-son As(e."," The central overpressure caused by the energy deposition blows a bubble in the SN remnant, similar to the dynamics studied in the context of pulsar wind nebulae \citep[e.g.,][]{Chevalier_77, Chevalier_92}. ."
"g., this bubble expands, it sweeps up ejecta 1992)..into a thin shell near the leading shock, leaving the hot, low density interior evident in the 1-D radiation hydrodynamical calculations of Figure 1.."," As this bubble expands, it sweeps up ejecta into a thin shell near the leading shock, leaving the hot, low density interior evident in the 1-D radiation hydrodynamical calculations of Figure \ref{Fig:ejecta}."
" In multi-dimensional calculations of pulsar wind nebulae, Rayleigh-Taylor instabilities broaden the shell and mix the swept-up material (Jun1998;Blondinetal.2001)."," In multi-dimensional calculations of pulsar wind nebulae, Rayleigh-Taylor instabilities broaden the shell and mix the swept-up material \citep{Jun_98, Blondin_01}."
. The bubble expansion will freeze out in Lagrangian coordinates when the leading shock velocity becomes comparable to the local velocity of theexpanding SN ejecta., The bubble expansion will freeze out in Lagrangian coordinates when the leading shock velocity becomes comparable to the local velocity of theexpanding SN ejecta.
" The postshock pressure is P=2ypov2/(1+7)(8/7)pov2 for a strong shock, and the pressure of the energized cavity is PzzE,/3V, where V is the volume, implying a shock velocity v?=pg."," The postshock pressure is $P = 2\gamma \rho_0 \vs^2 /(1+\gamma) = (8/7) \rho_0 \vs^2$ for a strong shock, and the pressure of the energized cavity is $P \approx \Ep/3 V$, where $V$ is the volume, implying a shock velocity $\vs^2 = 7 \Ep/ 32 \pi R^3 \rho_0$."
" The shock becomes weak when vszzR/t, which7Ej/32z determinesI? the final velocity coordinate of the dense shell The weak dependence onEy, vayοςEy!. for ὃ=1, places nnearΌε."," The shock becomes weak when $\vs \approx R/t$, which determines the final velocity coordinate of the dense shell The weak dependence on, $\vsh \propto \Ep^{1/4}$, for $\delta = 1$, places near."
. The total mass swept up in the shell is Msn=Moj(%/vsn3°., The total mass swept up in the shell is $\Msh = \Mej (\vt/\vsh)^{3-\delta}$.
" The magnetar does not affect the dynamics of the outer layers of the SN ejecta unless Ej2Esn, in which case the bubble expands beyond aand accelerates more rapidly down the steep outer density gradient."," The magnetar does not affect the dynamics of the outer layers of the SN ejecta unless $\Ep \ga \Esn$, in which case the bubble expands beyond and accelerates more rapidly down the steep outer density gradient."
 Essentially all of the ejecta is then swept up into the shell at a final shell velocity Both estimates for aassume no radiative losses., Essentially all of the ejecta is then swept up into the shell at a final shell velocity Both estimates for assume no radiative losses.
 The presence of a dense shell has consequences for the supernova spectra., The presence of a dense shell has consequences for the supernova spectra.
" Initially the photospheric velocity, Upn;, as measured from the Dopplershift of absorption line minima, decreases with time as the outer layers of"," Initially the photospheric velocity, , as measured from the Dopplershift of absorption line minima, decreases with time as the outer layers of"
wake.,wake.
 Detailed observations of the G-fragment impact were synthesised with their numerical simulations of the fireball (e.g. their Figure 7)., Detailed observations of the G-fragment impact were synthesised with their numerical simulations of the fireball (e.g. their Figure 7).
 Chyba et al. (, Chyba et al. (
1993). Zahnle and MacLow (1994) and MacLow and Zahnle (1994) showed that the arburst phase is dominated by ablation only in the upper atmosphere. and below that by catastrophic pressure-driven disruption.,"1993), Zahnle and MacLow (1994) and MacLow and Zahnle (1994) showed that the airburst phase is dominated by ablation only in the upper atmosphere, and below that by catastrophic pressure-driven disruption."
" In the Jupiter case. the latter is true for all impacting masses except those very much smaller than that (=10"" ο) of the Shoemaker-Levy 9 fragments in 1994."," In the Jupiter case, the latter is true for all impacting masses except those very much smaller than that $\approx 10^{15}$ g) of the Shoemaker-Levy 9 fragments in 1994."
 We summarize here the essence of the arguments leading to these conclusions so that we can adapt them to the somewhat different conditions of solar impacts., We summarize here the essence of the arguments leading to these conclusions so that we can adapt them to the somewhat different conditions of solar impacts.
 The relative importance of ablation. explosion and deceleration in the initial destruction of the nucleus depends on atmospheric scale height H as well as on the nucleus mass. speed and entry angle.," The relative importance of ablation, explosion and deceleration in the initial destruction of the nucleus depends on atmospheric scale height $H$ as well as on the nucleus mass, speed and entry angle."
 They can be expressed in terms of the atmospheric column densities N(cm?)=1nds traversed s of the nucleus after which each processes is complete. (, They can be expressed in terms of the atmospheric column densities $\cal N$ $^{-2})=\int_s nds$ traversed $s$ of the nucleus after which each processes is complete. (
The vertical column density N=cosó for constant path angle to the vertical).,The vertical column density $N={\cal N}\cos \phi$ for constant path angle $\phi$ to the vertical).
" First. to deliver enough energy to ablate the whole mass (without expansion) requires an encounter with column density Αρ such that NailyzRES£M,az where my, is the proton mass and ji the mean mass present per H (atoms and ions) in units of nj."," First, to deliver enough energy to ablate the whole mass (without expansion) requires an encounter with column density ${\cal N}_{abl}$ such that ${\cal N}_{abl}\mu m_pv_\odot^2/2\approx {\cal L}M_o/a_o^2$ where $m_p$ is the proton mass and $\mu$ the mean mass present per H (atoms and ions) in units of $m_p$."
" This implies Total deceleration of the nucleus (without expansion) requires Ny. such that A/g!my*M,faz, implying which 15 >>Na.", This implies Total deceleration of the nucleus (without expansion) requires ${\cal N}_{dec}$ such that ${\cal N}_{dec}\mu m_p \approx M_o/a_o^2$ implying which is $\gg {\cal N}_{abl}$.
" For internal strength (energy density) P.=10°P, dynefem*. nuclei will undergo hydrodynamic flow once they reach the depth where atmospheric ram pressure unyv/2 approaches the value P.."," For internal strength (energy density) $P_c=10^6P_6$ $^2$, nuclei will undergo hydrodynamic flow once they reach the depth where atmospheric ram pressure $\mu m_pv_\odot^2/2$ approaches the value $P_c$ ."
" This implies an atmospheric proton density n°>Hpyes=Mesx2P,fpna.", This implies an atmospheric proton density $n \succeq n_{pres}=n_{**}\approx 2P_c/\mu m_p v_\odot^2$.
" The corresponding column density A, along the path for atmospheric scale height H=10°Ης em is The ratio of critical | for ablation end to ram pressure onset is (Eqns. (5.. 7)))"," The corresponding column density ${\cal N}_{pres}$ along the path for atmospheric scale height $H=10^8H_8$ cm is The ratio of critical $\cal N$ for ablation end to ram pressure onset is (Eqns. \ref{Nabl}, \ref{Npres}) ))"
 where the last expression is for p.=0.5.4|.," where the last expression is for $\tilde{\rho}=0.5,\tilde{{\cal L}}=1$."
 There is thus a maximum mass M77(6) above which destruction of the nucleus 1s driven by explosion rather than ablation. namely (setting ratio (8)) to unity) the last value again for 6=0.5.£|. (," There is thus a maximum mass $M_{abl}^{max}(\phi)$ above which destruction of the nucleus is driven by explosion rather than ablation, namely (setting ratio \ref{NablovrNPres}) ) to unity) the last value again for $\tilde{\rho}=0.5,\tilde{{\cal L}}=1$. ("
Note - In the rest of this paper we use subscripts and .. respectively for values of quantities at the points where ablation first exceeds sublimation. and ... where explosion first becomes dominant.),"Note - In the rest of this paper we use subscripts $_*$ and $_{**}$ respectively for values of quantities at the points where ablation first exceeds sublimation, and $_{**}$ where explosion first becomes dominant.)"
 Thus although ram pressure sets in at a lower 7 for the sun than for Jupiter (because v is higher) the same v dependence occurs in the ablation expression so the ratio (8)) just scales as I/H and reaches unity only at a higher 2 in the solar case., Thus although ram pressure sets in at a lower $n$ for the sun than for Jupiter (because $v$ is higher) the same $v$ dependence occurs in the ablation expression so the ratio \ref{NablovrNPres}) ) just scales as $1/H$ and reaches unity only at a higher $n$ in the solar case.
" For solar Hy~0.5 we find n4,=nn,2.5x10UPs em™ while M,,=3xI0?P g and MM=ΑΧI0'(Psecby 5, In the case of impacts with Jupiter Hy~0.01—0.1 so. for relatively steep impacts (secó.<2). Eqns. (8.. 9))"," For solar $H_8\sim 0.5$ we find $n_{pres}=n_{abl}^{end} = n_{**}=2.5\times 10^{14}P_6$ $^{-3}$ while $M_{**}\approx 3\times 10^{10}P_6^3$ g and $M_{abl}^{max}\approx 3\times 10^{10}(P_6\sec\phi)^3$ g. In the case of impacts with Jupiter $H_8 \sim 0.01-0.1$ so, for relatively steep impacts $\sec \phi \preceq 2$ ), Eqns. \ref{NablovrNPres}, \ref{MAblmax}) )"
" show that ram pressure will cause explosion before ablation i$ complete for M,=2x103Ρὸ 5, Thus. as argued by Chyba et al (1993). Zahnle and MacLow (1994) and Carlson et al. ("," show that ram pressure will cause explosion before ablation is complete for $M_o\succeq 2\times 10^9P_6^3$ g. Thus, as argued by Chyba et al (1993), Zahnle and MacLow (1994) and Carlson et al. ("
1997). for steep planetary impacts the ablation process is secondary to pressure-driven explosion except for small masses (or extremely shallow entry) and during the initial entry phase.,"1997), for steep planetary impacts the ablation process is secondary to pressure-driven explosion except for small masses (or extremely shallow entry) and during the initial entry phase."
 In the case of solar chromospheric impacts. however. Hy=0.5 is larger so that the maximum ablation-dominated mass (after sublimation losses) is around 3x10ps £ for steep entry.," In the case of solar chromospheric impacts, however, $H_8=0.5$ is larger so that the maximum ablation-dominated mass (after sublimation losses) is around $3\times 10^{10}P_6^3$ g for steep entry."
 However. most sun-impactors have g=Rx» and secó>>10.," However, most sun-impactors have $q\approx R_\odot$ and $\sec \phi \gg 10$."
" Thus the maximum ablation-dominated mass for these is >>10""pS 5 as discussed further in Section 6.1.", Thus the maximum ablation-dominated mass for these is $\gg 10^{14}P_6^3$ g as discussed further in Section 6.1.
 In the planetary impact literature two factors are discussed which act to limit ablation and we show briefly here why they are less relevant to the solar case., In the planetary impact literature two factors are discussed which act to limit ablation and we show briefly here why they are less relevant to the solar case.
 The first is possible overestimation of the drag coefficient - we have taken the nucleus effective area to equal its geometrical area. which is an overestimation for fluid flow round an obstacle.," The first is possible overestimation of the drag coefficient - we have taken the nucleus effective area to equal its geometrical area, which is an overestimation for fluid flow round an obstacle."
 However. for the solar case the nucleus speed ts so high and the atmospheric density so low that we can reasonably use a kinetic description and treat the particles as impacting the nucleus directly rather than flowing around it.," However, for the solar case the nucleus speed is so high and the atmospheric density so low that we can reasonably use a kinetic description and treat the particles as impacting the nucleus directly rather than flowing around it."
 We note. however. that our kinetic approach may break down toward the lower chromosphere relevant to very heavy nuclei and steeper incidence angles.," We note, however, that our kinetic approach may break down toward the lower chromosphere relevant to very heavy nuclei and steeper incidence angles."
 There the high density results in a shorter collisional mean free path and more fluid like behavior with only indirect ablation via the radiation from a stand-off shock., There the high density results in a shorter collisional mean free path and more fluid like behavior with only indirect ablation via the radiation from a stand-off shock.
 This can decrease the drag coefficient considerably so that our ablation and deceleratior rates are overestimates there., This can decrease the drag coefficient considerably so that our ablation and deceleration rates are overestimates there.
 Secondly. it has been argued 1 the planetary case (e.g. Zahnle and MacLow 1994) that ablatior is limited by the loss of energy needed to ionize hydrogen.," Secondly, it has been argued in the planetary case (e.g. Zahnle and MacLow 1994) that ablation is limited by the loss of energy needed to ionize hydrogen."
 However. the solar atmosphere is around 100 times hotter than Jupiter's atmosphere and is already substantially ionizec at least 1n the upper chromopshere.," However, the solar atmosphere is around 100 times hotter than Jupiter's atmosphere and is already substantially ionized at least in the upper chromopshere."
 Further. solar impacting protons have keV kinetic energies - far higher than H ionizatior energy (13.6 eV) and than the 10 eVvalue for Jupiter so direct impact tonization is very effective in the sun but not Jupiter.," Further, solar impacting protons have keV kinetic energies - far higher than H ionization energy (13.6 eV) and than the 10 eVvalue for Jupiter so direct impact ionization is very effective in the sun but not Jupiter."
 Again. however. heavier and steeper entry nuclei reaching the deep chromosphere are in a cooler medium of lower ionization," Again, however, heavier and steeper entry nuclei reaching the deep chromosphere are in a cooler medium of lower ionization"
sone vears ago we proposed a new algorithm. the Implicit Integral Method (IAI). to solving those radiative transfer problems in which the specific source functions (one for each requency aud direction pair) depend linearly on the radiation field via a single quantity independent of both frequency. and direction.,"Some years ago we proposed a new algorithm, the Implicit Integral Method (IIM), to solving those radiative transfer problems in which the specific source functions (one for each frequency and direction pair) depend linearly on the radiation field via a single quantity independent of both frequency and direction."
 In the paradigm instance of radiative transfer through an ideal medium formed by atoms with only two energy levels (Two-Level Atom nodel). (his quantity is the inleeral over Irequencies of the mean specific intensity of the racialion field. weighted with the spectral profile. (," In the paradigm instance of radiative transfer through an ideal medium formed by atoms with only two energy levels (Two-Level Atom model), this quantity is the integral over frequencies of the mean specific intensity of the radiation field, weighted with the spectral profile. ("
See Simonneau and Crivellari. 1993. ereinafter Paper 1.) Because it is independent of both frequency and direction. such a quantity constitutes a single scalar coupling for all the specific RT equations. ancl can be chosen in a natural way as (he protagonist variable for the numerical solution of the RT problem.,"See Simonneau and Crivellari, 1993, hereinafter Paper I.) Because it is independent of both frequency and direction, such a quantity constitutes a single scalar coupling for all the specific RT equations, and can be chosen in a natural way as the protagonist variable for the numerical solution of the RT problem."
 This choice is the distinctive and essential feature of our LAL: to work with a quantitv which is independent of both frequency. aud direction brings about that the method does not require to store and invert huge matrices like in (he customary numerical algoritlims emploved in IT problenis., This choice is the distinctive and essential feature of our IIM: to work with a quantity which is independent of both frequency and direction brings about that the method does not require to store and invert huge matrices like in the customary numerical algorithms employed in RT problems.
 We have already remarked in Paper I that our algorithm is a mere phenomenological representation of the actual physical process., We have already remarked in Paper I that our algorithm is a mere phenomenological representation of the actual physical process.
 Because of (hat and due to the lack of a matricial structure. the advantages of the HAL in terms of reliability. accuracy and robustness sould be sell-evident. as well as the conspicous saving of both computational (ime and memory storage il makes possible.," Because of that and due to the lack of a matricial structure, the advantages of the IIM in terms of reliability, accuracy and robustness sould be self-evident, as well as the conspicous saving of both computational time and memory storage it makes possible."
 The aforesail advantages suggested us (he possibility to employ the LAI also in the computation of stellar atmospheres models. where we must solve many (some hundreds) RT equations. one lor each Irequency.," The aforesaid advantages suggested us the possibility to employ the IIM also in the computation of stellar atmospheres models, where we must solve many (some hundreds) RT equations, one for each frequency."
 The source function of each specilic RT equation is here the weighted mean of a term that includes the mean specific intensity of the radiation field through a scattering-like integral with a thermal contribution given by the Planck function BAT)., The source function of each specific RT equation is here the weighted mean of a term that includes the mean specific intensity of the radiation field through a scattering-like integral with a thermal contribution given by the Planck function $B_{\nu}(T)$.
 The paradigm problem of the sell-consistent temperature correction when computing stellar atinosphiere models was considered in Crivellari and Simonneau (1994)., The paradigm problem of the self-consistent temperature correction when computing stellar atmosphere models was considered in Crivellari and Simonneau (1994).
 First of all. we must recognize that the geometrical structure of the svstem. that is the sequence of the discrete atmospheric lavers. must be necessarily (he same for all the frequencies.," First of all, we must recognize that the geometrical structure of the system, that is the sequence of the discrete atmospheric layers, must be necessarily the same for all the frequencies."
 Bul we must also recognize that for any given frequency some lavers do nol contribute to the formation of the spectrum., But we must also recognize that for any given frequency some layers do not contribute to the formation of the spectrum.
" They co not take an effective part in the radiative transfer process because either (hey are exceedingly transparent 1) or they correspond to optically very deep regions exp[—.Nz,]c 0).", They do not take an effective part in the radiative transfer process because either they are exceedingly transparent $\exp \lbrace - \Delta\tau_{\nu}\rbrace \simeq 1$ ) or they correspond to optically very deep regions $\exp \lbrace - \Delta\tau_{\nu}\rbrace \simeq 0$ ).
 The layers intermediate between the above (wo groups constitute the specilic spectral formation region., The layers intermediate between the above two groups constitute the specific spectral formation region.
 llowever all the lavers of the structure must be taken into account in (he numerical algorithm. inrespectivelv of the frequency. considered.," However all the layers of the structure must be taken into account in the numerical algorithm, irrespectively of the frequency considered."
 Yet. due to the dramatic difference amongthe," Yet, due to the dramatic difference amongthe"
the cooling flux. of he burst decreases with time.,the cooling flux of the burst decreases with time.
 Such variations have been discussed in Damien et (1989) and iu Bhattacharyya. MIler. Calloway (2010).," Such variations have been discussed in Damen et (1989) and in Bhattacharyya, Miller, Galloway (2010)."
 Note that this svsteniatic effect can be corrected. in principle. if the data are fit directly with detailed models of ueutron-star atiiosplheres (see. e... Majezyua Made] 2005).," Note that this systematic effect can be corrected, in principle, if the data are fit directly with detailed models of neutron-star atmospheres (see, e.g., Majczyna Madej 2005)."
 A potentially imvortant source of unucertaintv in the measurements is introduced by the errors iu the absolute flux calibration of the RNTE/PCA detector., A potentially important source of uncertainty in the measurements is introduced by the errors in the absolute flux calibration of the RXTE/PCA detector.
 The current calibration of the PCA aud the cross-calibration between X-ray satellites lave boeji carried out using the Crab nebula as a standard caudle (Jahoda et 22006: see also Toor Seward 1971: Iil et 22005: Weisskopf et 22010)., The current calibration of the PCA and the cross-calibration between X-ray satellites have been carried out using the Crab nebula as a standard candle (Jahoda et 2006; see also Toor Seward 1974; Kirsch et 2005; Weisskopf et 2010).
 The uncertainties in the fux calibration can be due to a poteutial overall ofiset between the inferred aud the true flux of the Crab nebula. which may be as buge as (IKirsch et 22005: Weisssopf ct 22010).," The uncertainties in the flux calibration can be due to a potential overall offset between the inferred and the true flux of the Crab nebula, which may be as large as (Kirsch et 2005; Weisskopf et 2010)."
 This can only chanee the mean value of the ferred apparent area in cach source and does not alter the observed dispersion., This can only change the mean value of the inferred apparent area in each source and does not alter the observed dispersion.
 We will explore this issue as well as uncertainties related to the variability of the Crab nebula itself (Wilsou-Tlodge et 22011) in more detail iu Paper HT of this series., We will explore this issue as well as uncertainties related to the variability of the Crab nebula itself (Wilson-Hodge et 2011) in more detail in Paper III of this series.
 Figure EMi (loft pzanels) show the dependence of the ciucreie flux on color temperature for all the spectra in the cooling tails of ISS 260. IU οι and IU 536 that we cousider to be statistically acceptable (see 833).," Figure \ref{fig:kt_flux_detail} (left panels) show the dependence of the emerging flux on color temperature for all the spectra in the cooling tails of KS $-$ 260, 4U $-$ 34, and 4U $-$ 536 that we consider to be statistically acceptable (see 3)."
 We chose these thrce sources to use as detailed examples because of the high number spectra obtained for cach aud the fact that they :«xxui the range of behavior in cooling tracks that we will discuss below., We chose these three sources to use as detailed examples because of the high number spectra obtained for each and the fact that they span the range of behavior in cooling tracks that we will discuss below.
 If the whole neutron star is cluitting as a sinele-temperature blackxlv and the color correction factor is independent of color temperature. then Fro.) should scale as T1.," If the whole neutron star is emitting as a single-temperature blackbody and the color correction factor is independent of color temperature, then $F_{\rm cool}$ should scale as $T_{\rm c}^4$."
 Our aiii Liere is to investigate the systematic uncertainties ou the measurement of the apparent surface ai‘ca du each source at different flix levels., Our aim here is to investigate the systematic uncertainties on the measurement of the apparent surface area in each source at different flux levels.
 We. therefore. divided the data into a umuber of flux bius aud plotted i the same figure (ight panels) the distribution of the blackbody normalization values for some representative bius.," We, therefore, divided the data into a number of flux bins and plotted in the same figure (right panels) the distribution of the blackbody normalization values for some representative bins."
 The blackbody normalization for cach spectrum is defined formally as A=Foo/ospi!. although in practice this is one of the two measured parameters aud the flux is derived from the above definition.," The blackbody normalization for each spectrum is defined formally as $A\equiv F_{\rm
 cool}/\sigma_{\rm SB}T_{\rm c}^4$, although in practice this is one of the two measured parameters and the flux is derived from the above definition."
" According to equation (3 n. the1dlackbodsy normalization: is: equal to A=f£.⊽82,DIE'/D? aud we report itpo in units: of. (lau/10D kpc)."," According to equation \ref{eq:rapp}) ), the blackbody normalization is equal to $A=f_c^{-4}R_{\rm app}^2/D^2$ and we report it in units of (km/10 $^2$."
D If we do not correct for the dependence οf the color correction factx on the flux. we expect the normalization A to show a dependence on color teniperature hat is the ΗΤΟΙ svuuuetric of that shown in Fieure 2.," If we do not correct for the dependence of the color correction factor on the flux, we expect the normalization $A$ to show a dependence on color temperature that is the mirror symmetric of that shown in Figure 2."
 The flix-tempcrature diagrams of KS 260. [U 31. ane IU 536 share a number of similarities but are also distinguished by a umber of differences.," The flux-temperature diagrams of KS $-$ 260, 4U $-$ 34, and 4U $-$ 536 share a number of similarities but are also distinguished by a number of differences."
 In all three cases. the vast majority of data poiuts lie along a very well defined correlation.," In all three cases, the vast majority of data points lie along a very well defined correlation."
 This reproducibiitv of the coolue CHrVOR ο tens of X-ray bursts per source. combined witli he lack of large αιuplitude flux oscillations during coolue tails of |nrsts; provides the strougest aremucut that the hermal emission eugulfs the eutire neutron-star surface with very suall temperature variations at different latitudes and longitudes.," This reproducibility of the cooling curves of tens of X-ray bursts per source, combined with the lack of large amplitude flux oscillations during cooling tails of bursts, provides the strongest argument that the thermal emission engulfs the entire neutron-star surface with very small temperature variations at different latitudes and longitudes."
 Iu IU 31. a deviation of the data poiuts from he ἔωωι~T correlation is evident at high fixes aud color enperatures (ας=2.5 keV). which may be due to the evolution of the color correction factor at high Edcingtou fluxes (see the discussion in §l and Fig. 2)).," In 4U $-$ 34, a deviation of the data points from the $F_{\rm
cool}\sim T_{\rm c}^4$ correlation is evident at high fluxes and color temperatures $T_{\rm c}\ge 2.5$ keV), which may be due to the evolution of the color correction factor at high Eddington fluxes (see the discussion in 4 and Fig. \ref{fig:fcolor}) )."
 The same deviation is not evident iu NS 260 or IU 536. for which he highest temperatures cheouutered in the cooling tails were Z2.5 keV. Finally. in all three sources. a nuniber of outliers exist at the lowest flux levels. with uormalizationus that deviate youn the above correlation.," The same deviation is not evident in KS $-$ 260 or 4U $-$ 536, for which the highest temperatures encountered in the cooling tails were $\lesssim
2.5$ keV. Finally, in all three sources, a number of outliers exist at the lowest flux levels, with normalizations that deviate from the above correlation."
 In WS 260 aud IU 536. the outliers correspond to higher normalization values. whereas in LU 3l they correspond to lower normalization values with respect to the majority of the data poiuts.," In KS $-$ 260 and 4U $-$ 536, the outliers correspond to higher normalization values, whereas in 4U $-$ 34 they correspond to lower normalization values with respect to the majority of the data points."
 Any combination of the effects discussed earlier in this aud in the previous section may be responsible for the outliers., Any combination of the effects discussed earlier in this and in the previous section may be responsible for the outliers.
 Non-uuiforum coolitic of the neutron star surface will lead to a reduced inferred value for the apparent surface area., Non-uniform cooling of the neutron star surface will lead to a reduced inferred value for the apparent surface area.
 Reflection of the πιwace N-vav cussion off a ecometrically thin accretion disk will cause au increase in the inferred value for the apparent surface area., Reflection of the surface X-ray emission off a geometrically thin accretion disk will cause an increase in the inferred value for the apparent surface area.
 Finally. Comptonization of the surface eiission in a corona may have either effect. depending ou the €'ounrptou temperature of 1ο radiation.," Finally, Comptonization of the surface emission in a corona may have either effect, depending on the Compton temperature of the radiation."
 Our goal in this work is not to uudoerstau the origin of the outhers. but rather to ensure that their presence docs rot introduce any |jdjases in the measuremens of the apparent surface areas inferred from the vast majority of spectra.," Our goal in this work is not to understand the origin of the outliers, but rather to ensure that their presence does not introduce any biases in the measurements of the apparent surface areas inferred from the vast majority of spectra."
 Tudeed. includiug the outliers in a formal fi of the flux-temperatire correlation will cause a systematic iucrease of he apparent surfacὉ area with decreasius fius in WS 260 aud in 1U 536 and a systematic decrease iu IU 1728 31 Damen ct 11990: Bhaacharvya et 22010).," Indeed, including the outliers in a formal fit of the flux-temperature correlation will cause a systematic increase of the apparent surface area with decreasing flux in KS $-$ 260 and in 4U $-$ 536 and a systematic decrease in 4U $-$ 34 Damen et 1990; Bhattacharyya et 2010)."
 This issue was lareely avoided in our earlier work (Ozzel et 22009: Cüvvere 220102. 2tHb) by considering oulv the relatively carly time intervals in the cooling ail of cach burst. which conyond ouly to the brightest flux bins.," This issue was largely avoided in our earlier work (Özzel et 2009; Güvver et 2010a, 2010b) by considering only the relatively early time intervals in the cooling tail of each burst, which correspond only to the brightest flux bins."
 In order to eo bevoud this limitation here. we consider all spectred data for ich burst aud eniploy a Bayesian Caussian mixture algorithm. which is a standard xocedure for outlicy detectio 1 robust statistics (c.g.. Titterington. Smith. Makov 1985: McLachlan Peel 2000: IIuber Rouchet1 2009).," In order to go beyond this limitation here, we consider all spectral data for each burst and employ a Bayesian Gaussian mixture algorithm, which is a standard procedure for outlier detection in robust statistics (e.g., Titterington, Smith, Makov 1985; McLachlan Peel 2000; Huber Ronchetti 2009)."
 Our workine hvpothesis that the main peak of normalizatious in cach flux interval corresponds to he signal ancl the remaincer are outliers is shaped by two aspects of the observations., Our working hypothesis that the main peak of normalizations in each flux interval corresponds to the signal and the remainder are outliers is shaped by two aspects of the observations.
 First. at high fluxes. cach hisogra can be described Won sinele Gaussian with no evidence or room for a second distribution of what we would call outliers.," First, at high fluxes, each histogram can be described by a single Gaussian with no evidence or room for a second distribution of what we would call outliers."
 At lower fluxes. when the histograms can be decomposed into two Caussiaus. the distribution with the highest peak has properties that suxvothly connect to those of the single Gaussians at higher fiuxes.," At lower fluxes, when the histograms can be decomposed into two Gaussians, the distribution with the highest peak has properties that smoothly connect to those of the single Gaussians at higher fluxes."
 Second. the Gaussians that we call our sigual always coutain the majority of the data poiuts compared to the distribution of what we call the outliers.," Second, the Gaussians that we call our signal always contain the majority of the data points compared to the distribution of what we call the outliers."
 We take these as our criteria for defining our signal., We take these as our criteria for defining our signal.
 0.,"Merging this result with previous studies of mirror stellar evolution, we obtain a self consistent scenario for the mirror dark matter interpretation of the current experimental data."
8Γ Y(He, Our calculation of $Y_{He'}$ will play an important role in future precision tests of the mirror dark matter paradigm.
’), 1cm 0.2cm This work was supported by the Australian Research Council and by the Belgian Fund for Scientific Research (FNRS).
We have also investigated the behaviour od. 2IT(z) above Shun for other distributions of primary electrons over the polar cap.,We have also investigated the behaviour od $\del(\varepsilon)$ above $\et$ for other distributions of primary electrons over the polar cap.
" We have considered. intermediate cases between models A and D. (ie.wiith a uniform coverage of only an outer part of the polar cap area between some inner radius ri,re and the polar cap radius ric). ane models with uniformly filled interior of the polar cap surface but with increased electron density. alone the polar cap rim (cf."," We have considered intermediate cases between models A and D, (i.e.wiith a uniform coverage of only an outer part of the polar cap area between some inner radius $r_{\rm in} < \rpc$ and the polar cap radius $\rpc$ ), and models with uniformly filled interior of the polar cap surface but with increased electron density along the polar cap rim (cf."
Daaugherty Πανάς 1906).,Daaugherty Harding 1996).
 We conclude tha regardless the actual shape of the active part covered!’ with primary electrons) of the polar cap (either an outer rim. or an entire cap. or a ring. or entire cap/ring | increase rim density). one does expect in general strong changes in he peak separation to occur at photon energies close to igh-energv spectral cutoll due to magnetic absorption.," We conclude that regardless the actual shape of the active part `covered' with primary electrons) of the polar cap (either an outer rim, or an entire cap, or a ring, or entire cap/ring + increased rim density), one does expect in general strong changes in the peak separation to occur at photon energies close to high-energy spectral cutoff due to magnetic absorption."
" A word of technical comment seems to be appropriate or AUFL0,43 at ~300MeV.", A word of technical comment seems to be appropriate for $\del = 0.43$ at $\sim 300\MeV$.
 LI appears that a technique adopted: by Ixanbach. (1999). of fitting the observed. pulse shapes with asvmetric Lorentz profiles tends to overestimate he true value of APE by a few thousandth parts of phase., It appears that a technique adopted by Kanbach \shortcite{kanbach} of fitting the observed pulse shapes with asymetric Lorentz profiles tends to overestimate the true value of $\del$ by a few thousandth parts of phase.
 Vherefore the actual value may be closer to 0.42 than 0.43 (Maurice Ciros. private communication).," Therefore the actual value may be closer to 0.42 than 0.43 (Maurice Gros, private communication)."
 Nonetheless. this shift MMin APouk does not change any conclusions. of⋅ this. work.," Nonetheless, this shift in $\del$ does not change any conclusions of this work."
 In this paper we addressed a recent suggestion of Ixanbach (1000) that peak separation Apeak in the double-peak gamma-ray pulses of the Vela pulsar may monotonicallv decrease with increasing photon energy at a rate 0.025 phase per decade in energy over the range 50MeV to 9GeV. We calculated ganimia-ray pulses expected. in. polar-cap models with magnetospheric activity induced by curvature radiation of beam particles.," In this paper we addressed a recent suggestion of Kanbach \shortcite{kanbach} that peak separation $\del$ in the double-peak gamma-ray pulses of the Vela pulsar may monotonically decrease with increasing photon energy at a rate $\sim 0.025$ phase per decade in energy over the range $50\MeV$ to $9\GeV$, We calculated gamma-ray pulses expected in polar-cap models with magnetospheric activity induced by curvature radiation of beam particles."
 “Pwo types of geometry of magnetospheric column above the polar cap were assumed: a hollow-column associated with an outer rim of the polar cap and a filled column associated with a uniform. polar cap., Two types of geometry of magnetospheric column above the polar cap were assumed: a hollow-column associated with an outer rim of the polar cap and a filled column associated with a uniform polar cap.
 Four models were considered with two scenarios for the acceleration of beam particles., Four models were considered with two scenarios for the acceleration of beam particles.
 Pulsed emission in the models Was a superposition of curvature radiation due to beam particles and. svnehrotron radiation due to secondary. 67 pairs in magnetospheric cascades., Pulsed emission in the models was a superposition of curvature radiation due to beam particles and synchrotron radiation due to secondary $e^\pm$ pairs in magnetospheric cascades.
 Lhe changes in the peak separation were investigated. with Monte. Carlo numerical simulations., The changes in the peak separation were investigated with Monte Carlo numerical simulations.
" We found that regardless the dillerences in the mocels. the peak separation A?""K below a few GeV. where the emission is dominated. by synchrotron component. is either a weak decreasing function. of photon energy. 5, or remains constant."," We found that regardless the differences in the models, the peak separation $\del$ below a few GeV, where the emission is dominated by synchrotron component, is either a weak decreasing function of photon energy $\varepsilon$, or remains constant."
 Both variants may be considered. to be in agreement with the results of Wanbach (1909) for the latter are allected by large statistical errors.," Both variants may be considered to be in agreement with the results of Kanbach \shortcite{kanbach}
 for the latter are affected by large statistical errors."
 A particular behaviour of Alen epends on a combination of several factors. including strength. of magnetic field in the region of pair formation ancl mocel of electron acceleration (both of which determine spectral and directional. properties. of the radiation at cdillerent altitudes). as well as viewing ecometryv.," A particular behaviour of $\del$ depends on a combination of several factors, including strength of magnetic field in the region of pair formation and model of electron acceleration (both of which determine spectral and directional properties of the radiation at different altitudes), as well as viewing geometry."
 Essentially. in strong fields. Bij210076. PUE decreases with increasing photon energy 2. whereas or δωcLOM€. the peak separation Abe atavs at a constant level.," Essentially, in strong fields, $B_{\rm local} \ga 10^{12}\G$, $\del$ decreases with increasing photon energy $\varepsilon$, whereas for $B_{\rm local} < 10^{12}\G$, the peak separation $\del$ stays at a constant level."
 Aloreover. we found that due to the magnetic absorption (Dc ) there exists a critical energy. foun at which he peak separation. AP'Senl makes an abrupt. turn. and hen changes dramatically for 2Shun.," Moreover, we found that due to the magnetic absorption $\gamma \bld B \rightarrow e^\pm$ ) there exists a critical energy $\et$ at which the peak separation $\del$ makes an abrupt turn and then changes dramatically for $\varepsilon > \et$."
 Lt increases in he hollow-column. models CX. D. and. €) ancl decreases in the filled-column. model (D). at a rate 0.28. phase vcr decade of photon energy.," It increases in the hollow-column models (A, B, and C) and decreases in the filled-column model (D), at a rate $\sim 0.28$ phase per decade of photon energy."
 The numerical behaviour of Are jn this regime in the hollow-column models was easily reproduced to high accuracy with a simple analytical moclel of magnetospheric transparency. for a photon of energy ancl its momentum tangential to local dipolar magnetic field line at a site of its origin.," The numerical behaviour of $\Delta^{\rm peak}$ in this regime in the hollow-column models was easily reproduced to high accuracy with a simple analytical model of magnetospheric transparency for a photon of energy $\varepsilon$ , and its momentum tangential to local dipolar magnetic field line at a site of its origin."
"e An exact value of Shun 18 model-dependent but it is confined to a range between ~0.9GeV and ~4.5Gey,", An exact value of $\varepsilon_{\rm turn}$ is model-dependent but it is confined to a range between $\sim 0.9\GeV$ and $\sim 4.5\GeV$.
 To lind. such a hypothetical turnover of Are in real observational data would require. however. high-sensitivity detectors. since for 52soun the expected Hux of eamma-ravs crops significantly.," To find such a hypothetical turnover of $\Delta^{\rm peak}$ in real observational data would require, however, high-sensitivity detectors, since for $\varepsilon > \varepsilon_{\rm turn}$ the expected flux of gamma-rays drops significantly."
 LE detecb. this turnover would be an important signature of polar cap activity in eamma-pav pulsars.," If detected, this turnover would be an important signature of polar cap activity in gamma-ray pulsars."
 Ht would support the notion that high-energv cutolfs in gamma-ray spectra of pulsars are due to magnetic absorption., It would support the notion that high-energy cutoffs in gamma-ray spectra of pulsars are due to magnetic absorption.
 The CR-incluced cascades models. like those considered in this work. are not the only possibility for nearly aligned rotators to produce double-peak pulses with large phase separations.," The CR-induced cascades models, like those considered in this work, are not the only possibility for nearly aligned rotators to produce double-peak pulses with large phase separations."
 There exists an alternative class of models. - with pair cascades above polar cap induced by magnetic inverse Compton scatterings (C8) of primary electrons in the field of soft photons from the stellar surface - proposed in a series of papers (e.g. Sturmer Dermoer 19094. Sturner et al.," There exists an alternative class of models - with pair cascades above polar cap induced by magnetic inverse Compton scatterings (ICS) of primary electrons in the field of soft photons from the stellar surface - proposed in a series of papers (e.g. Sturner Dermer 1994, Sturner et al."
 1995)., 1995).
 In particular. Sturmer et al. (," In particular, Sturner et al. ("
1995) present the detailed Monte Carlo model spectra of the Vela pulsar.,1995) present the detailed Monte Carlo model spectra of the Vela pulsar.
 They also present pulse profiles at a fixed energy of 50 MeV. (for several viewing angles) but no word of Comment is given regarding the problem of Apeak versus photon energy., They also present pulse profiles at a fixed energy of 50 MeV (for several viewing angles) but no word of comment is given regarding the problem of $\Delta^{\rm peak}$ versus photon energy.
 We expect the outcome to be qualitatively similar to our results., We expect the outcome to be qualitatively similar to our results.
 First. the scatterings take place mostly within a very Limited height above the polar cap surface (below h~ fe) and the preferred. directions of propagation of the LCS photons will be fixed by magnetic field lines just above the surface.," First, the scatterings take place mostly within a very limited height above the polar cap surface (below $h \sim R_{\rm pc}$ ) and the preferred directions of propagation of the ICS photons will be fixed by magnetic field lines just above the surface."
 Therefore. AP due solely to LCS photons should. stay constant for a wide range of energy.," Therefore, $\Delta^{\rm peak}$ due solely to ICS photons should stay constant for a wide range of energy."
 Inclusion of svnchrotron photons due to pairs is unlikely to notably alfect Ares unless the pair formation front is vertically more extended wan for CR-incluced cascades., Inclusion of synchrotron photons due to pairs is unlikely to notably affect $\Delta^{\rm peak}$ unless the pair formation front is vertically more extended than for CR-induced cascades.
 Second. some turnover point ab fuga not exceeding 1 GeV should be present due to magnetic absorption.," Second, some turnover point at $\varepsilon_{\rm turn}$ not exceeding 1 GeV should be present due to magnetic absorption."
 Phe behaviour of AU for =>Shun garoulcl roughly follow the clashed lines in Fig., The behaviour of $\del$ for $\varepsilon > \et$ should roughly follow the dashed lines in Fig.
 2 (upper panel) and. Fig.3 as long as the assumption about photons (which are to beabsorbed) propagating tangentially to local ipolar magnetic field line at their site of origin remains valid for majority of LCS, 2 (upper panel) and Fig.3 as long as the assumption about photons (which are to beabsorbed) propagating tangentially to local dipolar magnetic field line at their site of origin remains valid for majority of ICS
lis an inensively studied: low-mass X-ray binary that was initially discovered. with the European X-ray Observatory SAXFellit«| (GEXOSAT)) in 1985. February 1985).,is an intensively studied low-mass X-ray binary that was initially discovered with the European X-ray Observatory SATellite ) in 1985 February .
 1owever. in retrospect the source already appeared active in sslow survey observations several times beginning 1984 July 1999). whereas the earliest detection dates back to 980 Alay. when wavas serendipitouslv observed. with the ssalellite 1986).," However, in retrospect the source already appeared active in slew survey observations several times beginning 1984 July , whereas the earliest detection dates back to 1980 May, when was serendipitously observed with the satellite ."
. Phe svstem exhibits irregular X-ray dips and displavs eclipses that last for 8.35 min and recur every 3.82 hr. which allow the unambiguous determination of the orbital period of the binary2009).," The system exhibits irregular X-ray dips and displays eclipses that last for $\sim 8.3$ min and recur every 3.82 hr, which allow the unambiguous determination of the orbital period of the binary."
. The «letection of tvpe-L X-ray bursts. conclusively identify the compact primary as a neuron star., The detection of type-I X-ray bursts conclusively identify the compact primary as a neutron star.
 A few X-ray bursts have been observed that exhibited photospjeric radius expansion (PRE). which indicates that the Edclineton luminosity is reached. near the burst »ealkk ancl allows for a distance estimate towards the source2008).," A few X-ray bursts have been observed that exhibited photospheric radius expansion (PRE), which indicates that the Eddington luminosity is reached near the burst peak and allows for a distance estimate towards the source."
. For a Ielium-dominated: photosphere. à. distance of D=Tdc0.0 kpe οui be derived. while assuming solar composition results in: Cdistance estimate of 2=5.90.9 kpe2008).," For a Helium-dominated photosphere, a distance of $D=7.4\pm0.9$ kpc can be derived, while assuming solar composition results in a distance estimate of $D=5.9\pm0.9$ kpc."
. The rise time ancl duration of the PRE bursts observed [rom ssugeest pure Helium ignition. rendering 1.4 kpe as the best," The rise time and duration of the PRE bursts observed from suggest pure Helium ignition, rendering 7.4 kpc as the best"
"For each of the 20 evolutionary curves we mark (he maximum values ,,; and Pg, of the simulated flare loops with a diamond svmbol in Fig.",For each of the 20 evolutionary curves we mark the maximum values $n_{max}$ and $T_{max}$ of the simulated flare loops with a diamond symbol in Fig.
 9 and compare it with the RTV law (Eq., 9 and compare it with the RTV law (Eq.
 21). which is indicated with a dashed line in Fig.," 21), which is indicated with a dashed line in Fig."
 9., 9.
" We can consider now the density ralio μιμμ. heatingal the [are maxinum temperature and find that (his ratio is à svstematic function of the time scale τω. but is almost d ol the heating rate 75,4 and maximum temperature ov "," We can consider now the density ratio $n_{max}/n_{RTV}$ at the flare maximum temperature and find that this ratio is a systematic function of the heating time scale $t_{heat}$, but is almost independent of the heating rate $E_{H0}$ and maximum temperature $T_{max}$ ."
"This ratio amounts to πμmn=1.0940.06 For /5,,,=328 SS μι=0.81dE0.07 lor liege=164 s. Dyeeμεν=57£0.06"" lor ενω=82 s. μιMery=0.35+0.06 dor linea=4l s. so the RIV overprecicts maximum densitv for small heating time scales. but agrees quite well for longer heating “ime scales and thus provides a good proxi to predict the maximum [lare densities."," This ratio amounts to $n_{max}/n_{RTV}=1.09\pm0.06$ for $t_{heat}=328$ s, $n_{max}/n_{RTV}=0.81\pm0.07$ for $t_{heat}=164$ s, $n_{max}/n_{RTV}=0.57\pm0.06$ for $t_{heat}=82$ s, $n_{max}/n_{RTV}=0.35\pm0.06$ for $t_{heat}=41$ s, so the RTV overpredicts the maximum density for small heating time scales, but agrees quite well for longer heating time scales and thus provides a good proxi to predict the maximum flare densities."
 Applying the RV law to scaling laws of flare parameters. one has to correct for a numerical factor for shorter heating limes. but (his factor islargely independent of the flare temperature.," Applying the RTV law to scaling laws of flare parameters, one has to correct for a numerical factor for shorter heating times, but this factor islargely independent of the flare temperature."
atmospieric Llow—asstuning local reheatiue trou dissipation—we note that oty result is «ependeut on the set-up we used. in paricular the spatial form used for the drag.,"atmospheric flow—assuming local reheating from dissipation—we note that our result is dependent on the set-up we used, in particular the spatial form used for the drag."
 As we have eiiphasizect th'oughlout. it is the spatial strιο of tie atmosphere tliat dictates the flow eiergetics. s»ecilically the cor'elations. between temperatures. ieating rates. aud clissipation rates.," As we have emphasized throughout, it is the spatial structure of the atmosphere that dictates the flow energetics, specifically the correlations between temperatures, heating rates, and dissipation rates."
" Our results here are depeucent ou the particular reation between the spatial function we have usec for the aluospheric drag aid the underlying atmospheric st""ucture.", Our results here are dependent on the particular relation between the spatial function we have used for the atmospheric drag and the underlying atmospheric structure.
 Mauy of the drag mechanuistIs Ineitloled earlier will have very clifferent spatia depeucerCles., Many of the drag mechanisms mentioned earlier will have very different spatial dependencies.
 The best way to estimate the euergeich iportance ol auy drag mechanisiu is to apply the Equatious (19--22)) to a particular atuosplere with some form of drag at work., The best way to estimate the energetic importance of any drag mechanism is to apply the Equations \ref{eq:finalu}- \ref{eq:finaleg}) ) to a particular atmosphere with some form of drag at work.
 We can further unclerstauc the energetics of the atmosphere by stidying its spatial properties., We can further understand the energetics of the atmosphere by studying its spatial properties.
 is will allow us to ideutify the regious of he atinosphere that are the nost relevant for APE eration. as deteriniued by the spatial correation between the atinosyheric quantities of interest.," This will allow us to identify the regions of the atmosphere that are the most relevant for APE generation, as determined by the spatial correlation between the atmospheric quantities of interest."
 Before we proceed. we need to hiehlight an important nuance in 1js analysis.," Before we proceed, we need to highlight an important nuance in this analysis."
 Although rations (19)) and (20)) give local expressiis lor UPE alπι APE geleration. note that the local tes only have meaulue in reference to the elobal total.," Although Equations \ref{eq:finalu}) ) and \ref{eq:finala}) ) give local expressions for UPE and APE generation, note that the local values only have meaning in reference to the global total."
" The quantiy T, is explicitly a global Va1ο aud is au inherent part of the definitiis of APE anc UPE. both of which were originally 1110ivated by the concept of a reference state with APE=0."," The quantity $T_r$ is explicitly a global value and is an inherent part of the definitions of APE and UPE, both of which were originally motivated by the concept of a reference state with $=0$."
 The reade ‘should keep iu imiud tliat the following plots of APE and UPE generajon are to be tasen as aul idication of which areas of the atinosphere (when integrated over) cont‘bute the most o the elobal total., The reader should keep in mind that the following plots of APE and UPE generation are to be taken as an indication of which areas of the atmosphere (when integrated over) contribute the most to the global total.
 In Figure 1 we plot the local values of UPE aud APEE generation as a function of longitiT ad pressure (the moclel’s vertical coordinate) Dor au equatorial slice hroug htie atiuosphere. (, In Figure \ref{fig:eq} we plot the local values of UPE and APE generation as a function of longitude and pressure (the model's vertical coordinate) for an equatorial slice through the atmosphere. (
CLE erey rates tend to be greatest at the equator.),The energy rates tend to be greatest at the equator.)
 Iu the left j»anel we show radiative generation for the Rauscher&Menou(2010) model at 10.000 clays. while in the t panel we plot dissing” frictional generation for the st‘OLigest-drag model rom Pernaetal.(2010a) at 2000 clay," In the left panel we show the radiative generation for the \citet{RM10} model at 10,000 days, while in the right panel we plot the “missing"" frictional generation for the strongest-drag model from \citet{Perna2010a} at 5000 days."
 je pattern of radiative UPE aud APE generation is simiar for all of 4e moclels. although he iplituces aud some of the detailed striοἱure is clilferent.," The pattern of radiative UPE and APE generation is similar for all of the models, although the amplitudes and some of the detailed structure is different."
" T πο ""UPE aud APE eeneratiO1 the Pλα and PMBb models is weaSer aud primarly ceonncentrated in he uppermost layers of the atinosphere. ou πο side."," The “missing"" UPE and APE generation in the PMRa and PMRb models is weaker and primarily concentrated in the uppermost layers of the atmosphere, on the night side."
 The OLeer drag times ii these models ueau that it is only he ieliest evels (where the applied drag is the strongest) wuch are able to siguilicantly clissipate kinetic energy. even though most of the unetic energy resicles in the deeper layers. (," The longer drag times in these models mean that it is only the highest levels (where the applied drag is the strongest) which are able to significantly dissipate kinetic energy, even though most of the kinetic energy resides in the deeper layers. ("
See Figure | of Rauscher Menou20) for a plot of kileic energy ve‘sus depth: it has a tmasximitun at 2 bar.),See Figure 4 of Rauscher Menou 2010 for a plot of kinetic energy versus depth; it has a maximum at $\sim2$ bar.)
Stellar lithium. abundances are continuing to attract strong interest among the astronomical community.,Stellar lithium abundances are continuing to attract strong interest among the astronomical community.
 In particular. many studies are devoted to constraining and finding explanations for the famous “Spite plateau” of warm metal-poor halo stars. first discovered by Spite&(1952).," In particular, many studies are devoted to constraining and finding explanations for the famous “Spite plateau” of warm metal-poor halo stars, first discovered by \citet{Spite82}."
 For recent studies of lithium in halo field stars see e.g. Ryanal. (2001).. Charbonnel&Primas(2005).. Asplundetal. (2006).. Bonifacioetal. (2007).. and Hosfordetal.(2009).," For recent studies of lithium in halo field stars see e.g. \citet{Ryan01}, , \citet{Charbonnel05a}, \citet{Asplund06}, \citet{Bonifacio07a}, and \citet{Hosford09}."
. To put constrainst on the primordial lithium abundance in the Universe. it is of fundamental importance to determine if and by how much the Li abundances increase with increasing metallicity on the Spite plateau.," To put constrainst on the primordial lithium abundance in the Universe, it is of fundamental importance to determine if and by how much the Li abundances increase with increasing metallicity on the Spite plateau."
 The behaviour of Li abundance with effective temperature is of equal importance. as it could lend support to the notion that atomic diffusion ts acting in metal-poor stellar atmospheres.," The behaviour of Li abundance with effective temperature is of equal importance, as it could lend support to the notion that atomic diffusion is acting in metal-poor stellar atmospheres."
 Lithium depletion through atomic diffusion has been suggested as a solution to the discrepancy between the Spite plateau abundance and the predicted value of the primordial lithium abundance (seee.g.Kornetal.2007:Lind 2009)..," Lithium depletion through atomic diffusion has been suggested as a solution to the discrepancy between the Spite plateau abundance and the predicted value of the primordial lithium abundance \citep[see e.g.][]{Korn07,Lind09b}."
" To accurately infer both the mean lithium abundance in the halo and the abundance behaviour with metallicity and effective temperature. it is crucial to have a realistic description of the Π line formation,"," To accurately infer both the mean lithium abundance in the halo and the abundance behaviour with metallicity and effective temperature, it is crucial to have a realistic description of the lithium line formation."
 Previous non-LTE analyses spanning a large stellar-parameter space (Carlssonetal.1994:Pavlenko&Maga- have shown that departures from LTE are generally small but significant at the required accuracy.," Previous non-LTE analyses spanning a large stellar-parameter space \citep{Carlsson94,Pavlenko96,Takeda05} have shown that departures from LTE are generally small but significant at the required accuracy."
 Barklemetal.(2003) applied quantum-mechanical calculations of cross sections for inelastic collisions with neutral hydrogen (Belyaev&Barklem2003:Croftetal.1999) in non-LTE analysis of the Sun and two metal-poor stars.," \citet{Barklem03b} applied quantum-mechanical calculations of cross sections for inelastic collisions with neutral hydrogen \citep{Belyaev03,Croft99} in non-LTE analysis of the Sun and two metal-poor stars."
 Their findings point to a negligible influence on the statistical equilibrium of lithium from collisional bound-bound transitions with hydrogen. but a significant influence from charge transfer reactions. specifically mutual neutralization and ion-pair production (Li?+H£5Li H).," Their findings point to a negligible influence on the statistical equilibrium of lithium from collisional bound-bound transitions with hydrogen, but a significant influence from charge transfer reactions, specifically mutual neutralization and ion-pair production $\rm Li^*+H \getsto Li^++H^-$ )."
 Otherinvestigations have relied on the classical Drawin recipes (Drawin1968).. as given by Steenbock&Holweger(1984) and Lambert(1993).. or the free electron model of Kaulakys(1985).. for estimates of collisions with neutral hydrogen.," Otherinvestigations have relied on the classical Drawin recipes \citep{Drawin68}, as given by \citet{Steenbock84} and \citet{Lambert93}, or the free electron model of \citet{Kaulakys85}, , for estimates of collisions with neutral hydrogen."
 None except Barklem et hhave included the influential charge transfer reaction., None except Barklem et have included the influential charge transfer reaction.
 We have extended the study by Barklemetal.(2003). to cover a large cool-star grid and calculate non-LTE abundance corrections for the lithium nnm (2s-2p) and nnm (2p-3d) lines., We have extended the study by \citet{Barklem03b} to cover a large cool-star grid and calculate non-LTE abundance corrections for the lithium nm (2s-2p) and nm (2p-3d) lines.
 We use the radiative-transfer code MULTI. version 2.3 (Carlsson1986.1992) to perform non-LTE calculations.," We use the radiative-transfer code MULTI, version 2.3 \citep{Carlsson86,Carlsson92} to perform non-LTE calculations."
 The model atom used includes the same 20 energy levels for neutral lithium as deseribed in Carlssonetal.(1994)... plus the Lill ground state.," The model atom used includes the same 20 energy levels for neutral lithium as described in \citet{Carlsson94}, plus the II ground state."
 The highest considered level in Lill has principle quantum number 2=9., The highest considered level in I has principle quantum number $n=9$.
 We have used TOPbase data (Peachetal.1988) for energy levels. oscillator strengths. and photo-ionization cross-sections for levels with orbital quantum number /< 3.," We have used TOPbase data \citep{Peach88} for energy levels, oscillator strengths, and photo-ionization cross-sections for levels with orbital quantum number $l\leq 3$ ."
 For the remaining levels. hydrogenie values are used.," For the remaining levels, hydrogenic values are used."
 For the resonance line at nnm. we adopt the oscillator strength. f=0.7468. as calculated by Yanetal.(1998)... and consider six hyperfine components in ‘Li (neglecting ?Li). with measured wavelengths givenby Sansonettietal.(1995).," For the resonance line at nm, we adopt the oscillator strength $f=0.7468$, as calculated by \citet{Yan98}, and consider six hyperfine components in $^7$ Li (neglecting $^6$ Li), with measured wavelengths givenby \citet{Sansonetti95}."
. For the subordinate line at nnm. we adopt f=0.6386. also determined by Yan et al. and account for the three fine-structure components with wavelengths determined by Lindgard&Nielson(1977).," For the subordinate line at nm, we adopt $f=0.6386$, also determined by Yan et al, and account for the three fine-structure components with wavelengths determined by \citet{Lindgard77}."
. For both lines. van-der-Waals-broadening parameters follow Anstee&O'Mara(1995) and Barklem&O'Mara(1997).," For both lines, van-der-Waals-broadening parameters follow \citet{Anstee95} and \citet{Barklem97}."
. Stark broadening is unimportant in the late-type atmospheres of interest here and is therefore neglected., Stark broadening is unimportant in the late-type atmospheres of interest here and is therefore neglected.
 Table | lists wavelenghts. oscillator strengths. and broadening data adopted for the nnm and nnm lines.," Table 1 lists wavelenghts, oscillator strengths, and broadening data adopted for the nm and nm lines."
 Cross-sections for collisional excitationby electrons are taken from Park(1971)and collisional ionization by electronsfrom Seaton(1962) as given by Allen (1976).., Cross-sections for collisional excitationby electrons are taken from \citet{Park71} and collisional ionization by electronsfrom \citet{Seaton62b} as given by \citet{Allen76}. .
 We add rate coefficients for excitation and de-excitation from collisions with neutral hydrogen atoms according to Belyaev&Barklem and Barklemetal. (2003).. as well as charge transfer," We add rate coefficients for excitation and de-excitation from collisions with neutral hydrogen atoms according to \citet{Belyaev03} and \citet{Barklem03b}, , as well as charge transfer"
in magnitudes (ο recshilts z>2 (also see Premacli&Martel(2005))).,in magnitudes to redshifts $z>2$ (also see \citet{premadimartel}) ).
 For the fiducial case of à SNAP-like SN survey. we expect more than a thousand SNe La in the range z=1.7—3 (see relsec:rates)) so treating lensing as an added dispersion is an excellent approximation. in contrast to cases where there might be only tens of sources in this range.," For the fiducial case of a SNAP-like SN survey, we expect more than a thousand SNe Ia in the range $z=1.7-3$ (see \\ref{sec:rates}) ) so treating lensing as an added dispersion is an excellent approximation, in contrast to cases where there might be only tens of sources in this range."
 We include lensing effects in this manner for the remainder of this article., We include lensing effects in this manner for the remainder of this article.
 Note that out to z=1.7. lensing has a «mall effect for a SNAP-like survey. with less than degradation of cosmological parameter estimation. for SN combined with the distance to the CAIB last scattering surface.," Note that out to $z=1.7$, lensing has a small effect for a SNAP-like survey, with less than degradation of cosmological parameter estimation, for SN combined with the distance to the CMB last scattering surface."
 In we consider the tails of the lensing amplification distribution where lensing can cause appreciable (de)magnilication. leading to possible confusion between SN types if categorized solely by Iuminositv.," In \\ref{sec:lenstype} we consider the tails of the lensing amplification distribution where lensing can cause appreciable (de)magnification, leading to possible confusion between SN types if categorized solely by luminosity."
 In the next section we examine cosmological bias. due to lensing in part but mostly from misestimation of redshift.," In the next section we examine cosmological bias, due to lensing in part but mostly from misestimation of redshift."
 The two kev quantities entering the distance probe are the fIux (as just discussed) aud the redshift., The two key quantities entering the distance probe are the flux (as just discussed) and the redshift.
 An important question is how much the redshift uncertaintv degrades the Lhibble diagram and cosmological parameter estimation., An important question is how much the redshift uncertainty degrades the Hubble diagram and cosmological parameter estimation.
 The great majority of the thousands of 5Ne discovered al 21.7. especially by. a survey dedicated to the follow up of zx1.1 events. will lack spectroscopic redshifts.," The great majority of the thousands of SNe discovered at $z>1.7$, especially by a survey dedicated to the follow up of $z\le1.7$ events, will lack spectroscopic redshifts."
 We examine the possibility of photometric redshifts in relsec:photoz:: here we investigate the effects in terms of a generic uncertainty oz., We examine the possibility of photometric redshifts in \\ref{sec:photoz}; ; here we investigate the effects in terms of a generic uncertainty $\delta z$.
 Propagating redshift uncertainties (through even a Fisher matrix approach is nontrivial due to correlations and dependencies., Propagating redshift uncertainties through even a Fisher matrix approach is nontrivial due to correlations and dependencies.
 A correct. technique would be to include the full covariance matrix (NauueeX Neouee} Of magnitude errors induced by redshift uncertainties.," A correct technique would be to include the full covariance matrix $N_{\rm source}\times 
N_{\rm source}$ ) of magnitude errors induced by redshift uncertainties."
 llowever. the standard practice in the literature citethint kim)) is to consider only uncorrelated. Gaussian. redshift bin centroid errors.," However, the standard practice in the literature \\citet{hutkim}) ) is to consider only uncorrelated, Gaussian, redshift bin centroid errors."
 While ]Iutereretal.(2004) implement this in terms of new fit parameters. to be marginalized over. we adopt (he approach of treating redshift uncertainties as an additional source of dispersion. approximating the [ull covariance matrix by diagonal entries.," While \citet{hutkim} implement this in terms of new fit parameters, to be marginalized over, we adopt the approach of treating redshift uncertainties as an additional source of dispersion, approximating the full covariance matrix by diagonal entries."
 In addition. though. we also consider redshift errors as a svstematic effect. and. as a function of cosmological parameter bias tolerated. bound Che allowed correlated error.," In addition, though, we also consider redshift errors as a systematic effect and, as a function of cosmological parameter bias tolerated, bound the allowed correlated error."
 In the statistical approach. the uncertaintv (Om/à0z)0z is added in quadrature with the other magnitude errors.," In the statistical approach, the uncertainty $(\partial m/\partial z)
\,\delta z$ is added in quadrature with the other magnitude errors."
 The typical amplitude of the distance modulus contribution to the source magnitude uncertaintyis 1.7-0.8 over the range 2=1.7— 3., The typical amplitude of the distance modulus contribution to the source magnitude uncertaintyis 1.7-0.8 over the range $z=1.7-3$ .
 This is not the full story.," This is not the full story,"
"the events and the 1FGL catalog in closer detail, we notice that the sample includes one pulsar, two unidentified sources, and eight AGN.","the events and the 1FGL catalog in closer detail, we notice that the sample includes one pulsar, two unidentified sources, and eight AGN."
 Four of the AGN have measured well beyond the GZK horizon with redshifts as high as z=1.843., Four of the AGN have measured well beyond the GZK horizon with redshifts as high as $z = 1.843$.
" The lowest redshifts correspond to Centaurus A (z=0.002) and NGC 4945 (z=0.002), as we shall discuss later."," The lowest redshifts correspond to Centaurus A $z = 0.002$ ) and NGC 4945 $z = 0.002$ ), as we shall discuss later."
 The remaining two AGN associations have no spectroscopic redshifts., The remaining two AGN associations have no spectroscopic redshifts.
 One could argue that the aforementioned AGN without spectroscopic redshifts could potentially form a parent population for UHECRs., One could argue that the aforementioned AGN without spectroscopic redshifts could potentially form a parent population for UHECRs.
" If it was, these AGN would most likely lie at distances closer than 100 Mpc (Abrahametal. 2008a)."," If it was, these AGN would most likely lie at distances closer than 100 Mpc \citep{abraham5}."
". Optical imaging studies of BL Lac reveal that the host galaxies of nearby AGN with jets pointed in our direction tend to cluster around a mean absolute magnitude in the R band equivalent to Mg=—22.8 (Sbarufatti,Treves&Falomo 2005).", Optical imaging studies of BL Lac reveal that the host galaxies of nearby AGN with jets pointed in our direction tend to cluster around a mean absolute magnitude in the $R$ band equivalent to $M_{R} = -22.8$ \citep{sba}.
". For distances closer than 100 Mpc, this would correspond to an apparent R-band magnitude R—12.2 or brighter."," For distances closer than 100 Mpc, this would correspond to an apparent $R$ -band magnitude $R = 12.2$ or brighter."
 A quick analysis of optical Digitized Sky Survey images centered on the positions of the two matched AGN without spectroscopic redshifts reveals no optical sources brighter than R=12.2 at the derived AGN positions (Abdoetal.2010b).," A quick analysis of optical Digitized Sky Survey images centered on the positions of the two matched AGN without spectroscopic redshifts reveals no optical sources brighter than $R =
12.2$ at the derived AGN positions \citep{abdo2}."
". Therefore, it is unlikely that any of the paired UHECRs/AGN are located within the hypothesized GZK horizon at 100 Mpc."," Therefore, it is unlikely that any of the paired UHECRs/AGN are located within the hypothesized GZK horizon at 100 Mpc."
" Interestingly, we do recover the two UHECRs previously associated with the radio galaxy Centaurus A (Abrahametal.20088) and one UHECR consistent with the position of the starburst/Seyfert 2 NGC 4945 (Moskalenkoetal."," Interestingly, we do recover the two UHECRs previously associated with the radio galaxy Centaurus A \citep{abraham5} and one UHECR consistent with the position of the starburst/Seyfert 2 NGC 4945 \citep{moskalenko}."
 2009).. Both Centaurus A (1FGL J1325.6—4300) and NGC 4945 (1FGL J1305.4—4928) are detected byFermi., Both Centaurus A (1FGL $-$ 4300) and NGC 4945 (1FGL $-$ 4928) are detected by.
". However, NGC 4945 only stands out as one of the two Seyfert 2 galaxies in the 1FGL catalog rather than for its qualifications as a UHECR accelerator (Abdoetal.2010b)."," However, NGC 4945 only stands out as one of the two Seyfert 2 galaxies in the 1FGL catalog rather than for its qualifications as a UHECR accelerator \citep{abdo2}."
". On the other hand, radio galaxies such as Centaurus A could potentially account for the acceleration of some UHECRs to energies >55 EeV (Dermeretal.2009)."," On the other hand, radio galaxies such as Centaurus A could potentially account for the acceleration of some UHECRs to energies $> 55$ EeV \citep{dermer}."
". As result, it is important to continue to explore the connectiona (if any) between UHECRs and Centaurus A. However, it is troubling that 24 UHECR events detected byAuger remain without an apparent ggamma-ray counterpart within 100 Mpc (we dub these “orphan” UHECRs)."," As a result, it is important to continue to explore the connection (if any) between UHECRs and Centaurus A. However, it is troubling that 24 UHECR events detected by remain without an apparent gamma-ray counterpart within 100 Mpc (we dub these “orphan” UHECRs)."
" In fact, there seems to be an apparent dearth of nearby AGN detected in gamma rays at distances less than 100 Mpc (z< 0.025)."," In fact, there seems to be an apparent dearth of nearby AGN detected in gamma rays at distances less than 100 Mpc $z \leq 0.025$ )."
" To see this more clearly, Figure 2 shows the distribution for AAGN in the redshift range z=0—0.1 (Abdoetal. 2010b)."," To see this more clearly, Figure \ref{figure2} shows the distribution for AGN in the redshift range $z=0-0.1$ \citep{abdo2}. ."
". In total, there are only five objects with z<0.025 that would be accessible to the southern site namely NGC 253 (z= 0.001), NGC 4945 (z= 0.002), Centaurus A (z= 0.002), M 87 (z= 0.004), and ESO 323-G77 (z= 0.015)."," In total, there are only five objects with $z \leq 0.025$ that would be accessible to the southern site namely NGC 253 $z= 0.001$ ), NGC 4945 $z = 0.002$ ), Centaurus A $z = 0.002$ ), M 87 $z = 0.004$ ), and ESO 323-G77 $z = 0.015$ )."
 NGC 4945 (being generous) and Centaurus A could potentially account for 3 of the 27 UHECR events as discussed above., NGC 4945 (being generous) and Centaurus A could potentially account for 3 of the 27 UHECR events as discussed above.
" However, we remark that there are no additional identified ssources in the 100 MeV-100 GeV energy band to account for the bulk of UHECRs seen byAuger so far."," However, we remark that there are no additional identified sources in the 100 MeV-100 GeV energy band to account for the bulk of UHECRs seen by so far."
" It is possible that other AGN remain unidentified in the 1FGL catalog, as implied by the observed north-south anisotropy of the associated sources, which might reflect the incompleteness of the existing AGN catalogs in other bands (Abdoetal.2010b)."," It is possible that other AGN remain unidentified in the 1FGL catalog, as implied by the observed north-south anisotropy of the associated sources, which might reflect the incompleteness of the existing AGN catalogs in other bands \citep{abdo2}."
". However, the remainder is expected to lie in the lower end of the gamma-ray flux distribution, and therefore in principle not represent the most efficient particle accelerators."," However, the remainder is expected to lie in the lower end of the gamma-ray flux distribution, and therefore in principle not represent the most efficient particle accelerators."
" In addition, some AGN are known to be highly variable, therefore we cannot fully discard that some sources have escaped ddetection and possible correlation with UHECRs because of a low state."," In addition, some AGN are known to be highly variable, therefore we cannot fully discard that some sources have escaped detection and possible correlation with UHECRs because of a low state."
" Accordingly, one must admit one or all of the following possibilities: 1) if UHECRs are truly accelerated within the current sample of detected ssources, their trajectories must experience deviations greater than 3.1° from their actual astrophysical origin, making the identification of UHECRs nearly intractable."," Accordingly, one must admit one or all of the following possibilities: 1) if UHECRs are truly accelerated within the current sample of detected sources, their trajectories must experience deviations greater than $3.1\!^\circ$ from their actual astrophysical origin, making the identification of UHECRs nearly intractable."
" In particular, the absence of astrophysical counterparts could be well justified if UHECRs are dominated by heavy nuclei including iron since the deviation scales linearly the atomic number Z, 2) there could potentially exist an undetected population of nearby AGN -- the so called proton blazars?"," In particular, the absence of astrophysical counterparts could be well justified if UHECRs are dominated by heavy nuclei including iron since the deviation scales linearly the atomic number $Z$, 2) there could potentially exist an undetected population of nearby AGN – the so called proton blazars?"
" (Mannheim1993) -- that only emit in a very narrow band above 100 GeV to be discovered by future VHE experiments such as the Cherenkov Telescope Array (CTA) or the Advanced GammaImaging System (AGIS), and 3) the absence of obvious particle accelerators in the MeV- energy band leaves some room for the exploration of"," \citep{mannheim} – that only emit in a very narrow band above 100 GeV to be discovered by future VHE experiments such as the Cherenkov Telescope Array (CTA) or the Advanced GammaImaging System (AGIS), and 3) the absence of obvious particle accelerators in the MeV-GeV energy band leaves some room for the exploration of"
constructed from (he resulis of the UCL Cosmic Dust Experiment (Perry&Price2003:Crei¢han.Perry&Price2006) and we compute LR emission spectra of Ils expected in interstellar clouds.,"constructed from the results of the UCL Cosmic Dust Experiment \citep{PD03,
C06} and we compute IR emission spectra of $_2$ expected in interstellar clouds."
 The Is formation pumping model is coupled with an accurate description ol both radiative transfer and chemistry in stratified clark clouds., The $_2$ formation pumping model is coupled with an accurate description of both radiative transfer and chemistry in stratified dark clouds.
 Section 2 contains a description of the formation pumping excitation model emploved in the caleulations., Section 2 contains a description of the formation pumping excitation model employed in the calculations.
 In section 3. we summarize the method and the procedures followed using dillerent formation pumping models to caleulate the complete vibrational distribution of Hs formed on dust erains.," In Section 3, we summarize the method and the procedures followed using different formation pumping models to calculate the complete vibrational distribution of $_2$ formed on dust grains."
 I» emission spectra are presented in Section 4., $_2$ emission spectra are presented in Section 4.
 We discuss observational implications and we present our conclusions in Section 5., We discuss observational implications and we present our conclusions in Section 5.
 The vibrational excitation of molecular hydrogen. desorbed [rom surlaces can be measured experimentally., The vibrational excitation of molecular hydrogen desorbed from surfaces can be measured experimentally.
 A quantitative partition of the excitation between vibrational states has been investigated by the UCL Cosmic Dust Experiment. which probes the distribution ol the rotational states within each vibrational manifold (Perry&Price2003:Creighan.Perry&Price 2006).," A quantitative partition of the excitation between vibrational states has been investigated by the UCL Cosmic Dust Experiment, which probes the distribution of the rotational states within each vibrational manifold \citep{PD03,C06}."
. The experiment studies (he formation of molecular hydrogen. primarily ILD. on a oriented pyroliic graphite (IIOPG) surface under ταΙσ vacuum. following continuous irradiation of the surface by H and D atoms.," The experiment studies the formation of molecular hydrogen, primarily HD, on a oriented pyrolitic graphite (HOPG) surface under ultrahigh vacuum, following continuous irradiation of the surface by H and D atoms."
 The nascent IID (or Hs) molecule will desorb in vibrational states (0.7) of the ground electronic state.," The nascent HD (or $_2$ ) molecule will desorb in vibrational states $(v, J)$ of the ground electronic state."
 IIvdrogen molecules are selectively ionised using induced resonance enhanced photon ionisation spectroscopy., Hydrogen molecules are selectively ionised using induced resonance enhanced photon ionisation spectroscopy.
 The relative populations of the vibrational states are then derived [rom ion vields., The relative populations of the vibrational states are then derived from ion yields.
 The experiment probes ILD preferentially to Ils (or Do) because there is a significant amount of undissociated I» originating from the atom source., The experiment probes HD preferentially to $_2$ (or $_2$ ) because there is a significant amount of undissociated $_2$ originating from the atom source.
" As studies of e=1 and 2 [or the isotopic species revealed very similar [lux densities aud rotational distributions for both HD ancl Ils molecules (Creighan.Perry&Price2006).. we extrapolate the ILD v=3-7 data to obtain the vibrational distribution of nascent Πα,"," As studies of $v = 1$ and 2 for the isotopic species revealed very similar flux densities and rotational distributions for both HD and $_2$ molecules \citep{C06}, we extrapolate the HD $v = 3 - 7$ data to obtain the vibrational distribution of nascent $_2$."
 However. nascent TID in the v=0 state could not be detected above the signal [rom background gas in the vacuum chamber (see Creighan.Perry&Price 2006)).," However, nascent HD in the $v = 0$ state could not be detected above the signal from background gas in the vacuum chamber (see \citealt{C06}) )."
 As a modest estimate of the internal excitation of IIl». we set the e=0 populations to be equal to the populations of the —1 states.," As a modest estimate of the internal excitation of $_2$, we set the $v = 0$ populations to be equal to the populations of the $v = 1$ states."
 Since (his is an arbitrary choice. IH» emission spectra have been generated with the e—0 states both signilicantlv more and less populated (han expected.," Since this is an arbitrary choice, $_2$ emission spectra have been generated with the $v = 0$ states both significantly more and less populated than expected."
 The sensitivity of the Ils spectra to the error in the population of the e=0 states is discussed further in Section 4., The sensitivity of the $_2$ spectra to the error in the population of the $v = 0$ states is discussed further in Section 4.
(Piran2000:Meészaros2002:Waxman2003) (Mészáros.Rees&Papathanassiou1,"\citep{Piran00,Meszaros02,Waxman03}
 \citep{MRP94}. \citep{BnM76},"
"994).. (Blandford&Mchee1976).. Pp=(10E/16zmmyc?)!?r7*7, E » T (Waxian2003): Zegp Ty Dj T. Τον)2Τι.r zr/4D7c (Waxman1997).. T (Waxman2003)E=10?E.4,»—0.1».pom7. 1+z z. "," $\Gamma_{\rm BM}=(17E/16\pi n m_p c^2)^{1/2}r^{-3/2}$ $E$ $n$ $T$ \citep{Waxman03}: $T_{\rm GRB}$ $T_{\Gamma}$ $\Gamma_{\rm BM}$ $\Gamma_i$ $\Gamma_{\rm BM}(T_{\Gamma})=\Gamma_i$$r$ $\approx r/4\Gamma^2 c$ \citep{W97}, $T$ \citep{Waxman03}
$E = 10^{53} E_{53}$$ n = 0.1 n_{-1} {\rm\, cm^{-3}}$ $1+z$ $z$ "
7 are the gas fraction ancl velocity dispersion at small racii. respectively).,"$\sigma$ are the gas fraction and velocity dispersion at small radii, respectively)."
 Although the AGN wind feedback has little elect on the early time star formation and DII accretion. Figure 1. shows that after final coalescence both the BIT accretion rate and the star formation rate decrease significantly. more rapidly with the inclusion of the AGN wind.," Although the AGN wind feedback has little effect on the early time star formation and BH accretion, Figure \ref{fig:mdotfid} shows that after final coalescence both the BH accretion rate and the star formation rate decrease significantly more rapidly with the inclusion of the AGN wind."
 The ellects of the Buls feedback are maximized at and after final coalescence because this is when the BIL reaches its final mass (to within a [actor⋅ of⋅ à lew)⋅ and when the DII accretion. rate is. largest., The effects of the BHs feedback are maximized at and after final coalescence because this is when the BH reaches its final mass (to within a factor of a few) and when the BH accretion rate is largest.
 Thus. the dynamics at and after final coalescence is the most sensitive to the details of the feedback physics.," Thus, the dynamics at and after final coalescence is the most sensitive to the details of the feedback physics."
 oPVo assess whydue the late-time; accretionqui and star ⋅⋅⋅formation: ⋜⊔⋅∢⊾⊳∖⊔↓≻↓≻↓⋅∢⋅≱∖≱∖∢⋅∠⇂∣⋡∙∖⇁↥↓↥∢⋅⊲↓↓⊔⇍⇂⊔≱∖⊲↓∪⊔∪⇂⋅⇂↓↕∢⋅⇀∖≺∶↓∖⊽∖∖⋰↓⊔∠⇂⊳↓⊲↝↓⋏∙≟⊔↓⋅⋖⊾⇉ shows the gas mass fraction (normalized to the gas mass at the start of the simulation) within 3 kpc of the DII as a function⋅. of D.time. for⋅ the same three simulations. as in. . qu . . . ↓⊲↓⋏∙≟⊔↓⋅∢⋅↓↦⊥," To assess why the late-time accretion and star formation are suppressed by the inclusion of the AGN wind, Figure \ref{fig:m3kpc} shows the gas mass fraction (normalized to the gas mass at the start of the simulation) within 3 kpc of the BH as a function of time, for the same three simulations as in Figure \ref{fig:mdotfid}."
↓⊔⇂⋜↧↓↓⊔↿⋖⊾↓⋅⋯⇍⊔∪⊔⊳∖∠∐⋅↓∖⇁⋖⋅⋏∙≟⋜↧⊳∖↓⊔↥∪↿⇂↥∢⋅≼∙∢⋅⊔↿⋖⋅↓⋅∪⇂ the two galaxies: (or the merged. remnant) after⋅ the. first⋅ close passage (/~0.5 Civr) and at final coalescence (f~1.0 Cyr). maintaining a large nuclear gas fraction in spite of the intense star formation.," Tidal interactions drive gas into the center of the two galaxies (or the merged remnant) after the first close passage $t \sim 0.5$ Gyr) and at final coalescence $t \sim 1.7$ Gyr), maintaining a large nuclear gas fraction in spite of the intense star formation."
 The resulting increase in the 1211 accretion rate (seen also in Fig. 1)), The resulting increase in the BH accretion rate (seen also in Fig. \ref{fig:mdotfid}) )
 significantly increases the streneth of the AGN feedback., significantly increases the strength of the AGN feedback.
 In the simulation including AGN winds. this feedback elicienthy removes material from the central regions of the galaxy. reducing the amount of gas inside the central 3 kpe of the merged system by a factor of ~LO relative to the simulation without explicit AGN winds.," In the simulation including AGN winds, this feedback efficiently removes material from the central regions of the galaxy, reducing the amount of gas inside the central 3 kpc of the merged system by a factor of $\sim 10$ relative to the simulation without explicit AGN winds."
" """"Igureiure :3 showstows LIA.ages οἱ/ theae proὉ]οσιοα. δαgas density--ps [or the fiducial simulation (feff) and the corresponding simulation:. without: the AGN MEEwind (righ/) at |=1371 BNvr. shortly alter∙ final∙ coalescence."," Figure \ref{fig:densplot} shows images of the projected gas density for the fiducial simulation ) and the corresponding simulation without the AGN wind ) at $t=1.71$ Gyr, shortly after final coalescence."
 The nnimages are 280 kpc on à side and show a roughly edee-on view of the orbital plane of the two galaxies., The images are 280 kpc on a side and show a roughly edge-on view of the orbital plane of the two galaxies.
 While there is material at large distances in the simulation without the AGN wind. these features‘ ‘are the tical “ostails 5generated‘ curing5 the merger© ‘ ‘ ⊳⋯∠⇂⊔↓∪≱∖↥∪⇂∎↓↥⊲↓⊳∖↓↥↓⊳↧↿⋖⊾↓⋅↕⊳↧↓↕≻⊔∢⋅⊳⊔⋅↿↓∐⋅∪↓⋅," While there is material at large distances in the simulation without the AGN wind, these features are the tidal tails generated during the merger and most of this material is near the orbital plane."
∣⋡⊲⊔⊳↧↓≻↓⊳⋯⋖⋅⊳↾∐⋯⊳∖| . m.. ∪⊔↓⋅⊳∖⊲↓⊔↓≻↓∢⊾↕↓↕↓≻↓⋖⋅↓↕↓∢⋅↓↕↿⊳↧⇂↕⋖≱↓↥∢⋟⇂⋅↓⋅⊳↧∠⇂↕⊳↧⇂↕⋖≱↓↥≻↓⋅∢⊾⊳∖⊳∖⊔↓⋅∢⋅⇂∎∢⋅⋖⋅∠∐⋡⊳⊔⇍↥⊳ow 1 .. is. not ellicientpe at unbinding gas [rom the galaxy. though it is very elfective at regulating the &rowth of the DII itself (DOM).," Thus our simple implementation of radiation pressure feedback is not efficient at unbinding gas from the galaxy, though it is very effective at regulating the growth of the BH itself (DQM)."
 By contrast. in the simulation with the AGN wind. Figure 3. (left panel) shows that there is significantly more material blown out of the galaxy. especially in the directions," By contrast, in the simulation with the AGN wind, Figure \ref{fig:densplot} (left panel) shows that there is significantly more material blown out of the galaxy, especially in the directions"
,$1\%$.
 50 Figure2: Dependenceofthe pa," However we still have to rely on extrapolations, and this produces some uncertainty, which may be systematic."
rtialwave functi, An interesting problem appears when $\xi=M_H/M_W > 2$.
onal deter, The nondiagonal terms $V_{34}=V_{43}$ behave as $\exp(-M_W r)$.
minant on the eutoll:large ," So if $f^{\alpha-}_3(r,\nu^2)\simeq exp(-\kappa_W r)$ it contributes with a behaviour $\exp(-(\kappa_W +M_W)r$ to $f^{\alpha-}_4(r,\nu^2)$ ."
"A: we displaythe quantity J,(0.", Now $f^{\alpha-}_4$ is supposed to behave as $\exp(-\kappa_H r)$ .
M) asa functionof |. The so, The cross term dominates over this behaviour for $\nu^2<M_H^2(M_H^2-4M_W^2)/4M_W^2$.
lidlines are theresu, This defines a real interval for $\nu$ if $M_H>2M_W$ .
ltsofGreen’s function approach., It found indeed that for $M_H>2M_W$ and small $\nu^2$ some of the functions $h^{\alpha-}_i$ increase exponentially with the expected behaviour.
 the svinhbols are those," One finds (numerically) that these contributions cancel in the determinant, both methods still produceconsistent results up to $\xi=M_H/M_W\simeq 2.5$."
ofthe Gel lancd-Yaglom method: squares:vy= 0: ," However, this cancellation is delicate numerically and for larger values $\xi$ the numerical procedure breaks down, the results become inconsistent."
diamonds: nv= 1. circles:n—2: triangles:η =5.crosses:n = 10.," So one would have to find a suitable modificationof the numerical procedure in order to maintain numerical reliability, herewelimit ourselves to the range$\xi < 2.5$ ."
The physical. conditions of the Intergalactic Medium (IGMD) have been under ciscussion ever since its existence was postulated.,The physical conditions of the Intergalactic Medium (IGM) have been under discussion ever since its existence was postulated.
 “Phe absence of strong absorption in the spectra of high-redshift objects showed that any homogeneous LGAL must be either. highly ionized or extremely tenuous (Gunn Peterson 1965. Bahcall Salpeter 1965).," The absence of strong absorption in the spectra of high-redshift objects showed that any homogeneous IGM must be either highly ionized or extremely tenuous (Gunn Peterson 1965, Bahcall Salpeter 1965)."
 Nevertheless the LGAL density and temperature have been poorly constrained for many vears., Nevertheless the IGM density and temperature have been poorly constrained for many years.
 oth a hot collisionally ionizecl LGOAL heated. by energy input due to supernovae ancl a warm intergalactic niccdium heated by photoionization have been discussed. extensively in the literature., Both a hot collisionally ionized IGM heated by energy input due to supernovae and a warm intergalactic medium heated by photoionization have been discussed extensively in the literature.
 Recent progress in numerical simulations including the relevant gas physics suggests that the low column clensity. forest. is due mainly {ο density luctuations of moderate amplitude in a warm photolonize hase of the LGAL (Cen et al., Recent progress in numerical simulations including the relevant gas physics suggests that the low column density forest is due mainly to density fluctuations of moderate amplitude in a warm photoionized phase of the IGM (Cen et al.
 1994: Petitjean. Müccket Μαιος 1995: Zhang. Anninos Norman 1995: LHernquis et al.," 1994; Petitjean, Müccket Kates 1995; Zhang, Anninos Norman 1995; Hernquist et al."
 1996: Miralda-Escudé et al., 1996; Miralda-Escudé et al.
 1996: but see also Di. Borrner Chu 1992 for carly analvtical predictions).," 1996; but see also Bi, Börrner Chu 1992 for early analytical predictions)."
 The ρανο content in this warm phase must be similar to tha »ecdieted by cosmic nucleosynthesis to produce the observec mean absorption in the spectra of high-redshift quasars., The baryon content in this warm phase must be similar to that predicted by cosmic nucleosynthesis to produce the observed mean absorption in the spectra of high-redshift quasars.
 A redshifts larger than two most of all barvons are therefore oobably in such a warm phase of the IGM and little room is left for an additional hot phase. (, At redshifts larger than two most of all baryons are therefore probably in such a warm phase of the IGM and little room is left for an additional hot phase. (
Rauch Llachnel 1995. Rauch οἱ al.,"Rauch Haehnelt 1995, Rauch et al."
 1997. Weinberg et al.," 1997, Weinberg et al."
 1997)., 1997).
 As pointe out bv a number of authors the recombination time scale of the LGAL becomes longer than the Hubble time at a redshift of about. five., As pointed out by a number of authors the recombination time scale of the IGM becomes longer than the Hubble time at a redshift of about five.
 As a result. the temperature of the warm LGAL has some memory for when and how it was (re-Jionized (Miralda-Escudé Rees 1994. Meiksin 1994. Hui CGnedin 1997).," As a result, the temperature of the warm IGM has some memory for when and how it was (re-)ionized (Miralda-Escudé Rees 1994, Meiksin 1994, Hui Gnedin 1997)."
 We have used a set of Biydrodyvnamical SPL simulations for different plausible reionization histories to demonstrate that these do indeed. result. in significant anc observable cdillerences in the thermal history of the IGM., We have used a set of hydrodynamical SPH simulations for different plausible reionization histories to demonstrate that these do indeed result in significant and observable differences in the thermal history of the IGM.
 In this letter we investigate the observable consequences on the Doppler parameter distribution of absorption lines anc speculate on the possibility to discriminate between dilleren reionization histories., In this letter we investigate the observable consequences on the Doppler parameter distribution of absorption lines and speculate on the possibility to discriminate between different reionization histories.
 We have calculated the thermal history of an IGA of primordial composition within a cosmological contex using SPILL simulations performed. with GRAPESPIL (see Steinmetz 1996 for a detailed. ceseription of the numerica techniques and Hauch. Haehnelt Steinmetz 1997. for a description of their. absorption properties)," We have calculated the thermal history of an IGM of primordial composition within a cosmological context using SPH simulations performed with GRAPESPH (see Steinmetz 1996 for a detailed description of the numerical techniques and Rauch, Haehnelt Steinmetz 1997 for a description of their absorption properties)."
 The code includes the relevant cooling and. heating processes anc follows self-consistenthy the non-equilibrium evolution of the barvonie species (Ll. LE. He. He. . and e ).," The code includes the relevant cooling and heating processes and follows self-consistently the non-equilibrium evolution of the baryonic species (H, $^{+}$ , He, $^{+}$, $^{++}$, and $e^-$ )."
 For the low-density regime relevant. in this paper the therma evolution of the LGAL israther simple. (Miralda-IZscudé, For the low-density regime relevant in this paper the thermal evolution of the IGM israther simple (Miralda-Escudé
FSRQ are the AGN class showing significantly harder spectral slopes. when compared io DLRG and SSRQ.,"FSRQ are the AGN class showing significantly harder spectral slopes, when compared to BLRG and SSRQ."
 According to the Ixolmogorov-Smirnov test. (he probability that the analvzed classes (FSRQ vs DLRG and FSRQ vs SSRQ) are drawn from the same distribution is verv small (κο—5xI0* and Pay=107. respectively).," According to the Kolmogorov-Smirnov test, the probability that the analyzed classes (FSRQ vs BLRG and FSRQ vs SSRQ) are drawn from the same distribution is very small $_{KS}=5\times10^{-3}$ and $_{KS}=10^{-3}$, respectively)."
 This is not a new result., This is not a new result.
 As FSRQ have (he larger core dominance value (110).>) a strong hard non-thermal contribution is expected to affect (heir X-ray continuum.," As FSRQ have the larger core dominance value $R$ ), a strong hard non-thermal contribution is expected to affect their X-ray continuum."
 What is instead interesting to note is that. taking into account the whole sample. the X-ray spectral slope does not correlate with J.," What is instead interesting to note is that, taking into account the whole sample, the X-ray spectral slope does not correlate with $R$."
 A non-parametric statistical test. (IXendall's Lau test) results in a correlation coefficient of 7;=—0.34. with a signilieance value ay=0.1. above the 0.05 limit that we have assumed (o accept (wo quantities as correlated.," A non-parametric statistical test (Kendall's tau test) results in a correlation coefficient of $r_{\rm s}=-0.34$ , with a significance value $\alpha_{\rm K}=0.1$, above the 0.05 limit that we have assumed to accept two quantities as correlated."
 The search for a correlation between the high οποίον eutoff and. the beaming indicator also gives a negalive result., The search for a correlation between the high energy cutoff and the beaming indicator also gives a negative result.
 In order (o assess (he significance of the correlation in the presence of lower limits. we applied the generalized IXendalls (an test in the ASURY package to the total sample.," In order to assess the significance of the correlation in the presence of lower limits, we applied the generalized Kendall's tau test in the ASURV package \citep{iso86} to the total sample."
 The resulüng significance. cTa=(Q.2. is again significantly hieher than the acceptance threshold.," The resulting significance, $\alpha_{\rm K}^{\rm ASURV}=0.2$, is again significantly higher than the acceptance threshold."
 For the reprocessing leatures. iron line ancl reflection component. the situation is slightly different.," For the reprocessing features, iron line and reflection component, the situation is slightly different."
 There is an apparent decreasing (rend of the equivalent width and the reflection component against the radio core dominance., There is an apparent decreasing trend of the equivalent width and the reflection component against the radio core dominance.
 The Ixendall's tau test returned significance values of atPY=0.07 and ad!=0.06 for the EW vs R and Ref vs R correlations. respectively. in both cases only slightly above the 0.05 acceptance (hresholcl.," The Kendall's tau test returned significance values of $\alpha_{\rm K}^{\rm ASURV}=0.07$ and $\alpha_{\rm K}^{\rm ASURV}=0.06$ for the EW vs $R$ and Ref vs $R$ correlations, respectively, in both cases only slightly above the 0.05 acceptance threshold."
 Figure 2 shows the EW against 2 lor FSRQ. BLRG and SSRQ.," Figure \ref{fig:EWvsR} shows the EW against $R$ for FSRQ, BLRG and SSRQ."
" Lowever. if we exclude the hiehlv absorbed DLRG 3C 445 (Vy,>1051 7). the already weak correlations disappear (αμRY= 0.2). casting farther doubts on the link between iron line EW ancl beaming."," However, if we exclude the highly absorbed BLRG 3C 445 $N_{\rm H}>10^{23}$ $^{-2}$ ), the already weak correlations disappear $\alpha_{\rm K}^{\rm ASURV}>0.2$ ), casting farther doubts on the link between iron line EW and beaming."
 The intense iron line (EW~170 eV) observed in 3C 445. instead of being the result of a low beaming contamination. could be due to the contribution of the transmission through a line-of-sight absorber. as discussed above.," The intense iron line $\sim170$ eV) observed in 3C 445, instead of being the result of a low beaming contamination, could be due to the contribution of the transmission through a line-of-sight absorber, as discussed above."
 The intrinsic column densitv is the only spectral parameter correlating with the radio core dominance., The intrinsic column density is the only spectral parameter correlating with the radio core dominance.
 The amount of gas along the line of sight decreases with increasing f (Figure 3))., The amount of gas along the line of sight decreases with increasing $R$ (Figure \ref{fig:NHvsR}) ).
 The generalized Ixendall's tau test gives a;ASURYN—(.02., The generalized Kendall's tau test gives $\alpha_{\rm K}^{\rm ASURV}=0.02$.
 If we exclude all upper limit values. the significance further improves to ay=0.01.," If we exclude all upper limit values, the significance further improves to $\alpha_{\rm K}=0.01$."
 This result is coherent with the unified model scenarios. which predict (he nuclear source to be more obscured and the jet enission stronelv de-amplilied at large viewing angles.," This result is coherent with the unified model scenarios, which predict the nuclear source to be more obscured and the jet emission strongly de-amplified at large viewing angles."
 Note that this correlation indirectly confirms the robustness of Ras indicator ofthe beaming orientation., Note that this correlation indirectly confirms the robustness of $R$ as indicator ofthe beaming orientation.
 The conclusion is that the lack of a strong; correlation between £2 aud (he X.ray spectral, The conclusion is that the lack of a strong correlation between $R$ and the X–ray spectral
For many astronomical measurements itis necessary {ο understand aud precisely measure the attenuation of lisht by Earth's atmosphere.,For many astronomical measurements it is necessary to understand and precisely measure the attenuation of light by Earth's atmosphere.
 While attempts have been made to cdo (his in an absolute sense (?7).. in general astronomical measurements al optical and NIB wavelengths are funcamentally dilferential. measuring [hix relative to either a standard star or comparison stars within a given [field of view.," While attempts have been made to do this in an absolute sense \citep{hayes1975}, in general astronomical measurements at optical and NIR wavelengths are fundamentally differential, measuring flux relative to either a standard star or comparison stars within a given field of view."
 As shown in Figure L.. across the optical and infrared Ravleigh scattering by molecules. Mie scattering by aerosols. and line absorption by molecules leads to atmospheric (ransniission (hat is a rapidly varving function of wavelength and can also vary rapiclly in (nme (7).," As shown in Figure \ref{fig8}, across the optical and infrared Rayleigh scattering by molecules, Mie scattering by aerosols, and line absorption by molecules leads to atmospheric transmission that is a rapidly varying function of wavelength and can also vary rapidly in time \citep{stubbs2007}."
 Massive miaging survevs like (he Large Synoptic Survey ((LSST). Dark Energv. Survev((DES).Pan-STARRS®.. and Ivper seek to produce broadband photometric measurements across the UV and near infrared (NUR) that are stable and uniform to better than 14. over periods of vears.," Massive imaging surveys like the Large Synoptic Survey (LSST), Dark Energy (DES), and Hyper seek to produce broadband photometric measurements across the UV and near infrared (NIR) that are stable and uniform to better than $1\%$ over periods of years."
 Even if the response of the telescope to monochromatic light of a given intensity is known verv precisely (e.g. ?.. 2)). global calibration of photometric measurements requires correcting flux for atmospheric variations.," Even if the response of the telescope to monochromatic light of a given intensity is known very precisely (e.g. \citealt{stubbs2006}, \citealt{stubbs2010}) ), global calibration of photometric measurements requires correcting flux for atmospheric variations."
 Similarly. precise radial velocity measurements al NI wavelengths require a detailed understanding of atmospheric (rausmission so (hal subtle Doppler shifts in stellar spectral leatures can be separated ου changes in telluric absorption leatures.," Similarly, precise radial velocity measurements at NIR wavelengths require a detailed understanding of atmospheric transmission so that subtle Doppler shifts in stellar spectral features can be separated from changes in telluric absorption features."
" The flux in ADU measured by a single-element detector from an astronomical source. fF. can be expressed as a function of the waveleneth-dependent properties of the telescope. ablmosphere. and source as where <1 is the elfective collecting area of the telescope. AT is the exposure time. 7(A) is (he (ransmiission of the atmosphere as a Iunction of wavelength. S(A) is (he response function ol the telescope in units of ADU/photon (including the elleets of optics. fillers. and detectors). and Fy is the flux density of the source in units of photons ten ""ο +."," The flux in ADU measured by a single-element detector from an astronomical source, $\digamma$, can be expressed as a function of the wavelength-dependent properties of the telescope, atmosphere, and source as where $A$ is the effective collecting area of the telescope, $\Delta T$ is the exposure time, $T(\lambda)$ is the transmission of the atmosphere as a function of wavelength, $S(\lambda)$ is the response function of the telescope in units of ADU/photon (including the effects of optics, filters, and detectors), and $F_{\lambda}$ is the flux density of the source in units of photons $^{-1}$ $^{-2}$ $^{-1}$."
" Historically. for broadband photometry the band-ntegrated value of the product 5(A)7(À) is measured Lor a given night using a combination of observations of a uniformly illuminated screen (the “flat field”) ancl observations of photometric ""standard stars."," Historically, for broadband photometry the band-integrated value of the product $S(\lambda)T(\lambda)$ is measured for a given night using a combination of observations of a uniformly illuminated screen (the “flat field”) and observations of photometric “standard” stars."
 Using this methodology. it, Using this methodology it
statistics (2 BSS and 0 HB stars).,statistics (2 BSS and 0 HB stars).
" In this respect Arp 2 seems more similar to old open clusters, like M 67 (Mathieu Gheller 2009), than to genuine globulars."," In this respect Arp 2 seems more similar to old open clusters, like M 67 (Mathieu Gheller 2009), than to genuine globulars."
" The higher concentration of BSS in the cluster internal region with respect to evolved stars we find is not new, and has been found already in several other globulars (Dalessandro et al 2008)."," The higher concentration of BSS in the cluster internal region with respect to evolved stars we find is not new, and has been found already in several other globulars (Dalessandro et al 2008)."
" At odds with what is found in other globulars, we do not see any outward increase of the BSS population, most probably because we are not sampling the cluster outskirts (see Section 3)in this study."," At odds with what is found in other globulars, we do not see any outward increase of the BSS population, most probably because we are not sampling the cluster outskirts (see Section 3)in this study."
" To better quantify the relationship between BSS and the other stars we make use of Kolmogorov-Smirnov (KS) statistics, in its one (1D) and two (2D) dimensional The 2D distributions of BSS with respect to other star samples in different evolutionary phases (HB, RGB or MS stars)"," To better quantify the relationship between BSS and the other stars we make use of Kolmogorov-Smirnov (KS) statistics, in its one (1D) and two (2D) dimensional The 2D distributions of BSS with respect to other star samples in different evolutionary phases (HB, RGB or MS stars)"
"This closely corresponds to the polarization of the two normal modes, the former corresponding to the O-mode, and the latter to the X-mode.","This closely corresponds to the polarization of the two normal modes, the former corresponding to the O-mode, and the latter to the X-mode."
" In addition, the following important property holds: irrespective of the pattern of change in plasma density and magnetic field along the trajectory, if two modes were orthogonally polarized in the beginning (OG?—O°)= 4/2, ES?=—e(?), then this property will also be retain subsequently, including the region where the geometrical optics approximation breaks down."," In addition, the following important property holds: irrespective of the pattern of change in plasma density and magnetic field along the trajectory, if two modes were orthogonally polarized in the beginning $\Theta_1^{(1)} - \Theta_1^{(2)} = \pi/2$ , $\Theta_2^{(1)} = - \Theta_2^{(2)}$ ), then this property will also be retain subsequently, including the region where the geometrical optics approximation breaks down."
" Hence, for high enough shear of the external magnetic field along the ray propagation when the derivative is high enough, the first term in the r.h.s."," Hence, for high enough shear of the external magnetic field along the ray propagation when the derivative is high enough, the first term in the r.h.s."
 of Eqn. (70)), of Eqn. \ref{t1}) )
 may be neglected., may be neglected.
" As for O5<1 we have sinh20»ctanh205, one can write down for V/I=tanh205 Here x=QI/c, and we used Ana; for the plasma number density."," As for $\Theta_2 \ll 1$ we have ${\rm sinh}2\Theta_2 \approx {\rm tanh}2\Theta_2$, one can write down for $V/I = {\tanh}2\Theta_2$ Here $x=\Omega l/c$, and we used $\lambda n_{GJ}$ for the plasma number density."
" 'Thus, the sign of the circular polarizationwill coincide with the sign of the derivative d(8p+ó)/dx for the O-mode and they must be opposite for the X-mode."," Thus, the sign of the circular polarizationwill coincide with the sign of the derivative ${\rm d}(\beta_{B}+\delta)/{\rm d}x$ for the O-mode and they must be opposite for the X-mode."
" This approximation can be used for large enoughderivative (i.e., for large enough total turn A(fjg+ within the light cylinder Ri= c/Q), and for small angle of propagation 0<1 through the relativistic plasma (vj/c~ 1)."," This approximation can be used for large enoughderivative (i.e., for large enough total turn $\Delta (\beta_{B}+\delta) \sim 1$ within the light cylinder $R_{\rm L} = c/\Omega$ ), and for small angle of propagation $\theta \ll 1$ through the relativistic plasma $v_{\parallel}/c \sim 1$ )."
 Both these conditions are valid in the magnetospheres of radio pulsars with a good accuracy., Both these conditions are valid in the magnetospheres of radio pulsars with a good accuracy.
" Indeed, assuming that U/c<1 and U;/c£zU/c©0 one can obtain So, the Stokes parameter V (77)) is to be much larger than Vo=+//Q resulting from standard evaluation (Ginzburg, 1961)."," Indeed, assuming that $U/c \ll 1$ and $U_x/c \approx U/c \approx \theta$ one can obtain So, the Stokes parameter $V$ \ref{twom}) ) is to be much larger than $V_0 = \pm I/Q$ resulting from standard evaluation (Ginzburg, 1961)."
"Xba. This fundamental property is well-known in plasma physics and crystal optics (see, e.g., Zheleznyakov et al."," This fundamental property is well-known in plasma physics and crystal optics (see, e.g., Zheleznyakov et al."
 1983; Czyz et al., 1983; Czyz et al.
" 2007), but up to now it was not used in connection with the pulsar radio emission."," 2007), but up to now it was not used in connection with the pulsar radio emission."
" Finally, our numerical simulations show that the sign of the derivative d(fp+ó)/dx is opposite to the sign of dp.a./dó."," Finally, our numerical simulations show that the sign of the derivative ${\rm d}(\beta_{B}+\delta)/{\rm d}x$ is opposite to the sign of ${\rm d}p.a./{\rm d}\phi$."
" As one can see from Eqn. (77)),"," As one can see from Eqn. \ref{twom}) ),"
" this results in an important prediction: 'This implies also that the effects of the particle drift motion, as was already found by Blaskiewicz et al. ("," this results in an important prediction: This implies also that the effects of the particle drift motion, as was already found by Blaskiewicz et al. ("
"1991) (see also Hibschman Arons 2001), shifts the curve to the trailingpart of the mean profile.","1991) (see also Hibschman Arons 2001), shifts the curve to the trailingpart of the mean profile."
" 'Thus, in this paper the arbitrary non-dipole magnetic field configuration, arbitrary number density profile within"," Thus, in this paper the arbitrary non-dipole magnetic field configuration, arbitrary number density profile within"
compajug these uutil we have estimated the systematic errors below.,comparing these until we have estimated the systematic errors below.
 The errors cau be roughly checked wine a bootstrap analysis. which agrees with those values quoted above.," The errors can be roughly checked using a bootstrap analysis, which agrees with those values quoted above."
 Our estimates for ihe nusauce parameter 1000xa are 6.150.31TESI for the SNLS sample and 2.65.(52EE for the nearby one., Our estimates for the nuisance parameter $1000 \times \alpha$ are $6.15^{+0.31}_{-0.31}$ for the SNLS sample and $5.65^{+0.25}_{-0.27}$ for the nearby one.
 a auc aale c ecoYelatecd. with a correlation coefficient of p=—0.6.," $\alpha$ and are quite correlated, with a correlation coefficient of $\rho = -0.6$."
" h order to test tlie adequacy of the quadratic rise-tinie model. we have also perforiued fits iu which t1e ex;»ouent is allowed to vary in the rise-tiuie relation (7 in. f.x/"")."," In order to test the adequacy of the quadratic rise-time model, we have also performed fits in which the exponent is allowed to vary in the rise-time relation $n$ in $f \propto t^n$ )."
" The best coustratut comes from le»w redshift. at which we fiud »=1.72:0.2. with the rise time reduced to 7,=18.500,4.0.32 days."," The best constraint comes from low redshift, at which we find $n=1.7 \pm 0.2$, with the rise time reduced to $\tr = 18.80^{+0.37}_{-0.32}$ days."
 Not surprisingly. the errors on aare considerably larger if 5 is uot ield. fixed.," Not surprisingly, the errors on are considerably larger if $n$ is not held fixed."
 The SNLS sample gives 1;−≻⋅∪∎∣⋮↜↓⋝⋅∖∖↽∐∐n ⋅ κοςlays.," The SNLS sample gives $n=2.0^{+0.4}_{-0.3}$, with $\tr = 19.39^{+1.07}_{-0.82}$ days."
" The combined value ls n=L5:0.2. essentially consistent wi∐↕∐↩⋜↕⊳∖⊳∖⋯∐≺↵≺⇂∖⇁⋜↕↥⋯↵ ο 5,"," The combined value is $n=1.8 \pm 0.2$, essentially consistent with the assumed value of $2$."
" In order ο test that our conclusious are nol too depeuder 1 the valie of n. we also re-fit the nearby ane distaut samples using a fixed vaue ο η=1.5. fiicing rise times of 7,=18.970.190.158 aid 18.19.|}5 days. respectively."," In order to test that our conclusions are not too dependent on the value of $n$, we also re-fit the nearby and distant samples using a fixed value of $n=1.8$, finding rise times of $\tr = 18.97^{+0.19}_{-0.18}$ and $18.49^{+0.17}_{-0.15}$ days, respectively."
 FisD> enw=1.5 5ἱjilfts the rise ti UL cloes 1οἱ appreciably affect the difference yetween the two samples., Fixing $n=1.8$ shifts the rise time but does not appreciably affect the difference between the two samples.
 This conc‘lusion also ho I1e when we compare subsets ο “the hiel-redshift sample. so we restrict the cliscson tonu-2 bsequentLy.," This conclusion also holds true when we compare subsets of the high-redshift sample, so we restrict the discussion to $n=2$ subsequently."
 I —-s also quite inte‘esting to compare the 1jse-tline ineasIreneits for different stubsamples of OUL ala., It is also quite interesting to compare the rise-time measurements for different subsamples of our data.
 Tle results of these fits are given 1 table 2.., The results of these fits are given in table \ref{tbl:fitvals}.
 We do uot cousicer these subsets of the nea‘by samdle because it is too small., We do not consider these subsets of the nearby sample because it is too small.
 First. we consider splitteeig by redshift to search [or evolution within our sample.," First, we consider splitting by redshift to search for evolution within our sample."
 The median recshill is Q.617., The median redshift is 0.647.
 This is quite close to the transition betwee observer frame mmatchinecl with rest-frame B and 7—B. whicl takes place at 2=0.589.," This is quite close to the transition between observer frame matching with rest-frame $B$ and $\ip \mapsto B$, which takes place at $z=0.589$."
 This test is therefore sensitive to two possible effects: evolution. and a ¢‘alibration mismatch between aand7.," This test is therefore sensitive to two possible effects: evolution, and a calibration mismatch between and."
 Because we cannot test each of these iuependeuly. we split at z=0.589. which results in 29 SNe Ia iu the iuterinediate-z sample ai| tl ju tle hieh-> sample X0.13 and. 0.7L. respectively).," Because we cannot test each of these independently, we split at $z=0.589$, which results in 29 SNe Ia in the $z$ sample and 44 in the $z$ sample $\left< z \right> = 0.43$ and $0.74$, respectively)."
 The photoimetric noise in the high-z portion o ‘the sample is much larger than in the interimediate-z portion. as shown in figure 3..," The photometric noise in the $z$ portion of the sample is much larger than in the $z$ portion, as shown in figure \ref{fig:lcoverplot}."
 Tlis figure als» demonstrates that data from different SNe Iacau be combined quite accurately usiug thie tecliuieqes described in 82., This figure also demonstrates that data from different SNe Iacan be combined quite accurately using the techniques described in \ref{sec:parameterization}.
 The data iu the rise-time regiou are shown iu [ie . again split into the two groups.," The data in the rise-time region are shown in figure \ref{fig:snlsrisetime}, again split into the two groups."
" Usit the same analysis. for 2<0.580 we measure 7,19.010.190.158 days aud for 2>O.589 we mneasur 2r.c19.60qe05155 days."," Using the same analysis, for $z \le 0.589$ we measure $\tr = 19.01^{+0.19}_{-0.18}$ days and for $z > 0.589$ we measure $\tr = 19.67^{+0.54}_{-0.49}$ days."
 The intermediate-2 portion «X te salije clearly dominates the fit to tl eh sample., The $z$ portion of the sample clearly dominates the fit to the full sample.
 These are statistically compatible (the «illerence is 1.2 σ]., These are statistically compatible (the difference is 1.2 ).
 Asa test of whether the stretch model works at such eary times. we split the sample by str," As a test of whether the stretch model works at such early times, we split the sample by stretch."
 Unlike the nearby sample. oue caunot measure tle rise time jorecisely for most of the distant SNe Ia iudividually. so this test must be doue ou a sauijxe-wide basis.," Unlike the nearby sample, one cannot measure the rise time precisely for most of the distant SNe Ia individually, so this test must be done on a sample-wide basis."
" Splitting around the mean str the low-stretch sample (s<0.99.(5j= 0.92) of 36 SNe gives y1=19.200.353ost flays aud the sample (5>0.99.(5j= 1.05. 37 SNe) οἱVes Tp=19.090.330,20 €lays."," Splitting around the mean stretch, the low-stretch sample $s \le 0.99, \left<s\right>=0.92$ ) of 36 SNe gives $\tr = 19.20^{+0.33}_{-0.34}$ days and the high-stretch sample $s > 0.99, \left<s\right>=1.05$ , 37 SNe) gives $\tr = 19.09^{+0.33}_{-0.20}$ days."
 These stracdle tL sample value. aud are statistically iudistiuguishable (0.3," These straddle the full sample value, and are statistically indistinguishable (0.3"
2) The obtained upper limit is consistent with the considerable drop in the NIR (4.94) intensity observed by COBE/DIRBE and disagrees with the much smaller decrease in the gamma-ray (above 100 MeV) background measured by EGRET.,2) The obtained upper limit is consistent with the considerable drop in the NIR $4.9\mu$ ) intensity observed by COBE/DIRBE and disagrees with the much smaller decrease in the gamma-ray (above $100$ MeV) background measured by EGRET.
" Therefore, the non-detection of the GRXE in the GA is consistent with the stellar mass distribution in the Galaxy, which does not contradict the stellar nature of GRXE, but is inconsistent with its CR. origin."," Therefore, the non-detection of the GRXE in the GA is consistent with the stellar mass distribution in the Galaxy, which does not contradict the stellar nature of GRXE, but is inconsistent with its CR origin."
 3) The developed background model potentially allows one to reach the statistically limited accuracy., 3) The developed background model potentially allows one to reach the statistically limited accuracy.
" However, the final uncertainty of the approach is associated with the source removal procedure, the systematic uncertainty of the method itself, and the CXB variance."," However, the final uncertainty of the approach is associated with the source removal procedure, the systematic uncertainty of the method itself, and the CXB variance."
" Nevertheless, the implemented method along with the special mode of observation is an optimal approach of modeling the ISGRI background and can be efficiently used for studying the Galactic hard X-ray background."," Nevertheless, the implemented method along with the special mode of observation is an optimal approach of modeling the ISGRI background and can be efficiently used for studying the Galactic hard X-ray background."
"(MPIA). the AMas-Plauck-Institute for Astroplysics (MPA). New Mexico State University. Olio State University, University of Pittsburgh. University of Portsmouth. Princeton University. the United States Naval Observatory. and the University of Washineton.","(MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
Using the new Alonso et al.,Using the new Alonso et al.
" scale we would have had lower abundances also on the CG97 scale, by about 0.09-0.10 dex for the more metal-poor GCs and by only 0.02-0.03 dex for the more metal-rich GCs, which accounts for most of the difference we found between the two scales."," scale we would have had lower abundances also on the CG97 scale, by about 0.09-0.10 dex for the more metal-poor GCs and by only 0.02-0.03 dex for the more metal-rich GCs, which accounts for most of the difference we found between the two scales."
 The metallicity found here for M 30 can be explained by the different temperatures: a star of magnitude similar to the two examined in CG97 has a temperature lower by about 140 K in our new analysis., The metallicity found here for M 30 can be explained by the different temperatures: a star of magnitude similar to the two examined in CG97 has a temperature lower by about 140 K in our new analysis.
" This difference, larger than the average one for metal-poor GCs (about 100 K), justifies the observed offset."," This difference, larger than the average one for metal-poor GCs (about 100 K), justifies the observed offset."
" However, even including this correction we would still have an offset of about 0.08-0.12 dex."," However, even including this correction we would still have an offset of about 0.08-0.12 dex."
 Part of this residual difference is due to the gf values: those currently adopted by us (see Gratton et al., Part of this residual difference is due to the $gf$ values: those currently adopted by us (see Gratton et al.
 2003) are larger on average by 0.026+0.005 dex., 2003) are larger on average by $0.026\pm 0.005$ dex.
" The remaining difference of 0.05-0.09 dex are probably due to the EWs, measured on higher quality and higher resolution spectra and in a more homogeneous way in the present project."," The remaining difference of 0.05-0.09 dex are probably due to the $EW$ s, measured on higher quality and higher resolution spectra and in a more homogeneous way in the present project."
" Finally, we note that most of the scatter around the best fit is given (with opposite sign) by two clusters, NGC 6121 (M 4) and NGC 7099 (M 30)."," Finally, we note that most of the scatter around the best fit is given (with opposite sign) by two clusters, NGC 6121 (M 4) and NGC 7099 (M 30)."
" For the first case, the difference is essentially due to the anomalous value of the ratio Ay/E(B—V), which was taken into account for this cluster in the present programme."," For the first case, the difference is essentially due to the anomalous value of the ratio $A_V/E(B-V)$, which was taken into account for this cluster in the present programme."
" The difference between the classical value 3.1 (CG97) and 4.0 (used in this work) implies a difference of 0.28 mag in the correction of V—K, which in turn translates in a difference of about 200 K in the effective temperatures, just enough to shift M 4 on the best fit regression given by the other GCs."," The difference between the classical value 3.1 (CG97) and 4.0 (used in this work) implies a difference of 0.28 mag in the correction of $V-K$, which in turn translates in a difference of about 200 K in the effective temperatures, just enough to shift M 4 on the best fit regression given by the other GCs."
" Concerning NGC 7099 (M 30), the original data adopted by CG97 were from Minniti et al. ("," Concerning NGC 7099 (M 30), the original data adopted by CG97 were from Minniti et al. ("
"1993) who derived the effective temperatures from the spectra; these temperature were adopted also in CG97, due to the lack of V—K colours for the two stars analysed.","1993) who derived the effective temperatures from the spectra; these temperature were adopted also in CG97, due to the lack of $V-K$ colours for the two stars analysed."
" The other most used metallicity scale for globular clusters is the one defined by Zinn and West (1984), based on a variety of integrated-light photometric and spectroscopic indices calibrated from the few echelle photographic spectra existing at the time."," The other most used metallicity scale for globular clusters is the one defined by Zinn and West (1984), based on a variety of integrated-light photometric and spectroscopic indices calibrated from the few echelle photographic spectra existing at the time."
" Its popularity was due to the ranking this scale provided for the first time for a large number of clusters; in fact, all 19 GCs of our project have a corresponding [Fe/H] entry in the ZW scale."," Its popularity was due to the ranking this scale provided for the first time for a large number of clusters; in fact, all 19 GCs of our project have a corresponding [Fe/H] entry in the ZW scale."
 The comparison between our average metallicities from UVES spectra and the ZW average [Fe/H] values is shown in Fig. 9.., The comparison between our average metallicities from UVES spectra and the ZW average [Fe/H] values is shown in Fig. \ref{f:zwu}.
 A least-squares linear fit (again obtained by exchanging the independent and dependent variables and averaging the results) gives the following relation to transform ZW values to our new scale: with σ=0.143 dex and correlation coefficient r=0.97 from 19 clusters., A least-squares linear fit (again obtained by exchanging the independent and dependent variables and averaging the results) gives the following relation to transform ZW values to our new scale: with $\sigma=0.143$ dex and correlation coefficient $r=0.97$ from 19 clusters.
" On average, the two scales formally differ by only 0.01 dex over the sampled metallicity range, but the scatter of the points is now clearly larger than in the case of the CG97 scale."," On average, the two scales formally differ by only 0.01 dex over the sampled metallicity range, but the scatter of the points is now clearly larger than in the case of the CG97 scale."
 By looking at Fig., By looking at Fig.
 9 a better transformation is obtained when a second-order polynomial is used: with σ=0.119 dex and correlation coefficientzw? r=0.98 from 19 clusters., \ref{f:zwu} a better transformation is obtained when a second-order polynomial is used: with $\sigma=0.119$ dex and correlation coefficient $r=0.98$ from 19 clusters.
 We made an analysis of variance and tested with an F-test the statistical significance of the regression with multiple components and the significance of the coefficient of the highest degree., We made an analysis of variance and tested with an F-test the statistical significance of the regression with multiple components and the significance of the coefficient of the highest degree.
" Both resulted highly significant, and the resulting scatter in the transformation is quite decreased by using the second order equation above."," Both resulted highly significant, and the resulting scatter in the transformation is quite decreased by using the second order equation above."
" We note however that metallicities on the ZW scale are obtained by averaging the results from a heterogeneous series of indices, calibrated on the best metal abundances from high dispersion spectroscopy available at the time (Cohen 1983; Frogel et al."," We note however that $metallicities$ on the ZW scale are obtained by averaging the results from a heterogeneous series of $indices$, calibrated on the best metal abundances from high dispersion spectroscopy available at the time (Cohen 1983; Frogel et al."
 1983)., 1983).
" A better approach is to directly calibrate the original and homogeneous Qo index by Zinn and West (1984), available for a large sample of GCs, on the new metallicities presented here and use it in addition to other homogeneous estimate of [Fe/H] to obtain an accurate ranking in metal abundance for Galactic GCs."," A better approach is to directly calibrate the original and homogeneous $Q_{39}$ index by Zinn and West (1984), available for a large sample of GCs, on the new metallicities presented here and use it in addition to other homogeneous estimate of [Fe/H] to obtain an accurate ranking in metal abundance for Galactic GCs."
 This is done in the Appendix., This is done in the Appendix.
 In Fig., In Fig.
 10 we show the comparison of our [Fe/H] values from UVES with those of the metallicity scale by Kraft and, \ref{f:ki03u} we show the comparison of our [Fe/H] values from UVES with those of the metallicity scale by Kraft and
While the central part of the tachocline (as define bv 7;) is clearly prolate in shape. the overall shape of the tachocline is more complex since the thickness changes.,"While the central part of the tachocline (as define by $r_t$ ) is clearly prolate in shape, the overall shape of the tachocline is more complex since the thickness changes."
 Given the model of the tachocline (Eq. 1)).," Given the model of the tachocline (Eq. \ref{eq:tach}) ),"
" we can assume that the tachocline is bounded. between lavers with radius r;—2w and r,+2:c.", we can assume that the tachocline is bounded between layers with radius $r_t-2w$ and $r_t+2w$.
 The rotation rate changed by about of 00 within these limits., The rotation rate changed by about of $\delta\Omega$ within these limits.
 The upper boundary. of the tachocline is clearly prolate since both r; and iw increase wilh latitude., The upper boundary of the tachocline is clearly prolate since both $r_t$ and $w$ increase with latitude.
 The lower boundary is another matter., The lower boundary is another matter.
 Given the tachocline model adopted in this work. the lower boundary is determined by the value of ry—210. which is —(0.0071220.0021)2... (GONG) or —(0.0066+0.0030)2... (MDID.," Given the tachocline model adopted in this work, the lower boundary is determined by the value of $r_{t3}-2w_3$ , which is $-(0.0071\pm0.0021)R_\odot$ (GONG) or $-(0.0066\pm0.0030)R_\odot$ (MDI)."
 Thus within 30. ihe lower boundary is close to being spherical.," Thus within $\sigma$, the lower boundary is close to being spherical."
 A similar conclusion was reached by Basu Antia (2003)., A similar conclusion was reached by Basu Antia (2003).
 It should be noted that the base of the convection zone is also spherical (Dasu Antia 2001)., It should be noted that the base of the convection zone is also spherical (Basu Antia 2001).
 The lower boundary of the tachocline is at about 0.6872. which is consistent with the extent of mixing required below the solar convection zone to match the solar sound-speed profile (e.e.. Brun οἱ al.," The lower boundary of the tachocline is at about $0.68R_\odot$ which is consistent with the extent of mixing required below the solar convection zone to match the solar sound-speed profile (e.g., Brun et al."
 2002)., 2002).
 The latitudinal variation in 90 is verv clear., The latitudinal variation in $\delta\Omega$ is very clear.
 There is very little change in the rotation rale across the (achocline at a latitude of about 30°., There is very little change in the rotation rate across the tachocline at a latitude of about $30^\circ$.
 At lower Iatitudes 0€) is positive (i.e.. higher rotation rate above the tachocline). while at higher latitudes (he sign of the difference is reversed.," At lower latitudes $\delta\Omega$ is positive (i.e., higher rotation rate above the tachocline), while at higher latitudes the sign of the difference is reversed."
 The latitude al which 90€? changes sign is of some interest., The latitude at which $\delta\Omega$ changes sign is of some interest.
 With the parameters ol tachocline as determined by us (his turns out to be about 297., With the parameters of tachocline as determined by us this turns out to be about $29^\circ$.
 If (0) were the only term present in the definition of 6Q. we would expect this number to be 26.6°.," If $P_3(\theta)$ were the only term present in the definition of $\delta\Omega$, we would expect this number to be $26.6^\circ$."
 Given that (0) is (he dominant term. it is not surprising that the latitude al which 00=0 is close to this value.," Given that $P_3(\theta)$ is the dominant term, it is not surprising that the latitude at which $\delta\Omega=0$ is close to this value."
 The latitude at which 0Q=0 can also be determined cuite easily by inspecting solar rotation profiles obtained from inversions., The latitude at which $\delta\Omega=0$ can also be determined quite easily by inspecting solar rotation profiles obtained from inversions.
 The numbers are quite similar., The numbers are quite similar.
 In order to detect tachocline variations linked to solar activity. we also average results that correspond to times of high and low activity we define (he period of high activity to be the one for which the 10.7 cm radio flux was greater (han 140 SEU and the period οἱ low aclivily is defined as the ones with 10.7 cm [lus less than 90 SEU.," In order to detect tachocline variations linked to solar activity, we also average results that correspond to times of high and low activity — we define the period of high activity to be the one for which the 10.7 cm radio flux was greater than 140 SFU and the period of low activity is defined as the ones with 10.7 cm flux less than 90 SFU."
 To study. possible difference between the (wo periods of minimum activity covered by the data sets we also Lake separate averages for the two periods of low activity., To study possible difference between the two periods of minimum activity covered by the data sets we also take separate averages for the two periods of low activity.
 These are also listed in Table 1., These are also listed in Table 1.
 We do not find any significant change in either the t(hiekness or the position of the Lachocline between the hieh and low activity periods., We do not find any significant change in either the thickness or the position of the tachocline between the high and low activity periods.
 However. ο) and OO; show dillerences al the level of 30.," However, $\delta\Omega_3$ and $\delta\Omega_5$ show differences at the level of $\sigma$."
 To check for anv solar cycle variation we also calculate the correlation coellicient between these parameters and (he 10.7 cm radio flux., To check for any solar cycle variation we also calculate the correlation coefficient between these parameters and the 10.7 cm radio flux.
" The correlationcoefficients for dQ, and dQ; are found. to be 0.40. aud. 0.20 respectively, when GONG data is usec."," The correlationcoefficients for $\delta\Omega_3$ and $\delta\Omega_5$ are found to be 0.40 and 0.20 respectively, when GONG data is used."
fact. theorists had. expected: gaseous massive disces around supermassive black holes to form stars or planets. (e.g...777777) lone before the properties of the voung stars in the GC became known.,"fact, theorists had expected gaseous massive discs around supermassive black holes to form stars or planets \citep[e.g.,][]{Paczynski78,Kolykhalov80,Shlosman89,Collin99,Gammie01,Goodman03}
 long before the properties of the young stars in the GC became known."
" Recently. 7). have numerically simulated the fragmentation process of a gcometrically thin gascous disc of mass z,105AL. for conditions appropriate for our GC (albeit with a rather simple cooling prescription) and found a top-heavy mass function for the stars formed there."," Recently, \cite{NayakshinEtal07} have numerically simulated the fragmentation process of a geometrically thin gaseous disc of mass $\simgt
10^4 \msun$ for conditions appropriate for our GC (albeit with a rather simple cooling prescription) and found a top-heavy mass function for the stars formed there."
 2). extended: numerical studies to the fragmentation of eccentric accretion disks., \cite{AlexanderEtAl08} extended numerical studies to the fragmentation of eccentric accretion disks.
 However. the in-situ model has not so far addressed in detail the origin of the gaseous dises themselves. although similar discs are believed. to exist. in AGN ancl quasars.," However, the in-situ model has not so far addressed in detail the origin of the gaseous discs themselves, although similar discs are believed to exist in AGN and quasars."
 Indeed. sub-parsec massive gaseous disces are invoked as means of [feeding supermassive black holes! immense appetite for gaseous fuel (c.g.2).," Indeed, sub-parsec massive gaseous discs are invoked as means of feeding supermassive black holes' immense appetite for gaseous fuel \citep[e.g.][]{Frank02}."
. Such dises are actually observed as the sites of maser emission in some of the nearby ealactic centres (7).., Such discs are actually observed as the sites of maser emission in some of the nearby galactic centres \citep{Miyoshi95}.
 )ut iis currently not a member of the AGN club. let alone the even more powerful and clitist quasar community.," But is currently not a member of the AGN club, let alone the even more powerful and elitist quasar community."
" Indeed. bholometric luminosity is around ~10."" of its Eddington limit (e.g.??).."," Indeed, bolometric luminosity is around $\sim 10^{-9}$ of its Eddington limit \citep[e.g.,][]{Narayan02,Baganoff03a}."
 The amount of ionised eas currently present in the inner parsec is estimated at perhaps a mere hundred solar masses (72).., The amount of ionised gas currently present in the inner parsec is estimated at perhaps a mere hundred solar masses \citep{Paumard04}.
 There are also tight observational constraints (77) on the presence of an optically thin disc on smaller scales (sav 0.01 pe or less).," There are also tight observational constraints \citep{Falcke97,Cuadra04} on the presence of an optically thin disc on smaller scales (say 0.01 pc or less)."
 We believe that the requisite gaseous. disces. could not have been assembled. by viscous transport of angular momentum. as is expected to be the case in other systems (ολο.," We believe that the requisite gaseous discs could not have been assembled by viscous transport of angular momentum as is expected to be the case in other systems \citep{Shakura73,Frank02}."
" Phe viscous. time seale of a thin. mareinally self-gravitating cise can be estimated as foi ὃν where Mig23,5109M. and M are masses of the blackhole and the disc. respectively. and Q is the Keplerian angular frequeney."," The viscous time scale of a thin, marginally self-gravitating disc can be estimated as $t_{\rm visc} =
(\mbh/M_{\rm disc})^2 (\alpha \Omega)^{-1}$ , where $\mbh \approx 3.5 \times
10^6\msun$ and $M_{\rm disc}$ are masses of the blackhole and the disc, respectively, and $\Omega$ is the Keplerian angular frequency."
 For a=0.1 ancl Maiiτε (reasons for this choice of mass are cliscussecl in 2))). lusso3y010 vears at pr=0.08 parsec. the inner edge of the disces (2).. and. 10 vears at 0.3 parsec.," For $\alpha=0.1$ and $M_{\rm disc} \approx 10^4 \msun$ (reasons for this choice of mass are discussed in \cite{NayakshinEtal06}) ), $t_{\rm visc} \sim 3 \times
10^7$ years at $r = 0.03$ parsec, the inner edge of the discs \citep{PaumardEtal06}, and $10^9$ years at 0.3 parsec."
 These times are very long compared with the age of the stellar svstemis., These times are very long compared with the age of the stellar systems.
 In fact. the gaseous disces would have to evolve even faster as the two stellar svstems are co-eval within one million vears (?)..," In fact, the gaseous discs would have to evolve even faster as the two stellar systems are co-eval within one million years \citep{PaumardEtal06}."
 lt is also not clear that one should. expect. any sort of a viscous quasi-steady state for the gas in our Galactic Centre., It is also not clear that one should expect any sort of a viscous quasi-steady state for the gas in our Galactic Centre.
 There is no strong evidence for other similar star ormation events in the central parsee within the last ~107 vears., There is no strong evidence for other similar star formation events in the central parsec within the last $\sim 10^8$ years.
 Therefore. the one-olf star formation event appears to oe best explained by a one-olf deposition of gas within the central parsec.," Therefore, the one-off star formation event appears to be best explained by a one-off deposition of gas within the central parsec."
 There are several ways in which this could rave happened. c.g. a Giant Molecular Cloud (GAC) with a sub-parsec impact parameter (in relation to A*)) could rave self-collided and become partially bound to the central xwsec (c.g.2).," There are several ways in which this could have happened, e.g. a Giant Molecular Cloud (GMC) with a sub-parsec impact parameter (in relation to ) could have self-collided and become partially bound to the central parsec \citep[e.g.,][]{NC05}."
 Alternatively. a GMC could have struck the Circumnuclear Disc (CND) located a few parsecs away fromA.. and then created. gas streams that settled into the central parsec.," Alternatively, a GMC could have struck the Circumnuclear Disc (CND) located a few parsecs away from, and then created gas streams that settled into the central parsec."
 Lt is very hard to estimate the probability of a cloud-cloud collision in the central part of the Galaxy., It is very hard to estimate the probability of a cloud-cloud collision in the central part of the Galaxy.
 ?) estimated the rate of GALC collisions as one per 20 Myrs in the central LOO pe. but the estimate is uncertain by at least an order of magnitude in either direction due to the lack of knowledge about the size distribution of GMCSs.," \cite{Hasegawa94} estimated the rate of GMC collisions as one per 20 Myrs in the central 100 pc, but the estimate is uncertain by at least an order of magnitude in either direction due to the lack of knowledge about the size distribution of GMCs."
 However. there is observational evidence that C:MCS can be put on orbits significantly dillerent from the simple circular orbits seen in the inner Galaxy.," However, there is observational evidence that GMCs can be put on orbits significantly different from the simple circular orbits seen in the inner Galaxy."
 2). recently measured the proper motions of the Arches star cluster.," \cite{Stolte08}
 recently measured the proper motions of the Arches star cluster."
 This star cluster is about 2.5 million. vears old and has a mass in excess of 103M. (c.g.7). ," This star cluster is about 2.5 million years old and has a mass in excess of $10^4 \msun$ \citep[e.g.,][]{Figer04}. ."
Lt ds presently about 30 pe away [rom lin projection., It is presently about 30 pc away from in projection.
 Its orbit must. be strongly non-circular as its 3D velocity is over 200 km/sec in the region where the circular velocity is only ~110 km/sec., Its orbit must be strongly non-circular as its 3D velocity is over 200 km/sec in the region where the circular velocity is only $\sim 110$ km/sec.
 Therefore. this sugeests that massive GAICs can be on highly non-circular orbits.," Therefore, this suggests that massive GMCs can be on highly non-circular orbits."
 Lt does not seem implausible that one of these clouds would pass within a few parsec ofA., It does not seem implausible that one of these clouds would pass within a few parsec of.
 In this paper we explore such a once-olf collision event in avery simple setup (822))., In this paper we explore such a one-off collision event in a very simple setup \ref{sec:ic}) ).
 We allow two massive. uniform ancl spherical clouds on. significantlv. οονομα orbits to collide with each other one parsec away fromA*.," We allow two massive, uniform and spherical clouds on significantly different orbits to collide with each other one parsec away from."
. The resulting gas chynamies. in particular the way in which gas settles into the inner. parsec. is the focus of our clfort here.," The resulting gas dynamics, in particular the way in which gas settles into the inner parsec, is the focus of our effort here."
 We find that the collision formis streams of gas with varving angular momentum. both in magnitude and direction.," We find that the collision forms streams of gas with varying angular momentum, both in magnitude and direction."
 Parts of these streams collicle ancl coalesce to form a. disc. and remaining parts form one or more orbiting filaments.," Parts of these streams collide and coalesce to form a disc, and remaining parts form one or more orbiting filaments."
 As the eas cools. it. becomes self-egravitating and stars are born. usually in both the disc and the filaments.," As the gas cools, it becomes self-gravitating and stars are born, usually in both the disc and the filaments."
 This overall picture Is discussed in 77..., This overall picture is discussed in \ref{sec:overall}.
" A reacler mainly interested in a comparison between the simulations and observations may find sections ?7?., 7? and ?? mostrelevant!"," A reader mainly interested in a comparison between the simulations and observations may find sections \ref{sec:radial}, \ref{sec:accretion} and \ref{sec:discussion} most."
 In £27. we argue that our simulation results sugeest that the location of the original impact was somewhat further away [rom tthan we assumed here. e.g. perhaps a few. parsecs.," In \ref{sec:obs} we argue that our simulation results suggest that the location of the original impact was somewhat further away from than we assumed here, e.g., perhaps a few parsecs."
 The numerical approach and the code usec in this paper is the same as in ?) with onlv slight mociilications., The numerical approach and the code used in this paper is the same as in \cite{NayakshinEtal07} with only slight modifications.
 We cmploy GADGET-2. a smoothecl particle hiverodyvnamic (SPLI)/N-bods code (?)..," We employ GADGET-2, a smoothed particle hydrodynamic (SPH)/N-body code \citep{Springel05}."
 The Newtonian N-bods gravitational interactions of particles in the code are calculated: via a tree algorithm., The Newtonian N-body gravitational interactions of particles in the code are calculated via a tree algorithm.
 Artificial viscosity is used. to resolve shocks in the gas., Artificial viscosity is used to resolve shocks in the gas.
 Gas cools according to dufdl=uflosaGR). where cooling time depends on radius as where Jaen=1/0 and Ὁ=(CMiB)? is the Ixeplerian angular velocity. ancl 347 is a dimensionless parameter.," Gas cools according to $du/dt
= - u/t_{\rm cool}(R)$, where cooling time depends on radius as where $t_{\rm dyn} = 1/\Omega$ and $\Omega = (GM_{\rm
bh}/R^{3})^{1/2}$ is the Keplerian angular velocity, and $\beta$ is a dimensionless parameter."
 This approach is motivated by simulations of mareinally stable sell-gravitating gaseous disces. where 3 is expected to be of the order of a few (22).," This approach is motivated by simulations of marginally stable self-gravitating gaseous discs, where $\beta$ is expected to be of the order of a few \citep{Gammie01,Rice05}."
 In the problem at hand. 3 would most likely be much smaller than this during the collision phase of the clouds as gas can heat up to high temperatures (sce £22).but it could be much Larger than unity when gas cools into a geometrically thin disc or a thin filament.," In the problem at hand, $\beta$ would most likely be much smaller than this during the collision phase of the clouds as gas can heat up to high temperatures (see \ref{sec:obs}) ),but it could be much larger than unity when gas cools into a geometrically thin disc or a thin filament."
 In the framework of our simplified approach. ancl given numerical limitations. we only consider two values for 3.," In the framework of our simplified approach, and given numerical limitations, we only consider two values for $\beta$ ."
 The 3=1 , The $\beta=1$ 
three categories: quasar. broad line radio galaxy (BLRG) or radio galaxy.,"three categories: quasar, broad line radio galaxy (BLRG) or radio galaxy."
 Given the relatively small number of reddened quasars found in the 6C. 7€-L and TC-LL samples. we do not believe the lack of nearinfrared. data for the 7C-HI sample has caused a statisticallysignificant number of quasars to be mis.classified.," Given the relatively small number of reddened quasars found in the 6C, 7C-I and 7C-II samples, we do not believe the lack of near–infrared data for the 7C-III sample has caused a statistically–significant number of quasars to be mis–classified."
 Spectra of the quasars and BLRCs in 7C- Land 7C-L are presented with a detailed. discussion of the classification scheme in Willott et al. (, Spectra of the quasars and BLRGs in 7C-I and 7C-II are presented with a detailed discussion of the classification scheme in Willott et al. (
199Sa),1998a).
 Our total saniple comprises 356 sources (one object. 3€ 200. is common to both the 3CRR and 7TC-LE samples).," Our total sample comprises 356 sources (one object, 3C 200, is common to both the 3CRR and 7C-II samples)."
 We do not exclude compact symmetric objects (CSOs) from our analysis as previous authors have (e.g. Barthel 1989. Singal 1996). since we have no a priori reason to expect them not to be part of the unified schemes.," We do not exclude compact symmetric objects (CSOs) from our analysis as previous authors have (e.g. Barthel 1989, Singal 1996), since we have no a priori reason to expect them not to be part of the unified schemes."
 In Section 3. we repeat our analysis excluding the CSOs to check this assumption., In Section \ref{pzplots} we repeat our analysis excluding the CSOs to check this assumption.
 Sources with FRI radio structure are however excluded., Sources with FRI radio structure are however excluded.
 The main reason for this was a practical one: most (24 out of 33) of the FRIs are from the 3CRR sample. ancl the existing spectrophotomoetry of these objects is inadequate for the quantitative investigations of this paper.," The main reason for this was a practical one: most (24 out of 33) of the FRIs are from the 3CRR sample, and the existing spectrophotometry of these objects is inadequate for the quantitative investigations of this paper."
 However. as discussed. fully in Section 4.5. the lack of quasars with FRI radio structure (e.g. Faleke. Gopal-Ixrishna 3icermann 1995). and the concentration of these objects at low Liz and low (sco Lig. 1))," However, as discussed fully in Section 4.5, the lack of quasars with FRI radio structure (e.g. Falcke, Gopal-Krishna Biermann 1995), and the concentration of these objects at low $L_{151}$ and low $z$ (see Fig. \ref{fig:qf1}) )"
 makes it easy to investigate the ellects of the exclusion. of FRIs on our study of the quasar fraction., makes it easy to investigate the effects of the exclusion of FRIs on our study of the quasar fraction.
 Although the radio structures of the quasars 5C'6.264 and 3€ 48 are arguably FRI. we count them in this paper as FRIIS because of their high radio luminosities.," Although the radio structures of the quasars 5C6.264 and 3C 48 are arguably FRI, we count them in this paper as FRIIs because of their high radio luminosities."
 Indeed. all objects that could. not unambiguously characterised as FRI or ΕΙ (e.g. the “DA sources in the notation of Blundell et al., Indeed all objects that could not unambiguously characterised as FRI or FRII (e.g. the `DA' sources in the notation of Blundell et al.
 in prep.), in prep.)
 are also counted as FRIIS in this stuck., are also counted as FRIIs in this study.
 Thus our total combined sample contains 323 objects deemed to be FItIs., Thus our total combined sample contains 323 objects deemed to be FRIIs.
 We now investigate the fraction. of objects which show broad emission lines (referred. to hereafter simply. as the quasar fraction) in the complete radio samples., We now investigate the fraction of objects which show broad emission lines (referred to hereafter simply as the quasar fraction) in the complete radio samples.
 Broad line radio galaxies (broad line objects with Mg — 23) are counted as quasars in this analvsis. because they are most likely to be weak quasars viewed: within the torus opening angle (Laing et al.," Broad line radio galaxies (broad line objects with $M_B>-23$ ) are counted as quasars in this analysis, because they are most likely to be weak quasars viewed within the torus opening angle (Laing et al."
 1994: Llarcdeastle ct al., 1994; Hardcastle et al.
 1998: Fig., 1998; Fig.
 6 of Willott ct al., 6 of Willott et al.
.. 1998a)., 1998a).
. However. objects with scattered. broad. lines seen only in polarised light (4. cases) are counted as galaxies. because the fact that the broad lines are not. observed. clirectly indicates that the nuclear regions are heavily obscured. most likely bv the torus.," However, objects with scattered broad lines seen only in polarised light (4 cases) are counted as galaxies, because the fact that the broad lines are not observed directly indicates that the nuclear regions are heavily obscured, most likely by the torus."
 Since spectropolarimetry is unavailable for. most. sources. there is clearly some danger that the split between weak and seattered-light quasars is imperfect.," Since spectropolarimetry is unavailable for most sources, there is clearly some danger that the split between weak and scattered-light quasars is imperfect."
 A further concern. is the question of the classification of lightly-reddened, A further concern is the question of the classification of lightly-reddened
 A further concern. is the question of the classification of lightly-reddened., A further concern is the question of the classification of lightly-reddened
J.NLC. acknowledges support from NSF eraut. AST-0507125.,J.M.C. acknowledges support from NSF grant AST-0507428.
 The work of R.L.G. aud D.S. was carried out at the Jet Propulsion Laboratory at the C'alifornia Iustitute of Technology. uuder a contract with NASA.," The work of R.L.G. and D.S. was carried out at the Jet Propulsion Laboratory at the California Institute of Technology, under a contract with NASA."
 D.F.G. acknowledges support by the U.S. Departineut of Enereyv under coutract uuuber DE-AC3-76SFO0515., B.F.G. acknowledges support by the U.S. Department of Energy under contract number DE-AC3-76SF00515.
 απο was supported by NASA through the Spitzer Space ΟνTelescope Fellowship program., M.C.C. was supported by NASA through the Spitzer Space Telescope Fellowship program.
 We thank Dai Fu for valuable assistance with the zCOSAMOS προς., We thank Hai Fu for valuable assistance with the zCOSMOS spectrum.
 This letter is partly based ou observations carried out using MegaPriue/MoegaCzu. a joint project of CFIIT and CEA/DAPNIA. at the CFUT which is operated bv the NRC of Canada. the Institut National des Science. de Univers of the CNRS of France. and the University of Hawaii:GALEN. a NASA funded Sul Explorer Mission: SDSS aud SDSS-IT. which are fuuded bv the Alfred P. Sloan Foundation. the Participating Tustitutions. the NSF. the US Dopartineut. of Eucrex. NASA. the Japanese Moubukagakusho. the Max Plauck Society. and the Digher Education Fuudiug Council for England: theTelescope. which is operated by the Jet Propulsion Laboratory. California Iustitute of Technology uuder a contract with NASA: the VLA. a facility of the NRAO. itself a facility of the NSF that is operated by Associated Universities. Πιο:Newton. au ESA scieuce mission with instruments and contributions directly funded by ESA Member States and the USA (NASA): and the W.M. Week Observatory.," This letter is partly based on observations carried out using MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at the CFHT which is operated by the NRC of Canada, the Institut National des Science de l'Univers of the CNRS of France, and the University of Hawaii;, a NASA funded Small Explorer Mission; SDSS and SDSS-II, which are funded by the Alfred P. Sloan Foundation, the Participating Institutions, the NSF, the US Department of Energy, NASA, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England; the, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA; the VLA, a facility of the NRAO, itself a facility of the NSF that is operated by Associated Universities, Inc.;, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA); and the W.M. Keck Observatory."
 We wish to recognize and acknowledge the very significant cultural role and reverence that the sunuuit of Alanna Kea has abwavs hac within the indigenous ΠασάΠα comunity., We wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.
 We are most fortunate to have the opportunity to conduct observations fron this iiouutaiu., We are most fortunate to have the opportunity to conduct observations from this mountain.
eas of similar surface density.,gas of similar surface density.
 In $66 we calculate the fraction of diffuse 21 jan enüssiou for our sample., In 6 we calculate the fraction of diffuse 24 $\mu$ m emission for our sample.
 We stuumarize our conclusions iu 8T., We summarize our conclusions in 7.
 We chose three galaxies. two erand-design (NGC οὓς and NCC 5191) aud one more floceulent (NCC 6916) 0n the basis of their proximity. orientation aud multibaud data.," We chose three galaxies, two grand-design (NGC 628 and NGC 5194) and one more flocculent (NGC 6946) on the basis of their proximity, orientation and multiband data."
 These have coverage in GALEN (NUV|FUV) (Cul de Paz et aL.," These have coverage in GALEX (NUV+FUV) (Gil de Paz et al.,"
 2007). TUOINGS (Walter et al.," 2007), THINGS (Walter et al.,"
 2008). SINGS (semnicutt et al.," 2008), SINGS (Kennicutt et al.,"
 2003) and IRAM 30 observations (Schuster et al., 2003) and IRAM 30 observations (Schuster et al.
 2007. Leroy et al..," 2007, Leroy et al.,"
 2009)., 2009).
" Iu order to quantify the amount of star formation iu the arm aud interarm regions we focus on a series of observables: 21,04 endssion. which traces young. dus eushrouded stars: UV cussion. which traces vouug. uuobseured stars: CO. which traces molecular gas that is prestunably orgauized iuto elant molecular clouds: are III sas. which is prestunably a imndxture of wari aux cold. atomic. diffuse. eas."," In order to quantify the amount of star formation in the arm and interarm regions we focus on a series of observables: $\mu$ m emission, which traces young, dust enshrouded stars; UV emission, which traces young, unobscured stars; CO, which traces molecular gas that is presumably organized into giant molecular clouds; and HI gas, which is presumably a mixture of warm and cold, atomic, diffuse, gas."
 Each of these probes a different stage iu the star formation process., Each of these probes a different stage in the star formation process.
 Towever. noue of them uniquely define the star formation rate as other sources nay contaminate or atteuuate the emission.," However, none of them uniquely define the star formation rate as other sources may contaminate or attenuate the emission."
 The tan enudsson. especially iu the iuterarii region. may also arise from diffuse cussion not associated witli recent star formation (curus 21514 cinissiou) (Ilelou. 1986. Calzetti et al.," The $\mu$ m emission, especially in the interarm region, may also arise from diffuse emission not associated with recent star formation (cirrus $\mu$ m emission) (Helou, 1986, Calzetti et al."
 2007)., 2007).
 The FUV. ou the other haud. may be atteuuated by dust in the ari region (Uscunicutt 1998. Calzetti et al.," The FUV, on the other hand, may be attenuated by dust in the arm region (Kennicutt 1998, Calzetti et al."
 2007)., 2007).
 Taken together. these two effects could potentially boost the relative amount of star formation iu the interarm region.," Taken together, these two effects could potentially boost the relative amount of star formation in the interarm region."
" Fortunately, CO neasurelents. which trace the molecular gas component. i.c. the fuel for star formation. can also be used to further xobe the amount star formation in the iterari region. if we asstune that the relations found by Bos aud Lüs indeed hold."," Fortunately, CO measurements, which trace the molecular gas component, i.e. the fuel for star formation, can also be used to further probe the amount star formation in the interarm region, if we assume that the relations found by B08 and L08 indeed hold."
 We frst describe the steps taken to το he images for our analysis as well as how the arm aud iuterarni regions are defined im the following subsections., We first describe the steps taken to render the images for our analysis as well as how the arm and interarm regions are defined in the following subsections.
 All of the images are aligned to the THINGS astrometric erid aud degraded to a coummon resolution of 13 FWIOIAL which is the resolution of the TERACLES CO images.," All of the images are aligned to the THINGS astrometric grid and degraded to a common resolution of $""$ FWHM, which is the resolution of the HERACLES CO images."
 Before this degradation. we remove foreground stars from the far-UW. tan and 3.6424 nuages uxing their UV colour.," Before this degradation, we remove foreground stars from the far-UV, $\mu$ m and $\mu$ m images using their UV colour."
 Pixels with an NUV-to-FUV intensity ratio between 9 aud 25. depending ou the galaxy. are blauked.," Pixels with an NUV-to-FUV intensity ratio between 9 and 25, depending on the galaxy, are blanked."
 We also require that the cut pixels have values greater than 5o in the NUV map., We also require that the cut pixels have values greater than $\sigma$ in the NUV map.
 The companion of NGC 5191 is removed by-eve. anc is bevond the radius considered in the analysis.," The companion of NGC 5194 is removed by-eye, and is beyond the radius considered in the analysis."
 All images are eveutually ceprojected to face-on according to the values found in Walter et al..," All images are eventually deprojected to face-on according to the values found in Walter et al.,"
 2008 aud listed in Table 1., 2008 and listed in Table 1.
 As in LOs and BOs. we remove a residual background frou the FUV and [jii images. measured as a median value in an offealaxy box.," As in L08 and B08, we remove a residual background from the FUV and $\mu$ m images, measured as a median value in an off-galaxy box."
 Stars are removed using the NUNV-to-FUV ratio aud the UV images are corrected for ealactic extinction (Schlegel et al., Stars are removed using the NUV-to-FUV ratio and the UV images are corrected for galactic extinction (Schlegel et al.
 1998) usine the E(D-V) values listed iu NED. which are 0.07. 0.035 aud 0.312 for NGC 628. NGC 5191 and NGC 6916. respectively.," 1998) using the E(B-V) values listed in NED, which are 0.07, 0.035 and 0.342 for NGC 628, NGC 5194 and NGC 6946, respectively."
 The UV and {μαι images are then combined to produce SER maps (see Appeudix of LU8): where Nery las units of M. 7 Land the FUV and 21402 intensity are cach iu My |., The UV and $\mu$ m images are then combined to produce SFR maps (see Appendix of L08): where $\Sigma_{SFR}$ has units of $_{\sun}$ $^{-2}$ $^{-1}$ and the FUV and 24 $\mu$ m intensity are each in MJy $^{-1}$.
 We use 21 cm line euissioun from THINGS to trace he atomic gas., We use 21 cm line emission from THINGS to trace the atomic gas.
 We convert the iuteerated intensity tfo a surface density aud include a factor of 1.36 to account for clin. following LOS.," We convert the integrated intensity to a surface density and include a factor of 1.36 to account for helium, following L08."
 We use integrated CO J—2»1 intensity maps from HERACLES (Leroy et aL.," We use integrated CO $J=2 \rightarrow 1$ intensity maps from HERACLES (Leroy et al.,"
 2009) ο estimate the distribution of To., 2009) to estimate the distribution of $_2$.
 For M51 this is a reprocessing of the maps preseuted by Schuster et al.," For M51 this is a reprocessing of the maps presented by Schuster et al.,"
 2007 and UWitschteld et al.," 2007 and Hitschfeld et al.,"
 2009 (the reprocessing docs rot significantly alter the map)., 2009 (the reprocessing does not significantly alter the map).
 To estimate the surface density of Πο from CO we use a constant conversion factor., To estimate the surface density of $_2$ from CO we use a constant conversion factor.
 As in LOS we adopt: The maps are deprojected aud added togetler to form total eas deusitv maps., As in L08 we adopt: The maps are deprojected and added together to form total gas density maps.
 The CO-to-IIo. couversion factor is a source of uncertaimty., The $_2$ conversion factor is a source of uncertainty.
 It shotId in principle be a fiction of (at. least) metallicity. radiation feld. density. aud temperature aud at least some of these conditions may change between the arm axd interarm regions.," It should in principle be a function of (at least) metallicity, radiation field, density, and temperature and at least some of these conditions may change between the arm and interarm regions."
 Towever. direct evidence for couversion factor variations In our tarects is mixed and oftei coutradictory (6.8. CGarcia-Duirillo et al..," However, direct evidence for conversion factor variations in our targets is mixed and often contradictory (e.g. Garcia-Burillo et al.,"
 1993) For a detailed. discussion see Schinnerer ct al..," 1993) For a detailed discussion see Schinnerer et al.,"
 2010., 2010.
 The fact that the studies of LOs and Bos showed striking agreement of different ealaxies in the Schuiüdt Law plot (plotting SER surface densities vs. gas surface densities) provides further confidence that the New conversion factor is roughly coustaut for the systems studied here., The fact that the studies of L08 and B08 showed striking agreement of different galaxies in the `Schmidt Law' plot (plotting SFR surface densities vs. gas surface densities) provides further confidence that the $_{\rm CO}$ conversion factor is roughly constant for the systems studied here.
 Thus. we use a constant New iu our analysis.," Thus, we use a constant $_{\rm CO}$ in our analysis."
 Tt is clear that we should define the ari and imterana regions. imn a wav that is least biased by the voung stellar population.," It is clear that we should define the arm and interarm regions, in a way that is least biased by the young stellar population."
 Ideally. we wish to define the spiral arms using the stellar mass deusity. or at least the old stellar population.," Ideally, we wish to define the spiral arms using the stellar mass density, or at least the old stellar population."
 Near-iufrared nuages have conunuoulv Όσοι used for this purpose (e.g. EhuegreeL& Ehneercen 198. Rix Zaritsky 1995. Seiegar James. 1998. Grosbol ot al.," Near-infrared images have commonly been used for this purpose (e.g. Elmegreen Elmegreen 1984, Rix Zaritsky 1995, Seigar James, 1998, Grosbol et al.,"
 2001 aud Ivencdall et al., 2004 and Kendall et al.
 2008)., 2008).
 We use the 3.6ju iniaees from the IRAC instrument o LSpitzer (IKeunicutt et al., We use the $\mu$ m images from the IRAC instrument on Spitzer (Kennicutt et al.
 2003) to trace the underlying od stellar population., 2003) to trace the underlying old stellar population.
 Tn this band most of the cussion is due to old stars although there is some patchy contamination from hot dust aud PATI features (e.g. Kendal et al., In this band most of the emission is due to old stars although there is some patchy contamination from hot dust and PAH features (e.g. Kendall et al.
 2008)., 2008).
 Zibetti, Zibetti
that weakly polarized radio cussion is produced by the evrosvuchnrotron incchauisu of a population of wildly relativistic electrons.,that weakly polarized radio emission is produced by the gyrosynchrotron mechanism of a population of mildly relativistic electrons.
 This mechanisia. however. eaunot account for polarizations exceeding about which may originate frou a coherent process.," This mechanism, however, cannot account for polarizations exceeding about which may originate from a coherent process."
 Circular polarizations of up to have been reported for YZ CAG at 20 cii and Ls cm (Lang&Wilson L986:: Beuz&Alef 1991)., Circular polarizations of up to have been reported for YZ CMi at 20 cm and 18 cm \cite{Lang}; ; \cite{B1} 1991).
 Tere we report on the results of VLBA experiments of YZ CMi and AD Leo at 3.6 cii., Here we report on the results of VLBA experiments of YZ CMi and AD Leo at 3.6 cm.
 These are well known. nearby. voung radio stars close to zero-age Mail sequenuce (ZANIS).," These are well known, nearby, young radio stars close to zero-age main sequence (ZAMS)."
 Some of their eeueral properties are listed in Table 1.., Some of their general properties are listed in Table \ref{tab:sum}.
 The single dMe stars YZ CAG and AD Leo were observed on 1997 April 18/19 and 1996 December 12. respectively. with the VLBA aud the phased-up VLA as a joint VLBI svsteni vieldiug an augular resolution of better than one nulharcsecond (aas).," The single dMe stars YZ CMi and AD Leo were observed on 1997 April 18/19 and 1996 December 12, respectively, with the VLBA and the phased-up VLA as a joint VLBI system yielding an angular resolution of better than one milliarcsecond (mas)."
 The VLÀ was also available im its uormal iuterferometer mode and was used for total flux ucasurements at ναν small baseline leneth which allowed us to monitor the changes iu total flux deusity and polarisation of both stars through the observations., The VLA was also available in its normal interferometer mode and was used for total flux measurements at very small baseline length which allowed us to monitor the changes in total flux density and polarisation of both stars through the observations.
 3C286 was used for the flux calibration of both stars., 3C286 was used for the flux calibration of both stars.
 The VLBI observations used a bandwidth of δ MIIZ in both left and right circulay polarization a sll Cz (3.6. cia) and two bit sample. eiving a data rate of 128 Mbit/s. Iu order to reliably nuage such weals SOULCOS (<3.0 ταν). the yhase referencing techuique was used (Beasley&Comvay 1995)) in which we switched between the tarect and a bright calibrator (quasar) in evcles of two to three miuuutes.," The VLBI observations used a bandwidth of 8 MHz in both left and right circular polarization at 8.41 GHz (3.6 cm) and two bit sampling, giving a data rate of 128 Mbit/s. In order to reliably image such weak sources $<3.0$ mJy) the phase referencing technique was used \cite{Beasley}) ) in which we switched between the target and a bright calibrator (quasar) in cycles of two to three minutes."
 For our targets YZ CMG and AD Leo these calibrators were 0736|O17 and 1022|19 respectively (sec Table 2)). at 2.1 and 1.1 degrees separation from the tarect.," For our targets YZ CMi and AD Leo these calibrators were 0736+017 and 1022+194 respectively (see Table \ref{tab:cali}) ), at 2.1 and 1.4 degrees separation from the target."
 The amplitude calibration was performed with AIPS. ‘ollowed bv. editing in DIFAIAP (in order to flag iore oeciselv single. bad visibilities of the VLBA data) iud by a continued aualvsis in AIPS (ucludiug all the mapping).," The amplitude calibration was performed with AIPS, followed by editing in DIFMAP (in order to flag more precisely single, bad visibilities of the VLBA data) and by a continued analysis in AIPS (including all the mapping)."
 We first mace livbrid naps of our calibrators using closure hase iiethods iu order to determine their structure., We first made hybrid maps of our calibrators using closure phase methods in order to determine their structure.
 We hen deteriuned the atmospheric contributions to the yhase of the calibrator data., We then determined the atmospheric contributions to the phase of the calibrator data.
 With hese solutions it was yosstble to shase-correct the stellar data., With these solutions it was possible to phase-correct the stellar data.
 We then made oxelininary wide maps in order to find the stars aud phase-rotated the data in order to bring the target stars to he inage phase ceuter., We then made preliminary wide maps in order to find the stars and phase-rotated the data in order to bring the target stars to the image phase center.
 Finally we deconvolved the stellar inaees using CLEAN., Finally we deconvolved the stellar images using CLEAN.
 As a check of the reliabilitv of the pliase calibration. our observations also included a second bright calibrator for cach star (0713-006 for YZ CAG. and 1013|208 for AD Leo. see Table 2)).," As a check of the reliability of the phase calibration, our observations also included a second bright calibrator for each star (0743-006 for YZ CMi, and 1013+208 for AD Leo, see Table \ref{tab:cali}) )."
 The maps of these secoudary calibrators. made using the phase solutious found towards the primary calibrators showed a dynamic range of about 20:1.," The maps of these secondary calibrators, made using the phase solutions found towards the primary calibrators showed a dynamic range of about 20:1."
 Since the separation between the two calibrators is approximately the same as that between the star aud cach calibrator. we cau conclude that our stellar images have also a 20:1 dynamic range aud heuce are in practice noise limited rather than dynamic range lanited.," Since the separation between the two calibrators is approximately the same as that between the star and each calibrator, we can conclude that our stellar images have also a 20:1 dynamic range and hence are in practice noise limited rather than dynamic range limited."
 We also produced: selfealibration maps and datasets of the secoudary calibrators., We also produced selfcalibration maps and datasets of the secondary calibrators.
 The plots of the friuge amplitude against gng-distenee show the sae fall off than the ones produced. from the phase referenced. datasets., The plots of the fringe amplitude against -distance show the same fall off than the ones produced from the phase referenced datasets.
 This confirms the value of the dvuamic range of the images of he targets., This confirms the value of the dynamic range of the images of the targets.
" The source centroid positions of the phase reference naps of the secoudary calibrators were within 0.5 ane 1,2 mas of the correlated positions (a-priori positions) for 713-006 aud. LOLS|208 respectively (see Table 21).", The source centroid positions of the phase referenced maps of the secondary calibrators were within 0.5 and 0.2 mas of the correlated positions (a-priori positions) for 0743-006 and 1013+208 respectively (see Table \ref{tab:cali}) ).
 This is consistent with he claimed accuracy of those correlate »ositious (0.5 mas. Jolinstonctal. 1995).," This is consistent with the claimed accuracy of those correlated positions (0.5 mas, \cite{Johnston} 1995)."
 This test gives us confidence iu he astrometric accuracy of our stellar )ositions (see Sect., This test gives us confidence in the astrometric accuracy of our stellar positions (see Sect.
 3.L)., 3.4).
 YZ CM was found at a surprisingly high. slowly decreasing flux density (average 2.9 imuJw) durius the whole observation (see Fie. l..," YZ CMi was found at a surprisingly high, slowly decreasing flux density (average 2.9 mJy) during the whole observation (see Fig. \ref{fig:light97},"
 left lot), left plot).
 As previously noted (see Sect. 1)).," As previously noted (see Sect. \ref{sec:uno}) ),"
 the eiuissio- X4 YZ CMi is often considerably polarized., the emission of YZ CMi is often considerably polarized.
 Throughout lis observation the polarization was predominantly left circular (about 60 circular polarisation durius the quiescent enüssiou. and 90 curing the flares).," Throughout this observation the polarization was predominantly left circular (about 60 circular polarisation during the quiescent emission, and 90 during the flares)."
 Two strone flares appear in the lighteurve obtained with the VLA., Two strong flares appear in the lightcurve obtained with the VLA.
 Their flux values reach up to 6.0 12Jy., Their flux values reach up to 6.0 mJy.
 The detection of YZ CAG with both iustrunents was clear. as the ris oise in the VLA image was 0.063 muJv/beam. aud iu our high scusitivity VLBI maps made using oulv VLBA-VLA baselines the noise was 0.13 udv/beam.," The detection of YZ CMi with both instruments was clear, as the rms noise in the VLA image was 0.063 mJy/beam, and in our high sensitivity VLBI maps made using only VLBA-VLA baselines the noise was 0.13 mJy/beam."
 (νοι these map noise values the peak intensity of the star was 17 0 in the VLA image. and 16 σ in the VLBA imaec.," Given these map noise values the peak intensity of the star was 47 $\sigma$ in the VLA image, and 16 $\sigma$ in the VLBA image."
 The mean flux deusitv value for AD Leo was at a nore typica level of 0.7 αν (see Fig. l. ," The mean flux density value for AD Leo was at a more typical level of 0.7 mJy (see Fig. \ref{fig:light97}, ,"
right plot)., right plot).
 Siuluw fluxes at 18 ci have previously been reported (ee. Jacksonetal.1989: Bengetal. 1995)., Similar fluxes at 18 cm have previously been reported (e.g. \cite{Jack}; \cite{B2} 1995).
 A weak flare appears in the VLA liglhiteurve (2.1 αν)., A weak flare appears in the VLA lightcurve (2.1 mJy).
 The star was detected bv both instruments despite its low flux., The star was detected by both instruments despite its low flux.
 The nius noise for the VLA map was 0.035 mJv/beam. while for the map mace with only the VLA-VLBA baselines it was 0.18 indy/beam.," The rms noise for the VLA map was 0.035 mJy/beam, while for the map made with only the VLA-VLBA baselines it was 0.18 mJy/beam."
 Peak map values were 15.6 σ for the VLA and 5 σ for the VLBA aps respectively., Peak map values were 15.6 $\sigma$ for the VLA and 5 $\sigma$ for the VLBA maps respectively.
 Tracing the lightcurve of AD Leo from the interferometric VLA data was a delicate task because of two stroug sources inthe field of view., Tracing the lightcurve of AD Leo from the interferometric VLA data was a delicate task because of two strong sources inthe field of view.
whic,magnetic field.
h inthe semiclassical limitbehaves as EE ox theuu) Hence, Similarly the creation of the generic dyons in the external field can be described in terms of the closed classical trajectories.
 density ofthe angular momentum atweak field and ει=—e», The radius of the trajectory is fixed by the external field moreover the configuration is unstable.
 =€reads as which indicates (he inst, The negative mode in the spectrum of the fluctuations implies the bounce interpretation of such Euclidean solution.
ability. of(he SYAL vacuum in (heexternal graviphoton RAR 1-formfield.A littlebit surprisingly the prepotential itself isproportional to ," The probability of the pair production in the unit time per unit volume is obtained from the action calculated at the bounce solution )) Since the classical trajectory is the circle with the fixed radius it implies the nontrivial angular momentum $<J>$ from the $R^4$ Euclidean viewpoint in the $x,t_E$ plane."
the densityof the angularmomentum. To complete this Section loex<JL— vae," The density of the loop trajectories in the $t_E,x$ plane corresponds to the angular momentum density from the $R^4$ viewpoint."
 (2.9) which implies thatiteffectively field.In what follows we shall suggest the degreesofTreeclom, Since we are interested in the derivative of the prepotential with respect to the external field let us examine the similar expression in QED <R^2> that is the area inside the semiclassical trajectory.
 whichcould serveas such rotators. 3., In the semiclassical approximation $<R^2>\propto R_{cr}^2 \propto E^{-2}$ and therefore diverges at weak field.
 example The toy, This behavior reflects the instability of the system.
 Considerthe simplest exampleof theQED theory is unstable wilh respect, The above picture can be generalized for the multiple pair production problem when the multiple Euclidean bounces have to be taken into account.
 tothe charged pairs creation. The tunneling can be described viat, Contrary to the single pair case the generic probability involves the account of the interaction between the bounces which makes the problem more complicated.
he first quantizedpicturein terms of the classical," The density of the bounces at the 2-dimensional plane selects the phase corresponding to ""gas"" or the ""liquid"" state of the bounce rotators system."
 particle trajectories in the Euclidean space-time|13].., To make the link with ${\cal N}=2$ SYM case let us represent the pair production in the brane terms \cite{gss}.
 The just the circlesin the lp planewhose radii are," The bounce corresponding to the generic $(p,q)$ dyon production is the tubular $(p,q)$ string stretched between two D3 branes in IIB setup."
 fixedbythe extremization trajectoriesofthe simple, In the field theory limit the deformation of the cylinder due to the finite string tension is neglected.
 effective action Sopp=2am— Ex R7(3.10) The coordinat, In what follows we shall consider very similar configuration of the tubular D2 branes between the parallel D4 branes in IIA setup.
ex is selected bythe directionofthe extern, The brane viewpoint allows to discuss the tunneling problem from the viewpoint within worldvolume theory of the created brane.
alfield. This semiclassical motionin the Euclidean space-time canbe interpreted as (he motionin (he con, Recall that the Schwinger -type process involves two steps - the tunneling described by Euclidean bounce and the Minkowski evolution upon the analytic continuation at the turning point.
stant," For instance,"
ie flow is aligned with the equatorial plane of the black role. while at larger distances it is smoothly tilted to its oeutial orbital plauc.,"the flow is aligned with the equatorial plane of the black hole, while at larger distances it is smoothly tilted to its initial orbital plane."
 The aligument radius (the radius)) was originally been assumed to be of jo same order as the radial distance. where the precession iue scale (~wit) ix comparable to the iufall time scale ~RV» lV ds the radial velocity of the eas).," The alignment radius (the ) was originally been assumed to be of the same order as the radial distance, where the precession time scale $\sim \omega _{\mathrm{LT}}^{-1}$ ) is comparable to the infall time scale $\sim R\,V_{\mathrm{R}}^{-1}$, $V_{\mathrm{R}}$ is the radial velocity of the gas)."
 When jo internal hydrodvuamics of the disk is fully aken iuto account it appears to be much siualler (Ini Prinele 1985)., When the internal hydrodynamics of the disk is fully taken into account it appears to be much smaller (Kumar Pringle 1985).
 Accretion disks spin up nonrotatiue. or slowly rotating. black holes because the augular uomentiu per nass unit of the accretiug gas at the iunernmost stable orbit exceeds that of the hole (Misner et al.," Accretion disks spin up nonrotating, or slowly rotating, black holes because the angular momentum per mass unit of the accreting gas at the innermost stable orbit exceeds that of the hole (Misner et al."
 1973: Shapiro Teukolsky 1983)., 1973; Shapiro Teukolsky 1983).
 The (electromagnetic extraction of the black hole spin eucergv) is probably not cficicut enough to reduce completely this acciunulated spin momentum (Modersky et al., The (electromagnetic extraction of the black hole spin energy) is probably not efficient enough to reduce completely this accumulated spin momentum (Modersky et al.
 1997. Livio ct al.," 1997, Livio et al."
 1998)., 1998).
 Iu other words. an ACN black hole is iiost probably fast rotating («& 1)," In other words, an AGN black hole is most probably fast rotating $a\simeq 1$ )."
 As a result. the accretion disk in ACN shoulk be nouplanar iu most cases; as there is no reason to suppose that the spin momentum of the accretiug ass aud the black hole spin momentum are always aligued.," As a result, the accretion disk in AGN should be nonplanar in most cases, as there is no reason to suppose that the spin momentum of the accreting mass and the black hole spin momentum are always aligned."
 Evidence or the existence of uouplanar disks in ACN has been recently ound by Nishiura et al. (, Evidence for the existence of nonplanar disks in AGN has been recently found by Nishiura et al. (
1998).Obviously. such a nonplanar ecometry of the disk males it possible for the central raciation to reach the outer parts. as these can be directly “seen” from the ceuter (Petterson 1977).,"1998).Obviously, such a nonplanar geometry of the disk makes it possible for the central radiation to reach the outer parts, as these can be directly ”seen” from the center (Petterson 1977)."
 The covering factor of a nouplanar disk in case of ceutral nvraciation is close to its initial inclination angle. measured iu lx units.," The covering factor of a nonplanar disk in case of central irradiation is close to its initial inclination angle, measured in $4\mathrm{\pi}$ units."
 Tn this paper. the broad Balmer cussion due to reprocessing of the central hieh euergv radiation bv a warped accretion disk is investigated.," In this paper, the broad Balmer emission due to reprocessing of the central high energy radiation by a warped accretion disk is investigated."
 Hore we present profiles., Here we present profiles.
 The profiles of other strong low ionization lines(IIo. for mstance) could be slishtlv. differecut because of different reprocessing properties of the medium.," The profiles of other strong low ionization lines, for instance) could be slightly different because of different reprocessing properties of the medium."
 Our main purpose here is o demonstrate the effect of disk twisting ou the line profiles. without making fits to observational data.," Our main purpose here is to demonstrate the effect of disk twisting on the line profiles, without making fits to observational data."
 That is why we choose the simplest model of a disk with simall inclination aux a point N-ray source located near the ceuter., That is why we choose the simplest model of a disk with small inclination and a point X-ray source located near the center.
 In the next sections we describe our method. (Sect., In the next sections we describe our method (Sect.
 2) and results (Sect., 2) and results (Sect.
 3)., 3).
 Discussion. coniparisou with the observations. aud couclusious are eiven in Sect.," Discussion, comparison with the observations, and conclusions are given in Sect."
 1 Iu our model a viscous. ecometrically thin. aud warped because of the Barcdecu-Petterson effect accretion disk is hradiated bv a point-lke central source.," 4 In our model a viscous, geometrically thin, and warped – because of the Bardeen-Petterson effect – accretion disk is irradiated by a point-like central source."
 We use a ονμπαΊσα] coordinate svsteii. centered at the black hole (Fie.," We use a cylindrical coordinate system, centered at the black hole (Fig."
 1)., 1).
 The imadiatiug source is located at Zy10Πο above the disk plane. along the black hole axis instead of at the exact center. as this is a more realistic situation.," The irradiating source is located at $Z_{\mathrm{S}}=10\, 
R_{\mathrm{G}}$ above the disk plane, along the black hole axis instead of at the exact center, as this is a more realistic situation."
 The profiles of cutission as a result of the reprocessing of the central hard N-xav raciatiou are obtained., The profiles of emission as a result of the reprocessing of the central hard X-ray radiation are obtained.
 The disk thickness is neglected iu the computations it docs not affect the profile shapes significantly as the semithickness to radius ratio is usually much less than the tilt angle., The disk thickness is neglected in the computations – it does not affect the profile shapes significantly as the semithickness to radius ratio is usually much less than the tilt angle.
 All relativistic corrections are also neglected. because they are snall while our goal is to obtain qualitative results ouly.," All relativistic corrections are also neglected, because they are small while our goal is to obtain qualitative results only."
 The disk is assumed to be optically thick to visual light and N-vavs o no cussion from the lower surface (probably also illuminated) can reach an observer located above the A-Y plane., The disk is assumed to be optically thick to visual light and X-rays – no emission from the lower surface (probably also illuminated) can reach an observer located above the X-Y plane.
 Since the precession velocity is small compared to the Iseplerian orbital velocity Vg=οοςΠ9 the disk is a stationary structure aud can be treated as clue coniposed of concentric rings. Lavine iu differcut planes.," Since the precession velocity is small compared to the Keplerian orbital velocity $V_{\mathrm{K}}=c\,R_{\mathrm{G}}^{0.5}R^{-0.5}$, the disk is a stationary structure and can be treated as being composed of concentric rings, laying in different planes."
 Warp waves can propagate through the disk surface if α«cl1 (Papaloizou Pringle 1983). where e is the dimensionless viscosity parameter (Shakura Suuvaev 1973). but this is not the case for ACN accretion disks. where a is usually assumed to be 0.1-1.," Warp waves can propagate through the disk surface if $\alpha<<1$ (Papaloizou Pringle 1983), where $\alpha$ is the dimensionless viscosity parameter (Shakura Sunyaev 1973), but this is not the case for AGN accretion disks, where $\alpha$ is usually assumed to be 0.1-1."
 Each disk rius is defined by two Eulerian augles ο aud > (Fig., Each disk ring is defined by two Eulerian angles $\beta$ and $\gamma$ (Fig.
 1) aud its radius A., 1) and its radius $R$.
 For a stationary twisted accretiou disk. 2 and 5 are slowly varviug functions of the radial distance R: ο)= XR). ~σι) and μις>Jy. SIey»0. Rox»0. as is ornginalhe shown bv Bardecu Petterson (1975).," For a stationary twisted accretion disk, $\beta $ and $\gamma $ are slowly varying functions of the radial distance $R$; $\beta =\beta (R)$ , $\gamma =\gamma (R)$ and $\beta \mathrm{|}_{\mathrm{R\rightarrow \infty}}\longrightarrow \beta _{0}$, $\beta \mathrm{|}_{\mathrm{R\rightarrow 0}}\longrightarrow 0$, $\gamma \mathrm{|}_{\mathrm{R\rightarrow \infty}}\longrightarrow 0$, as is originally shown by Bardeen Petterson (1975)."
 In this paper we use the steady solution of Scheucr Feiler (1996). derived following Prinele (1992).," In this paper we use the steady solution of Scheuer Feiler (1996), derived following Pringle (1992)."
 This solution. shown in Fig.," This solution, shown in Fig."
 2. cau be preseuted analytically bv:," 2, can be presented analytically by:"
(( oh (AEG. )) ((,( ) ( ) ( ( 6; ).
 Recent cosmic microwave background observations strongly unply a spatially flat universe (cle (2000).. Roos&Harun-or-Bashid (2000))). which naturally motivates application ol the new cistauce formulae presented in $22 to a set of staudard caudles to estimate v.," Recent cosmic microwave background observations strongly imply a spatially flat universe \cite{dB}, \cite{RH}) ), which naturally motivates application of the new distance formulae presented in 2 to a set of standard candles to estimate $\nu$."
 We use he 60 SNe Ia from the combined Calan/Tololo + Supernova Cosmology Project (CT+S5CP) as oesented in Riessetal.(1998).. Perlinutteretal. (1999).," We use the 60 SNe Ia from the combined Calan/Tololo + Supernova Cosmology Project (CT+SCP) as presented in \cite{RA}, \cite{PS1}."
. Combining these data with those from he High-z SN search (Schunictetal. (1998))) would bring the total to about 1060 SNe., Combining these data with those from the High-z SN search \cite{SB}) ) would bring the total to about 100 SNe.
 However. we shall see that this somewhat complicated task would uot increase the uumbers of SNe enough o noticeably improve the estimate.," However, we shall see that this somewhat complicated task would not increase the numbers of SNe enough to noticeably improve the estimate."
 Rather than subject the data to an in-depth Bayesian re-analysis with the ackditioual beam illiugH parameter 7 included.H we merely use a 47> gooduess-ol-lit⋅⋅ estimation.," Rather than subject the data to an in-depth Bayesian re-analysis with the additional beam filling parameter $\nu$ included, we merely use a $\chi^2$ goodness-of-fit estimation."
⋅ We- assume an intrinsicH SN absolute maguitude of A45=—19.33 aud Ay=65 kms 1 Hand follow the same procedure as Wang(2000a) to recover results consistent with Perlmutteretal.(1999) when v=0. Le.. [or observatious in a lLomogeneous universe.," We assume an intrinsic SN absolute magnitude of $M_B=-19.33$ and $H_0=65$ km $^{-1}$ $^{-1}$ and follow the same procedure as \cite{WY2000a} to recover results consistent with \cite{PS1} when $\nu=0$, i.e., for observations in a homogeneous universe."
" Then we employ the formulae presented here to obtain coulicdence contours iu the η plane. constraiulug octQ4 so that Qy=1.0—9,,."," Then we employ the formulae presented here to obtain confidence contours in the $\OM$ $\nu$ plane, constraining $\OL$ so that $\OL=1.0-\OM$."
Soap: By proceed[umi this way we are assiuning that iulioniogeneous matter. e.g.. galaxies. are sufficiently. removed from the lines of site of the 60 SNe ancl that lensing is negligible.," By proceeding this way we are assuming that inhomogeneous matter, e.g., galaxies, are sufficiently removed from the lines of site of the 60 SNe and that lensing is negligible."
the following telescopes.,the following telescopes.
 The ΠουΕαν Telescope (IIET) is a. joimt project of the University of Texas at Austin. the Penusylvania State University. Stautord Universitv.. Ludwie-Maxinilliaus-Universitàtt Ablüunchen. and CGeorg-August-Universitàtt Cotttingen.," The Hobby-Eberly Telescope (HET) is a joint project of the University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximillians-Universitätt Münnchen, and Georg-August-Universitätt Götttingen."
 The WET is named in honor of its principal benefactors. Willam DP. Dobby and Robart E. Eberh.," The HET is named in honor of its principal benefactors, William P. Hobby and Robert E. Eberly."
 The Marcario Low-Resolution Spectrograph is aimed for Mike Marcario of Wigh Lonesome Optics. who fabricated several optical clemeuts for the instrument but died before its completion: it is a joint project of the IIobby-Eberly Telescope partnership aud the Instituto de Astronomiaa de la Universidad Nacional Autéunomace Meéxsico.," The Marcario Low-Resolution Spectrograph is named for Mike Marcario of High Lonesome Optics, who fabricated several optical elements for the instrument but died before its completion; it is a joint project of the Hobby-Eberly Telescope partnership and the Instituto de Astronomíaa de la Universidad Nacional Autónnomade Méxxico."
 The Apache Point Observatory 3.5 11 telescope is owned aud operated bv the Astroplivsical Research Consortium., The Apache Point Observatory 3.5 m telescope is owned and operated by the Astrophysical Research Consortium.
 We thank the observatory director. Suzauue Hawley. aud site manager. Bruce Cuüllespie. for their support of this project.," We thank the observatory director, Suzanne Hawley, and site manager, Bruce Gillespie, for their support of this project."
 The Subaru Telescope is operated bv the National Astronomical Observatory of Japan., The Subaru Telescope is operated by the National Astronomical Observatory of Japan.
 The Willam IHerschel Telescope is operated bw the Isaac Newton Croup on the islaud of La Palma in the Spanish Observatorio del Roque de los Muchliachos of tle Tustituto de Astrofisica de Canarias., The William Herschel Telescope is operated by the Isaac Newton Group on the island of La Palma in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias.
 The W. M. Keck. Observatory is operated as a scieutific partnership among the Califormia Iustitute of Techuolosv. the University of California. and the National Acronautics anc Space Achuinistration: the observatory was nde possible by the eenerous financial support of the W. M. necks Foundation.," The W. M. Keck Observatory is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration; the observatory was made possible by the generous financial support of the W. M. Keck Foundation."
 This work was funded by the NASA ADP Program NNXNOOACT5C.Sloan.GALENX..UKIRT..," This work was funded by the NASA ADP Program NNX09AC75G.,."
production mechanism and/or ejecta distribution (e.g.. bipolar outflow,"production mechanism and/or ejecta distribution (e.g., bipolar outflow)."
" Looking). forward. the inclusion of spectroscopic diagnostics (e.g.. vyy measurements. He I feature intensities) would enable an extension of this systematic study to search for additional correlations and to properly break the model degeneracies between M, and Εκ."," Looking forward, the inclusion of spectroscopic diagnostics (e.g., $v_{\rm ph}$ measurements, He I feature intensities) would enable an extension of this systematic study to search for additional correlations and to properly break the model degeneracies between $M_{\rm ej}$ and $E_K$."
 Along this line. we reiterate that ten of the SNe in our study were observed as part of the CCCP and thus spectroscopic observations (often extensive) exist for this sub-sample.," Along this line, we reiterate that ten of the SNe in our study were observed as part of the CCCP and thus spectroscopic observations (often extensive) exist for this sub-sample."
 This project will be the focus of a separate study., This project will be the focus of a separate study.
 Furthermore. a direet comparison of host galaxy diagnostics with optical light curve properties (e.g.. peak luminosity) may reveal further clues on the nature of SN Ibe progenitors and their relation to those of engine-driven explosions.," Furthermore, a direct comparison of host galaxy diagnostics with optical light curve properties (e.g., peak luminosity) may reveal further clues on the nature of SN Ibc progenitors and their relation to those of engine-driven explosions."
 Since our sample is drawn from targeted surveys. itis likely biased toward metal-rich host galaxies.," Since our sample is drawn from targeted surveys, it is likely biased toward metal-rich host galaxies."
 Fortunately. new wide-field transient surveys such as Pan-STARRS (Kaiserefa£..2002) and the Palomar Transient Factory (Laweta£.2009) are discovering SNe in a broad range of host galaxy environments and will enable such comparisons to be made in an unbiased fashion.," Fortunately, new wide-field transient surveys such as Pan-STARRS \citep{kab+02} and the Palomar Transient Factory \citep{lkd+09} are discovering SNe in a broad range of host galaxy environments and will enable such comparisons to be made in an unbiased fashion."
 Moreover. thanks to the higher cadence of these optical surveys. SNe Ibe will be discovered earlier with respect to the explosion date. reducing the fraction of Bronze SNe in future follow-up samples.," Moreover, thanks to the higher cadence of these optical surveys, SNe Ibc will be discovered earlier with respect to the explosion date, reducing the fraction of Bronze SNe in future follow-up samples."
 Finally we note that the significant E(B—Vyo values inferred here motivates the regular use of near-IR facilities in future studies of SNe Ibc., Finally we note that the significant $E(B-V)_{\rm host}$ values inferred here motivates the regular use of near-IR facilities in future studies of SNe Ibc.
 M.R.D. and A.M.S. acknowledge support by the National Science Foundation Research Experiences for Undergraduates (REU) and Department of Defense Awards to Stimulate and Support Undergraduate Research Experiences (ASSURE) programs under Grant no., M.R.D. and A.M.S. acknowledge support by the National Science Foundation Research Experiences for Undergraduates (REU) and Department of Defense Awards to Stimulate and Support Undergraduate Research Experiences (ASSURE) programs under Grant no.
 0754568 and by the Smithsonian Institution., 0754568 and by the Smithsonian Institution.
 The work of A.G. is supported by grants from the Israeli Science Foundation (ISF). an EU/FP7 Marie Curie IRG Fellowship and a research grant from the Gruber Awards.," The work of A.G. is supported by grants from the Israeli Science Foundation (ISF), an EU/FP7 Marie Curie IRG Fellowship and a research grant from the Gruber Awards."
 S.B.C acknowledges generous support from Gary and Cynthia Bengier and the Richard and Rhoda Goldman Foundation., S.B.C acknowledges generous support from Gary and Cynthia Bengier and the Richard and Rhoda Goldman Foundation.
 The work of D.C.L. is supported by National Science Foundation (NSF) grant AST-1009571., The work of D.C.L. is supported by National Science Foundation (NSF) grant AST-1009571.
" D.C.L. is grateful for an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0401479, during which part of this work was completed."," D.C.L. is grateful for an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0401479, during which part of this work was completed."
SALA. the Weck IE and. Very Large Telescopes (VLT).,"SMA, the Keck II and Very Large Telescopes (VLT)."
 The time coverage on this day was excellent, The time coverage on this day was excellent
"For the developments of (2)?cos2A, the same conclusion is obtained.","For the developments of $(\frac{a}{r})^3 \cos 2\lambda$, the same conclusion is obtained."
" Fig.13 shows the discrepancies between the numerical and the semi analytical results using the developments of 23cos2A, to 3rd and 6th order."," \ref{fig9} shows the discrepancies between the numerical and the semi analytical results using the developments of $\frac{a}{r}^3 \cos 2\lambda$, to 3rd and 6th order."
" The difference is descreased from 0"".3 to 07.01 (peak to peak)."," The difference is descreased from 0"".3 to 0"".01 (peak to peak)."
" The developments are given by : To express (2)? and (2)?cos2A by purely analytical functions, Kinoshita replaced e,M,Ls by their linear mean values."," The developments are given by : To express $\frac{1}{2}(\frac{a}{r})^3$ and $(\frac{a}{r})^3 \cos 2\lambda$ by purely analytical functions, Kinoshita replaced $e, M, L_{S}$ by their linear mean values."
" For Venus, the same method is used in Cottereau and Souchay (2009), who replaced e,M,Ls by their linear mean values given by Simon et al (1994)."," For Venus, the same method is used in Cottereau and Souchay (2009), who replaced $e, M, L_{S}$ by their linear mean values given by Simon et al (1994)."
" Using the mean elements of the section (??)) and the developments given by (30)) and (31)), an analytical model is obtained."," Using the mean elements of the section \ref{2}) ) and the developments given by \ref{ar}) ) and \ref{arlam}) ), an analytical model is obtained."
 As explained in Kinoshita and Souchay (1990) and Cottereau et al. (, As explained in Kinoshita and Souchay (1990) and Cottereau et al. (
"2010) for the Earth and for Venus, the difference between the nutation computed by analytical tables and the numerical integration is small (at the order of 107?) and caused by the indirect planetary effects onthe planet.","2010) for the Earth and for Venus, the difference between the nutation computed by analytical tables and the numerical integration is small (at the order of $10^{-5}$ ) and caused by the indirect planetary effects onthe planet."
" For Phoebe, this difference is also"," For Phoebe, this difference is also"
with the major axis in this galaxy. the PV plot is alfected by the elliptical streaming of gas in the bar.,"with the major axis in this galaxy, the PV plot is affected by the elliptical streaming of gas in the bar."
 This explains why the rotation curve has a velocity at the bulge radius lower (han that seen in the PV diagram: this was also evident when we compared our rotation curve wilh the PV plot of Sakamoto. Baker and Scoville (2000).," This explains why the rotation curve has a velocity at the bulge radius lower than that seen in the PV diagram; this was also evident when we compared our rotation curve with the PV plot of Sakamoto, Baker and Scoville (2000)."
 However. the beam smearineg ellect is important for mainly the inner 2-3 beamwidths. which is well within (he bulge radius for this galaxy.," However, the beam smearing effect is important for mainly the inner 2-3 beamwidths, which is well within the bulge radius for this galaxy."
 Also. the effect of elliptical streaming is considerably reduced in deriving the rotation curve because velocities are azimuthallv averaged over annuli.," Also, the effect of elliptical streaming is considerably reduced in deriving the rotation curve because velocities are azimuthally averaged over annuli."
 We thus believe our isophote measurement of the bar ellipticity and (he rotation curve determination of aare both reasonable estimates for NGC 5005.3184:, We thus believe our isophote measurement of the bar ellipticity and the rotation curve determination of are both reasonable estimates for NGC 5005.:
 This galaxy has both a large bulge and a round bar., This galaxy has both a large bulge and a round bar.
 The bar is not prominent in the Ix band image., The bar is not prominent in the K band image.
 However the CO distribution shows the classic response of gas in a barred potential. with trailing spiral arms emereine from the ends of the bar.," However the CO distribution shows the classic response of gas in a barred potential, with trailing spiral arms emerging from the ends of the bar."
 The bulge is a little over half the bar size and is very. bright in the near-IR. which max indicate a large nass concentration in (he center of the galaxy.:," The bulge is a little over half the bar size and is very bright in the near-IR, which may indicate a large mass concentration in the center of the galaxy.:"
 The ellipticity of this egalaxy is so low. it mightoO be questionable whether it is in fact à barred galaxy.," The ellipticity of this galaxy is so low, it might be questionable whether it is in fact a barred galaxy."
 But first we note that this galaxy is classified as a barred galaxy in (he RC3 catalogue (Table 1)., But first we note that this galaxy is classified as a barred galaxy in the RC3 catalogue (Table 1).
 ILowever. the bar is not easy to distinguish.," However, the bar is not easy to distinguish."
 This is partly because of the low ellipticity of the bar but also because of the high inclination of the galaxy (627)., This is partly because of the low ellipticity of the bar but also because of the high inclination of the galaxy ).
 We found evidence for the bar in the Ix band photometry ancl also in the intensity profile along the major axis of the galaxy., We found evidence for the bar in the K band photometry and also in the intensity profile along the major axis of the galaxy.
 Also. Zeilinger et al. (," Also, Zeilinger et al. ("
2001) find evidence for the bar both in their B. band photometry ancl in the stellar kinematics derived fom (he long slit spectra obtaimed along different axes in the galaxy.,2001) find evidence for the bar both in their R band photometry and in the stellar kinematics derived from the long slit spectra obtained along different axes in the galaxy.
 This leads us to believe that there is a rather round bar in (his galaxy., This leads us to believe that there is a rather round bar in this galaxy.
 The bulge is verv bright ancl very. large: it il over 3/4 the size of the bar and this indicates that it may also be very. massive., The bulge is very bright and very large; it it over 3/4 the size of the bar and this indicates that it may also be very massive.
 It therelore appears that the three galaxies with low measured elliplicilies are indeed barred galaxies with relatively low ellipticity and with high central mass concentrations., It therefore appears that the three galaxies with low measured ellipticities are indeed barred galaxies with relatively low ellipticity and with high central mass concentrations.
 Nonetheless. it will be important to confirm the trend identified here using a larger sample ol galaxies.," Nonetheless, it will be important to confirm the trend identified here using a larger sample of galaxies."
" The last issue which should be discussed is the effect of elliptical streaming in the bar. due to gas moving on c, and s orbits."," The last issue which should be discussed is the effect of elliptical streaming in the bar, due to gas moving on $x_{1}$ and $x_{2}$ orbits."
 The wry orbits are aligned along the bar within corotation radius and (he c» orbils are elongated perpendicular to the leneth of the bar (Contopoulos Papavannopoulos 1980: Binney Tremaine 1987)., The $x_{1}$ orbits are aligned along the bar within corotation radius and the $x_{2}$ orbits are elongated perpendicular to the length of the bar (Contopoulos Papayannopoulos 1980; Binney Tremaine 1987).
 We have tried (o minimize the effect of elliptical streaming bv measuring velocities al (he bulge and bar radii., We have tried to minimize the effect of elliptical streaming by measuring velocities at the bulge and bar radii.
" We assume that we are measuring the or» orbits at the edge of the bulge and the xy, orbits at the edge of the bar.", We assume that we are measuring the $x_2$ orbits at the edge of the bulge and the $x_1$ orbits at the edge of the bar.
 For a bar aligned with the major axis of the galaxy. (jj 1s measured al the pericenter of its cc» orbit and (therefore larger (han a circular orbit consistent with the," For a bar aligned with the major axis of the galaxy, $v_{blg}$ is measured at the pericenter of its $x_2$ orbit and therefore larger than a circular orbit consistent with the"
eud. the rate of energy input actually varies little. when averaged over a sufficiently long perio.,"end, the rate of energy input actually varies little, when averaged over a sufficiently long period."
 What if the cluster is seeded with mauy binaries initially?, What if the cluster is seeded with many binaries initially?
 Figure LO shows the evolution of ie. top-level energy [or the Pleiacles simulation., Figure 10 shows the evolution of the top-level energy for the Pleiades simulation.
 There are now many Diaries even at the start. and thus uo initial period of constant. £p.," There are now many binaries even at the start, and thus no initial period of constant $E_{\rm top}$."
 Remarkably. however. the evolutiou is quite similar to the case of no primordial binaries.," Remarkably, however, the evolution is quite similar to the case of no primordial binaries."
 The top-level euergy is changed iu a few discrete jumps., The top-level energy is changed in a few discrete jumps.
 These [ew major intera‘ious always involve binary (or triple) svstems composed of the few most massive stars (see also clelaFuenteMarcos(1996b)), These few major interactions always involve binary (or triple) systems composed of the few most massive stars (see also \citet{fm96b}) ).
 The importaut lessou is that ouly a special subset of binaries strongly. influences a cluster’s evolution., The important lesson is that only a special subset of binaries strongly influences a cluster's evolution.
 These are systems which are relatively massive. wide euough to have a siguilicaut interaction Cross section with other stars. aud yet tight enough to be hard.," These are systems which are relatively massive, wide enough to have a significant interaction cross section with other stars, and yet tight enough to be hard."
 To be sure. the primordial binaries in the Pleiacles-like simulatiou shown in Figure 7 o halt the initial contraction.," To be sure, the primordial binaries in the Pleiades-like simulation shown in Figure 7 do halt the initial contraction."
 Relatively little energy iuput is required to do so., Relatively little energy input is required to do so.
 Virtually all primordial binaries are either of too low a ass. Or are so tight that they effectively interact as a single system.," Virtually all primordial binaries are either of too low a mass, or are so tight that they effectively interact as a single system."
 It dis the subsequent coupling ol relatively few iuassive stars that inject much greater euergy auc principally drive the clusters expansion., It is the subsequent coupling of relatively few massive stars that inject much greater energy and principally drive the cluster's expansion.
 We have seen how binary heating cau dominate a clusters evolution., We have seen how binary heating can dominate a cluster's evolution.
 For a realistic stellar Dass specirum. tie. effect begins very quickly. in less than a single relaxation time.," For a realistic stellar mass spectrum, the effect begins very quickly, in less than a single relaxation time."
 Under these circumstances. the cluster is still very far from the point of true core collapse.," Under these circumstances, the cluster is still very far from the point of true core collapse."
 Followiug ou: Pleiades study (Paper I). we liave set the maximum mass at μας=10M...," Following our Pleiades study (Paper I), we have set the maximum mass at $m_{\rm max} = 10\,\Msun$."
 The reasoniug heὁ was that more massive objects would have ionized the parent cloud. allowing the stars to clispe‘se before they could form a bouud cluster.," The reasoning here was that more massive objects would have ionized the parent cloud, allowing the stars to disperse before they could form a bound cluster."
 In any event. it is instructive when elucidating basic johysical principles. to relax this assumptiou aud gauge the effect.," In any event, it is instructive when elucidating basic physical principles, to relax this assumption and gauge the effect."
 We uow allow stars in our No—1096 cluster to be drawn from the same mass funetion as before. but with a nominal upper lirjt of naa=LOOAL..," We now allow stars in our $N = 4096$ cluster to be drawn from the same mass function as before, but with a nominal upper limit of $m_{\rm max} = 100\,\Msun$."
 In practice. no star ever realizes this mass: the largest eeneratec is abou 0AL...," In practice, no star ever realizes this mass; the largest generated is about $60\,\Msun$."
 Again. there are no primordial binaries.," Again, there are no primordial binaries."
 Based ou our earlier arguments. we would expect binary formation to begin even soouer.," Based on our earlier arguments, we would expect binary formation to begin even sooner."
 ludeed this is the case., Indeed this is the case.
 Tie first stable. hard binary forms atqui. a factor of 2 earlier in time.," The first stable, hard binary forms at, a factor of 2 earlier in time."
 Figure 11 shows tie evolution of both Lagrangian mass aud ΠΟΔΙ shells., Figure 11 shows the evolution of both Lagrangian mass and number shells.
 By either measure. the cluster undergoes elobal expausiou at all radii.," By either measure, the cluster undergoes global expansion at all radii."
 Not surprisingly. there are detailed differences from the μιας=10M. case.," Not surprisingly, there are detailed differences from the $m_{\rm max} = 10\,\Msun$ case."
 The apparent early contraction of the inueruiost mass shells is nowernlerely due to 1ass segregation., The apparent early contraction of the innermost mass shells is now due to mass segregation.
 Even the number shell with δη=0.03 expands from the start., Even the number shell with $N_r / N_0 = 0.03$ expands from the start.
 A closer analysis shows that there is agaiu never more than a single binary in the cluster at auy iustaut. although the specific pair cliauges identity in time.," A closer analysis shows that there is again never more than a single binary in the cluster at any instant, although the specific pair changes identity in time."
 Both components of this pair are always among the top 5 stars by mass., Both components of this pair are always among the top 5 stars by mass.
 These binaries generate especially strong heating., These binaries generate especially strong heating.
ralio (SNR) and peak magnitude: these factors ensure (hat it is almost always al a sullicientlv high elevation to minimize atmospheric opacity contributions.,ratio (SNR) and peak magnitude; these factors ensure that it is almost always at a sufficiently high elevation to minimize atmospheric opacity contributions.
 As phase self-calibration is used. our final VLBA images of the SiO maser emission toward TX Cam retain no absolute astrometric information or alignment.," As phase self-calibration is used, our final VLBA images of the SiO maser emission toward TX Cam retain no absolute astrometric information or alignment."
 In Paper I. zeroth Ireequency-monment total intensity images from successive epochs were iniGally aligned using spatial eross-correlation.," In Paper I, zeroth frequency-moment total intensity images from successive epochs were initially aligned using spatial cross-correlation."
 This initial alignment was then refined using feature-based registration in which the deviation in proper motion paths predicted for individual maser components was minimized., This initial alignment was then refined using feature-based registration in which the deviation in proper motion paths predicted for individual maser components was minimized.
 For the polarization and total intensity images produced and reported in the current reduction. we adopt the alignment frame developed in Paper I as an absolute relerence.," For the polarization and total intensity images produced and reported in the current reduction, we adopt the alignment frame developed in Paper I as an absolute reference."
 Each of the zeroth-moment total intensity images produced in the current work was reeistered relative to ils corresponding epoch in the Paper I frame using Fourier-based matched filtering., Each of the zeroth-moment total intensity images produced in the current work was registered relative to its corresponding epoch in the Paper I frame using Fourier-based matched filtering.
 A robust. weighted-sium estimator of the filter output. peak was used to measure (he (two-dimensional spatial offsets.," A robust, weighted-sum estimator of the filter output peak was used to measure the two-dimensional spatial offsets."
 This was found to out-perform an elliptical Gaussian fit to the peak. likely because it is less sensitive to small morphological differences in (he total intensity. images being aligned.," This was found to out-perform an elliptical Gaussian fit to the peak, likely because it is less sensitive to small morphological differences in the total intensity images being aligned."
 Based on manual cross-checks of the positions of individual SiO maser components. we believe the registration of our polarization images relative to the alignment. [rame determined in Paper I has an accuracy. of 0056s. or less.," Based on manual cross-checks of the positions of individual SiO maser components, we believe the registration of our polarization images relative to the alignment frame determined in Paper I has an accuracy of $100 \mu as$ or less."
 Note (that the images in Paper I used to produce the Stokes J movie published (here were enhanced in contrast by applying the transformation VI and a subsequent threshold. cutoll V1<0.7., Note that the images in Paper I used to produce the Stokes $I$ movie published there were enhanced in contrast by applying the transformation $\sqrt{I}$ and a subsequent threshold cutoff $\sqrt{I} < 0.7$.
 For consistency. the same transformation was applied to the total intensity images reduced here belore the matched filter alignment.," For consistency, the same transformation was applied to the total intensity images reduced here before the matched filter alignment."
 The calibrators. JO3859+509. 3C454.3. ancl JOGO9-157 were reduced. separately using continnuum polarization VLBI techniques for millimeter wavelength observations described bv Ixemball.Diaanond.&Paulinv-Toth(1996).," The calibrators, J0359+509, 3C454.3, and J0609-157 were reduced separately using continnuum polarization VLBI techniques for millimeter wavelength observations described by \citet{kemball96}."
. As the only calibrator source with sullicient parallaclic augle coverage. continuum imstrumental polarization leakage terms were determined from the observations of J03594-509 alone and applied to the other calibrator sources.," As the only calibrator source with sufficient parallactic angle coverage, continuum instrumental polarization leakage terms were determined from the observations of J0359+509 alone and applied to the other calibrator sources."
 Continuum polarization images in Stokes (/.Q.U) were (hen produced [or each calibrator at each epoch.," Continuum polarization images in Stokes $(I,Q,U)$ were then produced for each calibrator at each epoch."
 As noted above. the calibrators 3C454.3 and J0609-157 were only observed for a single 13-min scan al each epoch. primarily as visibili. calibrators. and have limited inge sensilivily and poor we-coverage as a result.," As noted above, the calibrators 3C454.3 and J0609-157 were only observed for a single 13-min scan at each epoch, primarily as visibility calibrators, and have limited image sensitivity and poor $uv$ -coverage as a result."
 The image quality Lor (hese sources was accordingly often poor. especially for 3454.3. which has more complex structure &Paulinv-Toth1996). than the more compact JO0609-157.," The image quality for these sources was accordingly often poor, especially for 3C454.3, which has more complex structure \citep{kemball96} than the more compact J0609-157."
 The absolute EVPA of anv linearlv-polarized emission measured in our VLBA observations, The absolute EVPA of any linearly-polarized emission measured in our VLBA observations
Nearly all known CRB with known red shifts lie either ou or to one side of the Amati relation. which is usually shown as a plot of CRB spectral peak euergy. E44. ou the v-axis. as a function of the isotropic equivaleut ΟΠΟΙΟΥ £j.,"Nearly all known GRB with known red shifts lie either on or to one side of the Amati relation, which is usually shown as a plot of GRB spectral peak energy $E_{peak}$, on the y-axis, as a function of the isotropic equivalent energy $E_{iso}$."
 Phi: can be understood as beige due to the observers light of sight Icing offset relative to the direction of motion of the material at which the ciission or last scattering took place (Eichler&Levinson2001:Yamazakietal. 2001).," This can be understood as being due to the observer's light of sight being offset relative to the direction of motion of the material at which the emission or last scattering took place \citealt{EL04, Yamazaki04}) )."
 However. the absence of CRB below the vedge” - defined by the Amati relation to be Eyο10rg]3 AIeV - ds curious.," However, the absence of GRB below the “edge"" - defined by the Amati relation to be $E_{peak} \sim [E_{iso}/10^{54}erg]^{1/2}$ MeV - is curious."
 For it states that the isotropic equivalent photon uuuber ima CRB decreases in proportion to the peak photon cucrey., For it states that the isotropic equivalent photon number in a GRB decreases in proportion to the peak photon energy.
 It thus strongly coustraius scatter iu the exteut of any process Within a GRD fireball that reduces photon cucrey without decreasingthe ummber of photous. c.g. adiabatic losses of photons enütted im au optically thick region.," It thus strongly constrains scatter in the extent of any process within a GRB fireball that reduces photon energy without decreasing the number of photons, e.g. adiabatic losses of photons emitted in an optically thick region."
 Simularly. it coustraims scatter iu the number of photous enütted iu CBR having a given spectral peal. c.g. the uunuber of energetic particles cutting iu an optically thin region ina given magnetic field.," Similarly, it constrains scatter in the number of photons emitted in GBR having a given spectral peak, e.g. the number of energetic particles emitting in an optically thin region in a given magnetic field."
 That the lower right side of the Amati eraph is free of any CRB of known redshift is particularly reniukable considering that it represeuts the brightest GRBs at a eiven spectral peak., That the lower right side of the Amati graph is free of any GRB of known redshift is particularly remarkable considering that it represents the brightest GRBs at a given spectral peak.
" It suggests that there is no such thing as a ""lightly. compromised CRB whose plotous are softened without beiug lost. c.g. as iu a dirty fireball."," It suggests that there is no such thing as a “slightly compromised"" GRB whose photons are softened without being lost, e.g. as in a dirty fireball."
 Iu this letter. however. we sueeest that the N-vay flash (NRF) 060218 is a dirty fireball.," In this letter, however, we suggest that the X-ray flash (XRF) 060218 is a dirty fireball."
 We note that. uulike most ARF. the low cnereyv componcut las. despite its ιο softer spectrum. a simular photon nuuber to classical GRB. as opposed to most NRF. which have far fewer photons.," We note that, unlike most XRF, the low energy component has, despite its much softer spectrum, a similar photon number to classical GRB, as opposed to most XRF, which have far fewer photons."
 NRF 060218 was detected with the Burst Alert Telescope (BAT) instrument onboard Swift spacecraft., XRF 060218 was detected with the Burst Alert Telescope (BAT) instrument onboard Swift spacecraft.
 The spectrum peaks below 10 keV. thus classifviug this transient as au NRF.," The spectrum peaks below 10 keV, thus classifying this transient as an XRF."
 ARF 060218 is distinguished by its unprecedenutedlv long duration (2000 sec) with a snooth lieht curve., XRF 060218 is distinguished by its unprecedentedly long duration $\sim 2000$ sec) with a smooth light curve.
 The lisht curves show a siguificaut spectral lag. with soft photons lageiug behind the hard photons. as usually seen iu loug CRBs (Norrisctal. 2000)).," The light curves show a significant spectral lag, with soft photons lagging behind the hard photons, as usually seen in long GRBs \citealt{Norris00}) )."
" The peak of the light curves are at LOS+25. T3HAYD, 91947 and 1082-123 seconds (Liangetal. 2006)) in the energy band (15-150) keV. (5-10) keV. (2-5) keV and (0.3-2) keV respectively,"," The peak of the light curves are at $405\pm25$, $735\pm9$, $919\pm7$ and $1082\pm13$ seconds \citealt{Liang06}) ) in the energy band (15-150) keV, (5-10) keV, (2-5) keV and (0.3-2) keV respectively."
 The NRT spectrum (0.5-10 keV) shows a thermal componucut in soft N-rav. with temperature ATz0.17 keV (Campanaetal.2006))., The XRT spectrum (0.3-10 keV) shows a thermal component in soft X-ray with temperature $KT\approx 0.17$ keV \citealt{Campana06}) ).
 The high euerev spectra (15-150 keV) from BAT show spectral softening with time., The high energy spectra (15-150 keV) from BAT show spectral softening with time.
 The optically discovered supernova 200G6a7 Is associated with NRF 060215., The optically discovered supernova 2006aj is associated with XRF 060218.
 The optical afterglow spectra aud the strong cussion lines from the host galaxy represcut a redshift of +1.033120.00002. corresponding to a distance of 110 Mpc (Pianetal.200633.," The optical afterglow spectra and the strong emission lines from the host galaxy represent a redshift of $z=0.03342\pm0.00002$, corresponding to a distance of $\sim 140$ Mpc \citealt{Pian06}) )."
 The isotropic equivalent prompt energv release is E;=(6.2to cre (Campanaetal. 2006))., The isotropic equivalent prompt energy release is $E_{iso}=(6.2\pm0.3)\times 10^{49}$ erg \citealt{Campana06}) ).
 The unusually long duration of this eveut raises the question of whether the central engine goes on this long. or whether the duration represents something downstream of the central cueine. e.g. shock breakout through the surface of the star or a light echo from slowly ovine. exteuded material.," The unusually long duration of this event raises the question of whether the central engine goes on this long, or whether the duration represents something downstream of the central engine, e.g. shock breakout through the surface of the star or a light echo from slowly moving, extended material."
 The chromaticity of the duration of the event aud the location of the peak is sienificaut., The chromaticity of the duration of the event and the location of the peak is significant.
 It challenges breakout and light eclio models. in which the duration is established by lydrocvuatniics.," It challenges breakout and light echo models, in which the duration is established by hydrodynamics."
 We focus on this issue., We focus on this issue.
 We propose that the radiation pressure of the photons ou the matter can accelerate it up to relativistic Lorentz factors. which makes the duration appear longer than the actual activity of the ceutral eugine.," We propose that the radiation pressure of the photons on the matter can accelerate it up to relativistic Lorentz factors, which makes the duration appear longer than the actual activity of the central engine."
 We show that the iutrinsic duration of central euge activity need not be pathologically long for au NRF. but appears longer because of the relativistic motion of the scatterer.," We show that the intrinsic duration of central engine activity need not be pathologically long for an XRF, but appears longer because of the relativistic motion of the scatterer."
 We also show that the apparent duration is expected to be wavelength dependent. aud we ft the observed elt curves.," We also show that the apparent duration is expected to be wavelength dependent, and we fit the observed light curves."
 We asstuuc that the primary raciation is scattered by an extended barvouic cloud of optical depth unity aud hat the scattered radiation is seen by an observer at au angele @ with the motion of the scatterer (Fig., We assume that the primary radiation is scattered by an extended baryonic cloud of optical depth unity and that the scattered radiation is seen by an observer at an angle $\theta$ with the motion of the scatterer (Fig.
 1)., 1).
 The cutive scattering cloud is within the cone of the primary jud radiation aud the observer sees the lard radiation due to scattering from the cloud whereas the extended soft component may also reach the observer directly (sce low)., The entire scattering cloud is within the cone of the primary hard radiation and the observer sees the hard radiation due to scattering from the cloud whereas the extended soft component may also reach the observer directly (see below).
 The scatterer itself is offset from the axis of the ormuary radiation and scatters oulv the soft frinecs of, The scatterer itself is offset from the axis of the primary radiation and scatters only the soft fringes of
The distribution of matter in the Universe ou larec scales is efficieutlv quantified by the power spectrmu of its overdensity fluctuations.,The distribution of matter in the Universe on large scales is efficiently quantified by the power spectrum of its overdensity fluctuations.
 This is because to a good approximation. the density field is a Caussian random field. at carly times (as established by observations of the cosmic nücrowave backeround. CMD). or on large scales at low redshift.," This is because to a good approximation, the density field is a Gaussian random field, at early times (as established by observations of the cosmic microwave background, CMB), or on large scales at low redshift."
 The problem of efficient structure quantification is interesting informatiou-theorctically. and also practically. for coustrainiug cosmological parameters.," The problem of efficient structure quantification is interesting information-theoretically, and also practically, for constraining cosmological parameters."
" While the power spectrum of the overdensity 6=(p.p)/p is the optimal statistic ou the largest scales. on translinear scales (0.2<Kk/[h tS0.8 at redshift 2=0 in a concordance ACDAL model). the darkanatter deusitv field departs substantially from Gaussianity. and the power spectrum covariance matrix develops a significant non-Gaussian component (2??).,"," While the power spectrum of the overdensity $\delta=(\rho-\bar\rho)/\bar\rho$ is the optimal statistic on the largest scales, on translinear scales $0.2\lesssim
k/[$ $]\lesssim 0.8$ at redshift $z=0$ in a concordance $\Lambda$ CDM model), the dark-matter density field departs substantially from Gaussianity, and the power spectrum covariance matrix develops a significant non-Gaussian component \citep{mw,szh,t09}."
 This leads to a plateau in Fisher information (?777): when measuring a parameter such as the initial power spectrum amplitude. modes in the translinear regime are lighly correlated. giving little additional coustraimine power when analyzed with larger-scale modes.," This leads to a plateau in Fisher information \citep{rh05,rh06,ns06,leepen}; when measuring a parameter such as the initial power spectrum amplitude, modes in the translinear regime are highly correlated, giving little additional constraining power when analyzed with larger-scale modes."
 This signals a failure of the power spectrum to describe the 8 field fully on these scales. and more practically. inuplies a substantial reduction in its power to constraiucosmological paraicters.," This signals a failure of the power spectrum to describe the $\delta$ field fully on these scales, and more practically, implies a substantial reduction in its power to constraincosmological parameters."
 The number of Fourier modes erows as k? for a 3D survey. aud so it would be a shame if these stualler-scale moces could not be used to constraiu cosimolosv.," The number of Fourier modes grows as $k^3$ for a 3D survey, and so it would be a shame if these smaller-scale modes could not be used to constrain cosmology."
 Methods lave been proposed that reduce the covariance on translinear scales to varying degrees., Methods have been proposed that reduce the covariance on translinear scales to varying degrees.
 These include pre-whitenius (C7): removing Óuege halos from a survey τς à Gaussianizing trausforui nonlinear wavelet Wiener filtering (λα and dividing à bv its dispersion oe1 few-Mpe cells (2).., These include pre-whitening \citep{ajsh2000}; removing large halos from a survey \citep{ns07}; a Gaussianizing transform; nonlinear wavelet Wiener filtering \citep{zhang11}; and dividing $\delta$ by its dispersion in few-Mpc cells \citep{n11a}.
" In this paper. by. ""Cuussianizig transform we mean a function applied to pixels iu a field 0 measured iu cells of some resolution) that increases the Gaussianity of l-poiut probability deusitv function (PDF) of this field."," In this paper, by “Gaussianizing transform” we mean a function applied to pixels in a field $\delta$ measured in cells of some resolution) that increases the Gaussianity of 1-point probability density function (PDF) of this field."
 Examples of Caussianiziug transforms for the cosmological density field iuclude: a logarithinic ὃ>In(l|à) transform (?7): rank-order Caussianization 6>»Gd). eiviue an exactly Gaussian distribution by mapping the 1-poiut PDF onto a Caussian of some width (277): aud a Dox-Cox ransform {οοον which can be considered a generalization of the logarithuuc transform with parameters tunable to eive vanishing skewness aud kurtosis.," Examples of Gaussianizing transforms for the cosmological density field include: a logarithmic $\delta\to\ln(1+\delta)$ transform \citep[]{nss09,seo}; rank-order Gaussianization $\delta\to G(\delta)$, giving an exactly Gaussian distribution by mapping the 1-point PDF onto a Gaussian of some width \citep{weinberg,nss09,yuyu}; and a Box-Cox transform \citep{boxcox,jtk11}, which can be considered a generalization of the logarithmic transform with parameters tunable to give vanishing skewness and kurtosis."
 A related statistic o the raulk-order-CGaussianized power spectrum is the copula (?).., A related statistic to the rank-order-Gaussianized power spectrum is the copula \citep{copula}.
 There are reasons to think that 44=lu(1|4) would )o niore appropriate to analvze than ὃν, There are reasons to think that $A=\ln(1+\delta)$ would be more appropriate to analyze than $\delta$.
 ? pointed out heoretically that a lognormal PDF emerges if peculiar volocities are assumed to grow according to linear theory., \citet{colesjones} pointed out theoretically that a lognormal PDF emerges if peculiar velocities are assumed to grow according to linear theory.
 Using a Schroddinger-equation framework. ? found that he variance in Ais ιο better-described iu trec-level erturbation theory than the variance in à. sugecstine hat Ais closer to linear theory.," Using a Schröddinger-equation framework, \citet{spt} found that the variance in $A$ is much better-described in tree-level perturbation theory than the variance in $\delta$ , suggesting that $A$ is closer to linear theory."
 Ina study of disereteness effects. 2 also observed in simulations that the first few noments of A have reduced fractional variance comparee ο à.," In a study of discreteness effects, \citet{romeo} also observed in simulations that the first few moments of $A$ have reduced fractional variance compared to $\delta$."
 Going bevoud the inherent statistics of A into xuwanmeter-dependeuce. 7?— found analytically that for a lognormal density field ὃ whose moments depen on a cosmological parameter. the underline Cassia Geld A is (often much) more informative about the xuwneter than the lognormal field à.," Going beyond the inherent statistics of $A$ into parameter-dependence, \citet{carron} found analytically that for a lognormal density field $\delta$ whose moments depend on a cosmological parameter, the underlying Gaussian field $A$ is (often much) more informative about the parameter than the lognormal field $\delta$ ."
" Also. ?— fou hat applving a log trausform to a simulatedweak-lensing convergence field. allows significantly tighter constraints i| a asg-Q,,, paraincter space."," Also, \citet{jtk11}
 found that applying a log transform to a simulatedweak-lensing convergence field allows significantly tighter constraints in a $\sigma_8$ $\Omega_m$ parameter space."
 Iowever. they foun hat adding realistic galaxy shape noise in the analysis," However, they found that adding realistic galaxy shape noise in the analysis"
"(Tout&Pringle|19921995) so that Combining equations [12], [11], [L0] and D] we find equation [8] becomes The second and considerable torque arises between star and disc by magnetic interaction.","\citep{tout1992,stepien1995} so that Combining equations \ref{meq5}, , \ref{meq4}, \ref{meq3} and \ref{meq2}
 we find equation \ref{meq1} becomes The second and considerable torque arises between star and disc by magnetic interaction."
 The framework of such an interaction was constructed by |Ghosh&Lamb|(1978)., The framework of such an interaction was constructed by \citet{ghosh1978}.
. Some of the stellar magnetic dipole flux connects to the accretion disc and transports angular momentum between star and disc., Some of the stellar magnetic dipole flux connects to the accretion disc and transports angular momentum between star and disc.
" We use the expression given by|Armitage &Clarke)(1996) and assume, as did |Stepiei|(2000,2002),, that the radius of magnetosphere is equal to the co-rotation radius, As they pointed out the disc torque does not depend on the mass of the disc so where ji,=B,R? is the magnetic moment of stellar magnetic field."," We use the expression given by \citet{armitage1996} and assume, as did \citet{stepien2000,stepien2002}, that the radius of magnetosphere is equal to the co-rotation radius, As they pointed out the disc torque does not depend on the mass of the disc so where $\mu=B_{\rm s}R^3$ is the magnetic moment of stellar magnetic field."
" Contrary to whoassumes magnetic flux is constant, we assume that the magnetic field strength D, remains constant because the mass of the accreting star increases at a substantial rate (Mace©1079?Ma yr-!)."," Contrary to \citet{stepien2002} whoassumes magnetic flux is constant, we assume that the magnetic field strength $B_{\rm s}$ remains constant because the mass of the accreting star increases at a substantial rate $ \dot{M}_{\rm acc} \approx 10^{-5}\,{\rm
 M_\odot\,yr^{-1}} $ )."
 We seek a solution for various fiducial values of Bg., We seek a solution for various fiducial values of $B_{\rm s}$ .
" The third torque, the accretion torque discussed in section B], has the effect of the mass-accreting star."," The third torque, the accretion torque discussed in section \ref{models}, , has the effect of spinning-up the mass-accreting star."
 Now we replace the radius of the star R in equation | with Reor because the stellar magnetic field disrupts the disc at Reor as constant at the value of co-rotation radius., Now we replace the radius of the star $R$ in equation \ref{Tacc} with $R_{\rm{cor}}$ because the stellar magnetic field disrupts the disc at $R_{\rm{cor}}$ as constant at the value of co-rotation radius.
 The wind mass-loss rate remains an unknown parameter., The wind mass-loss rate remains an unknown parameter.
 Although there is no a priori rate we can set an upper limit., Although there is no a priori rate we can set an upper limit.
 All observed Algols show reversed mass ratio so much of the material lost by the donor must be accreted by the gainer., All observed Algols show reversed mass ratio so much of the material lost by the donor must be accreted by the gainer.
 We may write and where 0«81., We may write and where $0<\beta <1$.
 For conservative evolution 6= 1., For conservative evolution $\beta = 1$ .
" |Matt&Pudritz(2005) claimed that the mass outflow rate in stellar winds is about pper cent of the accretion rate for the pre-mainsequence stars, corresponding to 3= 0.9."," \citet{matt2005} claimed that the mass outflow rate in stellar winds is about per cent of the accretion rate for the pre-mainsequence stars, corresponding to $\beta \approx 0.9$ ."
 This is whythe classical T Tauri starsspin at lessthan pper cent of their breakup velocity., This is whythe classical T Tauri starsspin at lessthan per cent of their breakup velocity.
 It is theoriginal suggestion by, It is theoriginal suggestion by
(?2)..,\citep{Deasy1988}.
 The analytic model was also applied to theoretical nodels of Cepheids with the metallicities of the Simall and Large Magellanic Clouds. and the Galaxy. aud mass oss was found to be significant m Alagellanic Cloud Cepheids as well (?)..," The analytic model was also applied to theoretical models of Cepheids with the metallicities of the Small and Large Magellanic Clouds, and the Galaxy, and mass loss was found to be significant in Magellanic Cloud Cepheids as well \citep{Neilson2008b}."
 This result. however. has not been ested by observations.," This result, however, has not been tested by observations."
 The dust shells that siurounud Coephleids are optically hin because they form at larec distances from the Cepheids., The dust shells that surround Cepheids are optically thin because they form at large distances from the Cepheids.
" For instance. a Cepheid with Zig=6000 A wil have a coudenusation radius of ,./R=LO(DTa)OF= 16. assuming a condensation cluperature of 1200 A."," For instance, a Cepheid with $T_{\rm{eff}} = 6000$ $K$ will have a condensation radius of $r_c/ R_* = 0.5(T_c/T_{\rm{eff}})^{-5/2} = 16$ , assuming a condensation temperature of $1500$ $K$."
 At the distance where dust orn. the dust shells are optically thin. aud do not contribute to the extinction of starlight.," At the distance where dust forms, the dust shells are optically thin, and do not contribute to the extinction of starlight."
 Mass loss παν also be a solution to the Cepheid mass discrepancy problem., Mass loss may also be a solution to the Cepheid mass discrepancy problem.
 The Cepheid mass discrepancy is the difference between Cepheid mass estimates based ou stellar evolution isochrones aud those based ou stellar pulsation calculations., The Cepheid mass discrepancy is the difference between Cepheid mass estimates based on stellar evolution isochrones and those based on stellar pulsation calculations.
 The lnass discrepancy is about 10 20% iu the Milky WavAC (?).. about 17 25% in thle LAIC' aud about 20% in the (?7).. ?.," The mass discrepancy is about $10$ $20\%$ in the Milky Way \citep{Caputo2005}, , about $17$ $25\%$ in the LMC and about $20\%$ in the SMC \citep{Brocato2004, Keller2006}."
 determined that the nass discrepancy inereascs as tlie iietallicitv decreases. and? argue that the discrepaucy may also be a functiou of mass. though ? provides evidence against this result.," \cite{Keller2006} determined that the mass discrepancy increases as the metallicity decreases, and \cite{Caputo2005} argue that the discrepancy may also be a function of mass, though \cite{Keller2008} provides evidence against this result."
 For mass loss to be a solution to the mass discrepancy. a Cepheid that starts ou its first crossing with mass of SAL. would need to lose about LAL... or have an average lassloss rate in the range of LO7 to LOCAL.urs," For mass loss to be a solution to the mass discrepancy, a Cepheid that starts on its first crossing with mass of $5M_\odot$ would need to lose about $1M_\odot$ or have an average mass–loss rate in the range of $10^{-8}$ to $10^{-6}M_\odot /yr$."
 The purpose of this work is to model infrared excess ini LAIC Cepheids using SAGE observations iu the IRAC bands combined with OGLE II observations of D.V. aud I(?77).," The purpose of this work is to model infrared excess in LMC Cepheids using SAGE observations in the IRAC bands combined with OGLE II observations of B,V, and I \citep{Udalski1999a, Udalski1999b}."
 The next section outlines the observations and the model describing circiunustellar dust created in a stellar wind that causes infrared excess., The next section outlines the observations and the model describing circumstellar dust created in a stellar wind that causes infrared excess.
 The process for determining the massloss rates is also described., The process for determining the mass–loss rates is also described.
 The results are given in Section 3 aud the predicted infrared PL relatious are described in Section 1., The results are given in Section 3 and the predicted infrared PL relations are described in Section 4.
 The fifth section will explore possible driving mechauisnis for mass loss in LAIC Cepheids. testing if the mass loss behavior is simular to that proposed in ?7..," The fifth section will explore possible driving mechanisms for mass loss in LMC Cepheids, testing if the mass loss behavior is similar to that proposed in \cite{Neilson2008, Neilson2008b}."
 We use OGLE II and SAGE observations of LAIC Cepheids to determine massloss rates., We use OGLE II and SAGE observations of LMC Cepheids to determine mass–loss rates.
 The OGLE II data are for D. V and I maguitudes while the SAGE magnitudes are in IRAC hands at wavelengths 3.6. 1.5. 5.8 and 8.0 ga.," The OGLE II data are for B, V and I magnitudes while the SAGE magnitudes are in IRAC bands at wavelengths 3.6, 4.5, 5.8 and 8.0 $\mu m$."
 The SAGE data we used are adopted from ?.. which consist of 730 OCLE II LAIC Cepheids with logP>0.1.," The SAGE data we used are adopted from \cite{Ngeow2008}, which consist of 730 OGLE II LMC Cepheids with $\log P > 0.4$."
 However we only use 188 of these Cepheids that have at least 3 IRAC bands and 2 of the BVI bauds from this dataset., However we only use 488 of these Cepheids that have at least 3 IRAC bands and 2 of the BVI bands from this dataset.
 The SAGE data are compiled bv matching the position of OGLE II Cepheids with positions of infrared sources in the SACE observations., The SAGE data are compiled by matching the position of OGLE II Cepheids with positions of infrared sources in the SAGE observations.
 ? match sources if they are within 3.5x5 areseconds of the position of the OGLE II Cepheids., \cite{Ngeow2008} match sources if they are within $3.5$ arcseconds of the position of the OGLE II Cepheids.
 Fiewe l1. show the infrared PL relations constructed using the IRAC magnitudes., Figure \ref{f1} show the infrared PL relations constructed using the IRAC magnitudes.
 A nuuber of (mostly short period) Cepheids appear to deviate from the infrared PL relations. implying there is some mfrared excess.," A number of (mostly short period) Cepheids appear to deviate from the infrared PL relations, implying there is some infrared excess."
 The are three possible causes of the infrared fiux excess: (1) blending of stars in SAGE observations. (2) false matches of the infrared sources to the OGLE II Cepheids. (3) or cmemustellar dust shells forming at a significant distance from the Cepheids in a stellar wind.," The are three possible causes of the infrared flux excess: (1) blending of stars in SAGE observations, (2) false matches of the infrared sources to the OGLE II Cepheids, (3) or circumstellar dust shells forming at a significant distance from the Cepheids in a stellar wind."
 It is possible that false matches contaminate the sample. and we check this in Figure 2/— where the magnitude residnials of the IR. PeriodΤαµουαν relations determined bv7 are shown as a function of the separation betweeu the OGLE II aud SAGE positions.," It is possible that false matches contaminate the sample, and we check this in Figure \ref{f1a} where the magnitude residuals of the IR Period–Luminosity relations determined by \cite{Ngeow2008} are shown as a function of the separation between the OGLE II and SAGE positions."
" Although the search radius used im ? is rather laree. most of the matched objects voustratedhave a separation less tla 0.77 arcseconds [as de in the Figure 1 aud Table 1 iu ο],"," Although the search radius used in \cite{Ngeow2008} is rather large, most of the matched objects have a separation less than $0.77$ arcseconds [as demonstrated in the Figure 1 and Table 1 in \cite{Ngeow2008}] ]."
 The potential for false matches is only significant for a πα faction of the total sample of and in the sample used here. only ten of the [55 Cepheid matches lave a separation greater than 1.3 arcsecouds. which is he separation where the matches almost all have large residuals to the fit of the PL relation.," The potential for false matches is only significant for a small fraction of the total sample of \cite{Ngeow2008}, and in the sample used here, only ten of the $488$ Cepheid matches have a separation greater than $1.3$ arcseconds, which is the separation where the matches almost all have large residuals to the fit of the PL relation."
 There is another asvuuuectiy apparent in Fieure 2. where the residuals lave a separation less than 0.1 arcsecouds., There is another asymmetry apparent in Figure \ref{f1a} where the residuals have a separation less than $0.4$ arcseconds.
 Iu this work. we keep the two samples. but ideutifv them iu the figures when imiportanut.," In this work, we keep the two samples, but identify them in the figures when important."
 Wecan test if the infrared excess is due to mass loss x caleulatiug the sim of the infrared luminosity of the Cepheid aud theluninosity of the dust that is generated in the wind. given by," Wecan test if the infrared excess is due to mass loss by calculating the sum of the infrared luminosity of the Cepheid and theluminosity of the dust that is generated in the wind, given by"
of their evolution. while the latter are. beamed objects.,"of their evolution, while the latter are beamed objects."
 The analysis of the multi-epoch spectral behaviour has ooved to be a powerful tool to discriminate between the wo populations Clorniainenetal.2005:TintiOrientietal. 2007).," The analysis of the multi-epoch spectral behaviour has proved to be a powerful tool to discriminate between the two populations \citep{torniainen05,tinti05,mo07}."
. In fact. beamed radio sources. although usually characterized by a flat ancl variable spectrum. may oe selected. in samples of high-frequeney. peaking objects when their emission is dominated by a Daring knot in a jet.," In fact, beamed radio sources, although usually characterized by a flat and variable spectrum, may be selected in samples of high-frequency peaking objects when their emission is dominated by a flaring knot in a jet."
 On the other hand. voung radio sources are known to be he least variable class of extragalactic radio sources (ODoa 1998).," On the other hand, young radio sources are known to be the least variable class of extragalactic radio sources \citep{odea98}."
. Llowever. it must be mentioned that in the voungest objects. substantial variability in the optically-thiek part of the spectrum. is expected. as a consequence of either the source. growthevolution (c.g.1459|3337.Orienti&Dallacasa—2008c) or changes in the possible absorber screen. or a combination of both (CTingay&deKool2003).," However, it must be mentioned that in the youngest objects, substantial variability in the optically-thick part of the spectrum is expected as a consequence of either the source growth/evolution \citep[e.g. J1459+3337,][]{mo08}, or changes in the possible absorber screen, or a combination of both \citep{tingay03}."
". In this paper we present a multi-epoch analysis based on simultaneous multi-frequency VLA. data of the racio spectra of 57 high frequcney peakers from. the ""faint"" LEP sample (Stanghellinictal.2009).", In this paper we present a multi-epoch analysis based on simultaneous multi-frequency VLA data of the radio spectra of 57 high frequency peakers from the “faint” HFP sample \citep{cs09}.
.. This sample was selected as the “bright” LIED sample (Dallacasaetal.2000) by cross-corrclating the 87GB survey at 4.9 Gllz with the NVSS at 1.4 Gllz. ancl including only sources fainter than 300 mJv at 4.9 Cll within a restricted area. around the northern galactic cap (lor details on the sample sce Stanghellini ct al.," This sample was selected as the “bright” HFP sample \citep{dd00} by cross-correlating the 87GB survey at 4.9 GHz with the NVSS at 1.4 GHz, and including only sources fainter than 300 mJy at 4.9 GHz within a restricted area around the northern galactic cap (for details on the sample see Stanghellini et al."
 2009)., 2009).
 The study of the radio properties of a sample of faint LPs is the first step in understanding the first stages of the radio source evolution., The study of the radio properties of a sample of faint HFPs is the first step in understanding the first stages of the radio source evolution.
 So far. spectral studies have been carried out for the bright LIP sources only. and an extension to fainter objects is necessary.," So far, spectral studies have been carried out for the bright HFP sources only, and an extension to fainter objects is necessary."
 The evolution models developed so far (c.g.Fantictal.1995:Snellenetal.2000:Waiser&AlexanderL997) preclict that in the earliest stage the raclio luminosity progressively increases. implving that the voungest objects are Likely to be found among faint sources.," The evolution models developed so far \citep[e.g.][]{cf95,snellen00,kaiser97} predict that in the earliest stage the radio luminosity progressively increases, implying that the youngest objects are likely to be found among faint sources."
" Furthermore. in. faint IHEPs. boosting elfects should be less relevant. making the contamination [roni blazars less severe than what found in samples of brighter ""Throughout this paper we assume the following cosmology: ily=Tikms!Mpe !. Ox,=027. and OX=0.73λ ina Hat Universe."," Furthermore, in faint HFPs, boosting effects should be less relevant, making the contamination from blazars less severe than what found in samples of brighter Throughout this paper we assume the following cosmology: $H_{0} = 71
{\rm km \; s^{-1} \; Mpc^{-1}}$ , $\Omega_{\rm M} = 0.27$, and $\Omega_{\Lambda} = 0.73$ in a flat Universe."
" The spectral index is defined as ο)xumm Simultaneous multi-frequency VLA observations of 57 out of the 61 sources from the ""faint HET sample (Stanghellinietal.2009) were carried out during cdillerent runs between September 2003 and April 2007 (Lable 13).", The spectral index is defined as $S{\rm (\nu)} \propto \nu^{- \alpha}$ Simultaneous multi-frequency VLA observations of 57 out of the 61 sources from the “faint” HFP sample \citep{cs09} were carried out during different runs between September 2003 and April 2007 (Table \ref{vla_oss}) ).
 Observations were performed. in L band (with the two LPs centered. at 1.415 and 1.665 112). € band (with the two LES centered at 4.565 and 4.935 CGllz). X band (with the two IEs centered at 5.055 and S465 CGllz). U band (14.940 112). Ix band (22.460 Gliz). and in Q band (43.340 11).," Observations were performed in L band (with the two IFs centered at 1.415 and 1.665 GHz), C band (with the two IFs centered at 4.565 and 4.935 GHz), X band (with the two IFs centered at 8.085 and 8.465 GHz), U band (14.940 GHz), K band (22.460 GHz), and in Q band (43.340 GHz)."
" At each frequeney, the target sources were observed. for about. 1 minute. evcling through frequencies."," At each frequency, the target sources were observed for about 1 minute, cycling through frequencies."
 During cach run. the primary [ux density calibrator either 2286 or 448 was observed for about 3 minutes at cach frequency.," During each run, the primary flux density calibrator either 286 or 48 was observed for about 3 minutes at each frequency."
 Secondary calibrators were chosen on the basis of their distance from the targets in order to minimize the telescope slewing time. and they were observed for 1.5 min at each frequency. every 20 The data reduction was carried. out following the standard. procedures. for the VLA. implemented. in. the NILAO ALPS package.," Secondary calibrators were chosen on the basis of their distance from the targets in order to minimize the telescope slewing time, and they were observed for 1.5 min at each frequency, every 20 The data reduction was carried out following the standard procedures for the VLA implemented in the NRAO AIPS package."
 Ehe flux density at cach frequency was measured on the final image produced. after a few phase-only. self-calibration iterations., The flux density at each frequency was measured on the final image produced after a few phase-only self-calibration iterations.
 In the L band it was ecncrally necessary to image a few confusing sources falling within the primary beam., In the L band it was generally necessary to image a few confusing sources falling within the primary beam.
 ALL the target sources appeared unresolved. at any frequency., All the target sources appeared unresolved at any frequency.
 During the observations of a lew sources. strong REI at 1420 and 1.665 Gllz was present. and in those cases the measurements of the fux density were not Uncertainties on the determination of the absolute Lux density scale are dominated by amplitude errors.," During the observations of a few sources, strong RFI at 1.420 and 1.665 GHz was present, and in those cases the measurements of the flux density were not Uncertainties on the determination of the absolute flux density scale are dominated by amplitude errors."
 Based on the variations of antenna complex gains curing the various observations. we can conservatively estimate an uncertainty of οί in L. €. and. X bands. ~5% in. U band. and. — in Ix and Q bands.," Based on the variations of antenna complex gains during the various observations, we can conservatively estimate an uncertainty of $\sim$ in L, C, and X bands, $\sim$ in U band, and $\sim$ in K and Q bands."
 The rms noise level on the image plane is about 0.1 mJv/beam in € and X. bands. and about 0.2-0.4 mJvfbeam in L. U. and Ix bands.," The rms noise level on the image plane is about 0.1 mJy/beam in C and X bands, and about 0.2-0.4 mJy/beam in L, U, and K bands."
 In the Q band it accounts. for 0.4-0.5. mvbeam. becoming comparable to the amplitude calibration errors for sources fainter than 20 mJ.," In the Q band it accounts for 0.4-0.5 mJy/beam, becoming comparable to the amplitude calibration errors for sources fainter than $\sim$ 20 mJy."
 Results are presented in Section Jo determine the optical properties. of the sources in the faint LIP sample. we complemented. the information available in the literature with that provided by the SDSS DRI (Abazajianetal.2009).," Results are presented in Section To determine the optical properties of the sources in the faint HFP sample, we complemented the information available in the literature with that provided by the SDSS DR7 \citep{sdss7}."
. The optical properties of each object (like source. extension ancl magnitude) have been carefully inspected bevond the automated. procedures in the SDSS. in order to unambiguously identify the host (Le. quasar. galaxy. or empty field) of each radio Of the 57 sources considered. in this paper. 12are identified. with galaxies with redshift between 0.03. and 0.6: 33 are quasars with a higher redshift. tvpicallv in the," The optical properties of each object (like source extension and magnitude) have been carefully inspected beyond the automated procedures in the SDSS, in order to unambiguously identify the host (i.e. quasar, galaxy, or empty field) of each radio Of the 57 sources considered in this paper, 12are identified with galaxies with redshift between 0.03 and 0.6; 33 are quasars with a higher redshift, typically in the"
al the surfaces of strange stars. which are higher than the nuclear densitv. the effect could reduce the bremsstrahlung emissivity of quark matter by an order ol magnitude.,"at the surfaces of strange stars, which are higher than the nuclear density, the Landau-Pomeranchuk effect could reduce the bremsstrahlung emissivity of quark matter by an order of magnitude."
 Photon emissivity of quark-gluon. plasma (QGP). which is conjectured to be Lormecl in ultrarelativistie heavy ion collisions. has been extensively investigated recently (Lor a recent. review of direct. photon. emission from QGDP. including comparisons of theoretical predictions wilh experiments see Peitznann Thoma (2002)).," Photon emissivity of quark-gluon plasma (QGP), which is conjectured to be formed in ultrarelativistic heavy ion collisions, has been extensively investigated recently (for a recent review of direct photon emission from QGP, including comparisons of theoretical predictions with experiments see Peitzmann Thoma (2002))."
 Photons and dilepton pairs onlv interact electromagnetically and their mean free paths are much larger than (he size of 1e QGP., Photons and dilepton pairs only interact electromagnetically and their mean free paths are much larger than the size of the QGP.
 Ποιος (hese electromagnetic probes leave the hot aud dense QGP? without further scaltering. (hus providing important informations about (he early stages of the collision and re structure of the QGP.," Hence these electromagnetic probes leave the hot and dense QGP without further scattering, thus providing important informations about the early stages of the collision and the structure of the QGP."
 As sources of direct photons one ean consider quark annihilation. ωνSomptlon scattering and bremsstrahlung following the initial hard scattering of partons οἱ je nuclei. as well as thermal photons from the QGP and trom hadronic interactions in 1e hol hadronic gas after the hadronization of the plasma (Peitzmann Thoma 2002).," As sources of direct photons one can consider quark annihilation, Compton scattering and bremsstrahlung following the initial hard scattering of partons of the nuclei, as well as thermal photons from the QGP and from hadronic interactions in the hot hadronic gas after the hadronization of the plasma (Peitzmann Thoma 2002)."
 The first observation of direct photon production in ultrarelativistic heavy ion collisions has been reported by the WA98 collaboration in 75Pb4-75 collisions at V/s=158 GeV αἱ ihe Super Proton Svnehrotron at CERN (Agearwal et al., The first observation of direct photon production in ultrarelativistic heavy ion collisions has been reported by the WA98 collaboration in $^{208}Pb+^{208}Pb$ collisions at $\sqrt{s}=158$ GeV at the Super Proton Synchrotron at CERN (Aggarwal et al.
 2000)., 2000).
 The results display a clear excess Of direct photons above the expected background [rom hadronic decays in (he range of (ransverse momentum py>1.5 GeVο im the most central collisions., The results display a clear excess of direct photons above the expected background from hadronic decays in the range of transverse momentum $p_{T}>1.5$ $/c$ in the most central collisions.
 These experimental lindings provide a confirmation of the feasibility of direct photons as reliable probes in heavy ion collisions aud may pave the way lor the understanding of the formation anc evolution of the QGD., These experimental findings provide a confirmation of the feasibility of direct photons as reliable probes in heavy ion collisions and may pave the way for the understanding of the formation and evolution of the QGP.
 The evaluation of the photo-emission rate from the QGP was initiated by Ixapusta et al. (, The evaluation of the photo-emission rate from the QGP was initiated by Kapusta et al. (
1991) and by Baier et al. (,1991) and by Baier et al. (
1992).,1992).
 However. as shown by Aurenche et al. (," However, as shown by Aurenche et al. ("
2000). these initial estimations were incomplete. since the bremisstrahlune ancl inelastic pair annihilation processes contains collinear enhancements which cause them (o contribute at the sime parametric order in (he coupling as the (wo to two processes even lor large photon energies.,"2000), these initial estimations were incomplete, since the bremsstrahlung and inelastic pair annihilation processes contains collinear enhancements which cause them to contribute at the same parametric order in the coupling as the two to two processes even for large photon energies."
 Bul these calculations are also incomplete. as they did not incorporate the suppression of photon emission due to multiple scatterings during the photon emission process. which limits the coherence length of the emitted. radiation (the. Landau-Pomeranchuk-Migeal. effect).," But these calculations are also incomplete, as they did not incorporate the suppression of photon emission due to multiple scatterings during the photon emission process, which limits the coherence length of the emitted radiation (the Landau-Pomeranchuk-Migdal effect)."
 The rate of photo-production by bremsstrahlung aid inelastic pair annihilation in a hot. equilibrated plasma at zero chemical potential. fully including the LPAI effect. was calculated. to leading orders in both the electromagnetic ancl strong coupling constants. by (Arnold. Moore Jaffe 2001. Arnold. Moore Jaffe 2002).," The rate of photo-production by bremsstrahlung and inelastic pair annihilation in a hot, equilibrated plasma at zero chemical potential, fully including the LPM effect, was calculated, to leading orders in both the electromagnetic and strong coupling constants, by (Arnold, Moore Jaffe 2001, Arnold, Moore Jaffe 2002)."
 The emission of hard photons from the QGP plasma. using a model based on the thermocdvuaniucs of QCD. and (he determination of the initial temperature of (he expanding fireball has been considered recently by Renk (2003).," The emission of hard photons from the QGP plasma, using a model based on the thermodynamics of QCD, and the determination of the initial temperature of the expanding fireball has been considered recently by Renk (2003)."
The use of adaptive optics iu astronomy with reasonable spatial order correction αμα temporal vancdwidths has been restricted. for παν vears to observatories with a high instrumentation budect.,The use of adaptive optics in astronomy with reasonable spatial order correction and temporal bandwidths has been restricted for many years to observatories with a high instrumentation budget.
 In curent versions. such svstenis use expensive deformable uirors. and digital signal processors to apply the recoustruction algorithius.," In current versions, such systems use expensive deformable mirrors, and digital signal processors to apply the reconstruction algorithms."
" This results in complex systems (οιο,[ippleretal.2000).. taking several vears o be developed."," This results in complex systems \citep[e.g.][]{hip2000}, taking several years to be developed."
 Nowadays. the total cost can be reduced w several orders of magnitude (e.c.Daintyetal.1999) thanks to the availability of relatively cheap memibraue deformable mirrors (Vdovin1995).. single photon detectors with a ligh frame rate based on the EMCCD technology (c.g.Dussault&IIoess2001) and low cost Tip-Tilt xvsteuis.," Nowadays, the total cost can be reduced by several orders of magnitude \cite[e.g.][]{dai1999} thanks to the availability of relatively cheap membrane deformable mirrors \cite{vdo1995}, single photon detectors with a high frame rate based on the EMCCD technology \citep[e.g.][]{dus2004} and low cost Tip-Tilt systems."
 Tn this paper we describe a System of Adaptive Optics with Lucky huaging (hereafter SAOLIND., In this paper we describe a System of Adaptive Optics with Lucky Imaging (hereafter SAOLIM).
 We present the design. construction and results of a low order AO syvsteu for 1-21n class telescopes. almost entirely developed with available commnercial compoucuts. aud with a total cost of ~ 35000 euros in hardware compoucuts.," We present the design, construction and results of a low order AO system for 1-2m class telescopes, almost entirely developed with available commercial components, and with a total cost of $\sim$ 35000 euros in hardware components."
 The optical desigu cuables a FOV of 90x90 arcsec? for the scientific camera., The optical design enables a FOV of 90x90 $^{2}$ for the scientific camera.
" SAOLIM is optically corrected aud transparent for a waveleugth rauge between 1.0-2.5¢nn. Qur system is based on a iiembraue deformable mirror (οιο,Patersonetal.2000).. aud a single PC to perform all the computationsτσ and hardware control. integrating everything iu a simple and compact design as we will see later."," SAOLIM is optically corrected and transparent for a wavelength range between $\mu$ m. Our system is based on a membrane deformable mirror \citep[e.g.][]{pat2000}, and a single PC to perform all the computations and hardware control, integrating everything in a simple and compact design as we will see later."
 The dual wireless/ethernet conunununication of the device allows au easy setup. because uo cabling has to be installed at the telescope. reducing a cousiderable amount of possible problems.," The dual wireless/ethernet communication of the device allows an easy setup, because no cabling has to be installed at the telescope, reducing a considerable amount of possible problems."
 This instrument can be used as an imput correction for applications where reaching the diffraction Πιτ of the telescope is required. such as a lucky nuagiug system. like ASTRALUN (Πωetal.2008)... or as a complement for the shift-aud-add approach (e.g.Dates&Cady 1980).," This instrument can be used as an input correction for applications where reaching the diffraction limit of the telescope is required, such as a lucky imaging system, like ASTRALUX \cite{hor2008}, or as a complement for the shift-and-add approach \citep[e.g.][]{bates1980}."
.. This device can take advantage of a low order AO svsteni in such a wav that the rate of the useful inages could be increased. and therefore the performance of the iustrunent can be improved.," This device can take advantage of a low order AO system, in such a way that the rate of the useful images could be increased, and therefore the performance of the instrument can be improved."
 Such innovative technique has been tested recently at the Palomar observatory. getting excellent results (6.8.Lawctal. 2008).," Such innovative technique has been tested recently at the Palomar observatory, getting excellent results \citep[e.g.][]{law2008}."
. Figure l shows a picture of the mstrumneut attached to the telescope., Figure 1 shows a picture of the instrument attached to the telescope.
 Labels indicating the main coupoucuts have been included., Labels indicating the main components have been included.
 A detail of the optical components is shown in Fieure 2 where a view of the inside of the instrument (located at the lab) is preseuted., A detail of the optical components is shown in Figure 2 where a view of the inside of the instrument (located at the lab) is presented.
 The main optical clemeuts have Όσοι labelled in this figure., The main optical elements have been labelled in this figure.
 A sketch of the optical setup is shown in Figure 3., A sketch of the optical setup is shown in Figure 3.
" The optical design is similar to that of ALFA. the AO system that was operated at the 23.51 telescope of the Calar Alto observatory since L997 πι. 2005 (οιο,Hippleretal.2000)."," The optical design is similar to that of ALFA, the AO system that was operated at the 3.5m telescope of the Calar Alto observatory since 1997 until 2005 \citep[e.g.][]{hip2000}."
. A Shack-Hartmann wavefrout sensor (hereafter SIIS) is placed optically conjugated with the membrane deformable mirror (MDMMS) and with the eutrauce pupil of the whole system., A Shack-Hartmann wavefront sensor (hereafter SHS) is placed optically conjugated with the membrane deformable mirror (MDMM) and with the entrance pupil of the whole system.
 Two achromatic doublets (El) separated by Suuu. with focal distances of 300nmun cach one conjugate the entrance pupil over a 201uu diameter circular area ou the deformable mirror.," Two achromatic doublets (E1) separated by 5mm, with focal distances of 300mm each one conjugate the entrance pupil over a 20mm diameter circular area on the deformable mirror."
 Another paix of achromatic doublet lenses (E23 and EG) is configured as a Ἱνορία telescope aud conjugates the membranes selected. zone, Another pair of achromatic doublet lenses (E3 and E6) is configured as a Kepler telescope and conjugates the membrane's selected zone
 6? (Raedleretal.1990:Durney1993).. (Dobleretal.2006).. (Duruevetal.1993)... (Alohautyetal.2002).," $\alpha\Omega$ \citep{rwm+90,ddr93}, \citep{dsb06}, \citep{ddr93}. \citep{mbs+02}."
. Ta (Saar&Liusky," $\alpha$ \citep{sl85,jv96}."
sources to Milagro diffuse emission the flix in Eq.(8)) has to be corrected for the cifferent threshold energy.,sources to Milagro diffuse emission the flux in \ref{eqnfluxtotal}) ) has to be corrected for the different threshold energy.
 All HESS sources are fitted bv a power law spectrum and the average spectral index is P—2.32., All HESS sources are fitted by a power law spectrum and the average spectral index is $\Gamma=2.32$.
 By correcting the flix in Exq.(8)) for the Milagro threshold energy the flix whieh IIESS-like sources contribute is The contribution of HESS source population amounts to al least 10 per cent of the diffuse {lux whieh Milagro measures above 3.5 TeV. This is a lower limit for the contribution of unresolved sources to Milagro diffuse emission. as only sources above 6 percent of the Crab flux were taken into account to estimate it.," By correcting the flux in \ref{eqnfluxtotal}) ) for the Milagro threshold energy the flux which HESS-like sources contribute is The contribution of HESS source population amounts to at least 10 per cent of the diffuse flux which Milagro measures above 3.5 TeV. This is a lower limit for the contribution of unresolved sources to Milagro diffuse emission, as only sources above 6 percent of the Crab flux were taken into account to estimate it."
" The Galactic cliffiise emission measured by Milagro can be used to constrain the minimum [Iux 55,5, below which the logN-log9 plot becomes flat in order not to overproduce the Milagro flix Sin cannot be less than 1.x10.photonssrtem?s! in order not to violate the constraint in Eq.(11))."," The Galactic diffuse emission measured by Milagro can be used to constrain the minimum flux $S_{min}$ below which the logN-logS plot becomes flat in order not to overproduce the Milagro flux $S_{min}$ cannot be less than $1 \times {10}^{-16} \, photons \, {\mathrm{{sr}^{-1}}} {\mathrm{{cm}^{-2}}} \,
{\mathrm {s^{-1}}}$ in order not to violate the constraint in \ref{eq:constraint}) )."
 From the number-intensityv relation for LESS source population it is possible to deduce ihe number of sources which ILESS. VERITAS. Milagro and ILAWC will detect.," From the number-intensity relation for HESS source population it is possible to deduce the number of sources which HESS, VERITAS, Milagro and HAWC will detect."
 The number ol SNRs and PWNs expected for ILESS if its entire field of view is scanned with a uniform sensitivity of 2 per cent of the Crab flix above 200 GeV is about 43410., The number of SNRs and PWNs expected for HESS if its entire field of view is scanned with a uniform sensitivity of 2 per cent of the Crab flux above 200 GeV is about $43\pm 10$.
" Η VERITAS will survey the Northern sky in the region between 30°«/220"" and —10*<b<10° reaching the level of 1 per cent of the Crab flix above 100. GeV. it should detect. approximately 35c10 sources."," If VERITAS will survey the Northern sky in the region between ${30}^{o}<l<{220}^{o}$ and ${-10}^{o}<b<{10}^{o}$ reaching the level of 1 per cent of the Crab flux above 100 GeV, it should detect approximately $35\pm10$ sources."
 Milagro has previously observed the Northern skv with a sensilivily equal to about 65 per cent the Crab Πιν at 3.5 TeV. At this level of sensitivity one should expect no detection of sources for Milagro and indeed no detection of sources was Claimed in, Milagro has previously observed the Northern sky with a sensitivity equal to about 65 per cent the Crab flux at 3.5 TeV. At this level of sensitivity one should expect no detection of sources for Milagro and indeed no detection of sources was claimed in
 The ζω ZZTodi». T; T... Tir) and fdr20.is ↸⊳∐⋅↸⊳∏↕⋜↧⊓∪∐∏∪↖↖↽↴∖↴↖↖⇁↸∖↸⊳∪∐↴∖↴↕≼∐∖↥⋅∐↸∖↥⋅↸∖∙↑↕∐↴∖↴↸⊳∪∪∐∐∶↴∙⊾↕↴∖↴⋜↧↴∖↴↴∖↴⋯⊔↸∖≼↧ creation , $\sim T_{virial}/3$ $T \lta T_{virial}/3$ $T_h$ $T_c$ $T_c(r)$ ${\dot M}$ $dT/dr > 0$ 
would have been contained within one svuthesized beam (together with any nuclear emission).,would have been contained within one synthesized beam (together with any nuclear emission).
 The luminosity limit suggestsMO that no comparable population of nearby 1uassive comipanious is present., The luminosity limit suggests that no comparable population of nearby massive companions is present.
 Moreover. no significant CO J=2-1 emission features are found within the entire fiel of view that spaus Lt Mpc. which suggests that such massive cold molecular gas conceutratious are rare.," Moreover, no significant CO J=2–1 emission features are found within the entire field of view that spans $\sim1.4$ Mpc, which suggests that such massive cold molecular gas concentrations are rare."
 Observations of SDSS 1011-0125 with better seusitivity are needed to explore whether smaller but still significant coucentrations of low excitation molecular gas are present iu tlie euviroument of this hieh redshift quasar., Observations of SDSS 1044-0125 with better sensitivity are needed to explore whether smaller but still significant concentrations of low excitation molecular gas are present in the environment of this high redshift quasar.
 The work was supported in part by NSF grant AST-21795 to the University of California., The work was supported in part by NSF grant AST-21795 to the University of California.
 We thank Leo Blitz lor scheduling additional observatious on this project alter revision cX the quasar redshift. ancl Alberto Bolatto for a generous donation of observing time.," We thank Leo Blitz for scheduling additional observations on this project after revision of the quasar redshift, and Alberto Bolatto for a generous donation of observing time."
cross product) defined by unit and VJ.,cross product) defined by ${\bf unit}$ and $\nabla J$.
 Lt is usually the that VJ40., It is usually the that $\nabla J \neq 0 $.
 The third order terms have the similar structure as the second order terms., The third order terms have the similar structure as the second order terms.
 The reason is because Ποια is analvtie and so is orientable., The reason is because $\kappa$ -field is analytic and so is orientable.
 In the third order. the phase is multiplied three (times. and so the unit vector is replaced by unit=(CS.Ss).," In the third order, the phase is multiplied three times, and so the unit vector is replaced by ${\bf unit} \equiv (C_3, ~S_3)$."
" In a tvpical Galactic bulge lensing where a star in (the direciton of) the bulge is lensed by a (faint) star or a planet svstem in the bulge or the disk. the stellar radius in units of the Einstein ring radius of the total mass of the lensing svstem is r,e107—D?5."," In a typical Galactic bulge lensing where a star in (the direciton of) the bulge is lensed by a (faint) star or a planet system in the bulge or the disk, the stellar radius in units of the Einstein ring radius of the total mass of the lensing system is $r_\ast \sim 10^{-3} - 10^{-2}$."
 In order to get an idea of the magnitude of the gradient. J on critical curves. let's consider the simplest case of a single lens.," In order to get an idea of the magnitude of the gradient $J$ on critical curves, let's consider the simplest case of a single lens."
 For a single lens. &=1/27: J=1-1/zkk VJ) ο ," For a single lens, $\kappa = 1/z^2$; $J = 1- 1/|z|^4$; $|\nabla J | = 4/|z|^5$ ."
the critical curve. [z|=lL. and |VJ|24>>r4.," On the critical curve, $|z|=1$, and $|\nabla J | = 4 >> r_\ast$."
 The gradient J is normal to the ring |z|=1and so is the positive eigenvector (4)., The gradient $J$ is normal to the ring $|z|=1$and so is the positive eigenvector $(\pm) E_+$.
 Thus. J—0. and J=(4)4.," Thus, $J_-=0$, and $J_+ = (\pm) 4$."
 The fact that J=0 evervwhere on the critical curve underlies that the point caustic is a degenerate cusp., The fact that $J_-=0$ everywhere on the critical curve underlies that the point caustic is a degenerate cusp.
 In a binary lens where its lens elements masses are comparable. we expect a simular range of values of |VJ|.," In a binary lens where its lens elements masses are comparable, we expect a similar range of values of $|\nabla J|$."
 We have shown the numerical values of JI| on the 4-cusped central caustie of a binary lens with the mass ratio of ~2:1 in Fig., We have shown the numerical values of $\sqrt{|J_-|}$ on the 4-cusped central caustic of a binary lens with the mass ratio of $\sim 2 : 1$ in Fig.
 1 of (rb99)., 1 of \citet{limb} (rb99).
 We also stated that (rioics have somewhat larger values., We also stated that trioids have somewhat larger values.
" The larger the value of [VJ around the eritical curve. the faster decreases the magnification value and smaller ihe ""width"" of the caustic curve."," The larger the value of $|\nabla J|$ around the critical curve, the faster decreases the magnification value and smaller the “width"" of the caustic curve."
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&b," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&ba," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&bar," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barp," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barpa," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barpar," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparl," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparli," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlia," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlial," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialb," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialba," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbar," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbark," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbarka," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbarkap," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbarkapp," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
 If the mass ratio is very small as in aplanetary. binary lens. then = 2partial = = 2&barparlialbarkappe," If the mass ratio is very small as in aplanetary binary lens, then = 2 J = - 2"
The NASA program in N-ray astronomy has two lone term goals: 1) to achieve sufficient aneular resolution to umaege the event horizon of a black hole (0.1 uucro arc sec) and 2) to achieve sufficieut. collecting area (50-150 sq m) aud angular resolution (0.1.1.0 arc sec) to observe in detail the first black holes aud ealaxies at hieh redshift.,The NASA program in X-ray astronomy has two long term goals: 1) to achieve sufficient angular resolution to image the event horizon of a black hole (0.1 micro arc sec) and 2) to achieve sufficient collecting area (50-150 sq m) and angular resolution (0.1–1.0 arc sec) to observe in detail the first black holes and galaxies at high redshift.
 These aimbitous goals cau be used to map out a series of nuissionis and a technology program., These ambitous goals can be used to map out a series of missions and a technology program.
 The next major mission will be Conustellation-N. which will be dedicated to high resolution. N-rav spectroscopy for lauuch iu ~ 2010., The next major mission will be Constellation-X which will be dedicated to high resolution X-ray spectroscopy for launch in $\sim$ 2010.
 This nüssion is a critical step im the roadmap to achieve these goals., This mission is a critical step in the roadmap to achieve these goals.
 Followiug Coustellation-X NASA is considering two very ambitious missions: MAXIM and Ceneration-X that will achieve the ultimate capabilities., Following Constellation-X NASA is considering two very ambitious missions: MAXIM and Generation-X that will achieve the ultimate capabilities.
 The modest miussious Astro-E2 and Switt address more focussed science goals on a rapid developiueut evcle and provide important pathfinders to the larger missions., The modest missions Astro-E2 and Swift address more focussed science goals on a rapid development cycle and provide important pathfinders to the larger missions.
" The current eeueratiou of larec X-ray observatories: the Chandra ταν Observatory frou: NASA aud the NMM-Neswtou Observatory from ESA are delivering an impressive array of results,", The current generation of large X-ray observatories: the Chandra X-ray Observatory from NASA and the XMM-Newton Observatory from ESA are delivering an impressive array of results.
 As we look to the future the requirements for future ταν observatories are driven by some basic questions: What is the nature of space and tine?, As we look to the future the requirements for future X-ray observatories are driven by some basic questions: What is the nature of space and time?
 When did the first black holes appear?, When did the first black holes appear?
 What is the Universe made of?, What is the Universe made of?
 A-rav observations are central to addressing these questions., X-ray observations are central to addressing these questions.
 Using these questions we can plan a roadmap for the future of N-ray. astronomy., Using these questions we can plan a roadmap for the future of X-ray astronomy.
 Black holes represeut the most extreme environnients kuown., Black holes represent the most extreme environments known.
 They provide the oulv known regime where the effects of General Relativity (GR) cau be tested iu the strong eravitv regime., They provide the only known regime where the effects of General Relativity (GR) can be tested in the strong gravity regime.
 Materi falliug iuto the black hole is heated to high temperatures aud the bulk of the euergv is emitted ia N-rvavs., Material falling into the black hole is heated to high temperatures and the bulk of the energy is emitted in X-rays.
 Results frou the Japanese-US ASCA iission have demonstrated that N-rav spectral features are comune from close to the event horizon of the black hole., Results from the Japanese-US ASCA mission have demonstrated that X-ray spectral features are coming from close to the event horizon of the black hole.
 These results sugeest that N-rav observations hold the promise of usiug black holes to test GR., These results suggest that X-ray observations hold the promise of using black holes to test GR.
 The most direct wav to do this is to resolve the event horizou of a black hole aud directly observe the predicted distortions of space time., The most direct way to do this is to resolve the event horizon of a black hole and directly observe the predicted distortions of space time.
 The best caudidates for this are, The best candidates for this are
"Four different central star types are considered in the model calculations: F2V, G2V, K2V, and M4.5V-type stars.","Four different central star types are considered in the model calculations: F2V, G2V, K2V, and M4.5V-type stars."
" The sample of stars used for the calculations are the F-dwarf σ Bootis 128167)), the Sun, the young active K-type star εEridani 22049)), and the M-type dwarf star AD Leo 388))."," The sample of stars used for the calculations are the F-dwarf $\sigma$ Bootis ), the Sun, the young active K-type star $\epsilon \ \mathrm{Eridani}$ ), and the M-type dwarf star AD Leo )."
" The spectra for the F, K, and M-type stars were obtained from measurements in the UV by the IUE satellite and were complemented at longer wavelengths with synthetic spectra (see Paper I for details)."," The spectra for the F, K, and M-type stars were obtained from measurements in the UV by the IUE satellite and were complemented at longer wavelengths with synthetic spectra (see Paper I for details)."
" In addition, the optical spectrum from ? between 335.5nm and 900nm and the NIR spectra of ? are used for the M-star."," In addition, the optical spectrum from \citet{Pettersen1989} between $335.5 \ \mathrm{nm}$ and $900 \ \mathrm{nm}$ and the NIR spectra of \citet{Leggett1996} are used for the M-star."
" For the Sun, a measured high-resolution spectrum is used (?).."," For the Sun, a measured high-resolution spectrum is used \citep{Gueymard2004}."
 We note that in general synthetic spectra of different stellar models may display considerable differences in the spectral lines (see e.g. ? for an intercomparison of stellar model spectra)., We note that in general synthetic spectra of different stellar models may display considerable differences in the spectral lines (see e.g. \citet{Sinclair2010MNRAS} for an intercomparison of stellar model spectra).
" The incident stellar fluxes are scaled by varying the orbital distances, such that the energy integrated over each incident stellar spectrum equals the present Total Solar Irradiance (TSI) at the top of the atmospheres."," The incident stellar fluxes are scaled by varying the orbital distances, such that the energy integrated over each incident stellar spectrum equals the present Total Solar Irradiance (TSI) at the top of the atmospheres."
 The corresponding stellar spectra were binned to the spectral intervals of the radiative transfer and are shown in Fig. 1.., The corresponding stellar spectra were binned to the spectral intervals of the radiative transfer and are shown in Fig. \ref{stellar_spectra}.
" According to Paper I, the orbital distances of the Earth-like planets are 1.89AU (F-type star), 1AU (G-type star), 0.61AU (K-type star), and 0.15AU (M-type star)."," According to Paper I, the orbital distances of the Earth-like planets are $1.89 \ \mathrm{AU}$ (F-type star), $1 \ \mathrm{AU}$ (G-type star), $0.61 \ \mathrm{AU}$ (K-type star), and $0.15 \ \mathrm{AU}$ (M-type star)."
" In our model, the measured terrestrial value of the global mean surface albedo (0.13) is used."," In our model, the measured terrestrial value of the global mean surface albedo (0.13) is used."
 A tuned surface albedo is sometimes used to mimic the effect of clouds on the surface temperature in cloud-free model calculations of Earth-like planetary atmospheres (seee.g.??)..," A tuned surface albedo is sometimes used to mimic the effect of clouds on the surface temperature in cloud-free model calculations of Earth-like planetary atmospheres \citep[see e.g.][]{Segura03, Segura05}."
" This adjustment is only feasible if the value of the resulting surface temperature is prescribed as, for example, the measured global mean surface temperature of Earth (288 K)."," This adjustment is only feasible if the value of the resulting surface temperature is prescribed as, for example, the measured global mean surface temperature of Earth $288 \ \mathrm K$ )."
" In contrast to this tuning approach, the planetary surface temperature is a result of the calculations in the case of an atmospheric model including clouds."," In contrast to this tuning approach, the planetary surface temperature is a result of the calculations in the case of an atmospheric model including clouds."
" To assess the applicability of our model for the calculation of reflection spectra of Earth-like planets around different types of central stars, we study at first the resulting spectra obtained for the Earth model introduced in Paper I. With the observed mean cloud coverages (39.5% low-level clouds, 15% high-level clouds, and a 7% overlap of both cloud layers), the measured global mean surface temperature of Earth (288K) is reproduced by our model."," To assess the applicability of our model for the calculation of reflection spectra of Earth-like planets around different types of central stars, we study at first the resulting spectra obtained for the Earth model introduced in Paper I. With the observed mean cloud coverages $39.5\%$ low-level clouds, $15\%$ high-level clouds, and a $7\%$ overlap of both cloud layers), the measured global mean surface temperature of Earth $288 \ \mathrm{K}$ ) is reproduced by our model."
" The clear sky calculation results in a surface temperature of 293K, which is clearly too high (see Paper I for a detailed discussion of the effects of clouds on the surface temperature)."," The clear sky calculation results in a surface temperature of $293 \ \mathrm{K}$, which is clearly too high (see Paper I for a detailed discussion of the effects of clouds on the surface temperature)."
" Mid-level clouds are omitted in the cloud model as most of them have been reported to be radiatively neutral in the Earth atmosphere, i.e. their greenhouse and albedo effect balance each other (see ? and Paper I&III for details)."," Mid-level clouds are omitted in the cloud model as most of them have been reported to be radiatively neutral in the Earth atmosphere, i.e. their greenhouse and albedo effect balance each other (see \citet{Poetzsch95} and Paper II for details)."
" However, the neglect of the mid-level clouds yields less back-scattered shortwave and more outgoing longwave radiation at the top of the atmosphere than the global energy budget of Earth including all cloud types."," However, the neglect of the mid-level clouds yields less back-scattered shortwave and more outgoing longwave radiation at the top of the atmosphere than the global energy budget of Earth including all cloud types."
" In this study, we focus on the effect of high and low-level clouds on the reflection spectra, corresponding to the model introduced in Paper I. These cloud types represent the extreme cases for the cloud effects on the reflection spectra."," In this study, we focus on the effect of high and low-level clouds on the reflection spectra, corresponding to the model introduced in Paper I. These cloud types represent the extreme cases for the cloud effects on the reflection spectra."
 The observations of e.g. ?? have shown that the effects of mid-level clouds on the reflection spectra of Earth are between these two extremes.," The observations of e.g. \citet{Tinetti2006a,Tinetti2006b} have shown that the effects of mid-level clouds on the reflection spectra of Earth are between these two extremes."
" Therefore, the important range of cloud effects on the reflection spectra are covered in our study of extrasolar planetary atmospheres."," Therefore, the important range of cloud effects on the reflection spectra are covered in our study of extrasolar planetary atmospheres."
 The corresponding spectra covering the wavelength range from the near UV to the IR are shown in Fig. 2.., The corresponding spectra covering the wavelength range from the near UV to the IR are shown in Fig. \ref{earth_ref_spectrum}.
" For a better visualisation, the IR part of the spectrum has been multiplied by a factor of 5."," For a better visualisation, the IR part of the spectrum has been multiplied by a factor of 5."
" In the visible range (<4 um) the spectrum is a reflection spectrum due to back-scattered, and partly absorbed, incident solar radiation."," In the visible range $< 4 \ \mathrm{\mu m}$ ) the spectrum is a reflection spectrum due to back-scattered, and partly absorbed, incident solar radiation."
" The emission spectrum in the IR wavelength range originates from the thermal emissions of the planetary surface and the atmosphere, affected by absorption and also scattering in the presence of clouds (see Paper II for details)."," The emission spectrum in the IR wavelength range originates from the thermal emissions of the planetary surface and the atmosphere, affected by absorption and also scattering in the presence of clouds (see Paper II for details)."
 The number of chemical species visible in a low-resolution spectrum is rather small., The number of chemical species visible in a low-resolution spectrum is rather small.
" In the IR range, CO; (15 um), Η2Ο. (vibrational and rotational at 5—8um and >15 pum), and the 9.6um Os band can be seen at low resolution."," In the IR range, $\mathrm{CO_2}$ $\sim 15 \ \mathrm{\mu m}$ ), $\mathrm{H_2 O}$ (vibrational and rotational at $5-8 \ \mathrm{\mu m}$ and $> 15 \ \mathrm{\mu m}$ ), and the $9.6 \ \mathrm{\mu m}$ $\mathrm{O_3}$ band can be seen at low resolution."
" The NUV-NIR range allows for the detection of O5, H2O, and in principle also O3 (see Fig. 2))."," The NUV-NIR range allows for the detection of $\mathrm{O_2}$, $\mathrm{H_2 O}$, and in principle also $\mathrm{O_3}$ (see Fig. \ref{earth_ref_spectrum}) )."
" Several H5O bands can be found in the reflection spectrum at for instance ~0.95um, ~1.1um, and ~1.4wm."," Several $\mathrm{H_2 O}$ bands can be found in the reflection spectrum at for instance $\sim 0.95 \ \mathrm{\mu m}$, $\sim 1.1 \ \mathrm{\mu m}$, and $\sim 1.4 \ \mathrm{\mu m}$."
 The small Fraunhofer A-band of O5 is visible at 0.76um., The small Fraunhofer A-band of $\mathrm{O_2}$ is visible at $0.76 \ \mathrm{\mu m}$.
" In principle, the broad Chappuis band of Os (0.55—0.70um) is also present, which, however, cannot be directly seen in the calculated low-resolution reflection spectrum of Earth (see Fig."," In principle, the broad Chappuis band of $\mathrm{O_3}$ $0.55-0.70 \ \mathrm{\mu m}$) is also present, which, however, cannot be directly seen in the calculated low-resolution reflection spectrum of Earth (see Fig."
 2 of Paper I for the atmospheric profiles of chemical species)., 2 of Paper I for the atmospheric profiles of chemical species).
" While for instance a spectral resolution of R=20 is sufficient to detect the Os absorption band at 9.6um in the IR (see ?),, it has to be R>50 for the detection of O at 0.76uum in the reflection spectrum."," While for instance a spectral resolution of $R=20$ is sufficient to detect the $\mathrm{O_3}$ absorption band at $9.6 \ \mathrm{\mu m}$ in the IR \citep[see][]{Kaltenegger2009}, , it has to be $R>50$ for the detection of $\mathrm{O_2}$ at $0.76 \ \mathrm{\mu m}$ in the reflection spectrum."
 We, We
of Fig.,of Fig.
 3 shows excellent agreement between the two described methods., \ref{fig:cool_smallg_safetyfactor} shows excellent agreement between the two described methods.
" We then increase the gravity parameter by a factor of 100, ie. g—0.100."," We then increase the gravity parameter by a factor of 100, i.e. $g=0.100$."
 The pressure and plasma beta are shown in Fig. 4.., The pressure and plasma beta are shown in Fig. \ref{fig:cool_intermediateg_pressure}.
" Since the fact that the parameter g is 100 times stronger, the pressure maximum is shifted downwards as expected."," Since the fact that the parameter $g$ is 100 times stronger, the pressure maximum is shifted downwards as expected."
" The plasma beta has values similar to those in the previous case, its maximum and minimum values being 0.105 and 0.014, respectively."," The plasma beta has values similar to those in the previous case, its maximum and minimum values being 0.105 and 0.014, respectively."
" 'The two-dimensional density profile, not shown here, has a similar behavior to that of the pressure profile."," The two-dimensional density profile, not shown here, has a similar behavior to that of the pressure profile."
 Its maximum is also shifted downwards as expected., Its maximum is also shifted downwards as expected.
 Fig., Fig.
 shows the safety factor q and the radial derivative of the Shafranov shift A(r)., \ref{fig:cool_intermediateg_safetyfactor} shows the safety factor $q$ and the radial derivative of the Shafranov shift $\Delta(r)$.
" Comparing the safety factor with the one of the previous equilibrium, one notices that there is hardly any difference."," Comparing the safety factor with the one of the previous equilibrium, one notices that there is hardly any difference."
 The two methods for computing the Shafranov shift now show a small difference., The two methods for computing the Shafranov shift now show a small difference.
" The second method, which using the Shafranov shift equation in Eq.(25),"," The second method, which using the Shafranov shift equation in Eq.,"
", underestimates the radial derivative for all radii.", underestimates the radial derivative for all radii.
 'The last cool solar prominence we discuss is the one for which the dimensionless gravity parameter is 1.000., The last cool solar prominence we discuss is the one for which the dimensionless gravity parameter is $g=1.000$ .
 The gravity is 1000 times stronger than in the first discussed equilibrium and 10 times stronger than the last one., The gravity is 1000 times stronger than in the first discussed equilibrium and 10 times stronger than the last one.
 The downward shift of the pressure is even stronger as shown in Fig. 6.., The downward shift of the pressure is even stronger as shown in Fig. \ref{fig:cool_strongg_pressure}.
" For this strong g case, the plasma beta varies from 0.013 at the edge up to 1.107 at the core of the plasma."," For this strong $g$ case, the plasma beta varies from 0.013 at the edge up to 1.107 at the core of the plasma."
 The latter value is larger than one would expect from observations., The latter value is larger than one would expect from observations.
" However, it is straightforward to compute an even more realistic equilibrium by adjusting the coefficient A» of the pressure equation in Eq."," However, it is straightforward to compute an even more realistic equilibrium by adjusting the coefficient $A_{2}$ of the pressure equation in Eq."
"(31).. For this equilibrium, the safety factor and the Shafranov shift are shown in Fig. 7.."," For this equilibrium, the safety factor and the Shafranov shift are shown in Fig. \ref{fig:cool_strongg_safetyfactor}."
" Owing to the strong gravity the safety factor is very different from its value for the two previous equilibria, particularly in terms of the multiple q=1 surfaces."," Owing to the strong gravity the safety factor is very different from its value for the two previous equilibria, particularly in terms of the multiple $q=1$ surfaces."
" In addition, the plot of the Shafranov shift shows a large discrepancy between the two methods, because of the strong gravity."," In addition, the plot of the Shafranov shift shows a large discrepancy between the two methods, because of the strong gravity."
 The first method of determining the Shafranov shift is superior in terms of accuracy and numerical cost., The first method of determining the Shafranov shift is superior in terms of accuracy and numerical cost.
" To illustrate the difference between the chosen flux function, we also computed the cases where the density or the entropy is a flux function for gravity g—1.000."," To illustrate the difference between the chosen flux function, we also computed the cases where the density or the entropy is a flux function for gravity $g=1.000$."
" For these choices, the pressure and plasma beta are plotted in Fig. 8.."," For these choices, the pressure and plasma beta are plotted in Fig. \ref{fig:cool_strongg_rho_S}."
" When density is a flux function, the plasma beta varies from 0.011 up to 0.254, and for the entropy case the range is 0.012 up to 0.395."," When density is a flux function, the plasma beta varies from 0.011 up to 0.254, and for the entropy case the range is 0.012 up to 0.395."
 Both plots are clearly distinct from the case where the temperature is chosen to be a flux function., Both plots are clearly distinct from the case where the temperature is chosen to be a flux function.
" From the equilibrium viewpoint, all three are realizable, but high resolution observations will be required to help us to select the properflux function for a particular filament."," From the equilibrium viewpoint, all three are realizable, but high resolution observations will be required to help us to select the properflux function for a particular filament."
" 'The second class of equilibria, based on the description presented by ?,, allows for multi-layered prominences."," The second class of equilibria, based on the description presented by \cite{Petrie_2007}, , allows for multi-layered prominences."
probably Galactic binarics. 2 are Galactic transient. objects and one could be associated either to a compact object in a supernova remnant or to a background AGN.,"probably Galactic binaries, 2 are Galactic transient objects and one could be associated either to a compact object in a supernova remnant or to a background AGN."
 For the SOULCCS in our sample. WO απο A-rav data acquired with the three A-ray CCD cameras (MOS 1. 2 and pn) comprising the EPIC instrument. on-board the spacecraft. (Struder et al.," For the sources in our sample, we use X-ray data acquired with the three X-ray CCD cameras (MOS 1, 2 and pn) comprising the EPIC instrument on-board the spacecraft (Struder et al."
 2001)., 2001).
 Phe EPLC cameras oller the possibility to. perform sensitive imaging observations over the telescopes field of view (FOV) of 30 arcmin and in the energy range from 0.15 to 15 keV with mocerate spectral (2/247~20. 50) and angular resoluion (PSE. 6 arcsec ENIM) and are therefore. ideal for our objectives of improving the source position ancl studving broad band X-ray spectra.," The EPIC cameras offer the possibility to perform sensitive imaging observations over the telescope's field of view (FOV) of 30 arcmin and in the energy range from 0.15 to 15 keV with moderate spectral $\Delta E \sim 20-50$ ) and angular resolution (PSF, 6 arcsec FWHM) and are therefore ideal for our objectives of improving the source position and studying broad band X-ray spectra."
 data were reprocessed using the Standard Analvsis Software (SAS) version 9.0.0 emploving the latest available calibration files., data were reprocessed using the Standard Analysis Software (SAS) version 9.0.0 employing the latest available calibration files.
 Only patterns corresponding to single ancl double events (VPATTERN<4) were taken into account and the stanclarel selection filter FLAC=0 was applied., Only patterns corresponding to single and double events $\leq$ 4) were taken into account and the standard selection filter FLAG=0 was applied.
 The observations were filtered for. periods of high ickeground and the resulting net exposures for each source and each camera are reported in table lL which also Lists he XMM edion date., The observations were filtered for periods of high background and the resulting net exposures for each source and each camera are reported in table 1 which also lists the XMM observation date.
 For each observation. we analvsed. the EPIC (MOS plus pn) images o search for X-ray sources which fall inside the IBIS error box and are therefore likely LBES counterparts.," For each observation, we analysed the EPIC (MOS plus pn) images to search for X-ray sources which fall inside the IBIS error box and are therefore likely IBIS counterparts."
 Next we obtained X-ray spectra in the 05-12. keV band of he likely associated. source., Next we obtained X-ray spectra in the 0.5-12 keV band of the likely associated source.
 Source counts were extracted ⋅rom circular. regions. of radiusD on20 centered on the source: vackerouncl spectra were extracted. [from circular regions close to the source or [rom source-free regions of SO” radius., Source counts were extracted from circular regions of radius $^{\prime\prime}$ centered on the source; background spectra were extracted from circular regions close to the source or from source-free regions of $^{\prime\prime}$ radius.
 The ancillary response matrices CABE) and the detector response malrices (ΗΔΗSs) were eenerated. using the SAS asks and while the spectral channels were rebinned in order to achieve a minimum of 20 counts in each bin., The ancillary response matrices (ARFs) and the detector response matrices (RMFs) were generated using the SAS tasks and while the spectral channels were rebinned in order to achieve a minimum of 20 counts in each bin.
" The 1«ΕςΔΙ, data reported here consist. of. all pointings performed (Winkler et al.", The INTEGRAL data reported here consist of all pointings performed (Winkler et al.
 2003) during 5 vears ol observations with typical exposures in the range S500 ksec (see 6th column of table 1): these are the same data used o obtain the 4th LDIS survey., 2003) during 5 years of observations with typical exposures in the range 3000--8500 ksec (see 6th column of table 1); these are the same data used to obtain the 4th IBIS survey.
 ΕΙ images for each available pointing were generated in various energv bands using the ISDC ollline scientific analvsis software OSA (€rolcwurnm et al., ISGRI images for each available pointing were generated in various energy bands using the ISDC offline scientific analysis software OSA (Goldwurm et al.
 2003) version 7.0., 2003) version 7.0.
 Count rates at the position of the source were extracted from individual images in order to provide light curves in various cnerev bands: from these light curves. average Ηχος were then extracted and combined to produce an average source spectrum (see Ard et al.," Count rates at the position of the source were extracted from individual images in order to provide light curves in various energy bands; from these light curves, average fluxes were then extracted and combined to produce an average source spectrum (see Bird et al."
 2010 for details)., 2010 for details).
 In the last column of table 1 the variability indicator as defined by Birel ct al. , In the last column of table 1 the variability indicator as defined by Bird et al. (
CC10) is also reported.,2010) is also reported.
 The data were then. fitted together with average spectra usingNSPEC 110v. 12.5.1 (Arnaucl 199deepa) to cover the broad band from NM05 to keV. A detailed ion of the results obtained this spectral analysis is given in a dedicated section for each source., The data were then fitted together with average spectra using v. 12.5.1 (Arnaud 1996) to cover the broad band from 0.5 to 110 keV. A detailed description of the results obtained by this spectral analysis is given in a dedicated section for each source.
 All together we analysed the data of 6 LBIS sources (see table 1)., All together we analysed the data of 6 IBIS sources (see table 1).
 ALL sources. discussed in this on appear in the fourth. IBIS catalogue 'Did et al., All sources discussed in this section appear in the fourth IBIS catalogue (Bird et al.
 menu) either as new detections (AX J1740.2-2903 and LOR 0102) Or as already MA hard. N-rav emitters (LGR 15359-5759N IC J17331-m AX J1739.3-2023 M IC J17445-2147m," 2010) either as new detections (AX J1740.2-2903 and IGR J18538-0102) or as already known hard X-ray emitters (IGR J15359-5750, IGR J17331-2406, AX J1739.3-2923 and IGR J17445-2747)."
 Apart from case of LGR. 11331-2406n for which we not find any XNMM counterpart w the LBIS positional uncertainty ancl that of LGR J17445-2741 where the only XMM detection inside the LBIS error box is possibly not the correct. counterpart of the high energy. source. for the remaining + sources we were able to find a convincing counterpart in the X-ray. band.," Apart from the case of IGR J17331-2406 for which we could not find any XMM counterpart within the IBIS positional uncertainty and that of IGR J17445-2747 where the only XMM detection inside the IBIS error box is possibly not the correct counterpart of the high energy source, for the remaining 4 sources we were able to find a convincing counterpart in the X-ray band."
 lor these sources db was also possible to extract the X-ray spectrum and to perform a broad. band analysis., For these sources it was also possible to extract the X-ray spectrum and to perform a broad band analysis.
 Table 2 lists the sources with X-ray detections together with their IBIS position ancl relative uncertainty as reported in Bird et al. , Table 2 lists the sources with X-ray detections together with their IBIS position and relative uncertainty as reported in Bird et al. (
).,2010).
 For cach of these eanima-ray οπίου. we provide the position and relative uncertainties (at level) of the sources detected by NATAL within theNM aMLBIS error counterpartscircle and considered. to be the likely (see section on individual sources): we also report their count rate in the 0.212 keV energy. band: and optical/infrared identification obtained. from. various on-line archives such as NED (NASA/IPAC Extragalactic Database). HEASARC (Ligh [ποιον Astrophysics. Science. Archive Research Center) and SLIAIBAD (Set. of. lelentifications. Aleasurements. ancl Bibliography for Astronomical Data).," For each of these gamma-ray emitters, we provide the position and relative uncertainties (at confidence level) of the sources detected by XMM within the IBIS error circle and considered to be the likely X-ray counterparts (see section on individual sources); we also report their count rate in the 0.2–12 keV energy band and optical/infrared identification obtained from various on-line archives such as NED (NASA/IPAC Extragalactic Database), HEASARC (High Energy Astrophysics Science Archive Research Center) and SIMBAD (Set of Identifications, Measurements, and Bibliography for Astronomical Data)."
 The results of the broad band spectral analysis relative to the 4 objects having good quality data are reported. in table 3 where we list the Galactic absorption. the column," The results of the broad band spectral analysis relative to the 4 objects having good quality data are reported in table 3 where we list the Galactic absorption, the column"
as — and its affinity for dust grains.,as – and its affinity for dust grains.
 Strong absorbers at OLSTube1.3 cause significant reddening of the background quasar light by dust. unlike DLAs and Mg I-selected absorption line systems (Murphy&Liske2004:Wild.Hewett.Pettini2006:Yorketal. 2006).. and an intriguing trend of increasing dust content with increasing rest-frame equivalent width (11°) is observed.," Strong absorbers at $0.8\la\zabs\la1.3$ cause significant reddening of the background quasar light by dust, unlike DLAs and Mg –selected absorption line systems \citep{2004MNRAS.354L..31M,2006MNRAS.367..211W,2006MNRAS.367..945Y}, and an intriguing trend of increasing dust content with increasing rest–frame equivalent width $W$ ) is observed."
 Within the Milky Way. Hunteretal.(2006) suggest that traces warm neutral gas clouds. thus strong lines ats 1 may probe he inner disks of chemically evolved galaxies.," Within the Milky Way, \citet{2006MNRAS.367.1478H} suggest that traces warm neutral gas clouds, thus strong lines at $z\sim1$ may probe the inner disks of chemically evolved galaxies."
 Alternatively. large abundances may arise when dust grains are destroyed. perhaps within shocks caused by major mergers (Wangetal.2005)..," Alternatively, large abundances may arise when dust grains are destroyed, perhaps within shocks caused by major mergers \citep{2005ApJ...622L.101W}."
 The shysics of dust is not well understood and destruction of a small Traction of the grains could lead to a large relative increase in gaseous.. consistent with the observed presence of dust in the systems.," The physics of dust is not well understood and destruction of a small fraction of the grains could lead to a large relative increase in gaseous, consistent with the observed presence of dust in the systems."
 A first step towards understanding the true nature of he systems is through imaging. which. for systems with redshifts below about unity. is well within the capabilities of 4m-class elescopes.," A first step towards understanding the true nature of the systems is through imaging, which, for systems with redshifts below about unity, is well within the capabilities of 4m-class telescopes."
 Aside from the objective of understanding the origins of the unusually strong lines. absorbers are expected to contain arge neutral hydrogen columns. with a significant number above he nominal DLA limit (Wild.Hewett.&Pettini2006).," Aside from the objective of understanding the origins of the unusually strong lines, absorbers are expected to contain large neutral hydrogen columns, with a significant number above the nominal DLA limit \citep{2006MNRAS.367..211W}."
. With 7-400 strong absorbers in the SDSS Data Release + (DR+4) quasar sample. compared to the ~40 Known DLAs at 2&1.65 (Rao.Turnshek.&Nestor2006). they potentially represent an important new method for the selection of galaxies by hydrogen cross section at low and intermediate redshifts where the Lyman-a line is currently observationally inaccessible.," With $\simeq$ 400 strong absorbers in the SDSS Data Release 4 (DR4) quasar sample, compared to the $\sim40$ known DLAs at $z\la1.65$ \citep{2006ApJ...636..610R}, they potentially represent an important new method for the selection of galaxies by hydrogen cross section at low and intermediate redshifts where the $\alpha$ line is currently observationally inaccessible."
 In this paper we present the results of A -band imaging of the fields around 30 absorbers with 0.7<z«1.2., In this paper we present the results of $K$ -band imaging of the fields around 30 absorbers with $0.7<z<1.2$.
 The typica number density of A -band objects is well known and low enough for single band imaging to be effective given the expected smal impact parameters between the absorbing galaxy and quasar., The typical number density of $K$ -band objects is well known and low enough for single band imaging to be effective given the expected small impact parameters between the absorbing galaxy and quasar.
 The sample selection is described in Section ??— and observations and data reduction in Section ??.., The sample selection is described in Section \ref{sec:sample} and observations and data reduction in Section \ref{sec:secobs}.
 In Section ?? we compare our A-band number counts to previous surveys and discuss the impact of potential contamination from the targetting of quasar fields (quasar host and companion galaxies)., In Section \ref{sec:results} we compare our $K$ -band number counts to previous surveys and discuss the impact of potential contamination from the targetting of quasar fields (quasar host and companion galaxies).
 Section ?? presents our results for he luminosity. impact parameters and morphology of the apsorber galaxies.," Section \ref{sec:props} presents our results for the luminosity, impact parameters and morphology of the absorber galaxies."
 In Section ??> we firs discuss these results in comparison to strong absorbers and DLAs. hen we consider the implications for the present day star formation rate derived in Wild.Hewett.& (2007).," In Section \ref{sec:disc} we first discuss these results in comparison to strong absorbers and DLAs, then we consider the implications for the present day star formation rate derived in \citet{2007MNRAS.374..292W}."
. Finally. in Sectioi1 2?.. we discuss viable models for the true nature of absorpion line systems.," Finally, in Section \ref{sec:concl}, we discuss viable models for the true nature of absorption line systems."
 Optical apparent and absolute magnitudes taken from the SDSS catalogues ure left in the AB-system. while the observed éy-band magnitudes are &iven on the Vega system. as employed by 2MASS.," Optical apparent and absolute magnitudes taken from the SDSS catalogues are left in the AB-system, while the observed $K$ -band magnitudes are given on the Vega system, as employed by 2MASS."
 Conversion between the systems can be achieved using the relations /XD)—i(Vega)|0.37. AostaXD)=Axoxpyss(Vega)|186 and. similarly. for the Mauna Kea Observatory. A-band.. Aano(AB)=Awuko(Vega)|1.90 (Hewettetal.2006.withniy-=+0.03forVega).," Conversion between the systems can be achieved using the relations $i{\rm (AB)}=i{\rm (Vega)}+0.37$, $K_{\rm 2MASS}{\rm (AB)}= K_{\rm
2MASS}{\rm (Vega)}+1.86$ and, similarly, for the Mauna Kea Observatory $K$ -band, $K_{\rm MKO}{\rm (AB)}=K_{\rm MKO}{\rm (Vega)}+1.90$ \citep[][with $m_V$=+0.03 for
Vega]{2006MNRAS.367..454H}."
 We use a flat cosmology with O4.=0.7. (day=(Q3 1005kms+Alpe* and f=0.7 throughout the paper.," We use a flat cosmology with $\Omega_\Lambda=0.7$, $\Omega_M=0.3$, $H_0=100\,h\,{\rm km\,s}^{-1}\,{\rm Mpc}^{-1}$ and $h=0.7$ throughout the paper."
 The absorber sample of 30 objects to be imaged was selected from the catalogue of 345 absorbers described in Wild.Hewett.&Pettini(2007)., The absorber sample of 30 objects to be imaged was selected from the catalogue of 345 absorbers described in \citet{2007MNRAS.374..292W}.
.. The absorber redshift range 0.N. was chosen to probe the old stellar populations associated with absorbers at a significant lookback-time., The absorber redshift range $0.8 \la z_{abs} \la 1.1$ was chosen to probe the old stellar populations associated with absorbers at a significant lookback-time.
 Three absorbers in the catalogue with measurements from Rao.Turnshek.&Nestor(2006) were preferentially targeted. including two with 0.7«σος0.8.," Three absorbers in the catalogue with measurements from \citet{2006ApJ...636..610R}
 were preferentially targeted, including two with $0.7 < z_{abs} < 0.8$."
 To avoid complications resulting from the detection of galaxies associated with the background quasar. targets with zstubs|0.2 were selected. with the majority of targets possessing differences of Az70.5.," To avoid complications resulting from the detection of galaxies associated with the background quasar, targets with $z_{quasar} \ga z_{abs}+0.2$ were selected, with the majority of targets possessing differences of $\Delta z > 0.5$."
 Otherwise. the make up of the imaging sample was determined by observing constraints. including accessibility of the fields from the United Kingdom Infra-Red Telescope (UKIRT) and the availability of a suitable guide star.," Otherwise, the make up of the imaging sample was determined by observing constraints, including accessibility of the fields from the United Kingdom Infra-Red Telescope (UKIRT) and the availability of a suitable guide star."
 The sample includes the full range of quasar brightness and equivalent width within the catalogue of 345 absorbers., The sample includes the full range of quasar brightness and equivalent width within the catalogue of 345 absorbers.
 To summarise. the 30 absorbers were selected according to the following criteria: In addition. a small sample of tive “control” quasars were imaged.," To summarise, the 30 absorbers were selected according to the following criteria: In addition, a small sample of five `control' quasars were imaged."
 The quasars were selected from (Schneideretal.2005) to have zs01.5. mic18.0 and to possess no detected absorption systems with cup.«1.2.," The quasars were selected from \citep{2005AJ....130..367S} to have $z_{quasar} \simeq 1.5$, $m_i \simeq 18.0$ and to possess no detected absorption systems with $z_{abs}
< 1.2$."
 Finally. a single “blank tield was also imaged.," Finally, a single `blank field' was also imaged."
 Table | includes the observing log for the target and control sample. along with the positions and redshifts of the targets.," Table 1 includes the observing log for the target and control sample, along with the positions and redshifts of the targets."
 The key aim of our observation and data reduction strategy was to obtain a well characterised point spread function (PSF) to allow accurate subtraction of the quasar image and recover candidate absorber host galaxies as faint and close to the quasar as possible., The key aim of our observation and data reduction strategy was to obtain a well characterised point spread function (PSF) to allow accurate subtraction of the quasar image and recover candidate absorber host galaxies as faint and close to the quasar as possible.
 Imaging was undertaken on the 3.8m UKIRT using the Fast-Track Imager CUFTT) near-infared imager (Rocheetal.2003) and a Mauna Kea Observatory Av-band filter. over the nights of 2006 April HI and 14-17 UT.," Imaging was undertaken on the 3.8m UKIRT using the Fast-Track Imager (UFTI) near-infared imager \citep{2003SPIE.4841..901R} and a Mauna Kea Observatory $K$ -band filter, over the nights of 2006 April 11 and 14-17 UT."
 The transparency on 2006 April 11 was variable. but otherwise transparency was good and conditions photometric for a significant fraction of the observations.," The transparency on 2006 April 11 was variable, but otherwise transparency was good and conditions photometric for a significant fraction of the observations."
 The median on-chip seeing was 07662 full width at half maximum (FWHM). with a full range of 07440-07885 FWHM. which was well sampled by the UFTI image scale of 00091 per pixel.," The median on-chip seeing was 62 full width at half maximum (FWHM), with a full range of 85 FWHM, which was well sampled by the UFTI image scale of 091 per pixel."
 The on-chip seeing generally improved over the first few hours of observation each night., The on-chip seeing generally improved over the first few hours of observation each night.
 An isolated Asst012.3 star was observed before and after each target to provide a reference PSF for use in the subsequent quasar-subtractions., An isolated $K_{\rm 2MASS}\simeq 12.3$ star was observed before and after each target to provide a reference PSF for use in the subsequent quasar-subtractions.
 The default observing sequence for a target consisted of a single 9-point dither of the PSF-star with 35s exposures. followed by a series of seven 9-point dithers for the target with 40ss exposures. finally. the single 9-point dither," The default observing sequence for a target consisted of a single 9-point dither of the PSF-star with s exposures, followed by a series of seven 9-point dithers for the target with s exposures, finally, the single 9-point dither"
power spectrum of the SW signal in cach slice. together with the predicted. linear ISW signal.,"power spectrum of the ISW signal in each slice, together with the predicted linear ISW signal."
" For these predictions. we have accounted for the smearing effects of the photometric redshift data by assuming that photo-z's apply Gaussian smoothing along the radial axis with o,—90. Mpe."," For these predictions, we have accounted for the smearing effects of the photometric redshift data by assuming that $z$ 's apply Gaussian smoothing along the radial axis with $\sigma_{r}=90\mpcoh$ ."
 Fig., Fig.
 2 also shows the cllect of this radial smearing on the total predicted ISN signal to saunas=0.3., \ref{fig:ISW_angpow} also shows the effect of this radial smearing on the total predicted ISW signal to $z_{\mathrm{max}}=0.3$.
 Overall. the agreement between observed and. predicted. power spectra is) good. which illustrates that the elfect of photo-z imperfection is relatively minor for f=20.," Overall, the agreement between observed and predicted power spectra is good, which illustrates that the effect of $z$ imperfection is relatively minor for $\ell\lsim 20$."
 We note that the estimated ISW signal in the lower two redshilt slices is slightly larger than expected at (6z 20: this is particularly noticeable in the z«0.1 slice., We note that the estimated ISW signal in the lower two redshift slices is slightly larger than expected at $\ell\gsim20$ ; this is particularly noticeable in the $z<0.1$ slice.
 The nonlinear ]tees-Sciama effect is expected to increase the power at high f. doubling the power at £z200 (e.g.2)..," The nonlinear Rees-Sciama effect is expected to increase the power at high $\ell$, doubling the power at $\ell\simeq 200$ \citep[e.g.][]{Cooray_and_Sheth2002}."
 Since the mean depth of the lowest-redshift shell is of order one tenth of the distances that dominate the total ISW effect. it is plausible hat we are secing sonie local Rees-Sciama effect.," Since the mean depth of the lowest-redshift shell is of order one tenth of the distances that dominate the total ISW effect, it is plausible that we are seeing some local Rees-Sciama effect."
 In. any case. the main focus of t10 present. paper is at larger angular scales.," In any case, the main focus of the present paper is at larger angular scales."
 We also note tha the estimated ISW power is larger han expected for 6S3 in the 0.1<τς0.2 shell and. for (<=5 in the 02«z0.3 shell. which leads to a larger han average total signa for τμ=0.3 on such scales.," We also note that the estimated ISW power is larger than expected for $\ell\lsim 3$ in the $0.1<z<0.2$ shell and for $\ell\lsim5$ in the $0.2<z<0.3$ shell, which leads to a larger than average total signal for $z_{\mathrm{max}}=0.3$ on such scales."
 Such small-£ modes are potentiallv sensitive to the corrections or mask incompleteness near the plane. but. tests of our method on mock data slow no tendeney to bias the power ugh in this regime (altrough it does increase the already considerable cosmic variance).," Such $\ell$ modes are potentially sensitive to the corrections for mask incompleteness near the plane, but tests of our method on mock data show no tendency to bias the power high in this regime (although it does increase the already considerable cosmic variance)."
" Any discrepancy. is relative ο our standard mocdel o ο0,,20.3. which is perhaps on the uigh-clensity side given current data (2).."," Any discrepancy is relative to our standard model of $\Omega_m=0.3$, which is perhaps on the high-density side given current data \citep{KomatsuWMAP09}."
 Deercasing ελ ο 0.25 with other paraneters held fixed raises the expected cquacirupole power for cau;x00:3 by a factor 1.5: our adopted normalization ofax=0.75 might also be increased. so that he predicted low-£ LS\W| power could. plausiblyincrease in," Decreasing $\Omega_m$ to 0.25 with other parameters held fixed raises the expected quadrupole power for $z_{\mathrm{max}}=0.3$ by a factor 1.5; our adopted normalization of$\sigma_8=0.75$ might also be increased, so that the predicted $\ell$ ISW power could plausiblyincrease in"
narrower passbands have been introduced since. but most photometric svstems use fillers significantly broader (han Stromgrens. and therefore tend to provide lower sensitivitv to the almospheric parameters.,"narrower passbands have been introduced since, but most photometric systems use filters significantly broader than Stromgren's, and therefore tend to provide lower sensitivity to the atmospheric parameters."
 Until recently. (he most widely-used photomeltric svstem was the broad-band Cousins UBVRI. but with the advent of the Sloan Digital Skv Survey (SDSS). which includes CCD-based photometry for 357 million unique optical sources over more (han 11.000 square degrees CAbazajianetal.2009).. and 2\LASS (Skrutskieοἱal.2006).. which includes nearly half a million near-IR. sources over the entire sky. these new svstems have taken over.," Until recently, the most widely-used photometric system was the broad-band Johnson-Cousins $UBVRI$, but with the advent of the Sloan Digital Sky Survey (SDSS), which includes CCD-based photometry for 357 million unique optical sources over more than 11,000 square degrees \citep{abazajian09}, and 2MASS \citep{skrutskie06}, which includes nearly half a million near-IR sources over the entire sky, these new systems have taken over."
 The heeemony of the SDSS svstem in the optical is illustrated by Che fact that new and [future insiruments. such as the Large Svnoptic Survey Telescope(LSST)!.. the Dark Energy Survey (DES)camera?.. or OSIRIS (Cepaetal.2000) on the Gran Telescopio Canarias (GTC) are adopting the same system.," The hegemony of the SDSS system in the optical is illustrated by the fact that new and future instruments, such as the Large Synoptic Survey Telescope, the Dark Energy Survey (DES), or OSIRIS \citep{cepa00} on the Gran Telescopio Canarias (GTC) are adopting the same system."
 Despite the widths of the SDSS [fillers are 36 times larger (han the Strómmeren passbands. Iveziéetal.(2008) have shown Chat. if the reddening is known with sullicient accuracy. db is possible to estimate stellar effective temperatures for late-tvpe stus wilh a (vpical precision of 100 Ix. and metallicities with a precision of ~0.2 dex [from SDSS photometry alone for stars of moderate metallicities.," Despite the widths of the SDSS filters are 3–6 times larger than the Strömmgren passbands, \cite{ivezic08} have shown that, if the reddening is known with sufficient accuracy, it is possible to estimate stellar effective temperatures for late-type stars with a typical precision of $\sim 100$ K, and metallicities with a precision of $\sim 0.2$ dex from SDSS photometry alone for stars of moderate metallicities."
 The sensitivitv of SDSS photometry to surface gravitv is much weaker. and disentangling this parameter [rom (he other two. even in (he absence of reddening. may not be possible. but in any case we are not aware of any sucessful calibration.," The sensitivity of SDSS photometry to surface gravity is much weaker, and disentangling this parameter from the other two, even in the absence of reddening, may not be possible, but in any case we are not aware of any sucessful calibration."
 Because of the vast number of sources with available SDSS photometry. it is desirable (o ensure (hat we are extracting all possible information captured by (his svstem.," Because of the vast number of sources with available SDSS photometry, it is desirable to ensure that we are extracting all possible information captured by this system."
 The most straightforward techniques for mapping photometric indices into (he quantities of interest. such as stellar atmospheric parameters. have provided limited success. and (he time is ripe to explore new avenues.," The most straightforward techniques for mapping photometric indices into the quantities of interest, such as stellar atmospheric parameters, have provided limited success, and the time is ripe to explore new avenues."
 This work examines the possibility of using a simple technique inspired on the concept of compressed sensing (CS:Candésetal.2006:Donoho2006) {ο reconstruct spectroscopic data of stellar objects [rom SDSS photometry.," This work examines the possibility of using a simple technique inspired on the concept of compressed sensing \citep[CS;][]{candes06,donoho06} to reconstruct spectroscopic data of stellar objects from SDSS photometry."
 If there is an accessible mapping between photometry ancl intermediate-dispersion spectra. (he available tools for spectroscopic analvsis could then be applied to the reconstructed spectra in order to recover (he parameters of interest.," If there is an accessible mapping between photometry and intermediate-dispersion spectra, the available tools for spectroscopic analysis could then be applied to the reconstructed spectra in order to recover the parameters of interest."
 The recent work on SDSS spectra by MeGurketal.(2010). indicates Chat lor stars in narrow color bands (about 0.02 mag wide in g— r). principal component analvsis (PCA)," The recent work on SDSS spectra by \cite{mcgurk10}
 indicates that for stars in narrow color bands (about 0.02 mag wide in $g-r$ ), principal component analysis (PCA)"
actors in favor of usine ATLASCAL: (1) catalog size. (2) a broad range of source positions and velocities. aud (3) a range of evolutionary states.,"factors in favor of using ATLASGAL: (1) catalog size, (2) a broad range of source positions and velocities, and (3) a range of evolutionary states."
 ATLASGAL provides a iuuch larecr catalog than ΠΟΡΟ or the ? survey of IRAS sources;, ATLASGAL provides a much larger catalog than HOPS or the \citet{Beltran:2006} survey of IRAS sources.
 2? report from the initial results of the survey about. 6000 sources Inighter than 0.25 Jy in 95 dee? in the Galactic range AQP<JALS and 1157x1€21° with |b)<1, \citet{Schuller:2009} report from the initial results of the survey about 6000 sources brighter than 0.25 Jy in 95 $^2$ in the Galactic range $-30^\circ \leq l \leq +11.5^\circ$ and $+15^\circ \leq l \leq +21^\circ$ with $|b| \leq 1^\circ$.
" Th coutrast the 7. survey contaims 235 |sources aud TOPS (?) is expected to contain a few hundred bright ssources,", In contrast the \citet{Beltran:2006} survey contains 235 sources and HOPS \citep{Walsh:2008} is expected to contain a few hundred bright sources.
 Neither TOPS nor IRAS coutains a sufficieut uunber of sources for the scicuce goals of NALT90., Neither HOPS nor IRAS contains a sufficient number of sources for the science goals of MALT90.
" ATLASGAL sources appear to sample many Galactic structures,", ATLASGAL sources appear to sample many Galactic structures.
 Figure 10 shows the Calactic longitude and velocity of sources withdetections (using hancd-determined velocities) plotted ou the ? CO map., Figure \ref{lvdistribution} shows the Galactic longitude and velocity of sources withdetections (using hand-determined velocities) plotted on the \citet{Dame:2001} CO map.
 The ? CO map is preseuted as a longitude-velocity diagram integrated over ο<b<27 and has units of I arceee., The \citet{Dame:2001} CO map is presented as a longitude-velocity diagram integrated over $-2^\circ < b < 2^\circ$ and has units of K arcdeg.
 The positious of our sources in this plot all fall within the 0.3 I& arcdeg CO contour and most cluster iu the portions of stronger CO cussion. as expected for deuse. star-forming eas.," The positions of our sources in this plot all fall within the 0.3 K arcdeg CO contour and most cluster in the portions of stronger CO emission, as expected for dense, star-forming gas."
 The preseuce of ATLASCAL sources at inanv positious in the loneitude-velocity diagram indicates that this catalog is detecting sources in a range of Galactic locations., The presence of ATLASGAL sources at many positions in the longitude-velocity diagram indicates that this catalog is detecting sources in a range of Galactic locations.
 ATLASGAL sources cover a range of evolutionary states., ATLASGAL sources cover a range of evolutionary states.
 The initial results of ATLASQAL found thirds of the sources do not have a inid- or far-infrared counterpart in the Midcouse Space Experiment (MSX) and IRAS catalogs., The initial results of ATLASGAL found two-thirds of the sources do not have a mid- or far-infrared counterpart in the Midcourse Space Experiment (MSX) and IRAS catalogs.
 A’ closer inspection of TRAS/AISN. dark sources iu dore scusitive GLIMPSE/MIPSQAL images reveals anv of them associated with weaker iufrared sources but still a considerable fraction of the subiülliueter ciaission appear dark in the nuages (c.g...Fig.125audllof7) and work on the first compact source release catalog demonstrates that ATLASGAL will. provide enough sources iu cach evolutionary state as assessed * the ocnmuüssion morphology scheme shown im Figure 1..," A closer inspection of IRAS/MSX dark sources in more sensitive GLIMPSE/MIPSGAL images reveals many of them associated with weaker infrared sources but still a considerable fraction of the submillimeter emission appear dark in the images \citep[e.g., Fig. 13 and 14 of][]{Schuller:2009} and work on the first compact source release catalog demonstrates that ATLASGAL will provide enough sources in each evolutionary state as assessed by the emission morphology scheme shown in Figure \ref{spitzer}."
 We thus choose ATLASQAL as the sole input catalog or the MALT90./ survey because if meets all our requirements for a source list., We thus choose ATLASGAL as the sole input catalog for the MALT90 survey because it meets all our requirements for a source list.
 ATLASCAL sources had Heh detection rates in this pilot) survey. include a diversity of evolutionary states. aud cover a broad rauge of Galactic positions.," ATLASGAL sources had high detection rates in this pilot survey, include a diversity of evolutionary states, and cover a broad range of Galactic positions."
 Choosing ATLASCAL as the sole input catalog also provides the benefits of having a single uniform catalog when selecting sources., Choosing ATLASGAL as the sole input catalog also provides the benefits of having a single uniform catalog when selecting sources.
 We cau choose our sources with a single uniform criteria and compare our MALT90 incasurements against ATLASGAL catalog properties such as the flux aud extent of STU CCLUSSION., We can choose our sources with a single uniform criteria and compare our MALT90 measurements against ATLASGAL catalog properties such as the flux and extent of 870 emission.
" The second goal MALT90 pilot survey was to test the observing set up aud verify that it allows us to achieve our science goals,", The second goal MALT90 pilot survey was to test the observing set up and verify that it allows us to achieve our science goals.
 NLAET90 is fundamentally a mapping survey: although some science goals (such as determining distances to clumps} could be achieved with a single pointing. the majority of our science Goals rely on maps.," MALT90 is fundamentally a mapping survey; although some science goals (such as determining distances to clumps) could be achieved with a single pointing, the majority of our science goals rely on maps."
 Our configuration allows us to map sources in multiple lues at high spectral resolution., Our configuration allows us to map sources in multiple lines at high spectral resolution.
 Maps of multiple lines allows us to study the chemical variation within a chunp. which is most useful if clumps are typically spatially resolved aud at least sometimes exhibit strong chemical variation.," Maps of multiple lines allows us to study the chemical variation within a clump, which is most useful if clumps are typically spatially resolved and at least sometimes exhibit strong chemical variation."
 \lapping at high velocity resolution (0.11 kin/s) allows us to study spatial variation iu line profiles which may iudicate changes iu the strength of turbulence. large scale motions (rotation. shear. iufall or outflow). or multiple velocity componcuts.," Mapping at high velocity resolution (0.11 km/s) allows us to study spatial variation in line profiles which may indicate changes in the strength of turbulence, large scale motions (rotation, shear, infall or outflow), or multiple velocity components."
 The pilot survey allowed us to verity that our ou-the-fiv maps had sufficient seusitivitv aud that we could make maps without sguificaut artifacts., The pilot survey allowed us to verify that our on-the-fly maps had sufficient sensitivity and that we could make maps without significant artifacts.
 We see strong cliceiical variation in the ALALT90 pilot sources Which validates our decision to map multiple lines in these sources., We see strong chemical variation in the MALT90 pilot sources which validates our decision to map multiple lines in these sources.
 All four of our main lines aand IENC)) ave ground state transitions of molecules with similar critical densities., All four of our main lines and ) are ground state transitions of molecules with similar critical densities.
 As tracers of deuse gas. the enuüssion from these lines typically show similar morphologics in ALALT9O pilot sources. but this is not always the case.," As tracers of dense gas, the emission from these lines typically show similar morphologies in MALT90 pilot sources, but this is not always the case."
 Figure 11 shows two example sources where yvarics significantly with respect to the other species., Figure \ref{nthp_odd} shows two example sources where varies significantly with respect to the other species.
 Figure HH. shows one source where the iiaxinnuu inteerated intensity is a factor of LS times weaker thu the other three main lines., Figure \ref{nthp_odd} shows one source where the maximum integrated intensity is a factor of 4-8 times weaker than the other three main lines.
 Conversely. the other source in Figure 11. las ecluission Which is twice as strong as that ofICN... oorICO!.," Conversely, the other source in Figure \ref{nthp_odd} has emission which is twice as strong as that of, or."
. Large variation iu the Huteerated intensity ratio iu ποπας star-forming regions has been noted before (e.g. 77)..," Large variation in the integrated intensity ratio in high-mass star-forming regions has been noted before \citep[e.g.,][]{Turner:1977,Walsh:2006}. ."
 Thus. the colubination of mapping in several lines sinmltaneouslv," Thus, the combination of mapping in several lines simultaneously"
"filtered to accept only good events (via ""PATTERN4"" and ""PLAG=0"").","filtered to accept only good events (via $\leq 4$ "" and $=$ 0"")."
" The MOS event lists were also filtered to accept only good events (via ""PATTERNx12"" and ""FLAG=0"")."," The MOS event lists were also filtered to accept only good events (via $\leq 12$ "" and $=$ 0"")."
 Inspection of the EPIC lightcurves revealed strong flaring at the end of the observation: therefore. only the first 48.1 ksec of the pn exposure was used to create Images and spectra. and only the first 53.8 ksee of the MOS exposures were used.," Inspection of the EPIC lightcurves revealed strong flaring at the end of the observation; therefore, only the first 48.1 ksec of the pn exposure was used to create images and spectra, and only the first 53.8 ksec of the MOS exposures were used."
 Circles with radit of 30 areseconds centered on 3C 287 were used as source extraction. regions., Circles with radii of 30 arcseconds centered on 3C 287 were used as source extraction regions.
 Background regions of equivalent size were taken nearby to the source. on the same CCD.," Background regions of equivalent size were taken nearby to the source, on the same CCD."
 The pn source and background spectra were generated by grouping channels 0-20.479 by a factor of 5: the MOS spectra were generated by grouping channels 0-11.999 by a factor of 15.," The pn source and background spectra were generated by grouping channels 0–20,479 by a factor of 5; the MOS spectra were generated by grouping channels 0–11,999 by a factor of 15."
" Custom response files for the EPIC spectra were generated using the tools ""rmfgen"" and ""arfgen""."," Custom response files for the EPIC spectra were generated using the tools ""rmfgen"" and ""arfgen""."
" To ensure valid constraints from the use of the  fitting statistic. the EPIC spectra were binned to require at least 20 counts per bin using the FTOOL ""grppha""."," To ensure valid constraints from the use of the $\chi^{2}$ fitting statistic, the EPIC spectra were binned to require at least 20 counts per bin using the FTOOL ""grppha""."
 Spectra were then fit using XSPEC version 11.3.2 (Arnaud 1996)., Spectra were then fit using XSPEC version 11.3.2 (Arnaud 1996).
 All EPIC spectra were fit in the 0.3-10.0 keV band as the cameras are best calibrated in this range., All EPIC spectra were fit in the 0.3–10.0 keV band as the cameras are best calibrated in this range.
" All errors in this work are confidence errors as derived using the XSPEC ""error? command. unless otherwise noted."," All errors in this work are confidence errors as derived using the XSPEC “error” command, unless otherwise noted."
 Model spectra were fitted to the data., Model spectra were fitted to the data.
 The pn and MOS spectral data were fit jointly. allowing an overall normalization constant to float between the spectra.," The pn and MOS spectral data were fit jointly, allowing an overall normalization constant to float between the spectra."
 This normalization constant was within of unity., This normalization constant was within of unity.
" All models were multiplied by the XSPEC model ""phabs"" to account for photoelectric absorption by neutral hydrogen along the line of sight. assuming default abundances."," All models were multiplied by the XSPEC model ""phabs"" to account for photoelectric absorption by neutral hydrogen along the line of sight, assuming default abundances."
 The X-ray spectra were jointly fitted using simple power-law models. where X statistics determined the goodness of each fit.," The X-ray spectra were jointly fitted using simple power-law models, where $\chi^{2}$ statistics determined the goodness of each fit."
 The three models we report are spectral fits to a power-law. broken power-law. and line functions incorporated into the power-law.," The three models we report are spectral fits to a power-law, broken power-law, and line functions incorporated into the power-law."
 We added line functions to the power-law model over a restricted energy range of 3.1-3.4 keV which is where we expect to locate an Fe K emission line based on a redshift of z21.055 for 3C 287.," We added line functions to the power-law model over a restricted energy range of 3.1–3.4 keV which is where we expect to locate an Fe K emission line based on a redshift of $z =
1.055$ for 3C 287."
 Spectral parameter values and fluxes. along with the corresponding upper limits. are recorded for fits to each of the models.," Spectral parameter values and fluxes, along with the corresponding upper limits, are recorded for fits to each of the models."
 Investigation of our orientation relative to the jet or disk comes from equivalent width methods. determined by the upper limit of the normalization for the line functions.," Investigation of our orientation relative to the jet or disk comes from equivalent width methods, determined by the upper limit of the normalization for the line functions."
 Both narrow and broad line fits were made to ensure robust limits on the equivalent width., Both narrow and broad line fits were made to ensure robust limits on the equivalent width.
 We present images as observed by and spectral fits along with parameters obtained using XSPEC., We present images as observed by and spectral fits along with parameters obtained using XSPEC.
 Figures | and 2 show an X-ray and V-band image of 3C 287 respectively. as observed byXMM-Newton.," Figures 1 and 2 show an X-ray and V-band image of 3C 287 respectively, as observed by."
 The results from the joint fits of MOS-1 and MOS-2 spectra to the pn spectra of 3C 287 over the 0.3-10.0 keV energy band are listed in Table I., The results from the joint fits of MOS-1 and MOS-2 spectra to the pn spectra of 3C 287 over the 0.3–10.0 keV energy band are listed in Table 1.
 Both a simple power-law model and power-law with a break in the 0.3-10.0 keV range yield acceptable fits based on X statistics., Both a simple power-law model and power-law with a break in the 0.3–10.0 keV range yield acceptable fits based on $\chi^{2}$ statistics.
 The best-fit power-law model gives V/r=650.3/598= 1.09. where v is the number of degrees of freedom.," The best-fit power-law model gives $\chi^{2}/\nu = 650.3/598 = 1.09$ , where $\nu$ is the number of degrees of freedom."
 The best-fit broken power-law model result is v/v2643.4/596=1.08: this is not a statistically significant improvement over the simple power-law model., The best-fit broken power-law model result is $\chi^2/\nu = 643.4/596 = 1.08$; this is not a statistically significant improvement over the simple power-law model.
 Spectral fits using a power-law model are shown in Figure 3., Spectral fits using a power-law model are shown in Figure 3.
 The simple power-law model gives a flux of F=7.1+0.2« eres. cm7s7! over the 0.3-10.0 keV band., The simple power-law model gives a flux of $F = 7.1\pm 0.2 \times 10^{-13}$ ergs ${\rm cm}^{-2} {\rm s}^{-1}$ over the 0.3–10.0 keV band.
 We calculate a corresponding bolometric luminosity of L=5.830.2.&10%ergss! based on the stated redshift z21.055. allowing us to place a minimum mass estimate on the central engine.," We calculate a corresponding bolometric luminosity of $L = 5.8\pm 0.2 \times
10^{45}~{\rm ergs}~ {\rm s}^{-1}$ based on the stated redshift $z =
1.055$, allowing us to place a minimum mass estimate on the central engine."
 The lower mass limit of 3C 287 is Mprnin74.6«107M... assuming that the X-ray luminosity accounts for 10% of the bolometric luminosity and radiationat the Eddington luminosity limit.," The lower mass limit of 3C 287 is $M_{BH min} \geq 4.6 \times
10^{7}~ {\rm M_{\odot}}$, assuming that the X-ray luminosity accounts for $\%$ of the bolometric luminosity and radiationat the Eddington luminosity limit."
" We calculate a recession velocity of v4, from the redshift. corresponding to a distance of d= assuming Hy=71kms! Mpe7!."," We calculate a recession velocity of $v_{rec} =
(0.62)c$ from the redshift, corresponding to a distance of $d = 2.6~
{\rm Gpc}$ assuming $H_0 = 71 ~{\rm km}~ {\rm s}^{-1}~ {\rm Mpc}^{-1}$ ."
 This is the distance we assume when calculating the luminosity., This is the distance we assume when calculating the luminosity.
RAV Επ is an eclipsing nova-like cataclysmic variable (CV) system with an orbital period of 5 hours 34 mins.,RW Tri is an eclipsing nova-like cataclysmic variable (CV) system with an orbital period of 5 hours 34 mins.
 πα systems contain a non-magnetic white dwarf primary and a [ate type (Ix-M) secondary star that fills its Roche lobe. transferring mass to the white dwarf via an accretion clisc," Nova-like systems contain a non-magnetic white dwarf primary and a late type (K-M) secondary star that fills its Roche lobe, transferring mass to the white dwarf via an accretion disc."
 Since first being observed as an eclipsing variable (Protiteh. 1937). multi-waveleneth studies of RAV Tri have provided much insight into the svstem. but have vet to vield an accurate measurement. of the masses of. the two component stars.," Since first being observed as an eclipsing variable (Protitch 1937) multi-wavelength studies of RW Tri have provided much insight into the system, but have yet to yield an accurate measurement of the masses of the two component stars."
 By observing the velocities of the component stars. their masses can be calculated: using Ixepler's laws.," By observing the velocities of the component stars, their masses can be calculated using Kepler's laws."
 A number of authors have attempted to measure the radial velocity of the white clwarl (c.g. Iwaitchuck Loss. Still 1995 ancl Mason. 2002). but thus far little work has been done on the secondary star.," A number of authors have attempted to measure the radial velocity of the white dwarf (e.g. Kaitchuck 1983, Still 1995 and Mason 2002), but thus far little work has been done on the secondary star."
 Aleasurement of the secondary star velocity is hampoered bv the large contrast in optical brightness between it and the accretion disk., Measurement of the secondary star velocity is hampered by the large contrast in optical brightness between it and the accretion disk.
" Weak secondary star absorption features were detected in the I-band spectrum of RAW Tri by. Friend (LOSS). and later used by Smith (1993) to estimate the secondary. star velocity as 250km/s using a ""skew-mapping technique (cf."," Weak secondary star absorption features were detected in the I-band spectrum of RW Tri by Friend (1988), and later used by Smith (1993) to estimate the secondary star velocity as $\sim250$ km/s using a `skew-mapping' technique (cf."
 Smith 1998a)., Smith 1998a).
 Dhillon (2000) have detected secondary. star features in low-resolution W-hancl spectra of RW Pri., Dhillon (2000) have detected secondary star features in low-resolution K-band spectra of RW Tri.
 “Phe relative contribution of the secondary. star is higher in the Ix-band and the equivalent width of the absorption features ercaler as a Consequence., The relative contribution of the secondary star is higher in the K-band and the equivalent width of the absorption features greater as a consequence.
 In this paper we report phasc-resolved spectral observations of RAV Tri in the WK-bancl. covering the binary orbit.," In this paper we report phase-resolved spectral observations of RW Tri in the K-band, covering the binary orbit."
 We use these observations to determine the racial velocity of the secondary star in RW Tri., We use these observations to determine the radial velocity of the secondary star in RW Tri.
 We also re-analyse the original I-band spectra of Smith (1993)., We also re-analyse the original I-band spectra of Smith (1993).
 Based on these results we revisit the mass ratio and the masses of the component stars in RAV Fri and discuss their implications., Based on these results we revisit the mass ratio and the masses of the component stars in RW Tri and discuss their implications.
 We have re-analvsed the RW Tri near I-band data of Smith (1993) using the back-projection routines in the software package MOLLY written by Tom. Marsh. (see also, We have re-analysed the RW Tri near I-band data of Smith (1993) using the back-projection routines in the software package MOLLY written by Tom Marsh (see also
where £ is the total length of DNA.,where $L$ is the total length of DNA.
 Hence. in order to find the partition function one must solve the Schrodinger-like differential equation (23)) and integrate the solution over all O values.," Hence, in order to find the partition function one must solve the Schrodinger-like differential equation \ref{sch2}) ) and integrate the solution over all $\Theta$ values."
 To solve equation (23)). we rewrite the Hamiltonian in (lie form where and Furthermore. we decompose fy lo its aud parts where and I is the Hamiltonian of a quantum top.," To solve equation \ref{sch2}) ), we rewrite the Hamiltonian in the form where and Furthermore, we decompose $H_0$ to its and parts where and $H_0^{\,R}$ is the Hamiltonian of a quantum top."
" It commutes with both J, and J.. where J. is the third component of (he angular momentum operator in the laboratory coordinate system |?].."," It commutes with both $J_3$ and $J_z$, where $J_z$ is the third component of the angular momentum operator in the laboratory coordinate system \cite{Landau}."
" Since J, and J. also commute with each other. one can find the simultaneous eigenvectors of these three operators."," Since $J_3$ and $J_z$ also commute with each other, one can find the simultaneous eigenvectors of these three operators."
 We denote these simultaneous eigenvectors by |n.&.m) where 7 and m are inleger numbers relerring (o the eigenvalues of Jy ancl J..respectively.," We denote these simultaneous eigenvectors by $|n,\,k,\,m\rangle$ where $k$ and $m$ are integer numbers referring to the eigenvalues of $J_3$ and $J_z$,respectively."
 The «quantum number n» distinguishes between the eigenvectors with identical / and i» numbers: From equations (27)). (322)). and (33)). it can further be seen Chat the eigenvectors of Lf) are also eigenvectors of 1H: where," The quantum number $n$ distinguishes between the eigenvectors with identical $k$ and $m$ numbers: From equations \ref{H0}) ), \ref{eig1}) ), and \ref{eig2}) ), it can further be seen that the eigenvectors of $H_0^{\,R}$ are also eigenvectors of $H_0$ : where"
Amone the available data. those collected with aare the best suited to search for X-ray emission. [rom bbecause of the superb angular resolution.,"Among the available data, those collected with are the best suited to search for X-ray emission from because of the superb angular resolution."
 Five observations were carried out between 2000 and 2006 (see Table 1)) with the source positioned on the back-illuminated ACTS-83 chip., Five observations were carried out between 2000 and 2006 (see Table \ref{tab:obs}) ) with the source positioned on the back-illuminated ACIS-S3 chip.
 We again reduced. the cata usingCLXO. applving aspect corrections and filtering the data to reject time intervals of ματς background as necessary: see and for detailed information about the datasets.," We again reduced the data using, applying aspect corrections and filtering the data to reject time intervals of flaring background as necessary; see and for detailed information about the datasets."
 Figure shows the image from. the combined data., Figure \ref{fig:image} shows the image from the combined data.
 Source detection was performed with the and routines in over numerous energy bancs on each individual dataset and on a cumulative image that was obtained from the 2006 with a total exposure of kks., Source detection was performed with the and routines in over numerous energy bands on each individual dataset and on a cumulative image that was obtained from the 2006 with a total exposure of ks.
" No X-ray source was found. near the rracio position (J2000: Ie""Μο4+004. 178)."," No X-ray source was found near the radio position (J2000: $18^{\rm h}46^{\rm m}15\fs49\pm0\fs04$, $-02\degr57\arcmin36\farcs0\pm1\farcs8$ )."
 As with wave used a generous search radius of 22resec. which combines the radio position uncertainty. the pposition uncertainty and the ppoint-spread. function.," As with we used a generous search radius of $2\,$ arcsec, which combines the radio position uncertainty, the position uncertainty and the point-spread function."
 Based on the S events within this circle. (hackeround expectation is 7). this upper limit is 6.810.7countsks tin the 0.3.8 keV energy band. at 3 confidence level (again based on 1986)).," Based on the 8 events within this circle (background expectation is 7), this upper limit is $\expnt{6.8}{-2}\,{\rm counts\,ks}^{-1}$ in the 0.3–8 keV energy band, at $\sigma$ confidence level (again based on )."
 To interpret the count-rate limits derived in the previous sections. we need a model for the X-ray spectrum. of the RRALs. as well as the distance and interstellar absorption along each line of sight.," To interpret the count-rate limits derived in the previous sections, we need a model for the X-ray spectrum of the RRATs, as well as the distance and interstellar absorption along each line of sight."
 For wave take the distance of 3.4454 kkpe based on the observed dispersion measure and the electron density mocdel., For we take the distance of $d_{3.4}$ kpc based on the observed dispersion measure and the electron density model.
 Using the 3-climensional Cialactic extinction model of(2003).. we find an optical extinction. of chy& 3mag.," Using the 3-dimensional Galactic extinction model of, we find an optical extinction of $A_V\approx 3\,$ mag."
 This translates to a hyverogen column density of Ny=5.I0em7 (based on 1995)).," This translates to a hydrogen column density of $\NH=5\times10^{21}\,\percmsq$ (based on )."
" ‘Vhis is similar to the value of 9]Ocm? based on the observed. dispersion measure of 203pcem "" and assuming 10 per cent ionisation (similar to the ΔΗΛ ratio of 1458)). or the same value from the data-set (likely to be an upper limit. since it is integrated through the Galaxy)."," This is similar to the value of $\expnt{9}{21}\,\percmsq$ based on the observed dispersion measure of $293\,{\rm pc\,cm}^{-3}$ and assuming 10 per cent ionisation (similar to the $\NH/DM$ ratio of ), or the same value from the data-set (likely to be an upper limit, since it is integrated through the Galaxy)."
 Overall. values of (5.10).107em7? seem likely. but the uncertainty is large given the poorly known distance.," Overall, values of $\expnt{(5-10)}{21}\,\percmsq$ seem likely, but the uncertainty is large given the poorly known distance."
 For0257.. we use a distance of 5.2d*2 kpe estimated from the dispersion measure.," For, we use a distance of $5.2d_{5.2}\,$ kpc estimated from the dispersion measure."
 The model gives an extinction of elyz Image which implies Ng=L6107em7.," The model gives an extinction of $A_V\approx 9\,$ mag which implies $\NH=1.6\times10^{22}\,\percmsq$."
 We compare this with the DAl-derived value of 0.7.1077em.7 or the E-cerived value of 2«107em7.," We compare this with the DM-derived value of $\expnt{0.7}{22}\,\percmsq$ or the -derived value of $\expnt{2}{22}\,\percmsq$."
 We also have the value of 41077em7 measured for PSI. 0258/Ixes 75. which is likely at a comparable or slightly larger distance 7.5kkpe: 2008)) as0257.," We also have the value of $\expnt{4}{22}\,\percmsq$ measured for PSR $-$ 0258/Kes 75, which is likely at a comparable or slightly larger distance kpc; ) as."
. This gives us a wide range to consider. (1.4)«1077emE7. again with the understanding that the large uncertainty on the distance eives an additional contribution.," This gives us a wide range to consider, $\expnt{(1-4)}{22}\,\percmsq$, again with the understanding that the large uncertainty on the distance gives an additional contribution."
 As an aside. we consider the possibility that PSH 0258 and wwere members of a binary system that was. disrupted bv a supernova.," As an aside, we consider the possibility that PSR $-$ 0258 and were members of a binary system that was disrupted by a supernova."
 At Skkpe their transverse separation. of 25aremin is ppc.," At kpc their transverse separation of $2.5\,$ arcmin is pc."
 In this scenario. the explosion that produced 00.4Myr ago cjected the progenitor of PSR 0258. which would have been moving at a transverse velocity of ~9kms4 before the second. supernova explosion that created the Wes 75 remnant.," In this scenario, the explosion that produced Myr ago ejected the progenitor of PSR $-$ 0258, which would have been moving at a transverse velocity of $\sim 9\,{\rm
km\,s}^{-1}$ before the second supernova explosion that created the Kes 75 remnant."
 Phis. velocity is actually. very low for an object ejected out of à. binary. which tvpicalle move at several hundred. | 2004).. although. we do not know the racial velocity of either object.," This velocity is actually very low for an object ejected out of a binary, which typically move at several hundred $^{-1}$ , although we do not know the radial velocity of either object."
 Waiting ~0.5 Myr for the second supernova seems. plausible. at. least. based: on some studies of binary evolution1995).," Waiting $\sim 0.5\,$ Myr for the second supernova seems plausible, at least based on some studies of binary evolution."
. IC PSR 0258 were actually the second supernova of the svstem it could have interesting implications. as the progenitors of magnetars are typically thought of as quite massive and PSR 0258 shares some (but not all) characteristies of magnetars2008).. but being the second. supernova would suggest it is less massive than its companion.," If PSR $-$ 0258 were actually the second supernova of the system it could have interesting implications, as the progenitors of magnetars are typically thought of as quite massive and PSR $-$ 0258 shares some (but not all) characteristics of magnetars, but being the second supernova would suggest it is less massive than its companion."
 Mass transfer or interactions during binary evolution could likely resolve the discrepancy., Mass transfer or interactions during binary evolution could likely resolve the discrepancy.
 We determine luminosity limits in two ways., We determine luminosity limits in two ways.
 First. we assume that the entire surface. of the RRATS cmits as a blackbody. and we use our count-rate limits to set. upper limits to the surface temperature. (measured. at infinity. like all blackhody parameters discussed. below).," First, we assume that the entire surface of the RRATs emits as a blackbody, and we use our count-rate limits to set upper limits to the surface temperature (measured at infinity, like all blackbody parameters discussed below)."
 Blackbodvemission is not. necessarily the best assumption as many other atmosphere models are more realistic (e.g... found hydrogen. atmospheres gave acceptable its for 1458)). but we keep it for simplicity.," Blackbodyemission is not necessarily the best assumption as many other atmosphere models are more realistic (e.g., found hydrogen atmospheres gave acceptable fits for ), but we keep it for simplicity."
 In Figure 2 we give temperature limits both for a canonical kkm neutron star racius and a kknm radius more similar o that inferred lor06.. covering the wide range in possible for cach source. ," In Figure \ref{fig:thermal} we give temperature limits both for a canonical km neutron star radius and a km radius more similar to that inferred for, covering the wide range in possible for each source. ["
The larger radius is not. likely yhvsical. and. may rellect a combination of errors in the distance or unrealistic models. for. the surface. emission.,"The larger radius is not likely physical, and may reflect a combination of errors in the distance or unrealistic models for the surface emission."
 Llowever. more realistic models (ee. model atmospheres) vpically have areas than blackbodies.]," However, more realistic models (.e.g, model atmospheres) typically have areas than blackbodies.]"
 As an alternate method. since assuming a constant emission area is not always consistent with observations of strongly pulsed. neutron. stars. we take the thermal spectrum measured. for1458: a blackbody with AY= I40eV. and unabsorbed luminosity 410eres+ 2007)..," As an alternate method, since assuming a constant emission area is not always consistent with observations of strongly pulsed neutron stars, we take the thermal spectrum measured for: a blackbody with $kT=140\,$ eV and unabsorbed luminosity $\expnt{4}{33}\,\ergsec$ ."
 Again considering the wide range in possible for each source. we plot the limits on the unabsorbed thermal luminosity in Figure 2— for three rlackbody temperatures: the nominal value. 1400eV. and +50 per cent of that value.," Again considering the wide range in possible for each source, we plot the limits on the unabsorbed thermal luminosity in Figure \ref{fig:thermal} for three blackbody temperatures: the nominal value eV, and $\pm50$ per cent of that value."
 The main goal of the range in emperature is to show the elfects of interstellar absorption and the limited energy window of oon the luminosity limits. although we physically motivate hem in Section 3..," The main goal of the range in temperature is to show the effects of interstellar absorption and the limited energy window of on the luminosity limits, although we physically motivate them in Section \ref{sec:disc}. ."
 Assuming blackbocly emission from a kkm radius and, Assuming blackbody emission from a km radius and
strip.,strip.
 All the temperature indicators adopted in this paper agree very well one each other and indicate a Z;4r 27500 Ex. i.c. a bit too hot fora pure > Dor variable.," All the temperature indicators adopted in this paper agree very well one each other and indicate a $T_{\rm eff}>$ 7500 K, i.e. a bit too hot for a pure $\gamma$ Dor variable."
 We conclude that the variability classification of this star is uncertain since it could be a 5 Dor ὁ Set Livbrict.180177):, We conclude that the variability classification of this star is uncertain since it could be a $\gamma$ Dor – $\delta$ Sct Hybrid.:
 Vhis star is the hottest star in the group. and Lies outside the instability strip (see Fie 4)).," This star is the hottest star in the group, and lies outside the instability strip (see Fig \ref{fig3}) )."
 Phe periodogram shows significant power at. very low frequeney., The periodogram shows significant power at very low frequency.
 The light curve of 0085837710 is presented in reffig:Ixep.., The light curve of 08583770 is presented in \\ref{fig:Kep}.
 No specific period can be deduced. but it is clear that there is something peculiar about this star. which needs further investigation.(HDISSLLG):," No specific period can be deduced, but it is clear that there is something peculiar about this star, which needs further investigation.:"
 Vhe frequenev. spectrum of 009775454. shows one dominant. peak in the ὁ SSct region. and several - seemingly eequidistant - peaks at lower frequencies.," The frequency spectrum of 09775454 shows one dominant peak in the $\delta$ Sct region, and several - seemingly equidistant - peaks at lower frequencies."
 Further investigation of theAepler light curves will clarify if this star is a dSSet star with rotational modulation cllects.234999):, Further investigation of the light curves will clarify if this star is a $\delta$ Sct star with rotational modulation effects.:
 ‘The lisht curves. of 111973705 LL97S7I05show a periodic long-term. behaviour. with Pox tdd (Vig. 6)).," The light curves of 11973705 show a periodic long-term behaviour, with $P \approx 4$ d (Fig. \ref{fig:Kep}) )."
 Phe spectral type BO. recorded in the Lenry Draper catalogue. is a full spectral class too early compared to our classification of 9.5V. It is most likely a 8 ος star in a binary svstem.," The spectral type B9, recorded in the Henry Draper catalogue, is a full spectral class too early compared to our classification of A9.5V. It is most likely a $\delta$ Sct star in a binary system."
 We presented a spectroscopic analysis of 19. candidate OSSet variables observed byAWepler both in. long and short. cadence mode., We presented a spectroscopic analysis of 19 candidate $\delta$ Sct variables observed by both in long and short cadence mode.
 The analysis is based on mecdium- to high-resolution spectra obtained at the. Loiano and OOACT observatories., The analysis is based on medium- to high-resolution spectra obtained at the Loiano and OACT observatories.
 For cach star we derived Tour. logg and esin; by matching the observed: spectra with synthetic spectra computed from theSYNTH code (Ixurucz&Averett1981). and using the L'TIZ atmospheric models calculated by ATLAS9 (Ixurucz1993a).," For each star we derived $T_{\rm eff}$, $\log g$ and $v\sin i$ by matching the observed spectra with synthetic spectra computed from the code \citep{kur81} and using the LTE atmospheric models calculated by ATLAS9 \citep{kur93}."
. The typical errors are about lxIx.. ddex. and 10 | for Tir. ogg. and esin/. respectively.," The typical errors are about K, dex, and 10 $^{-1}$ for $T_{\rm eff}$, $\log g$, and $v 
\sin i$, respectively."
 Equivalent spectral types and uninositv classes were also derived., Equivalent spectral types and luminosity classes were also derived.
 “Phe luminosities of he stars were obtained using the tables of Sehmiclt-haler (1982)., The luminosities of the stars were obtained using the tables of \citet{schmidt}.
 For ten stars we used Strommeren photometry fron he literature to estimate the reddening., For ten stars we used Strömmgren photometry from the literature to estimate the reddening.
 For seven stars or which 3 photometry was also available. Zi and logg could be obtained for comparison with our spectroscopic values.," For seven stars for which $\beta$ photometry was also available, $T_{\rm eff}$ and $\log g$ could be obtained for comparison with our spectroscopic values."
 In addition. Vdv colours. the IETUM. method and values listed in the KIC were used to obtain independent estimates of Zar.," In addition, $V-K$ colours, the IFRM method and values listed in the KIC were used to obtain independent estimates of $T_{\rm eff}$."
 We find a general good agreement between photometric and. spectroscopic. results., We find a general good agreement between photometric and spectroscopic results.
 Four stars. with significant parallaxes ancl four cluster member objects were used. to check our estimate of the luminosities., Four stars with significant parallaxes and four cluster member objects were used to check our estimate of the luminosities.
 We obtain consistent results for. all the stars. with the exception of 003429637. 1178875). which is a wide binary and mav have an erroneus parallax determination.," We obtain consistent results for all the stars, with the exception of 03429637 178875), which is a wide binary and may have an erroneus parallax determination."
 Moreover. for 005724440 1187234) we suspect a problem with the ucby: photometry. since Zi derived from the WoA index is in agreement with our estimate within the errors.," Moreover, for 05724440 187234) we suspect a problem with the $uvby\beta$ photometry, since $T_{\rm eff}$ derived from the $V-K$ index is in agreement with our estimate within the errors."
 Finally. we present the periodograms for the 19 investigated stars. based on the Wepler satellite photomoetry.," Finally, we present the periodograms for the 19 investigated stars, based on the satellite photometry."
 These beautiful cata allowed: us to classify the type of variability of cach star. including 005296877. which is a high5 amplitude ellipsoidal variable candidate.," These beautiful data allowed us to classify the type of variability of each star, including 05296877, which is a high amplitude ellipsoidal variable candidate."
 As a result. we find six pure à Sct. 3 1pure > Dor and nine hybrid. 1pulsators.," As a result, we find six pure $\delta$ Sct, 3 pure $\gamma$ Dor and nine hybrid pulsators."
 This classification is consistent with the derived: physical parameters and their position in the LRdiagram., This classification is consistent with the derived physical parameters and their position in the HRdiagram.
 As already noted by Cirigahceneetal. (2010).we were surprised by the," As already noted by \citet{griga10}, ,we were surprised by the"
eclipse. we do not expect to be able to fit the light curve of with high accuracy.,"eclipse, we do not expect to be able to fit the light curve of with high accuracy."
 Nevertheless. in the absence of longer-teriii photometric monitoring. the data of. Douauos(2007)- allow reasonable constraints. to be placed on the orbital inclination.," Nevertheless, in the absence of longer-term photometric monitoring, the data of \cite{bonanos} allow reasonable constraints to be placed on the orbital inclination."
 The was used to model the light. curve of., The was used to model the light curve of.
"W13.. The: effective""+ temperature ofB the emission. line. object. was fixeda - ∐∐∐⊢≼↧⋜∐⋅↨↘↽↸∖↕∏∐≼↴⋁⋜∏↖↽⋜↧∐≼↧↸⊳↕↥⋅↸⊳∏↕⋜∐⋅∪↥⋅↴⋝↕↑↴∖↴↖↖⇁↸∖↥⋅↸∖⋜↕↴∖↴↴∖↴∏⋯↸∖≼↧∙- "" - - EN"," The effective temperature of the emission line object was fixed at 25kK, appropriate for its spectral type, and a linear limb-darkening law and circular orbits were assumed."
 The mass ratio“ was derived∙ from. the RV.. curve. aud the inclination.∙∙∙ Roc;1¢ lobe filliug∙∙ factors. for. both objects. aud temperature of the optical secondary were allowed to vary.," The mass ratio was derived from the RV curve, and the inclination, Roche lobe filling factors for both objects and temperature of the optical secondary were allowed to vary."
 The code rapidly converges to a near-coutact configuration in which the enuüsson-ine object has almost filled its Roche lobe (filling factor ~0.93+0.05) and the other star ias a somewhat lower filling factor (0.7 150.1)., The code rapidly converges to a near-contact configuration in which the emission-line object has almost filled its Roche lobe (filling factor $\sim$ $\pm$ 0.05) and the other star has a somewhat lower filling factor $\sim$ $\pm$ 0.1).
 The best-fit. model has an iuclination. ¢.=62°n and providesi a close natch to the light curve from o~0.5 to o~l. although he region around o=0.2 0.1 is less well reproduced.," The best-fit model has an inclination $i=62^\circ$ and provides a close match to the light curve from $\phi$$\sim$ 0.5 to $\phi$$\sim$ 1, although the region around $\phi=0.2$ $0.4$ is less well reproduced."
 Although the code supports additional features such as hot spots that may provide a better fit to this yortion of the light curve (sec. for example. the model of presented by Linderetal. 20093). we consider ≯↿∐⋅↕∐∖↥⋅↥⋅↸∖∱∎⊔↸∖⋯↸∖∐↑∪↕⋟↑↕∐∖↕⊔∪≼∐∖↕↕∐⋜↧⋯∐⋅∪↻↥⋅↕⋜↧↑↸∖∶↴∙⊾↕↖⇁↸∖∐↑∐↸∖ ↕⋯↕↑⋜↧↑↕∪∐↴∖↴∪↕≯↑∐↸∖↻∐∪↑∪∐∐," Although the code supports additional features such as `hot spots' that may provide a better fit to this portion of the light curve (see, for example, the model of presented by \citealt{linder}) ), we consider further refinement of the model inappropriate given the limitations of the photometric dataset used."
∖⊓⋅↕↸⊳≼↧⋜↧↑⋜↧↴∖↴↸∖↑∏↴∖↴↸∖≼↧∙↕⋟⋜∐⋅⋜∐⊔↸∖↑↸∖↥⋅↴∖↴ ⋜↧↕↑↕∪∏∶↴⋁∐↖↖↽↸∖↴∖↴⊓⋅↸∖↴∖∷∖↴↑∐⋜↧↑≺∏∐⋅↻↥⋅∐⊔⋜∐⋅⋅↖⇁∶↴⋁∪⋜↧↕↕↴∖↴↑∪↸⊳∪∐↴∖↴⊓⋅⋜⊔∐ he inclination⋅⋅⋅ of. the svstem aud other parameters should )o regarded as provisional⋅⋅ pending. acquisition⋅⋅⋅ of. louecr-ena photometry.," Parameters derived from the light curve model are listed in Table \ref{tab:results}, although we stress that our primary goal is to constrain the inclination of the system and other parameters should be regarded as provisional pending acquisition of longer-term photometry."
" Aum ↴∪↸∖↼↽⋪⋅↸∖↸∖⋅⋅∪⋅ he loris TLlinati ↸∖∩∖⋅∎↽↸∖≺⋅⊳⋅⋪⋅∪↖↽↸∖iy investigated.To models in orswhichm / is fixed at values from 55 o 687ANS Oxvelualewhile AllingAlline factorsfac uo:anc FomaperaB*uuresCOR areATO allowed:val ο vary as before: the best fit /=62° curve aud models with ἐξ 59°.)=65"" aud ;=68"" are plotted in Figure LL. Tuclinat"," To examine errors in the derived inclination we investigated models in which $i$ is fixed at values from $55^\circ$ to $68^\circ$ while filling factors and temperatures are allowed to vary as before; the best fit $i=62^\circ$ curve and models with $i=59^\circ$, $i=65^\circ$ and $i=68^\circ$ are plotted in Figure \ref{fig:phot}."
"nclinations ercaterereatcr tlthan our best-fitt-fit model] lead to shehtlysielth {wed«ων Shine factorsfac uoaud lohoer fenrperafὰBpures,Qs v“hilehile theὰ converse is fillingtrue for lower hisinchuatious,"," Inclinations greater than our best-fit model lead to slightly lower filling factors and higher temperatures, while the converse is true for lower inclinations."
"i The depth of the eclipseClipse provides the strongeststrong constraintconstiüut oon theο model withvith inclinationsiucliuatious greater παtli 765 65"" stronglySugly disfavouredcistavourec."," The depth of the eclipse provides the strongest constraint on the model, with inclinations greater than $\sim$ $^\circ$ strongly disfavoured."
. I -bxux; pliototomietrywdvv isOus also presentuntepreseutedx bv~ DRon:οuutosP9007(2007‘)).. but audcousiderablv greater scatter is present in the data at iud-eclipse. making it less suitable for modelling.," $I$ -band photometry is also presented by \cite{bonanos}, but considerably greater scatter is present in the data at mid-eclipse, making it less suitable for modelling."
" Nevertheless.: these data also support /:<Γ65"". favouring.' a ""lue although 62"". the;oue dege degreeFeeatter of rendersscatter reudersHus 4444this uncertain."," Nevertheless, these data also support $i\le65^\circ$, favouring a value $\sim$ $^\circ$ , although the degree of scatter renders this uncertain."
" Therefore. taking τς657 places robust lower of Myc21.14 2.6M ind Maji,232.84 LOX. for linüts enission-ine oject andaud its companion.conrpauion. rising,rng tto 232!thesu und 35.1!0 GAL... for. our preferred. inclination. j=M627]."," Therefore, taking $i\le65^\circ$ places robust lower limits of $M_{\text{em}}\ge21.4\pm2.6 $ $_\odot$ and $M_{\text{abs}}\ge32.8\pm4.0 $ $_\odot$ for the emission-line object and its companion, rising to $23.2^{+3.3}_{-3.0}$ $_\odot$ and $35.4^{+5.0}_{-4.6}$ $_\odot$ for our preferred inclination $i=62^{+3}_{-4}$$^\circ$."
" For the purposes of discussion we take the lower""OY massesHRSÓMS derivedvo fromn the iow;/65° GMSlinüt.", For the purposes of discussion we take the lower masses derived from the $i\le65^\circ$ limit.
Jays The↴ short orbital⋅ period.⋅ mear-coutact configuration⋅ and evolved. mass-cepleted nature of. the enmuüssion-liue⋅⋅. object. all iuplv. that the two components of. the ↴∖↴⋅↖⇁↴∖↴↑↸∖⋯⋯∏↴∖↴↑∐⋜↧↖⇁↸∖∏∐≼∐∖↥⋅∶↴∙⊾∪∐↸∖↴∖↴∏⋅∪∐∶↴∙⊾↕∐↑↸∖↥⋅⋜↕↸⊳↑↕∪∐≼⊔∐⋅∐↓∶↴∙⊾ ⋅ ⋅ their⋅ evolution.," The short orbital period, near-contact configuration and evolved, mass-depleted nature of the emission-line object all imply that the two components of the system must have undergone strong interaction during their evolution."
"⋅ The↴ 9.3-dayv∙ orbital⋅ period⋅ suggestsus a late-Caseate-C'ase ὁA or Caseas B scenario,:ario. with ass.as trausferal besiunins: near the ouset of: shell lydrogen burniug1 (Petrovieotal.2005.. hereafter. POS}."," The 9.3-day orbital period suggests a late-Case A or Case B scenario, with mass transfer beginning near the onset of shell hydrogen burning \citealt{petrovic}, hereafter P05)."
 The presence. of unevolved late-O stars 30M.) in Wal suggests a mini initial mass —35ML.. for shell buruing to have conuuenced. with: in: excess of ~LOAL.. lost ounce mass transfer beeius.," The presence of unevolved late-O stars ) in Wd1 suggests a minimum initial mass $\sim$ $_\odot$ for shell burning to have commenced, with in excess of $\sim$ $_\odot$ lost once mass transfer begins."
 However. trausfer of augular moment model:i. expectedto lead to the accretor rapidly: reaching ∙⋡∙∙critical rotation: (Packet↽∖↽1981:Lauger.etal.:∖2008)... while: rapid: rotation: will: also ereatly. increase: wind: nass loss rates (Langer1998).. and a fullv-conservative transfer scenario appears unlikely.," However, transfer of angular momentum is expected to lead to the accretor rapidly reaching critical rotation \citep{packet,
 langer}, while rapid rotation will also greatly increase wind mass loss rates \citep{langer98}, and a fully-conservative transfer scenario appears unlikely."
 Dudeced. models of short-period WR|O binaries by POS sugsests that imass-trauster is highly dn ∙∙such scenarios. with: oulv: 10 of ↸⋞↴↕∏∿ transferred mass being retained bv companion star.," Indeed, models of short-period WR+O binaries by P05 suggests that mass-transfer is highly -conservative in such scenarios, with only $\sim$ of transferred mass being retained by companion star."
 The current 2UAL.. B3\L. (1uniniuni) mass ratio in is consistent: with: the POS mo of lnte-Case A/Case D evolution at low accretion efficieucy. with higher accretion --- leading to a more unequal lass ratiothan we obscrve.," The current $_\odot$ $_\odot$ (minimum) mass ratio in is consistent with the P05 model of late-Case A/Case B evolution at low accretion efficiency, with higher accretion efficiencies leading to a more unequal mass ratiothan we observe."
 POS estimate a relationship between initial AIS tnass and final WR μας for Case D svstclus of, P05 estimate a relationship between initial MS mass and final WR mass for Case B systems of
traces In P.,traces in $P$.
 The secoud cause for canals is depth depolarization., The second cause for canals is depth depolarization.
" If the oenuitting aud rotating mediun is nuiform. P=sn(2RAPASMASRAIAS) so that ""nulls of jiinuinmunu polarization exist at positious where 2RALA= υπ."," If the emitting and rotating medium is uniform, $P =
\sin(2 RM \lambda^2) / (2 RM \lambda^2)$ so that “nulls” of minimum polarization exist at positions where $2 RM \lambda^2 = n
\pi$ ."
 However. there is no reason why depolarization canals would be one beam wide in this model. aud low-P canals onlv exist for à uniformi medium.," However, there is no reason why depolarization canals would be one beam wide in this model, and $P$ canals only exist for a uniform medium."
 Furthermore. the second model predicts the position of the canals to shift with wavelength. which is not observed.," Furthermore, the second model predicts the position of the canals to shift with wavelength, which is not observed."
 Therefore we believe that the depolkwization canals are predonmünantlv the result of scant depolarization., Therefore we believe that the depolarization canals are predominantly the result of beam depolarization.
 For au extended discussion of the causes of depolarization canals. see IHaverkorn et (9000. 20024).," For an extended discussion of the causes of depolarization canals, see Haverkorn et (2000, 2003a)."
 If the eracient iu RAL is constant. depolarization canals due to bean depolarization would not be possible.," If the gradient in $RM$ is constant, depolarization canals due to beam depolarization would not be possible."
 So the RAL must increase discoutiunnouslv. as is visible in Fie. 12..," So the $RM$ must increase discontinuously, as is visible in Fig. \ref{f5:fil}."
 I£ ARAL over one beam is between 2.1 aud a2.5 (for between 6.3⋅∙ aud —7.5Γ7.. ete.).," If $\Delta RM$ over one beam is between 2.1 and 2.5 (or between 6.3 and 7.5, etc.),"
 depolarizationB. canals occur., depolarization canals occur.
 Fie., Fig.
 14. shows P and RAL over a narrow slit (— 1.57)) across the canals plotted agaist b. located within the ος In Fie. 12..," \ref{f5:fil_n} shows $P$ and $RM$ over a narrow slit $\sim$ ) across the canals plotted against $b$ , located within the box in Fig. \ref{f5:fil}."
 The top panel shows plots for P at al 5 frequencies. oversauupled by a factor of 5.," The top panel shows plots for $P$ at all 5 frequencies, oversampled by a factor of 5."
" The ceutral pane shows RAL. where RAL values are only plotted for. bemus forJ whichH \j,,D<a2 aud (P;/720 1iJvοσα,"," The central panel shows $RM$, where $RM$ values are only plotted for beams for which $\chi_{red}^2 < 2$ and $\left<P\right> > 20$ mJy/beam."
 Polarization augle o is given in the bottom panel., Polarization angle $\phi$ is given in the bottom panel.
 The two sharpest depolarization canals visible at all frequencics. at pixel uunubers 113 aud 115 correspond to abrupt changes in. RAL of. about .2 aand angle chauges around 90°... Abrupt RAL changes with other magnitudes are present as well. but these do not create depolarization canals.," The two sharpest depolarization canals visible at all frequencies, at pixel numbers 113 and 145 correspond to abrupt changes in $RM$ of about 2 and angle changes around Abrupt $RM$ changes with other magnitudes are present as well, but these do not create depolarization canals."
 The ring-like structure in polarized intensity P with a radius of about sshows a regular increase in polarization angle frou its center out to  suggesting tha he structure is a disk mstead of a ring. aud extends to larger radii than the ving in P.," The ring-like structure in polarized intensity $P$ with a radius of about shows a regular increase in polarization angle from its center out to $\sim$, suggesting that the structure is a disk instead of a ring, and extends to larger radii than the ring in $P$."
 The rotation measure is shehtly positive outside the rine. reverses sign inside the re arc decreases alinost coutiuuouslv to RALxNradaiu* aat the ceuter.," The rotation measure is slightly positive outside the ring, reverses sign inside the ring and decreases almost continuously to $RM~\approx -8$ at the center."
 This property rules out its production by any spherically svuumetric structure. such as a supernova remnant or a wind-blown bubble.," This property rules out its production by any spherically symmetric structure, such as a supernova remnant or a wind-blown bubble."
 The RAL structure. and the observation that the coherence in angle slowly disappears beyond a radius of ~ iudicates a mnienetic origiu for the polarization ving. probably accompanied bv anu electrou density enhancement.," The $RM$ structure, and the observation that the coherence in angle slowly disappears beyond a radius of $\sim$ indicates a magnetic origin for the polarization ring, probably accompanied by an electron density enhancement."
 We propose that the rime is produced by a predominantly magnetic Πο)κο structure. iu which the parallel magnetic field streueth is maximal at the center aud directed away from us.," We propose that the ring is produced by a predominantly magnetic funnel-like structure, in which the parallel magnetic field strength is maximal at the center and directed away from us."
 The enhancement in P at radius ~ lis caused by a lack of depolarization duc to the relative constancy of RAL at that radius., The enhancement in $P$ at radius $\sim$ is caused by a lack of depolarization due to the relative constancy of $RM$ at that radius.
 A filamentary pattern of parallel. narrow depolarization canals indicates structure inRAL which is aligned with Calactic longitude.," A filamentary pattern of parallel, narrow depolarization canals indicates structure in$RM$ which is aligned with Galactic longitude."
The depolarization canals are probably created by beau depolarization due to abrupt spatia eradicuts in RAL.,The depolarization canals are probably created by beam depolarization due to abrupt spatial gradients in $RM$ .
"contributes (to7. assuming Ie Tat z>:3 andan additional 7,mοσο0.002 for He III at z<3).","contributes to$\tau_e$, assuming He II at $z > 3$ andan additional $\tau_e \approx 0.002$ for He III at $z \leq 3$ )."
" second. we investigate (he effects ofIGM partia ionization al 2>z,."," Second, we investigate the effects ofIGM partial ionization at $z > z_r$."
 Finally. in computing the contribution of residual electrons at high recshiis. we adopt the concordance parameters from the WMAD-3 data set.," Finally, in computing the contribution of residual electrons at high redshifts, we adopt the concordance parameters from the WMAP-3 data set."
" The CMD optical depth is formally only a 36 result. which may change. as WAIAP producesbetter determinations of tje matter density. O,,. and the parameters. o and ης. (hat govern small-scale power."," The CMB optical depth is formally only a $3 \sigma$ result, which may change, as WMAP producesbetter determinations of the matter density, $\Omega_m$, and the parameters, $\sigma_8$ and $n_s$, that govern small-scale power."
 Both ox and n have well-known degeneracies with 7. in CAMB parameter extraction (Spergel 2006]., Both $\sigma_8$ and $n$ have well-known degeneracies with $\tau_e$ in CMB parameter extraction (Spergel 2006).
" Fherefore. (heir derived values may change in [future CAIB data analvses. especially as (he constraints on 7, continue to evolve."," Therefore, their derived values may change in future CMB data analyses, especially as the constraints on $\tau_e$ continue to evolve."
 In addition. inaccuracies in (he incomplete recombination epoch and residual ionization history. οὐ(). add uncertainties to the CAIB radiative transler. the ciunping of (-modes. and the polarization signal used to derive an overall 7.We now discuss the contribution of residual electrons in the IGM following the recombination epoch al z21000.," In addition, inaccuracies in the incomplete recombination epoch and residual ionization history, $x_e(z)$, add uncertainties to the CMB radiative transfer, the damping of $\ell$ -modes, and the polarization signal used to derive an overall $\tau_e$.We now discuss the contribution of residual electrons in the IGM following the recombination epoch at $z \approx 1000$."
 Scattering [rom (hese electrons is significant and is normally accounted Lor in CMD transport codes such as CAIBFAST (Seljak Zaldarriaga 1996) through the recombination IGM ionization history. οἴ(5).," Scattering from these electrons is significant and is normally accounted for in CMB transport codes such as CMBFAST (Seljak Zaldarriaga 1996) through the post-recombination IGM ionization history, $x_e(z)$."
 However. a number of past papers are vague on 10w the ionization history is treated. which has led to confusion in how much residual optical depth and power-damping has been subtracted from the CMD signal.," However, a number of past papers are vague on how the ionization history is treated, which has led to confusion in how much residual optical depth and power-damping has been subtracted from the CMB signal."
 Modern calculations of row the IGM. became neutral have been done by Seager (2000). although their code (RECEAST) continues to be modified to deal with subtle effects of the recombination epoch and the atomic physics of hydrogen (25— 15) two-photon transitions (W. Wong D. Scott. private communication).," Modern calculations of how the IGM became neutral have been done by Seager (2000), although their code (RECFAST) continues to be modified to deal with subtle effects of the recombination epoch and the atomic physics of hydrogen $2s \rightarrow 1s$ ) two-photon transitions (W. Wong D. Scott, private communication)."
" To illustrate the potential effects of hieh-z residual electrons. we have used numbers [rom Figure 2 of Seager (2000). (he top-panel model. which assumed a cosmology wilh Qi=1. OQ,=0.05. h=0.5. Y=0.24. and Tesis=2.728 Ik. At low redshifts. zmοσοp. just. before reionizalion. they find a residual electron [raction c,2210. 57."," To illustrate the potential effects of $z$ residual electrons, we have used numbers from Figure 2 of Seager (2000), the top-panel model, which assumed a cosmology with $\Omega_{\rm tot} = 1$, $\Omega_b = 0.05$, $h = 0.5$, $Y = 0.24$, and $T_{\rm CMB} = 2.728$ K. At low redshifts, $z \approx z_r$, just before reionization, they find a residual electron fraction $x_{e,0} \approx 10^{-3.3}$ ."
" We fitted their curve Lor log zr,out to z500 to the formula: where a52.303x10"".", We fitted their curve for log $x_e$out to $z \approx 500$ to the formula: where $\alpha \approx 2.303 \times 10^{-3}$.
" More recent recombination calculations (W. Y. \Wong D. Scott. private communication) using WALAP-32 parameters (Q,= 0.04. h=0.73.) 3. €O4,= 0.24. Q4= 0.76. Y= 0.244) find somewhat lower values. ir.9210.* with azm2.12x 10.7. We attribute the lower 7,9 lo the faster recombination rates arising from their higher assumedbarvon censity. Ομ= 0.0213. compared to Qj?=0.0125 in Seager (2000)."," More recent recombination calculations (W. Y. Wong D. Scott, private communication) using WMAP-3 parameters $\Omega_b = 0.04$ , $h = 0.73$ , $\Omega_m = 0.24$ , $\Omega_{\Lambda} = 0.76$ , $Y = 0.244$ ) find somewhat lower values, $x_{e,0} \approx 10^{-3.67}$ with $\alpha \approx 2.12 \times 10^{-3}$ We attribute the lower $x_{e,0}$ to the faster recombination rates arising from their higher assumedbaryon density, $\Omega_b h^2 = 0.0213$ , compared to $\Omega_b h^2 = 0.0125$ in Seager (2000)."
"where we have assumed a constant circular speed. so v.dR,.","where we have assumed a constant circular speed, so $\d L_z=\vc\d\Rg$ ."
" Combining equations (6). (C7)) and (91). we have Our assumption ofconstant v, allows us to evaluate A(R). because then so the normalisation condition |=27dRRP reads where and So As here we are aiming at velocity distributions and not stellar densities at a certain position. we will henceforth use the normalised velocity distribution at a fixed radius R where N normalises the integral of 7 in v, to unity."," Combining equations \ref{eq:givesP}) ), \ref{eq:givesnv}) ) and \ref{eq:givesN}) ), we have Our assumption ofconstant $\vc$ allows us to evaluate $K(\Rg)$ , because then so the normalisation condition $1=2\pi\int\d R\,RP$ reads where and So As here we are aiming at velocity distributions and not stellar densities at a certain position, we will henceforth use the normalised velocity distribution at a fixed radius R where ${\mathcal N}$ normalises the integral of $n$ in $v_\phi$ to unity."
 Note that on the right side of eq. (16) , Note that on the right side of eq. \ref{eq:nvphi}) )
"the dependence on v, is carried by the instances of A.=Rr;/v, and by c(Ry).", the dependence on $v_\phi$ is carried by the instances of $\Rg=Rv_\phi/\vc$ and by $c(\Rg)$.
 Fromequation (6)) we expect A (which has the units of a frequency) to be constant when 7«v. and indeed from an asymptotic expansion of equation (12)) for large e we obtain etc)« rand K2vumΑΟ).," Fromequation \ref{eq:givesP}) ) we expect $K$ (which has the units of a frequency) to be constant when $\sigma\ll\vc$, and indeed from an asymptotic expansion of equation \ref{eq:normc}) ) for large c we obtain $g(c)\propto
c^{-1/2}\propto \sigma$ and $K = \vc / (2\pi^{3/2} R_g)$."
 shows that e(c) satisties this expectation throughout the entire parameter range of interest., shows that $g(c)$ satisfies this expectation throughout the entire parameter range of interest.
 For (c> 25). the direct computation of etc) becomes impractical so apart from the advantage of a numerically less costly formula a reasonable approximation must be found.," For $c > 25$ ), the direct computation of $g(c)$ becomes impractical so apart from the advantage of a numerically less costly formula a reasonable approximation must be found."
 The dashed green line in demonstrates that forc>2 this is achieved to high precision For c we adopt the radial dependence Above we have restricted ourselves to stars with E.=0., The dashed green line in demonstrates that for $c>2$ this is achieved to high precision For $\sigma$ we adopt the radial dependence Above we have restricted ourselves to stars with $E_z=0$.
 However. to the extent that the motion in A of a star is unattected by its motion perpendicular to the plane. the distribution we have derived will apply to the population formed by all stars that now ie in the solar cylinder (the region restricted in radius to but unrestricted in z).," However, to the extent that the motion in $R$ of a star is unaffected by its motion perpendicular to the plane, the distribution we have derived will apply to the population formed by all stars that now lie in the solar cylinder (the region restricted in radius to $R\simeq R_0$ but unrestricted in $z$ )."
" From our formulae we have that the shape of this velocity distribution is controlled by four parameters: he galactocentric radius of the measurement Ay. the scale-length Ry of the young disc. the local velocity dispersion c. and the scale-length A,. on which the velocity dispersion varies."," From our formulae we have that the shape of this velocity distribution is controlled by four parameters: the galactocentric radius of the measurement $R_0$, the scale-length $R_\d$ of the young disc, the local velocity dispersion $\sigma_0$, and the scale-length $R_\sigma$ on which the velocity dispersion varies."
 The first wo parameters are generally well-known and for fits of the solar neighbourhood can be set to Ry=Skpe and Ay=2.5kpe., The first two parameters are generally well-known and for fits of the solar neighbourhood can be set to $R_0=8\kpc$ and $R_\d=2.5\kpc$.
" The value of &,. is less clear and will be discussed below.", The value of $R_\sigma$ is less clear and will be discussed below.
" The dependence of ntv,[R4) on o and &,. is shown in2.", The dependence of $n(v_\phi|R_0)$ on $\sigma_0$ and $R_{\sigma}$ is shown in.
. In the upper panel we hold the dispersion scale-length constant at 7.5kpe and show the velocity distributions for local dispersion values of c=20.25.3040.50kms'.," In the upper panel we hold the dispersion scale-length constant at $7.5\kpc$ and show the velocity distributions for local dispersion values of $\sigma_0=20,25,30,40,50\kms$."
. As ce increases. the distribution becomes wider and the low-velocity tail rises much faster than does the high-velocity tail. corresponding to an increasing asymmetric drift.," As $\sigma_0$ increases, the distribution becomes wider and the low-velocity tail rises much faster than does the high-velocity tail, corresponding to an increasing asymmetric drift."
 Simultaneously. the peak slowly shifts to lower velocities.," Simultaneously, the peak slowly shifts to lower velocities."
" The lower panel shows the velocity distributions for fixed συ=30kms !. butdifferent scale-lengths of the velocity dispersion. A,=5.A 10.7.5."," The lower panel shows the velocity distributions for fixed $\sigma_0=30\kms$ , butdifferent scale-lengths of the velocity dispersion, $R_\sigma=10,7.5,5,4\kpc$ ."
 Smaller values of A. imply higher dispersion in the inner kpe.regions of the Galaxy and lower dispersion in the outskirts., Smaller values of $R_\sigma$ imply higher dispersion in the inner regions of the Galaxy and lower dispersion in the outskirts.
" Thus for small &,- stars from the inner dise can more easily reach the Ry compared to their counterparts", Thus for small $R_\sigma$ stars from the inner disc can more easily reach the $R_0$ compared to their counterparts
The production of thiorine in AGB stars is of interest also in the light of the Galactic chemical evolution.,The production of fluorine in AGB stars is of interest also in the light of the Galactic chemical evolution.
 In and. we present vields for cecalculated [or the different model shown inl., In and we present yields for calculated for the different model shown in.
. Yields are a direct function of the amount of TDU., Yields are a direct function of the amount of TDU.
 They are calculated as net vields: Af=IMBY—NU)SLdl where 7 is the total lifetime of the star. «AZ/dl is (he mass-loss rate and X and Xy reler to the current aud initial mass fraction ofPF.," They are calculated as net yields: $M = \int_0^{\tau} (X - 
X_0) \frac{dM}{dt} dt$ where $\tau$ is the total lifetime of the star, $dM/dt$ is the mass-loss rate and $X$ and $X_0$ refer to the current and initial mass fraction of."
. The vield is positive if iis produced and negative if it is destroved., The yield is positive if is produced and negative if it is destroyed.
 The vvields are (vpically negative. decreasing from ~ 0 for stars of 1 (to ~—2x10? [or stars of GM.," The yields are typically negative, decreasing from $\sim$ 0 for stars of 1 to $\simeq - 2 \times 10^{-5}$ for stars of 6."
.. This means that this isotope is destroved in all the models. except those with Z=0.0001 and mass higher than 2.25 Al.," This means that this isotope is destroyed in all the models, except those with $Z$ =0.0001 and mass higher than 2.25 ."
.. The vvield reaches a positive maximum of 4xLO° for the highest mass model computed as this mnelallicily (5 )).," The yield reaches a positive maximum of $4 \times 
10^{-6}$ for the highest mass model computed as this metallicity (5 )."
 This is due to a combination of different factors: (i) the temperature al the base of the convective envelope is as high as 9.7 x10' Ix in this model. at. which temperature the δρ. ΙΟ reaction becomes as important as the δρα rreaction and ecan actually be produced by proton captures during hot bottom burning. (ii) (he abundance of iis extremely. high because of the operation of strong TDU and hot bottom burning. and (iii) the initial aabundance (Ny in the formula above) is verv low.," This is due to a combination of different factors: (i) the temperature at the base of the convective envelope is as high as 9.7 $\times 10^7$ K in this model, at which temperature the $p,\gamma)^{15}$ O reaction becomes as important as the $p,\alpha$ reaction and can actually be produced by proton captures during hot bottom burning, (ii) the abundance of is extremely high because of the operation of strong TDU and hot bottom burning, and (iii) the initial abundance $X_0$ in the formula above) is very low."
 The initial aabundance is also very small and hence the vields for this metallicity are less negative lor masses above 4. ccompared to more metal-rieh models of the same mass., The initial abundance is also very small and hence the yields for this metallicity are less negative for masses above 4 compared to more metal-rich models of the same mass.
 In we compare some selected model predictions with the observations by (1992)., In we compare some selected model predictions with the observations by \citet{jorissen:92}.
. The metallicitv. of the observed stars ranges [rom about Z—0.006 to about Z=0.04 with an average of 0.016., The metallicity of the observed stars ranges from about $Z$ =0.006 to about $Z$ =0.04 with an average of 0.016.
 Hence the 2.5 Z = 0.004modelhasametatlicil yloolowlobeconsideredtomatehlheabservationsandilisincludedinthe figure , Hence the 2.5 $Z$ =0.004 model has a metallicity too low to be considered to match the observations and it is included in the figure only to illustrate the trend of our results with metallicity.
Z=0.008 model. which has the highest final aabundance in (he intershell. does not represent a good match to the stellar data.," The 3 $Z$ =0.008 model, which has the highest final abundance in the intershell, does not represent a good match to the stellar data."
 This is because the final C/O abundance in this model is 5.6. while the stellar data have C/O up to about 1.5.," This is because the final C/O abundance in this model is 5.6, while the stellar data have C/O up to about 1.5."
 It follows that since the large aabundance in this model is a consequence of the large aabundance in (he envelope. we cannot take this model to explain the highest observed values. (," It follows that since the large abundance in this model is a consequence of the large abundance in the envelope, we cannot take this model to explain the highest observed values. ("
We note though that stars with the high C/O ratio and high aabundance produced by this model may in principle exist but. be obscured by their dusty envelopes).,We note though that stars with the high C/O ratio and high abundance produced by this model may in principle exist but be obscured by their dusty envelopes).
 It shouldalso be considered (hat the observational data regarding SC stus, It shouldalso be considered that the observational data regarding SC stars
stars at intermediate ages. we find the intermediate values for ML.,"stars at intermediate ages, we find the intermediate values for $M/L$."
 The interpretation of the dynamics under the proposed model thus accords with the natural increase iu the ML ratios of stellar populations due to the ageing of stars aud the consequent build up of black holes. neutron stars. and white dwarfs.," The interpretation of the dynamics under the proposed model thus accords with the natural increase in the $M/L$ ratios of stellar populations due to the ageing of stars and the consequent build up of black holes, neutron stars, and white dwarfs."
 This correspondence is natural oei anv modified eravity scheme where the stars alouc are responsible for the dviauics. bu has to be thought of as a fortuitous coincidence under the dark matter hypothesis.," This correspondence is natural in any modified gravity scheme where the stars alone are responsible for the dynamics, but has to be thought of as a fortuitous coincidence under the dark matter hypothesis."
 As a secoud consistency check. we now show the resulting projected surface deusity mass profile for a model of the Leo II eilaxy. normalised by the total huuinosity of hat svsteni.," As a second consistency check, we now show the resulting projected surface density mass profile for a model of the Leo II galaxy, normalised by the total luminosity of that system."
" We coustruc the surface brightuess profile using our model for an equilibriumu isothermal syste raving velocity dispersion aud projected A, (assumed equal to the observed A). equal to those observed. for Leo IL"," We construct the surface brightness profile using our model for an equilibrium isothermal system having velocity dispersion and projected $R_{\rm hm}$ (assumed equal to the observed $R_{\rm hl}$ ), equal to those observed for Leo II."
 This is giveu in Figure (2). where we have also jotted the observation for the actual surface density light xofile of Leo IL. from the star couuts analvsis of Coleman ct al. (," This is given in Figure (2), where we have also plotted the observation for the actual surface density light profile of Leo II, from the star counts analysis of Coleman et al. ("
2007). out to the radius where measurements fall rclow the backgrouud noise level. both normalised to the sale total Iunünositv.,"2007), out to the radius where measurements fall below the background noise level, both normalised to the same total luminosity."
 We uote that the error bars of Coleman et al. (, We note that the error bars of Coleman et al. (
2007) are oulv a lower estimate of the confidence intervals for this comparison. as they ouly refer to the errors in the measured star counts aud not to full surface density profile inferences.,"2007) are only a lower estimate of the confidence intervals for this comparison, as they only refer to the errors in the measured star counts and not to full surface density profile inferences."
 A very. good agreement is evident. showing the proposed models are a good. fully selt-cousisteut representation of the dynamics aud both the integral and spatially resolved light distribution iu- the well-studied local dSphli galaxies.," A very good agreement is evident, showing the proposed models are a good, fully self-consistent representation of the dynamics and both the integral and spatially resolved light distribution in the well-studied local dSph galaxies."
 We now turn to the scalines shownby the dSph galaxies iu our saüuple., We now turn to the scalings shown by the dSph galaxies in our sample.
" Eistlv. we show iu Figure(3) thebehaviour of the equilibrium isothermal coufigurations we are solving for. in terms of the resulting projected Ayy, as a function of the input value of py. at a constant value of 6= lOkmi/s."," Firstly, we show in Figure (3) the behaviour of the equilibrium isothermal configurations we are solving for, in terms of the resulting projected $R_{\rm hm}$ as a function of the input value of $\rho_{0}$ , at a constant value of $\sigma = 10 \rm{km/s}$ ."
" The remarkable feature of Figure (3) is that. after decreasing as pu increases. when py reaches a value of about O.3AL.pe7. the resulting Py, stops chaueine aud settles at relatively constant value of around 150pc."," The remarkable feature of Figure (3) is that, after decreasing as $\rho_{0}$ increases, when $\rho_{0}$ reaches a value of about $0.3 M_{\odot} \rm{pc}^{-3}$, the resulting $R_{\rm hm}$ stops changing and settles at relatively constant value of around $150\rm{pc}$,"
Obreschkow et al. (,Obreschkow et al. (
2009d) that are available as part of the SKA Simulated Skieswebsite*.,2009d) that are available as part of the SKA Simulated Skies.
. This simulation is based on the galaxy catalog from De Lucia & Blaizot (2007) which post-processed the Millennium dark matter simulation with semi-analytical models of galaxy formation., This simulation is based on the galaxy catalog from De Lucia $\&$ Blaizot (2007) which post-processed the Millennium dark matter simulation with semi-analytical models of galaxy formation.
" It captures the gas astrophysics, especially HI and H» and allow for the cosmic evolution of these two of the cold hydrogen gas (see, Obreschkow et al."," It captures the gas astrophysics, especially HI and $_2$ and allow for the cosmic evolution of these two phases of the cold hydrogen gas (see, Obreschkow et al."
 2009b; phases2009c for details)., 2009b; 2009c for details).
" They also constructed a virtual sky field with HI and CO line luminosities for each of the galaxies, calculated from a combination of the neutral and molecular gas masses and the metallicity of each galaxy."," They also constructed a virtual sky field with HI and CO line luminosities for each of the galaxies, calculated from a combination of the neutral and molecular hydrogen gas masses and the metallicity of each galaxy."
 Using the hydrogenonline tools we made a query with a four degree observing cone containing more than 5x10° galaxies at each of z=7., Using the online tools we made a query with a four degree observing cone containing more than $5\times 10^5$ galaxies at each of $z=7$.
 The Lco(M) from every halo at z7 from this simulation is shown in the left panel of Figure 1 with the shade of gray darker in the regions where the halo density is higher., The $L_{\rm CO}(M)$ from every halo at $z \sim 7$ from this simulation is shown in the left panel of Figure \ref{fig:L_CO_M_HI} with the shade of gray darker in the regions where the halo density is higher.
 Typically we find one galaxy per halo at these redshifts., Typically we find one galaxy per halo at these redshifts.
 The lines show the result of averaging the total Lco from the halos for each of 150 logarithmic bins of halo mass., The lines show the result of averaging the total $L_{\rm CO}$ from the halos for each of 150 logarithmic bins of halo mass.
" Motivated by parameterization of optical luminosity and halo mass 2005a), we use relation of the form (CoorayLco(M)Milosavljevic=L((M/M.)'(14- to adescribe the mean relation and determine the four free M/M.)parameters? A, b, c and d."," Motivated by parameterization of optical luminosity and halo mass (Cooray Milosavljevic 2005a), we use a relation of the form $L_{\rm CO}(M) = L_0(M/M_c)^b(1+M/M_c)^{-d}$ to describe the mean relation and determine the four free parameters $A$, $b$, $c$ and $d$."
" At z26,7, and 8, these parameters take the values of Lo24.3x10°,6.2106,4.010° Lo, 5—2.4,2.6,2.8, M.=3.5x101,3.010!,2.010!! Mo, and d=2.8,3.4,3.3, respectively."," At $z=6,7,$ and 8, these parameters take the values of $L_0=4.3\times 10^{6},6.2\times 10^{6}, 4.0\times 10^{6}$ $_{\sun}$, $b=2.4,2.6,2.8$, $M_c= 3.5\times 10^{11}, 3.0\times 10^{11}, 2.0\times 10^{11}$ $_{\sun}$, and $d=2.8,3.4,3.3$, respectively."
" Using the above relation between Lco and halo mass M as a function of redshift, we can write the mean intensity of the CO(1-0) line as Lo dM να”, where we take Minin=10°(Mo as the minimum mass of the dark matter halos that can host /h)galaxies (to be consistent with mass scale of atomic Hydrogen cooling or a virial temperature of 104K (Loeb Barkana 2001), dn/dM is the halo mass function (Sheth Tormen 1999), D; is the luminosity distance, and Dy is the comoving angular diameter distance."," Using the above relation between $L_{\rm CO}$ and halo mass $M$ as a function of redshift, we can write the mean intensity of the CO(1-0) line as = dM ^2, where we take $M_{\rm min} = 10^8 (M_{\odot}/h)$ as the minimum mass of the dark matter halos that can host galaxies (to be consistent with mass scale of atomic Hydrogen cooling or a virial temperature of $^4$ K (Loeb Barkana 2001), $dn/dM$ is the halo mass function (Sheth Tormen 1999), $D_{L}$ is the luminosity distance, and $D_{\rm A}$ is the comoving angular diameter distance."
" In above y(z)=dy/dvAco(1+z)*/H(2), where x is the comoving distance, v is the observed frequency, Aco=2.6 mm is the rest frame wavelength of the CO(1-0) line."," In above $y(z)={d\chi}/{d\nu}={\lambda_{\rm CO}(1+z)^2}/{H(z)}$, where $\chi$ is the comoving distance, $\nu$ is the observed frequency, $\lambda_{\rm CO}=2.6$ mm is the rest frame wavelength of the CO(1-0) line."
 The signal is expected to contain spatial variations around the average intensity due to poisson fluctuations in the number, The signal is expected to contain spatial variations around the average intensity due to poisson fluctuations in the number
and ten are the percentages of predicted ages that are within a factor of 3 of the published ages for the solar abundance model. the 0.2 solar abundance model. and the 2.5 solar abundance model. respectively.,"and ten are the percentages of predicted ages that are within a factor of 3 of the published ages for the solar abundance model, the 0.2 solar abundance model, and the 2.5 solar abundance model, respectively."
 This percentage is determined for the best-fit ages without consideration of the predicted or measured age uncertainties., This percentage is determined for the best-fit ages without consideration of the predicted or measured age uncertainties.
 Unfortunately. the only samples of OCs large enough to draw any significant interpretations are the By. (BV) and Byand(B V).," Unfortunately, the only samples of OCs large enough to draw any significant interpretations are the $-$ B), $-$ V) and $-$ B) and $-$ V)."
 Results from the other combinations of colours are merely suggestive., Results from the other combinations of colours are merely suggestive.
 We see from Table 2 that for several colour combinations the uncertainties in predicted age are relatively large., We see from Table 2 that for several colour combinations the uncertainties in predicted age are relatively large.
 This is because we have assumed the maximum uncertainty in measured colours for the OCs., This is because we have assumed the maximum uncertainty in measured colours for the OCs.
 Using the actual uncertainties in integrated colours for this sample would not change the predicted ages but would result in better age constraints., Using the actual uncertainties in integrated colours for this sample would not change the predicted ages but would result in better age constraints.
 Unfortunately the uncertainties in integrated colours of each of these OCs are not available., Unfortunately the uncertainties in integrated colours of each of these OCs are not available.
 However. the relative predicted age uncertainties provide a good measure of the precision afforded by each colour and colour combination for age-dating.," However, the relative predicted age uncertainties provide a good measure of the precision afforded by each colour and colour combination for age-dating."
 Looking at the mean uncertainties (column 7 in Table, Looking at the mean uncertainties (column 7 in Table
 (hereafter CJF) sour'ces and their analysis.,Flat-Spectrum (hereafter CJF) sources and their analysis.
of 11622 (Casassus et al.,of 1622 (Casassus et al.
 2006) to show a clistinet excess of microwave emission., 2006) to show a distinct excess of microwave emission.
" LDN11III (a=2140™30°.5|57748'00"" 2000) lies on the inside edge of the heel of the large (23 cldegree) horseshoe shaped region 11396 (Sh 2-131: Sharpless 1959)."," 1111 $\alpha = 21^{\rm{h}} 40^{\rm{m}} 30^{\rm{s}}, \delta = 
+57^{\circ} 48' 00''$ J2000) lies on the inside edge of the heel of the large $\approx 3$ degree) horseshoe shaped region 1396 (Sh 2-131; Sharpless 1959)."
 It has no known LRAS association (Parker 1988) but has been studied in the submillimetre at yam with the SCUBA instrument (Visser. Richer Chandler 2001).," It has no known IRAS association (Parker 1988) but has been studied in the submillimetre at $\mu$ m with the SCUBA instrument (Visser, Richer Chandler 2001)."
 An opacity class 6 object. Gl. 5pmmaeg: Lynds. 1962). it is one of the most opaque. dark nebulae.," An opacity class 6 object $A_{\nu} \ge 5$ mag; Lynds 1962), it is one of the most opaque dark nebulae."
sources. and filled svimbols. unresolved sources.,"sources, and filled symbols, unresolved sources."
 As expected. those sources exhibiting resolved structure are also those which show Large galaxy fractions and are at relatively low redshilts.," As expected, those sources exhibiting resolved structure are also those which show large galaxy fractions and are at relatively low redshifts."
 We now investigate whether emission from the host galaxies of Parkes quasars can significantly. contribute to their observed. ByA colours., We now investigate whether emission from the host galaxies of Parkes quasars can significantly contribute to their observed $B_{J}-K$ colours.
 We do this by computing the Bydv colour of the hypothesised underlying “quasar”. (D—BIN). we would expect if contribution from the host ealaxy was absent in each source.," We do this by computing the $B_{J}-K$ colour of the hypothesised underlying “quasar”, $(B_{J}-K)_{q}$, we would expect if contribution from the host galaxy was absent in each source."
" Lo the observed. colours were entirely due to galactic emission. then we expect. the distribution in (D;—A), to show a relatively small scatter. i.c. similar to that observed for opticallv-selected: quasars where typically (23,A),NEZο24."," If the observed colours were entirely due to galactic emission, then we expect the distribution in $(B_{J}-K)_{q}$ to show a relatively small scatter, i.e. similar to that observed for optically-selected quasars where typically $(B_{J}-K)_{q}\simeq2.5$."
" The colour of an underlying quasar. (2,όν. can be written in terms of the observed colour (2,fv)an. ancl the galaxy fractional contributions £5,,:GD,) ancl PisI) as follows:"," The colour of an underlying quasar, $(B_{J}-K)_{q}$, can be written in terms of the observed colour $(B_{J}-K)_{obs}$ and the galaxy fractional contributions $F_{gal}(B_{J})$ and $F_{gal}(K)$ as follows:"
measure is no longer fair when the X is small but. the curvature of the likelihood function is very big. meaning wt the constraints on the parameters are loose although the resulting X7 is small.,"measure is no longer fair when the $\chi^2$ is small but the curvature of the likelihood function is very big, meaning that the constraints on the parameters are loose although the resulting $\chi^2$ is small."
 In this paper. we introduce another figure of merit in nalogous to (Albrechtefαἱ2006:Albrecht&Bern-gaein 2007).. the area of the ws)> band to evaluate the »erformance of dilferent. parameterizations.," In this paper, we introduce another figure of merit in analogous to \citep{DETF1,DETF2}, the area of the $w(z)-z$ band to evaluate the performance of different parameterizations."
 The justification {this measure lies in that our ultimate goal is to constrain re shape of w(z2) as much as we can from the data., The justification of this measure lies in that our ultimate goal is to constrain the shape of $w(z)$ as much as we can from the data.
 In our analysis. we will compare the /parameterizations with identical number of parameters.," In our analysis, we will compare the parameterizations with identical number of parameters."
 Note that comparing »;wameterizations with different number of. parameters owed on Akaike’s Information Criterion (AIC). Bayesian Information Criterion (BIC) or other criteria is arbitrary in he sense of the criteria one chooses.," Note that comparing parameterizations with different number of parameters based on Akaike's Information Criterion (AIC), Bayesian Information Criterion (BIC) or other criteria is arbitrary in the sense of the criteria one chooses."
 In a sensible Bayesian method. a model is penalized for having a larger number of parameters that gives a reasonable fit (not the best it compared to models with more parameters) is awarded with increased. evidence for that model. (see for example. (Liddleefa£ 2006))).," In a sensible Bayesian method, a model is penalized for having a larger number of parameters that gives a reasonable fit (not the best fit compared to models with more parameters) is awarded with increased evidence for that model, (see for example, \citep{Liddle06}) )."
 Llowever. this is not what we are concerned and the purpose of this paper is Just to show what is the best way to parameterize the equation of state of dark energy for a variety of prevalent models with number of parameters.," However, this is not what we are concerned and the purpose of this paper is just to show what is the best way to parameterize the equation of state of dark energy for a variety of prevalent models with number of parameters."
 Our results show that the widely used parameterization. (ín)—ang|owzf£(l2). 1s not the one that can tell us most of the information of (2) in two-parameter parameterization family based on the SNla data.," Our results show that the widely used parameterization, $w(z)=w_0+w_1z/(1+z)$, is not the one that can tell us most of the information of $w(z)$ in two-parameter parameterization family based on the SNIa data."
 Amonge the many observations that can help to constrain the shape of οσο). SNla cata provide most sensitive and struehtlorwarel constraints.," Among the many observations that can help to constrain the shape of $w(z)$, SNIa data provide most sensitive and straightforward constraints."
 Therefore. in this paper. we will study the cllects of dilferent. parameterizations on our understanding of the evolution of dark energy. based on ολα data.," Therefore, in this paper, we will study the effects of different parameterizations on our understanding of the evolution of dark energy based on SNIa data."
 The outline of the paper is as follows., The outline of the paper is as follows.
 In section LL. we discuss the expansion history of the universe and. the observational variables from the supernova experiments.," In section II, we discuss the expansion history of the universe and the observational variables from the supernova experiments."
 In section LLL two parameterization Families and some prevalent. mioclels of equation. of state of dark. energy. is introduced.," In section III, two parameterization families and some prevalent models of equation of state of dark energy is introduced."
 Throughout the paper. we only consider. two-parameter models anc parameterizations.," Throughout the paper, we only consider two-parameter models and parameterizations."
 In. section. IV. the method ancl results of the analysis is. presented.," In section IV, the method and results of the analysis is presented."
 ba the last. we conclude with some remarks on the choice of parameterization of dark energy.," In the last, we conclude with some remarks on the choice of parameterization of dark energy."
 In the framework of standard cosmological model. assuming a spatially [at (4= 0) Friedmann universe. the equations governing the expansion of the universe are and where Lf=ἁα is the Hubble parameter. q=fad? is the deceleration parameter. 3;=ϱρε. ds the cosmic matter density. parameter and ws) is the dark energy equation of state. which is defined hy The dark energv2 density parameter {por evolves as PpE(S)=Poet(2) and To date. Supernovae Type la (SNla) provide the most direct. indication of the accelerating expansion of the universe.," In the framework of standard cosmological model, assuming a spatially flat $k = 0$ ) Friedmann universe, the equations governing the expansion of the universe are and where $H\equiv \dot{a}/a$ is the Hubble parameter, $q\equiv -\ddot{a}/aH^2$ is the deceleration parameter, $\Omega_M\equiv \rho_M/\rho_C$ is the cosmic matter density parameter and $w(z)$ is the dark energy equation of state, which is defined by The dark energy density parameter $\rho_{DE}$ evolves as $\rho_{DE}(z)=\rho_{DE}^0f(z)$ and To date, Supernovae Type Ia (SNIa) provide the most direct indication of the accelerating expansion of the universe."
 For the distant SNla. one can directly observe their apparent magnitude m and redshift ο. because the absolute magnitude M of them are assumed to be constant. i0. SNla are standard candles.," For the distant SNIa, one can directly observe their apparent magnitude $m$ and redshift $z$, because the absolute magnitude $M$ of them are assumed to be constant, i.e., SNIa are standard candles."
" The luminosity distance d;(z) is the ""meeting point” between the observed m(z) and theoretical prediction (2): and Although //(2) is more cürectlv related to. the observable Luminosity distance and then is easier to measure more accurately. in order to investigate the evolution of dark energy with time and the scale factor. constraints on (eC) is essentially equivalent to that of ἐς) and is also crucial for understanding the nature of dark energy (Liuterer&Starkman 2003)."," The luminosity distance $d_L(z)$ is the ""meeting point"" between the observed $m(z)$ and theoretical prediction $H(z)$: and Although $H(z)$ is more directly related to the observable luminosity distance and then is easier to measure more accurately, in order to investigate the evolution of dark energy with time and the scale factor, constraints on $w(z)$ is essentially equivalent to that of $H(z)$ and is also crucial for understanding the nature of dark energy \citep{Huterer}."
. Since ws) is a continuous function with an infinite number of values at a finite redshift range. ες). must be modeled. using just a few parameters whose values are determined by fitting to observations.," Since $w(z)$ is a continuous function with an infinite number of values at a finite redshift range, $w(z)$ must be modeled using just a few parameters whose values are determined by fitting to observations."
 A merit of using ος) with a particular parametrization is to compare the performance of dillerent. experiments., A merit of using $w(z)$ with a particular parametrization is to compare the performance of different experiments.
 Note here that no single parameterization can represent all possibilities for eC)., Note here that no single parameterization can represent all possibilities for $w(z)$.
 AX reasonable parameterization must be accorded with the demand that dark energy is important at late times and insignificant at carly times., A reasonable parameterization must be accorded with the demand that dark energy is important at late times and insignificant at early times.
 There exist plenty of parameterizations for the equation of state (a) (Johri&Rath2006.2007:Johri2004) where à=(1|2)Fl. but. most. of them. are purely phenomenological.," There exist plenty of parameterizations for the equation of state $w(a)$ \citep{p1,p2,p3} where $a=(1+z)^{-1}$, but most of them are purely phenomenological."
 Maybe. we should consider some of them in the sense that they are generalized [from the behavior of physically motivated sets of models (Lincer2007).," Maybe, we should consider some of them in the sense that they are generalized from the behavior of physically motivated sets of models \citep{Linder07}."
. For single parameter. models. e.g. a= costant. no dvnamies is embodied: and can not. parameterize the rate of change of « and then high fine-tuning is needed.," For single parameter models, e.g. $w=costant$, no dynamics is embodied and can not parameterize the rate of change of $w$ and then high fine-tuning is needed."
 The physical svnimetry motivated: one-parameter. mocels. such as topological defects. are not consistent with the observation data.," The physical symmetry motivated one-parameter models, such as topological defects, are not consistent with the observation data."
 More. parameters mean more degrees of freedom for adaptability το Observations. at the same time more degeneracies in the determination of parameters.," More parameters mean more degrees of freedom for adaptability to observations, at the same time more degeneracies in the determination of parameters."
 For models with more than two parameters. they lack predictability ancl even the next generation of experiments will not be able to constrain stringently (Linder&Llutercr 2005)..," For models with more than two parameters, they lack predictability and even the next generation of experiments will not be able to constrain stringently \citep{LH05}. ."
 Therefore. we only consider the two-parancter," Therefore, we only consider the two-parameter"
the same number of times on average.,the same number of times on average.
 Moreover. given that we expect correlations on the size of the simulation box to be negligible. there is no dillerence between this statistical measurement and one with only a single skewer and multiple realizations of the simulation volume are produced: using different initial concitions.," Moreover, given that we expect correlations on the size of the simulation box to be negligible, there is no difference between this statistical measurement and one with only a single skewer and multiple realizations of the simulation volume are produced using different initial conditions."
 Figure 1. compares the contributions from cosmic and Poisson variance as calculated by linear theory (, Figure \ref{fig:sigfov} compares the contributions from cosmic and Poisson variance as calculated by linear theory (Eq.
"leq. 1. and 2)) lor a,= as a function of the opening angle of the survey, @=a,/(2)."," \ref{eq:var} and \ref{eq:cosvar}) ) for $a_x=a_y$ as a function of the opening angle of the survey, $\theta=a_x/\chi(z)$."
 This diagram can be used to calculate the effectiveness of future surveys with Large fields of view., This diagram can be used to calculate the effectiveness of future surveys with large fields of view.
 Given the limiting size of our simulation volume. a sullicient number of skewers is not obtainable to perform comparable numerical calculations for large survevs.," Given the limiting size of our simulation volume, a sufficient number of skewers is not obtainable to perform comparable numerical calculations for large surveys."
 However. figures and 3 show the agreement for the LLUDE field-of-view.," However, figures \ref{fig:var6-8} and \ref{fig:var8-10} show the agreement for the HUDF field-of-view."
 Figures 2. and 3. compare the contributions to the variance in the simulations with calculations based on linear theory for the >=6S and z=58Ü ranges. respectively.," Figures \ref{fig:var6-8} and \ref{fig:var8-10} compare the contributions to the variance in the simulations with calculations based on linear theory for the $z=6-8$ and $z=8-10$ ranges, respectively."
 They. suggest that the observed. statistics are well approximated by the analvtie caleulations at the Low luminosity limits., They suggest that the observed statistics are well approximated by the analytic calculations at the low luminosity limits.
 However. while linear theory predicts a Gaussian probability cistribution of the count of halos with variance given by equation 2.. the simulated: probability istribution of bright LBGs has a non-Gaussian shape.," However, while linear theory predicts a Gaussian probability distribution of the count of halos with variance given by equation \ref{eq:cosvar}, the simulated probability distribution of bright LBGs has a non-Gaussian shape."
 This can be clearly seen. by eve in. Figure d. where he solid histogram representing the measured probability istribution is best represented by a log-normal clistribution.," This can be clearly seen by eye in Figure \ref{fig:pdist8-10}, where the solid histogram representing the measured probability distribution is best represented by a log-normal distribution."
 The skewness as a function of minimum luminosity is presented. in the bottom. panels of figures 2 and 3..., The skewness as a function of minimum luminosity is presented in the bottom panels of figures \ref{fig:var6-8} and \ref{fig:var8-10}.
 We cline the skewness as the third moment of the probability istribution normalized bv the variance to the 3/2 power: The seemingly large amplitude variations in the skewness at low luminosity for z=6 SN are due to small numerical Uuctuations around the nearly zero skewness from the simulation. plotted on a log scale.," We define the skewness as the third moment of the probability distribution normalized by the variance to the $3/2$ power: The seemingly large amplitude variations in the skewness at low luminosity for $z=6$ $8$ are due to small numerical fluctuations around the nearly zero skewness from the simulation, plotted on a log scale."
 Deviations between the analvtie and simulation values of the sample variance grow as the skewness becomes more significant., Deviations between the analytic and simulation values of the sample variance grow as the skewness becomes more significant.
 This behavior is a manifestation of nonlinear clustering on the small scales probed by the narrowness of the skewer., This behavior is a manifestation of nonlinear clustering on the small scales probed by the narrowness of the skewer.
 Finally. the figures show sharp regimes where the field- variance is dominated. either by cosmic or Poisson variance.," Finally, the figures show sharp regimes where the field-to-field variance is dominated either by cosmic or Poisson variance."
" For the two redshift ranges ο=6 S and 2=8 10 this division comes at around L(Al=510?AL.)-6«LOecres/s/liz and L(AM-—2.10""AM.)=3107 ergs/s/LMz.. respectively."," For the two redshift ranges $z=6$ $8$ and $z=8$ $10$ this division comes at around $L\left(M=5\times10^{10}\msun\right)=6\times10^{28}\,\rm{ergs/s/Hz}$ and $L\left(M=2\times10^{10}\msun\right)=3\times10^{28}\,\rm{ergs/s/Hz}$ , respectively."
 Phe nonlinearity at higher luminosities causes the sample variance to dominate at masses which are ~50 1004 Lweer than otherwise expected., The nonlinearity at higher luminosities causes the sample variance to dominate at masses which are $\sim 50$ $100\%$ larger than otherwise expected.
 However. since the Poisson variance is unallected by nonlinearities. there remains only a specific mass range (shown in the figures) over which the total variance is different than expected.," However, since the Poisson variance is unaffected by nonlinearities, there remains only a specific mass range (shown in the figures) over which the total variance is different than expected."
continuum curvature would have on the inferred X-ray reflection as a function of radius in the clisk.,continuum curvature would have on the inferred X-ray reflection as a function of radius in the disk.
 For this study. we examine the Κον EDPIC-pn data supplemented by the kkeW IUNTE-PC data.," For this study, we examine the keV EPIC-pn data supplemented by the keV RXTE-PCA data."
 Adding the 2kkeV. EPIC-pn. data complicates the spectral fitting (since one must account for the soft. X-rav absorptionemission) without improving the constraints., Adding the keV EPIC-pn data complicates the spectral fitting (since one must account for the soft X-ray absorption/emission) without improving the constraints.
 The major limitation in the study of à curved continuum is the lack of reacily available reflection models that can handle non-power law input spectra., The major limitation in the study of a curved continuum is the lack of readily available reflection models that can handle non-power law input spectra.
 We use the mocel (Vitarchuk 1994). in NSPEC to describe the continuum. resulting from thermal C'omptonisation., We use the model (Titarchuk 1994) in XSPEC to describe the continuum resulting from thermal Comptonisation.
 The seed. photons are assumed to be characterized by a Wien spectrum with KT= 50ceV (typical of the optically-thick part of AGN disk)., The seed photons are assumed to be characterized by a Wien spectrum with $kT=50$ eV (typical of the optically-thick part of AGN disk).
 Phe temperature Z ancl optical depth 7 of the corona are left as free parameters., The temperature $T$ and optical depth $\tau$ of the corona are left as free parameters.
 We then employ. the model (Alagdziarz Zdziarski 1995) in NSPEC to model the Compton reflection of this continuum spectrum from the disk surface., We then employ the model (Magdziarz Zdziarski 1995) in XSPEC to model the Compton reflection of this continuum spectrum from the disk surface.
 AX narrow Gaussian emission line with a rest-frame energy in the range kkeV is added: to model the iron IIuorescence line. ancl the whole spectrum is convolved with the (Laor 1991) kernel to describe the Doppler and eravitational redshift effects associated with the accretion disk.," A narrow Gaussian emission line with a rest-frame energy in the range keV is added to model the iron fluorescence line, and the whole spectrum is convolved with the (Laor 1991) kernel to describe the Doppler and gravitational redshift effects associated with the accretion disk."
 This procedure. constrains the temperature of the corona to be greater than AY~ 5.2kkeV: for coronal temperatures exceeding this. there is an almost perfect cancellation between spectral curvatures resulting from different coronal temperatures. reflection. fractions and relativistic smearings.," This procedure constrains the temperature of the corona to be greater than $kT\sim 5.2$ keV; for coronal temperatures exceeding this, there is an almost perfect cancellation between spectral curvatures resulting from different coronal temperatures, reflection fractions and relativistic smearings."
 However. for all allowable coronal temperatures. a very steep emissivitv index is required. jc 3.9.," However, for all allowable coronal temperatures, a very steep emissivity index is required, $\beta>3.9$ ."
 Freezing the emissivity index to be 3= resulted in a worsening of the goodness of fit. parameter by Ay?=400., Freezing the emissivity index to be $\beta=3$ resulted in a worsening of the goodness of fit parameter by $\Delta\chi^2=400$.
 Thus. our principal result is secure against the continuum curvature introduced by standard thermal C'omptonisation moclels.," Thus, our principal result is secure against the continuum curvature introduced by standard thermal Comptonisation models."
 In the above discussion. we have examined these data with a variety of spectral models.," In the above discussion, we have examined these data with a variety of spectral models."
 We have found that the need for extreme relativistic cllects is robust to dilferent. treatments of the soft X-ray complexity. complex absorption. ancl the use of a Comptonisation model instead. of a simple aw to describe the primary X-ray continuum.," We have found that the need for extreme relativistic effects is robust to different treatments of the soft X-ray complexity, complex absorption, and the use of a Comptonisation model instead of a simple power-law to describe the primary X-ray continuum."
" Until now. we have emploved a phenomenological mocel or the radial dependence of the disk emissivity. assuming hat it can be deseribed by a power-law form οxr runcatecd by inner and outer radii ri, and rau."," Until now, we have employed a phenomenological model for the radial dependence of the disk emissivity, assuming that it can be described by a power-law form $\epsilon\propto r^{-\beta}$ truncated by inner and outer radii $r_{\rm in}$ and $r_{\rm out}$."
 While his is an extremely useful. parameterization. it does. not correspond to any particular physical disk model.," While this is an extremely useful parameterization, it does not correspond to any particular physical disk model."
 Given the quality of these data. we can go bevone these simple power-aw emissivitv profiles and attempt to constrain physical relativistic smearing models.," Given the quality of these data, we can go beyond these simple power-law emissivity profiles and attempt to constrain physical relativistic smearing models."
 Firstly. we shall review some of the pertinent theory related o the gcometrically-thin accretion disks of the type thought o be operating in σον[ο nuclei such as Ας6-30-15.," Firstly, we shall review some of the pertinent theory related to the geometrically-thin accretion disks of the type thought to be operating in Seyfert nuclei such as MCG–6-30-15."
 The standard thin disk model of black hole accretion was developed in a Newtonian setting by Shakura Sunvacy (1973). and extended. into a fully relativistic theory by Ὀνίκον Thorne (1974) and Page Vhorne (1974: yoreafter PT).," The standard thin disk model of black hole accretion was developed in a Newtonian setting by Shakura Sunyaev (1973), and extended into a fully relativistic theory by Novikov Thorne (1974) and Page Thorne (1974; hereafter PT)."
 In this moclel. the accretion disk is assumed to x &cometricallv-thin. raciatively-cllicicnt. and in a steacly-state.," In this model, the accretion disk is assumed to be geometrically-thin, radiatively-efficient, and in a steady-state."
 Furthermore. it is postulated that the disk experiences zero torque at the radius of marginal stability.," Furthermore, it is postulated that the disk experiences zero torque at the radius of marginal stability."
 With these assuniptions. one can compute the dissipation rate. ancl hence total radiative Hux as a functionof radius and black hole spin: where we have defined the function JF ," With these assumptions, one can compute the dissipation rate, and hence total radiative flux as a functionof radius and black hole spin: where we have defined the function ${\cal F}$ "
since theoretical wavelengths do not reach spectroscopic accuracy. statistical comparisons with measurements are essential inasmuch as thev enable small empirical adjustments {ο make (he computed values astrophysically useful.,"Since theoretical wavelengths do not reach spectroscopic accuracy, statistical comparisons with measurements are essential inasmuch as they enable small empirical adjustments to make the computed values astrophysically useful."
 It is worth pointing out that. due to the lack of experimental data for ionic species of the second row. this procedure is only currently possible for svstems with electron number V<9.," It is worth pointing out that, due to the lack of experimental data for ionic species of the second row, this procedure is only currently possible for systems with electron number $N\leq 9$."
 Experimental wavelengths for Ar ions with 3€NN<9 have been reported by as well as theoretical values computed with the aud codes., Experimental wavelengths for Ar ions with $3\leq N\leq 9$ have been reported by \citet{bie00} as well as theoretical values computed with the and codes.
 The two latter data sets are to be hereafter referred to as the IIER2 and MCDEL. respectively.," The two latter data sets are to be hereafter referred to as the HFR2 and MCDF1, respectively."
 As shown in Fig. T..," As shown in Fig. \ref{wl},"
" a statistical comparison with the measured values can be carried oul where the average theoryexperiment wavelength difference. AA,. is plotted as afunction of the ion electron number NV."," a statistical comparison with the measured values can be carried out where the average theory–experiment wavelength difference, $\overline{\Delta
\lambda}_{\rm e}$, is plotted as afunction of the ion electron number $N$."
 It may be seen that theoretical wavelengths appear to be always shorter than the measured values., It may be seen that theoretical wavelengths appear to be always shorter than the measured values.
 In the case of our HERI data. .NÀ.| increases with V from under 1 [for NV=3 to around 6 [for σσ.Nx9.," In the case of our HFR1 data, $|\overline{\Delta \lambda}_{\rm e}|$ increases with $N$ from under 1 for $N=3$ to around 6 for $8\leq
N\leq 9$."
" A similar behavior is displaved by the MCDEI values except for 8<<N«9 where .NA,|<LmAs this seems to indicate that the transition energies in these (wo ions have been adjusted with the experimentalvalues."," A similar behavior is displayed by the MCDF1 values except for $8\leq N\leq 9$ where $|\overline{\Delta \lambda_{\rm
e}}|< 1$; this seems to indicate that the transition energies in these two ions have been adjusted with the experimentalvalues."
" This is certainlv also the case with the ΠΕΙΣ results for 3<Nx9 where AA,<|mA.."," This is certainly also the case with the HFR2 results for $3\leq
N\leq 9$ where $|\overline{\Delta \lambda_{\rm e}}| < 1$."
 However. as shown in Fig. 7..," However, as shown in Fig. \ref{wl},"
 the standard deviations for all three theoretical data sets are awound 23 wwhich gives an indication of the level of accuracy that ean be attained theoretically., the standard deviations for all three theoretical data sets are around 2–3 which gives an indication of the level of accuracy that can be attained theoretically.
" For the computed MCDE wavelengths listed by Chen&Crasemann(1988). and al.(1997). for D- and C-like Ar. |AA.)22x2 aand [NA,|=446 mA.. respectively."," For the computed MCDF wavelengths listed by \citet{che88} and \citet{che97} for B- and C-like Ar, $|\overline{\Delta
\lambda_{\rm e}}|=2\pm 2$ and $|\overline{\Delta
\lambda_{\rm e}}|=4\pm 6$ , respectively."
 The somewhat large standard. deviation in the svstem is caused by (ransitions involving heavily mixed Ix-vacancey states. in particular I5C8$)2s2p!CD)!Da. but the level of agreement is in general similar to that discussed above [or IIFRI.," The somewhat large standard deviation in the C-like system is caused by transitions involving heavily mixed K-vacancy states, in particular ${\rm
1s(^2S)2s2p^4(^2D)\ ^1D_2}$, but the level of agreement is in general similar to that discussed above for HFR1."
 Diflerences between the IERI. wavelengths and the measurements listed by Faenov(1994). for ions with 12<Z16 and 4xN<9 displav a similar behavior {ο that found in the Ar ions: for each isomiclear sequence. [NA| increases with VV.," Differences between the HFR1 wavelengths and the measurements listed by \citet{fae94} for ions with $12\leq Z\leq 16$ and $4\leq
N\leq 9$ display a similar behavior to that found in the Ar ions: for each isonuclear sequence, $|\overline{\Delta \lambda_{\rm
e}}|$ increases with $N$."
 Also. as expected. il decreases wilh Z along the isoelectronic sequence: for instance. in the case of NV— 9. [AAz18 (for Z=14 while .NA.|z10 for Z= 16.," Also, as expected, it decreases with $Z$ along the isoelectronic sequence; for instance, in the case of $N=9$ , $|\overline{\Delta \lambda_{\rm e}}|\approx 18$ for $Z=14$ while $|\overline{\Delta \lambda_{\rm e}}|\approx
10$ for $Z=16$ ."
 On the otherhand. we have found three transitions in N-like ions CN= 7) with unusually large wavelength differences which are listed in Table 3..," On the otherhand, we have found three transitions in N-like ions $N=7$ ) with unusually large wavelength differences which are listed in Table \ref{w-faenov}. ."
range of 75°<γης1057.,range of $75^\circ < \gamma_\mathrm{He} < 105^\circ$.
 Azimuth differences~.. dye—si. in between the white contours are presented in the histograms of Figs. 13((," Azimuth , $\phi_\mathrm{He} - \phi_\mathrm{Si}$, in between the white contours are presented in the histograms of Figs. \ref{Fig:histograms}( ("
c.e). revealing that: To understand why the azimuth of the chromospheric vector field always pointed in a slightly conterclockwise direction with respect to the photospheric one. we provide the inset of Fig. 13((,"c,e), revealing that: To understand why the azimuth of the chromospheric vector field always pointed in a slightly conterclockwise direction with respect to the photospheric one, we provide the inset of Fig. \ref{Fig:histograms}( ("
d).,d).
 On July 3rd. the chromospheric magnetic field was aligned with the spine. while the photospheric one displayed an inverse configuration. thus. having smaller azimuth angles (angles are measured from the .-axis and increase counterclockwise).," On July 3rd, the chromospheric magnetic field was aligned with the spine, while the photospheric one displayed an inverse configuration, thus, having smaller azimuth angles (angles are measured from the $x$ -axis and increase counterclockwise)."
 We interpret this as an indication that the field lines were wrapped around the filament axis., We interpret this as an indication that the field lines were wrapped around the filament axis.
 On July Sth. the situation was reversed: the filament axis was delineated by the photospheric field lines while the chromospheric ones showed a normal polarity configuration. thus mapping a higher part of the structure than two days earlier.," On July 5th, the situation was reversed: the filament axis was delineated by the photospheric field lines while the chromospheric ones showed a normal polarity configuration, thus mapping a higher part of the structure than two days earlier."
 Again the chromospheric azimuth angles were larger than the photospherte ones., Again the chromospheric azimuth angles were larger than the photospheric ones.
 In this paper. we present an extensive study of the vector magnetic field at two different heights in a filament that lies on top of the polarity inversion line of an active region.," In this paper, we present an extensive study of the vector magnetic field at two different heights in a filament that lies on top of the polarity inversion line of an active region."
 This study is mostly based on TIP-II data in the sspectral region. but also uses other multiwavelength data sources.," This study is mostly based on TIP-II data in the spectral region, but also uses other multiwavelength data sources."
 Helium intensity core images clearly showed the presence of a thin filament spine. that we identify with the filament axis. on top of the PIL on both days.," Helium intensity core images clearly showed the presence of a thin filament spine, that we identify with the filament axis, on top of the PIL on both days."
 Our observations reveal that this spine ts dominated by strong homogeneous horizontal fields. typically in the range of 500GG im the chromosphere and 100GG smaller in. the photosphere underneath.," Our observations reveal that this spine is dominated by strong homogeneous horizontal fields, typically in the range of G in the chromosphere and G smaller in the photosphere underneath."
 Vertical gradients of the field strength with stronger fields in the upper layers. have already been observed (e.g..inpolarcrownprominencesasshownby?) and modelled in the past (2) and have been ascribed to the presence of dips (?2)..," Vertical gradients of the field strength with stronger fields in the upper layers, have already been observed \citep[e.g., in
polar crown prominences as shown by][]{leroy83} and modelled in the past \citep{aulanier03} and have been ascribed to the presence of dips \citep{anzer69,demoulin89}."
 The vector magnetic field in the filament spine at chromospheric heights is highly sheared (1.e.. parallel to the PIL). while the photospherie vector field below the spine shows a uniformly inverse polarity.," The vector magnetic field in the filament spine at chromospheric heights is highly sheared (i.e., parallel to the PIL), while the photospheric vector field below the spine shows a uniformly inverse polarity."
 Such a configuration naturally suggests a flux rope topology., Such a configuration naturally suggests a flux rope topology.
 This scenario is also favored by the at 171AA.. which showed the filament with an inverse-S shape structure.," This scenario is also favored by the at $171$, which showed the filament with an inverse-S shape structure."
 Such a shape is thought to be an indication of the existence of an underlying flux rope (?).., Such a shape is thought to be an indication of the existence of an underlying flux rope \citep{gibson02}.
 This configuration of the spine with inverse orientation in the photosphere and shearedfield lines in the chromosphere. shows no significant," This configuration of the spine with inverse orientation in the photosphere and shearedfield lines in the chromosphere, shows no significant"
"noisy due to the low oscillatory power of the two velocity signals, and we do not find any clear relation between them.","noisy due to the low oscillatory power of the two velocity signals, and we do not find any clear relation between them."
" In the spectral range from 2 mHz to 4 mHz, the phase difference takes a constant value around zero, while for frequencies above 4 mHz, Ad decreases with increasing frequency."," In the spectral range from 2 mHz to 4 mHz, the phase difference takes a constant value around zero, while for frequencies above 4 mHz, $\Delta \phi$ decreases with increasing frequency."
" However, the low power at the photosphere at frequencies between 6 and 8 mHz (Fig. 8))"," However, the low power at the photosphere at frequencies between 6 and 8 mHz (Fig. \ref{fig:power_spectra}) )"
 results in a noisy phase spectrum in this band., results in a noisy phase spectrum in this band.
 Comparing the phase difference diagrams between the same two lines inside the umbra of the sunspot (top panel in Fig. 12)), Comparing the phase difference diagrams between the same two lines inside the umbra of the sunspot (top panel in Fig. \ref{fig:dfase_FeSi}) )
" and in the quiet Sun, we find that in quiet Sun the negative slope of the phase difference is less steep."," and in the quiet Sun, we find that in quiet Sun the negative slope of the phase difference is less steep."
 We also fitted the quiet Sun phase difference and amplification spectra with our wave propagation model., We also fitted the quiet Sun phase difference and amplification spectra with our wave propagation model.
" Inside the umbra of the sunspot, the magnetic field imposes wave propagation along field lines, and thus the hypothesis of vertical propagation is justified."," Inside the umbra of the sunspot, the magnetic field imposes wave propagation along field lines, and thus the hypothesis of vertical propagation is justified."
" However, in the quiet Sun waves can propagate in different directions and this approximation may not be true."," However, in the quiet Sun waves can propagate in different directions and this approximation may not be true."
" Still, we were able to find a fit matching reasonably the phase spectrum."," Still, we were able to find a fit matching reasonably the phase spectrum."
 We retrieve a lower value of formation height differences and a higher temperature in the quiet Sun compared to the umbra (see Table 3))., We retrieve a lower value of formation height differences and a higher temperature in the quiet Sun compared to the umbra (see Table \ref{tb:parameters}) ).
 The bottom panel of Fig., The bottom panel of Fig.
 15 shows the amplification spectrum for this case., \ref{fig:dfase_FeSi_quiet} shows the amplification spectrum for this case.
" For frequencies in the band of 2-7 mHz, the theoretical amplification matches the observational one."," For frequencies in the band of 2-7 mHz, the theoretical amplification matches the observational one."
" From the fit to the phase and amplitude spectra of all the pairs of lines we obtain the temperature, the difference between the formation heights of two lines, and the cooling time that best match the observational data."," From the fit to the phase and amplitude spectra of all the pairs of lines we obtain the temperature, the difference between the formation heights of two lines, and the cooling time that best match the observational data."
" Obviously, our model is simplified and has several important limitations."," Obviously, our model is simplified and has several important limitations."
" It only describes linear wave propagation in an isothermal stratified atmosphere, not taking into account a realistic temperature stratification of the sunspot atmosphere."," It only describes linear wave propagation in an isothermal stratified atmosphere, not taking into account a realistic temperature stratification of the sunspot atmosphere."
" In a real sunspot, the temperature varies with height, so the temperature that we obtain from the fit represents a mean value between two heights of formation."," In a real sunspot, the temperature varies with height, so the temperature that we obtain from the fit represents a mean value between two heights of formation."
" In the case of spectral lines formed at a similar height(i.e.,, the llines), the temperature that we find should be close to the real temperature in the layer."," In the case of spectral lines formed at a similar height, the lines), the temperature that we find should be close to the real temperature in the layer."
" However, when the formation height difference (Az) obtained from the fit is larger(i.e, Sir-He1)) we can not assign the temperature to a certain height."," However, when the formation height difference $\Delta z$ ) obtained from the fit is larger, ), we can not assign the temperature to a certain height."
" The presence of shocks in the temporal evolution of the aand vvelocities (Fig. 7)),"," The presence of shocks in the temporal evolution of the and velocities (Fig. \ref{fig:velocity_x}) ),"
" together with the difficulties of the model to fit the amplification spectra between aand ((Fig. 14)),"," together with the difficulties of the model to fit the amplification spectra between and (Fig. \ref{fig:dfase_CaHe}) ),"
 show that the propagation at these heights is non-linear., show that the propagation at these heights is non-linear.
 Thus our determination of the height difference between these two layers must be somewhat affected by the deviations from the linear regime., Thus our determination of the height difference between these two layers must be somewhat affected by the deviations from the linear regime.
" Despite these limitations, we plotted the deduced temperature values over the temperature stratification of the sunspot model of Maltbyetal.(1986) (Fig. 16))."," Despite these limitations, we plotted the deduced temperature values over the temperature stratification of the sunspot model of \citet{Maltby+etal1986} (Fig. \ref{fig:strat_T}) )."
 The symbols in this figure mark the values retrieved from the fit to our observations., The symbols in this figure mark the values retrieved from the fit to our observations.
" From the fit, we obtain the relative difference between the formation heights of spectral lines, not the absolute values."," From the fit, we obtain the relative difference between the formation heights of spectral lines, not the absolute values."
" As reference point, we set the height of the lline to z=308 km, as given by Bard&Carlsson for the sunspot atmosphere model of Maltbyetal. (1986)."," As reference point, we set the height of the line to z=308 km, as given by \citet{bard+carlsson2008} for the sunspot atmosphere model of \citet{Maltby+etal1986}."
. The formation heights of the other lines then follow from their relative distance to the lline (Table 3))., The formation heights of the other lines then follow from their relative distance to the line (Table \ref{tb:parameters}) ).
 Table 4 lists the resulting height values and the corresponding temperature., Table \ref{tb:formation} lists the resulting height values and the corresponding temperature.
" In the case of the llines, we obtained from the low phase difference (top panel of Fig. 11))"," In the case of the lines, we obtained from the low phase difference (top panel of Fig. \ref{fig:dfase_Fe2Fe5}) )"
" that the distances between them are very small, but the amplification spectra (bottom panel of Fig. 11))"," that the distances between them are very small, but the amplification spectra (bottom panel of Fig. \ref{fig:dfase_Fe2Fe5}) )"
 or the histograms of velocities (Fig. 6)), or the histograms of velocities (Fig. \ref{fig:velocity_hist}) )
 show a certain increase of the amplitude with height., show a certain increase of the amplitude with height.
 This amplification in spite of the small geometrical distance indicates a low pressure scale height Ho and consequently a low temperature T'., This amplification in spite of the small geometrical distance indicates a low pressure scale height $H_0$ and consequently a low temperature $T$.
 The height of the lline close to the temperature minimum agrees with the results from the fit of the model to the phase and amplification spectra., The height of the line close to the temperature minimum agrees with the results from the fit of the model to the phase and amplification spectra.
" In the case of the chromospheric signals ccore and we assign to both formation heights the lline),temperature retrieved from the fit of the phase spectra between aand vvelocities."," In the case of the chromospheric signals core and line), we assign to both formation heights the temperature retrieved from the fit of the phase spectra between and velocities."
" According to the geometrical differences of Table 3,, there is an uncertainty in the formation height of these two lines."," According to the geometrical differences of Table \ref{tb:parameters}, there is an uncertainty in the formation height of these two lines."
" On the one hand, we have set them from their Az with respect to the Iline, locating the formation height of the ccore at 958 km and the formation height of the lline at 1208 km."," On the one hand, we have set them from their $\Delta z$ with respect to the line, locating the formation height of the core at 958 km and the formation height of the line at 1208 km."
" On the other hand, we have obtained an alternative height for the ccore by considering its Az to the lline, and subtracting it from the previous estimate of"," On the other hand, we have obtained an alternative height for the core by considering its $\Delta z$ to the line, and subtracting it from the previous estimate of"
DM Tau is a pre-main sequence star located —110 pe away [rom the Sun. surrouided by a eas-rich Wepleriau circumstellar disk first discovered by Guilloteau&Dutrey(199D).,"DM Tau is a pre-main sequence star located $\sim$ 140 pc away from the Sun, surrounded by a gas-rich Keplerian circumstellar disk first discovered by \citet{Guilloteau94}."
. With a mass of 0.01 M... the disk is one of the more massive protoplanetary clisks known (Audrewsetal.2011).," With a mass of 0.04 $_\odot$, the disk is one of the more massive protoplanetary disks known \citep{Andrews11}."
. Π is a transition disk with an imaged central hole of radius 19 AU (Calvet 2011).," It is a transition disk with an imaged central hole of radius 19 AU \citep{Calvet05,Andrews11}."
. The DM Tau cisk shows strong emission [rom species proposed to trace cold gas (Obereetal.2010).. in particular — and N5H.. . and has previously been detected iu a variety oL molecular lines using single-dish telescopes (Dutreyetal.1996)..," The DM Tau disk shows strong emission from species proposed to trace cold gas \citep{Oberg10c}, in particular $^+$ and $_2$ $^+$, and has previously been detected in a variety of molecular lines using single-dish telescopes \citep{Dutrey96}. ."
 It is therefore an especially £&ood target to search lor HoD eimission., It is therefore an especially good target to search for $_2$ $^+$ emission.
 Ceccarellietal.(2001) reported a detection of the H»D lio—lyy line at 372 GHz at the L7 6 level toward DML Tau using the Caltech Subimillimeter Observatory (C50).," \citet{Ceccarelli04} reported a detection of the $_2$ $^+$ $1_{1,0}-1_{1,1}$ line at 372 GHz at the 4.7 $\sigma$ level toward DM Tau using the Caltech Submillimeter Observatory (CSO)."
" However. Guilloteauetal.(2006) suggestMOD that the DM Tau result is ""at best a 2o cletection” when analyzed with a proper model."," However, \citet{Guilloteau06} suggest that the DM Tau result is “at best a $\sigma$ detection” when analyzed with a proper model."
 To help to resolve these conflicting claims. we have used the Submillimeter Array to observe this HoD liue toward DAL Tau with improved sensitivity.," To help to resolve these conflicting claims, we have used the Submillimeter Array to observe this $_2$ $^+$ line toward DM Tau with improved sensitivity."
 We analyze these results. together with observatious of other molecular ious. to provide new constraints on the ionization fraction in different regions of the clisk.," We analyze these results, together with observations of other molecular ions, to provide new constraints on the ionization fraction in different regions of the disk."
 The paper is orgauized as follows., The paper is organized as follows.
 In 822. we describe the SMLA observations aud complementary IRAM. 30 meter telescope observations of molecular ious in the DAL Tau disk.," In 2, we describe the SMA observations and complementary IRAM 30 meter telescope observations of molecular ions in the DM Tau disk."
 Iu. 833. we present ipper limits on HeD and NoH J—153. emission aud detectious of the CO J=3-2 aud HYCO J=3-2 lines aud other resolved observations of molecular ious from the literature.," In 3, we present upper limits on $_2$ $^+$ and $_2$ $^+$ J=4–3 emission and detections of the CO J=3–2 and $^{13}$ $^+$ J=3–2 lines and other resolved observations of molecular ions from the literature."
 In 811. we use he line intensities together with an empirically derived disk structure to coustrain the midplane emiperaturere (S111) and to calculate upper limits on the o—HeD — coliunu (811.2). the total H3D columa 11.3). the total midplaue ion column ($11.1). the midplane ionization fraction (811.5). aud he ionization fraction in liigher disk lavers (81L6).," In 4, we use the line intensities together with an empirically derived disk structure to constrain the midplane temperature 4.1) and to calculate upper limits on the $o-$ $_2$ $^+$ column 4.2), the total $_2$ $^+$ column 4.3), the total midplane ion column 4.4), the midplane ionization fraction 4.5), and the ionization fraction in higher disk layers 4.6)."
 In 855. we compare the results with predictions rom a generic disk chemistry moclel.," In 5, we compare the results with predictions from a generic disk chemistry model."
 DAL Tau was observed with the Submillimeter Array on top of Mauna Ixea. Hawaii. on 2010 October 7.," DM Tau was observed with the Submillimeter Array on top of Mauna Kea, Hawaii, on 2010 October 7."
 Six of the array antennas were available lor these observations. arrauged in a compact couliguration spanning baselines of 16 to 77 meters.," Six of the array antennas were available for these observations, arranged in a compact configuration spanning baselines of 16 to 77 meters."
" Using the dual receiver mode. we simultaneously targeted the orfho-H»D — 144— liue at 372.1213 GHz aud the ΝΟΗ 24-3 at 372.6725 GHz with the ""high frequency” receiver. aud the 315.796 GHz CO J—3-2 liue with the “low frequeucy"" receiver."," Using the dual receiver mode, we simultaneously targeted the $ortho$ $_2$ $^+$ $1_{1,0}-1_{1,1}$ line at 372.4213 GHz and the $_2$ $^+$ J=4–3 at 372.6725 GHz with the “high frequency” receiver, and the 345.796 GHz CO J=3–2 line with the “low frequency” receiver."
 Table 1 sunuuarizes the spectral line setup., Table 1 summarizes the spectral line setup.
 At these frequencies. the synthesized beam sizes were 272xUF and 274x2/0. respectively.," At these frequencies, the synthesized beam sizes were $2\farcs2\times1\farcs7$ and $2\farcs4\times2\farcs0$, respectively."
" The observing conditions were generally very good. with atmospheric opacity τουσ, measured atthe nearby Caltech Subinillimeter Observatory. aud stable atinospheric phase."," The observing conditions were generally very good, with atmospheric opacity $\tau_{\rm 225{\:GHz}}\sim0.04$ measured atthe nearby Caltech Submillimeter Observatory, and stable atmospheric phase."
 This opacity corresponds to a zenith transmission of 50% at the HeD line lrequency., This opacity corresponds to a zenith transmission of $\sim50$ at the $_2$ $^+$ line frequency.
 Short observations of the two calibrators 3C120 aud JOLLO+1L13 were interleaved with the observations of, Short observations of the two calibrators 3C120 and J0449+113 were interleaved with the observations of
measuring the width of emission lines of à comparison spectrum after the wavelength calibration) was 3.10 (CENIIND.,measuring the width of emission lines of a comparison spectrum after the wavelength calibration) was $3.10$ (FWHM).
 This corresponds to an instrumental dispersion ol TingcSÜ aat 5170A.," This corresponds to an instrumental dispersion of $\sigma_{\it
 inst}\approx80$ at 5170."
. Phe spatial scale along the slit was 0.275 arcesec pixel, The spatial scale along the slit was $0.275$ arcsec $^{-1}$.
 For cach galaxy. we took a spectrum with the slit along the major axis. as well as two or more ollsct spectra with the slit parallel to the major axis on each side of the center.," For each galaxy, we took a spectrum with the slit along the major axis, as well as two or more offset spectra with the slit parallel to the major axis on each side of the center."
 The major axis position angles were chosen according to the values obtained from the surface photometry. as described in Sect.," The major axis position angles were chosen according to the values obtained from the surface photometry, as described in Sect."
 4.1 and listed in Table 3...," \ref{sec:photometrical_parameters}
 and listed in Table \ref{tab:measured_pa}."
 The location of the offset spectra was chosen. using the photometry. to give larec .V (and therefore large Y) and high SYN.," The location of the offset spectra was chosen, using the photometry, to give large $\pin$ (and therefore large $\kin$ ) and high $S/N$."
 Integration time of the galaxy spectra was split into exposures of 2700/3600 sec to deal with cosmic ravs., Integration time of the galaxy spectra was split into exposures of 2700/3600 sec to deal with cosmic rays.
 A log of the observations is given in Table 5.., A log of the observations is given in Table \ref{tab:log_spec}.
 In each observing run we obtained spectra of twelve giant stars with spectral type ranging from late-G to earlv-Ix. selected. from Faber et al. (," In each observing run we obtained spectra of twelve giant stars with spectral type ranging from late-G to early-K, selected from Faber et al. ("
1985). lo use as templates in measuring stellar kinematies.,"1985), to use as templates in measuring stellar kinematics."
 Are lamp spectra were taken before and/or alter every object exposure to allow an accurate wavelength calibration., Arc lamp spectra were taken before and/or after every object exposure to allow an accurate wavelength calibration.
 The range of the secing FWHLIAIL during the observing runs was 0.7. 1.6 arcsec in February 2001 (as measured by fitting a two-dimensional Gaussian to the PNG guide star). 0.5 2.2 aresec in May," The range of the seeing FWHM during the observing runs was $0.7$ $1.6$ arcsec in February 2001 (as measured by fitting a two-dimensional Gaussian to the TNG guide star), $0.5$ $2.2$ arcsec in May"
The binary [Traction is an essential component in the formation and evolution ofdynamically active svstems. like globular clusters (GCs).,"The binary fraction is an essential component in the formation and evolution of dynamically active systems, like globular clusters (GCs)."
 In such dense environments. where stellar eravitational interactions are very [requent. binaries can exert a sienilicant influence on both the dynamical evolution of the svstem and the properties of its stellar populations.," In such dense environments, where stellar gravitational interactions are very frequent, binaries can exert a significant influence on both the dynamical evolution of the system and the properties of its stellar populations."
 Being. on average. more massive than the other stars. binaries tend to sink into the hiehlv crowded cluster centers. because of equipartiüon.," Being, on average, more massive than the other stars, binaries tend to sink into the highly crowded cluster centers, because of equipartition."
 The characteristic timescale of this process (the relaxation time) depends on the cluster structure and can be even longer (han a IIubble time in the outskirts., The characteristic timescale of this process (the relaxation time) depends on the cluster structure and can be even longer than a Hubble time in the outskirts.
 Hence. in the outer regions of GC's we essentially expect {ο observe primordial binary svstenis. ie.. binaries created as part of the star formation process and evolving undisturbed.," Hence, in the outer regions of GCs we essentially expect to observe primordial binary systems, i.e., binaries created as part of the star formation process and evolving undisturbed."
 In the cluster core. on the other hand. a variety of dynamical processes (exchange interactions. three-body encounters. tidal captures. etc.)," In the cluster core, on the other hand, a variety of dynamical processes (exchange interactions, three-body encounters, tidal captures, etc.)"
 can take place. wilh competing effects on the binary population: binaries can be destroyed. created or just moclified (e.g.Iliοἱal.1992)... with relative efficiencies that still are a matter of debate in the literature (e.g.Ivanovaetal.2005:IIurlev2007:Sollima2008: 2009).," can take place, with competing effects on the binary population: binaries can be destroyed, created or just modified \citep[e.g.][]{hut92}, with relative efficiencies that still are a matter of debate in the literature \citep[e.g.][]{ivanova05, hurl07,
 sol08, freg09}."
. In general. however. (he gravitational encounters occurring in the core tend to make (he binaries harder (more tightly bound). (hus providing a central energy source able to slow down (he cluster core collapse (Goodman&πας1959).," In general, however, the gravitational encounters occurring in the core tend to make the binaries harder (more tightly bound), thus providing a central energy source able to slow down the cluster core collapse \citep{goodhut89}."
. Following the N-body simulations of Gilletal.(2003).. (his energy source could also suppress the mass segregation process. will a detectable effect on the radial behavior of the mass function of main-sequence (AIS) stars.," Following the N-body simulations of \citet{gill08}, this energy source could also suppress the mass segregation process, with a detectable effect on the radial behavior of the mass function of main-sequence (MS) stars."
 Since the same ellect could alternatively be due to a central intermeciate-mass black hole (IMDIIseealsoPasquatoοἱal.2009:Deccariet2010.herefaterDIO).. if we can measure the fraction of binaries. (hen we can sav whether or not we need an IMDII to explain the low level of mass segregation that has been observed.," Since the same effect could alternatively be due to a central intermediate-mass black hole \citep[IMBH; see also][herefater B10]{pasqua09, bec10}, if we can measure the fraction of binaries, then we can say whether or not we need an IMBH to explain the low level of mass segregation that has been observed."
 Hence. the empirical estimate of the binary fraction in a sample of GC's representative of different environments is a prime ingredient [or dvnamical models. which help us understixd (he internal cluster dynamics.," Hence, the empirical estimate of the binary fraction in a sample of GCs representative of different environments is a prime ingredient for dynamical models, which help us understand the internal cluster dynamics."
 The knowledge of the binary fraction is also crucial for understanding the properties of puzzling objects like blue stragelers. millisecond pulsars and cataclysmic variables. which are all thought to be the by-products of binary evolution (e.g..MeCrea1964:Romanietal.1987:Ferraroetal.2001:Leieh2011.andreferences therein).," The knowledge of the binary fraction is also crucial for understanding the properties of puzzling objects like blue stragglers, millisecond pulsars and cataclysmic variables, which are all thought to be the by-products of binary evolution \citep[e.g.,][and references
 therein]{mcrea64,romani87,fe01,leigh11}."
. In particular. (he analvsis of the bimodal radial distribution of blue stragglers observed in a number of GCs (e.g.FerraroDalessandroetal.2008) suggests that a non-negligible Traction of these stars 1s generated by primordial binaries. which still orbit in isolation in (he cluster outskirts and produce (he observed rising branch of the distribution (Mapelli et al.," In particular, the analysis of the bimodal radial distribution of blue stragglers observed in a number of GCs \citep[e.g., Ferraro et al. 1997, 2004;][]{ema_6388}
 suggests that a non-negligible fraction of these stars is generated by primordial binaries, which still orbit in isolation in the cluster outskirts and produce the observed rising branch of the distribution (Mapelli et al."
 2004. 2006: Lanzoni et al.," 2004, 2006; Lanzoni et al."
 2007a.b).," 2007a,b)."
 The interpretation of the double blue strageler sequence recently discovered in (he core of M30. also requires a significant Iraction of primordial binaries, The interpretation of the double blue straggler sequence recently discovered in the core of M30 also requires a significant fraction of primordial binaries
With this assumption the equilibrium couditiou becomes gj. where a prime denotes au ;c-derivative.,"With this assumption the equilibrium condition becomes _0 = x +, where a prime denotes an $x$ -derivative."
 One can regard the background flow as providing au effective shear rate that varies with or. in which case vj=—[q6s)dsOg.," One can regard the background flow as providing an effective shear rate that varies with $x$, in which case $\bv_0 = -\int^x \qe(s) ds \, \Omega
\ey$."
 Localized ou this background flow we will cousider a shearing wave witli ZJiLx1., Localized on this background flow we will consider a shearing wave with $k_y L \gg 1$.
 That is. we will cousicer nouaxisymanetrie short-waveleugth Eulerian perturbations with spacetime cdepeucdence o(l)expti[πο...Hp ο constants aud t AREE -- ]t may not be unmediately. obvious hat this is a valid expansion since the sliwaves sit on top of a radially-varyiug background. (see Toomre1969 for a discussion of waves in a background)," That is, we will consider nonaxisymmetric short-wavelength Eulerian perturbations with spacetime dependence $\delta(t) \exp(i\int^x \tilde{k}_x(t,s) ds + ik_y y + ik_z z)$, where $k_y$ and $k_z$ are constants and _x(t,x) + k_y t. It may not be immediately obvious that this is a valid expansion since the shwaves sit on top of a radially-varying background (see \citealt{toom69} for a discussion of waves in a slowly-varying background)."
 But his is an orclitary WEB expansiou iu disguise., But this is an ordinary WKB expansion in disguise.
 To see this. oue ueed ouly rauslorin to cOMOVILD>oO coordinates oSn g(s)dsQf. f=? (this procedure may be more familiar in a cosinological context: as Balbussí(1988 has pointed ott. this is possible for auy flow in which the velocities depend linearly on tle spatial οςordinates).," To see this, one need only transform to “comoving” coordinates $x' = x$, $y'
= y + \int^x \qe(s) ds \Omega t$ , $t' = t$ (this procedure may be more familiar in a cosmological context; as \cite{bal88} has pointed out, this is possible for any flow in which the velocities depend linearly on the spatial coordinates)."
 Ii this frame the wavevect«y eiven above is transornmecd to a time-inde»endent waverrector., In this frame the time-dependent wavevector given above is transformed to a time-independent wavevector.
" The price paid for tlis is that ὃν—Ow+gOlO,. so new explicit tirje clepenclehees appear oi the right haud side of the perturbed eqalious of motion. auc the perturbed variabes uo longer have time cepeucdence expel’)."," The price paid for this is that $\partial_x \rightarrow
\partial_{x'} + \qe \Omega t \partial_{y'}$, so new explicit time dependences appear on the right hand side of the perturbed equations of motion, and the perturbed variables no longer have time dependence $\exp(i\omega t')$."
 Instead. we must solve an ODE for δή.," Instead, we must solve an ODE for $\delta(t')$."
" The y ependence ca1 be decomposed as explifyy"").", The $y'$ dependence can be decomposed as $\exp(ik_y y')$.
 The wr ependeuce cau be treated via V‘WB. since tlie perturbation may be assumed to have the form Wer’.ef)explik’τα].," The $x'$ dependence can be treated via WKB, since the perturbation may be assumed to have the form $W(\epsilon x', \epsilon t')\, \exp(i\bk' 
\cdot \bx')$."
" This “neary cdiagonalizes 1le operator Q,».", This “nearly diagonalizes” the operator $\partial_{x'}$.
" Γιs we are cousidering the evoltion of a wavepacket in comovirig coordinates— a ""sInvavepacket"".", Thus we are considering the evolution of a wavepacket in comoving coordinates— a “shwavepacket”.
 For this procedure to be valid two e«ocditions 1just be met., For this procedure to be valid two conditions must be met.
" Firs the usual WISB coudition iust apply. kyLz»1.Secoud. the parameters of the flow that are ""seen by the shwavepacket must Change little ou the characteristic timescale for variation of 0(/). which isϱ1 for the incompressive"," First the usual WKB condition must apply, $k_y L \gg 1$.Second, the parameters of the flow that are “seen” by the shwavepacket must change little on the characteristic timescale for variation of $\delta(t)$ , which is$\Omega^{-1}$ for the incompressive"
"R,—6 but this is inconsistent with other measured values, and with the value of R, found to reconcile the FUV, Ha and [Ori] measurements in our analysis.","$R_{v}=6$ but this is inconsistent with other measured values, and with the value of $R_{v}$ found to reconcile the FUV, $\alpha$ and ] measurements in our analysis."
 The fact that the NUV SFRs are overestimated more than the FUV SFRs in all MW obscuration curves and not in the Calzetti (2001) curve suggest that this effect is due to the ffeature being included in the extinction curve., The fact that the NUV SFRs are overestimated more than the FUV SFRs in all MW obscuration curves and not in the Calzetti (2001) curve suggest that this effect is due to the feature being included in the extinction curve.
 The NUV being the only SFR indicator that otherwise does not agree with the other indicators lends weight to this argument as none of those wavelengths fall within the range of the ffeature., The NUV being the only SFR indicator that otherwise does not agree with the other indicators lends weight to this argument as none of those wavelengths fall within the range of the feature.
" We show that the Ho and NUV derived SFRs agree very well if the ffeature is eliminated in an FD05 curve with an R, value of 4.5.", We show that the $\alpha$ and NUV derived SFRs agree very well if the feature is eliminated in an FD05 curve with an $R_{v}$ value of 4.5.
 This provides us with a consistent extinction correction curve for all the SFR indicators., This provides us with a consistent extinction correction curve for all the SFR indicators.
 The role of the ffeature will require further analysis in future work., The role of the feature will require further analysis in future work.
 The use of the 9 parameter in estimating obscuration corrections for our sample of largely quiescent galaxies shows that intrinsic SFRs will be overestimated., The use of the $\beta$ parameter in estimating obscuration corrections for our sample of largely quiescent galaxies shows that intrinsic SFRs will be overestimated.
 Kong et al. (, Kong et al. (
2004) attribute this to a secondary effect based on SFR histories.,2004) attribute this to a secondary effect based on SFR histories.
" We applied galaxy evolutionary synthesis models to the comparison between Ha and FUV derived SFRs, and there was good agreement for all SFRs."," We applied galaxy evolutionary synthesis models to the comparison between $\alpha$ and FUV derived SFRs, and there was good agreement for all SFRs."
" The evolutionary paths were compared according to galaxy mass, characteristic decay time and IMF (assuming an exponentially decaying SFR)."," The evolutionary paths were compared according to galaxy mass, characteristic decay time and IMF (assuming an exponentially decaying SFR)."
co-operation with the Centre National dEtudes Spatiales of France and the Korean Ministry of Science and Technology.,co-operation with the Centre National d'Etudes Spatiales of France and the Korean Ministry of Science and Technology.
 Finally. the WigeleZ survey would. not. be possible without the dedicated work of the stall of the Xnglo-Australian Observatory in the development and support of the AXOmoeea spectrograph. and the running of the AAT.," Finally, the WiggleZ survey would not be possible without the dedicated work of the staff of the Anglo-Australian Observatory in the development and support of the AAOmega spectrograph, and the running of the AAT."
he break in the observed data.,the break in the observed data.
 Àn alternative explanation for the break. as will be shown in this paper. is that after acceleration the ugh energy electrons might have suffered signiticant losses within he sources itself before they are released into the Galaxy.," An alternative explanation for the break, as will be shown in this paper, is that after acceleration the high energy electrons might have suffered significant losses within the sources itself before they are released into the Galaxy."
" On the other hand. measurements of CR secondary-to-primary (s/p) ratios like the boron-to-carbon (B/C) up to £6=100 GeV/n by several independent experiments (see the experiments listed in Stephens & Streitmatter 1998) show that above 1. GeV/n. the ratio decreases with energy as ££"" with the index 9.~6 implying an energy dependent CR propagation path length in the Galaxy."," On the other hand, measurements of CR secondary-to-primary (s/p) ratios like the boron-to-carbon (B/C) up to $E\approx 100$ GeV/n by several independent experiments (see the experiments listed in Stephens $\&$ Streitmatter 1998) show that above $1$ GeV/n, the ratio decreases with energy as $\sim E^{-\delta}$ with the index $\delta\approx 0.6$ implying an energy dependent CR propagation path length in the Galaxy."
 However. these measurements seem to indicate a yossible hardening at higher energies (see also Swordy et al.," However, these measurements seem to indicate a possible hardening at higher energies (see also Swordy et al."
 1990)., 1990).
 Recent data from the CREAM balloon-borne experiment also show Hattening at energies around 1 TeV/n CAhn et al., Recent data from the CREAM balloon-borne experiment also show flattening at energies around $1$ TeV/n (Ahn et al.
 2008)., 2008).
 One way to explain this is to assume that CRs traverse some minimum amount of matter of approximately 0.3 & > in the Galaxy ¢Ave et al., One way to explain this is to assume that CRs traverse some minimum amount of matter of approximately $0.3$ g $^{-2}$ in the Galaxy (Ave et al.
 2009)., 2009).
 This implies an energy dependent escape path length of CRs up to some certain energy and beyond that energy. the path length becomes constant in energy.," This implies an energy dependent escape path length of CRs up to some certain energy and beyond that energy, the path length becomes constant in energy."
 This view is. in fact. supported by the observed CR anisotropy which remain almost constant above around 100 GeV energies.," This view is, in fact, supported by the observed CR anisotropy which remain almost constant above around $100$ GeV energies."
 But such model need to assume a break in the source spectrum in order to explain the observed CR spectrum which closely follow a pure power-law without any break up to the knee., But such model need to assume a break in the source spectrum in order to explain the observed CR spectrum which closely follow a pure power-law without any break up to the knee.
 Another possible explanation is that some amount of secondaries might be produced inside the sources due to the nuclear interaction of the primaries with the matter (Berezhko et al., Another possible explanation is that some amount of secondaries might be produced inside the sources due to the nuclear interaction of the primaries with the matter (Berezhko et al.
 2003)., 2003).
 Under such models. the CR anisotropy is expected to increase with energy as Z which does not agree well with the observed data.," Under such models, the CR anisotropy is expected to increase with energy as $E^\delta$ which does not agree well with the observed data."
 However. it is quite possible that at higher energies the anisotropy might not be determined by the global diffusion leakage of CRs from the Galaxy.," However, it is quite possible that at higher energies the anisotropy might not be determined by the global diffusion leakage of CRs from the Galaxy."
 Rather. it might be the effects of high energy CRs coming from nearby sources because of their faster diffusion in the Galaxy (Ptuskin et al.," Rather, it might be the effects of high energy CRs coming from nearby sources because of their faster diffusion in the Galaxy (Ptuskin et al."
 2005. Thoudam 2007).," 2005, Thoudam 2007)."
 One more possible explanation which we will not discuss here is the possible reacceleration of some fraction of the background CRs by strong SNR shocks while propagating in the Galaxy (Wandel et al., One more possible explanation which we will not discuss here is the possible reacceleration of some fraction of the background CRs by strong SNR shocks while propagating in the Galaxy (Wandel et al.
 1987. Berezhko et al.," 1987, Berezhko et al."
 2003)., 2003).
 In this paper. we present one possible explanation for the observed break in the electron spectrum at around | TeV and its possible correlation with the flattening in the B/C ratio at energies above around. 100 GeV/n. In our model. we assume that CRs are accelerated by SNR shock waves by diffusive shock acceleration (DSA) mechanism (Bell 1978. Blandford & Eichler 1987).," In this paper, we present one possible explanation for the observed break in the electron spectrum at around $1$ TeV and its possible correlation with the flattening in the B/C ratio at energies above around $100$ GeV/n. In our model, we assume that CRs are accelerated by SNR shock waves by diffusive shock acceleration (DSA) mechanism (Bell 1978, Blandford $\&$ Eichler 1987)."
 Under DSA theory. suprathermal particles in the tail of the Maxwell-Boltzmann distribution present in the interstellar medium (ISM) are infected into the supernova shock front.," Under DSA theory, suprathermal particles in the tail of the Maxwell-Boltzmann distribution present in the interstellar medium (ISM) are injected into the supernova shock front."
 They are then reflected back and forth several times across the shock by the magnetic turbulence generated on either side of the shock and in each crossing. the particles gain energy by tirst order Fermi acceleration.," They are then reflected back and forth several times across the shock by the magnetic turbulence generated on either side of the shock and in each crossing, the particles gain energy by first order Fermi acceleration."
 In a simple planar shocks model. such a mechanism naturally leads to a power-law spectrum of the form f°U with P=2 for strong shocks which is in good agreement with radio observations of several SNRs (Green 2009).," In a simple planar shocks model, such a mechanism naturally leads to a power-law spectrum of the form $E^{-\Gamma}$ with $\Gamma=2$ for strong shocks which is in good agreement with radio observations of several SNRs (Green 2009)."
 One important consideration of the standard DSA theory is that most of the particles do not escape upstream of the shock. they can only escape downstream mainly due to advection by the bulk flow and remain confined within the remnant due to the strong magnetic turbulence generated by the CRs themselves.," One important consideration of the standard DSA theory is that most of the particles do not escape upstream of the shock, they can only escape downstream mainly due to advection by the bulk flow and remain confined within the remnant due to the strong magnetic turbulence generated by the CRs themselves."
 Though it is still not fully understood. it is generally considered that the particles remain inside the remnant until the shock remains strong.," Though it is still not fully understood, it is generally considered that the particles remain inside the remnant until the shock remains strong."
 During the continement period. the high energy electrons may suffer from radiative energy losses while the nuclear components suffer from nuclear fragmentations.," During the confinement period, the high energy electrons may suffer from radiative energy losses while the nuclear components suffer from nuclear fragmentations."
 At later stages when the shock slows down and does not efficiently accelerate the CRs. the magnetic turbulence level goes down and the particles can no longer be contined effectively within the remnant.," At later stages when the shock slows down and does not efficiently accelerate the CRs, the magnetic turbulence level goes down and the particles can no longer be confined effectively within the remnant."
 At such stage. all the particles present inside the remnant escape and they are injected into the ISM.," At such stage, all the particles present inside the remnant escape and they are injected into the ISM."
 Fora typical ISM density of 1 Hem. . this happens at around 10° vr after the supernova explosion (Berezhko & ΝΟΚ 2000).," For a typical ISM density of $1$ H $^{-3}$, this happens at around $10^5$ yr after the supernova explosion (Berezhko $\&$ Völlk 2000)."
 The scenario described above might be true mostly for the lower energy particles because it is quite possible that some of the highest energy particles may escape the remnant already at the start of the sedov phase itself due to their faster diffusion., The scenario described above might be true mostly for the lower energy particles because it is quite possible that some of the highest energy particles may escape the remnant already at the start of the sedov phase itself due to their faster diffusion.
 Such an energy dependent escape of particles has been discussed in a number of literatures (Ptuskin & Zirakashvili 2005. Gabici. Aharonian & Casanova 2009).," Such an energy dependent escape of particles has been discussed in a number of literatures (Ptuskin $\&$ Zirakashvili 2005, Gabici, Aharonian $\&$ Casanova 2009)."
 Recent high energy -ray observations of SNRs associated with molecular clouds also suggest that some of the high energy particles might have already escaped the remnant at much early times ¢(Aharonian et al., Recent high energy $\gamma$ -ray observations of SNRs associated with molecular clouds also suggest that some of the high energy particles might have already escaped the remnant at much early times (Aharonian et al.
 2007. Aharonian et al.," 2007, Aharonian et al."
 20082)., 2008a).
 But. as mentioned above. it is still not clear how and when exactly the particles escape the remnant.," But, as mentioned above, it is still not clear how and when exactly the particles escape the remnant."
 For simplicity. we consider an energy independent scenario for our present work and assume that all the particles are released into the ISM at some characteristic time after the supernova explosion.," For simplicity, we consider an energy independent scenario for our present work and assume that all the particles are released into the ISM at some characteristic time after the supernova explosion."
 Once the CRs are released into the ISM. we assume that they undergo diffusive propagation in the Galaxy where the electrons again suffer from radiative losses and the nuclei from nuclear spallation with the interstellar matter.," Once the CRs are released into the ISM, we assume that they undergo diffusive propagation in the Galaxy where the electrons again suffer from radiative losses and the nuclei from nuclear spallation with the interstellar matter."
 Therefore. in our model the CR spectra that we tinally observe in the Solar System are moditied from their original source spectrum (the one generated by the SNR shocks) due to the various interactions or energy loss processes occurring not only during their propagation in the Galaxy. but also within the SNRs itself.," Therefore, in our model the CR spectra that we finally observe in the Solar System are modified from their original source spectrum (the one generated by the SNR shocks) due to the various interactions or energy loss processes occurring not only during their propagation in the Galaxy, but also within the SNRs itself."
 Our paper is planned as follows., Our paper is planned as follows.
 In section 2. we present our calculations for the CR spectra inside the SNRs and in section 3. we calculate the spectra in the Galaxy.," In section 2, we present our calculations for the CR spectra inside the SNRs and in section 3, we calculate the spectra in the Galaxy."
 Then. in section 4. we compare our results with the observed data and also give a short discussions about our results.," Then, in section 4, we compare our results with the observed data and also give a short discussions about our results."
 Finally in section 5. we give a brief conclusion of our study.," Finally in section 5, we give a brief conclusion of our study."
 We assume that CR acceleration by SNR shock waves begins at the time of the supernova explosion itself and the CRs then escape downstream of the shock and remain confined within the remnant., We assume that CR acceleration by SNR shock waves begins at the time of the supernova explosion itself and the CRs then escape downstream of the shock and remain confined within the remnant.
 We further assume that at some later stage of the SNR evolution characterize by age /=7 when the shock slows down. particle acceleration stops and also all the CRs present inside the remnant are released into the ISM.," We further assume that at some later stage of the SNR evolution characterize by age $t=T$ when the shock slows down, particle acceleration stops and also all the CRs present inside the remnant are released into the ISM."
 In our study. we do not consider the expansion of the remnant and hence neglect the effect of adiabatic energy losses and other related effects.," In our study, we do not consider the expansion of the remnant and hence neglect the effect of adiabatic energy losses and other related effects."
 Detailed study including the evolution of the remnant. the weakening of the shock with time and the possible energy dependent escape of CRs will be presented elsewhere.," Detailed study including the evolution of the remnant, the weakening of the shock with time and the possible energy dependent escape of CRs will be presented elsewhere."
 Under the assumption that the acceleration time of CRs is much less than their continement time within the SNR. the CR electron spectrum inside an SNR can be described by the following equation.," Under the assumption that the acceleration time of CRs is much less than their confinement time within the SNR, the CR electron spectrum inside an SNR can be described by the following equation,"
Reverberation mapping of active galactic nuclei (AGNs) exploits the light-travel-time delay between continuum variations and the response of broad emission lines to measure the sizes of AGN broad-line regions (BLRs).,Reverberation mapping of active galactic nuclei (AGNs) exploits the light-travel-time delay between continuum variations and the response of broad emission lines to measure the sizes of AGN broad-line regions (BLRs).
 This size can then be used. together with the estimated velocities of the BLR gas to derive masses for the central black holes (See Netzer Peterson 1996. for a review.)," This size can then be used, together with the estimated velocities of the BLR gas to derive masses for the central black holes (See Netzer Peterson 1996, for a review.)"
 The size-luminosity and mass luminosity relations of AGNs may shed new light on understanding these objects and the connections between the black holes in AGNs and those found in normal nearby galaxies., The size-luminosity and mass luminosity relations of AGNs may shed new light on understanding these objects and the connections between the black holes in AGNs and those found in normal nearby galaxies.
 Kaspi et al. (, Kaspi et al. (
2000) measured reverberation sizes. Razr. for 17 PG quasars between 1991—1998 using the Wise Observatory Im telescope and the Steward Observatory 2.3m telescope. with acombination of CCD spectrophotometry and photometry.,"2000) measured reverberation sizes, $R_{BLR}$, for 17 PG quasars between 1991–1998 using the Wise Observatory 1m telescope and the Steward Observatory 2.3m telescope, with a combination of CCD spectrophotometry and photometry."
 The time averaged flux for each quasar was derived from 20-70 epochs per quasar., The time averaged flux for each quasar was derived from 20-70 epochs per quasar.
 The rest-frame 5100 ccontinuum flux densities were measured from the ~10 rresolution spectra while minding potential systematics such as broad-emission-line wings and atmospheric absorptions in the higher-redshift objects., The rest-frame 5100 continuum flux densities were measured from the $\sim 10$ resolution spectra while minding potential systematics such as broad-emission-line wings and atmospheric absorptions in the higher-redshift objects.
 Following Galactic extinction corrections (small for most of the objects) the absolute fluxes were used to calculate 5100 lluminosities. ALA(5100).," Following Galactic extinction corrections (small for most of the objects) the absolute fluxes were used to calculate 5100 luminosities, $\lambda L_{\lambda}(5100 {\rm \AA})$."
 The sizes and luminosities of Seyfert galaxies with reverberation measurements were compiled from previously published works., The sizes and luminosities of Seyfert galaxies with reverberation measurements were compiled from previously published works.
" A regression analysis taking into account the errors in both Rare and ALA(S100) showed that Rarex[ALAS100)|S, ", A regression analysis taking into account the errors in both $R_{BLR}$ and $\lambda L_{\lambda}(5100{\rm \AA})$ showed that $R_{BLR} \propto [\lambda L_{\lambda}(5100{\rm \AA})]^{0.70\pm 0.03}$.
"""This was a surprising result. since it has long been speculated that the dependence would be to the power 0.5."," This was a surprising result, since it has long been speculated that the dependence would be to the power 0.5."
 Such a scaling would lead to an tonization parameter (ratio of ionizing photon density to electron density) at the surface of the BLR clouds that is independent of luminosity. and would explain the similarity of AGN spectra over many orders of magnitude in luminosity.," Such a scaling would lead to an ionization parameter (ratio of ionizing photon density to electron density) at the surface of the BLR clouds that is independent of luminosity, and would explain the similarity of AGN spectra over many orders of magnitude in luminosity."
 In a recent paper. MeLure Jarvis (2002) examine to what degree UV observables. namely luminosities and Mg II line velocities. can serve the purpose of the optical observables — 5100 lluminosities and H.>j widths — used to date in AGN black hole estimates.," In a recent paper, McLure Jarvis (2002) examine to what degree UV observables, namely luminosities and Mg II line velocities, can serve the purpose of the optical observables – 5100 luminosities and $\beta$ widths – used to date in AGN black hole estimates."
 In the course of their work. they re-analyze the data presented by Kaspi et al. (," In the course of their work, they re-analyze the data presented by Kaspi et al. ("
2000) and conclude that actually Rpigκ[ALA(GI00)/79*97* is the best fit. contrary to the result of Kaspi et al..,"2000) and conclude that actually $R_{BLR} \propto [\lambda L_{\lambda}(5100{\rm \AA})]^{0.50\pm0.02}$ is the best fit, contrary to the result of Kaspi et al.,"
 and consistent with the expectations for a constant ionization parameter., and consistent with the expectations for a constant ionization parameter.
 Vestergaard (2002) has carried out an analysis along similar lines., Vestergaard (2002) has carried out an analysis along similar lines.
" After studying the various systematics that can affect the slope determination. she concludes that the best estimate (her equation AS) is Rarex[ALAS100)""OF"" consistent with the result of Kaspi et al. ("," After studying the various systematics that can affect the slope determination, she concludes that the best estimate (her equation A5) is $R_{BLR} \propto [\lambda L_{\lambda}(5100)]^{0.66\pm0.09}$, consistent with the result of Kaspi et al. ("
2000).,2000).
 The statement by McLure Jarvis (2002). that Vestergaard (2002) found a slope consistent with 0.5 but decided to adopt a slope of 0.7 anyway. is incorrect.," The statement by McLure Jarvis (2002), that Vestergaard (2002) found a slope consistent with 0.5 but decided to adopt a slope of 0.7 anyway, is incorrect."
 In this Note. I investigate the source of the discrepancy between Kaspi et al.," In this Note, I investigate the source of the discrepancy between Kaspi et al."
 and MeLure Jarvis., and McLure Jarvis.
 | show that it arises. first. due to the adoption by MeLure Jarvis of old. single-epoch. and systematically higher fluxes. but only for the high luminosity part of the sample: and second. due to incorrect conversion from flux to luminosity.," I show that it arises, first, due to the adoption by McLure Jarvis of old, single-epoch, and systematically higher fluxes, but only for the high luminosity part of the sample; and second, due to incorrect conversion from flux to luminosity."
 This is confirmed and corrected m à revised version of their paper (R. McLure. private communication).," This is confirmed and corrected in a revised version of their paper (R. McLure, private communication)."
 MeLure Jarvis chose to base their luminosities for the PG quasars in the Kaspr et al. (, McLure Jarvis chose to base their luminosities for the PG quasars in the Kaspi et al. (
2000) sample on the measurements by Neugebauer et al. (,2000) sample on the measurements by Neugebauer et al. (
1987).,1987).
 These measurements were obtained in. 1980 with a multichannel spectrophotometer mounted on the Palomar 5m telescope., These measurements were obtained in 1980 with a multichannel spectrophotometer mounted on the Palomar 5m telescope.
 The measurements had coarse resolution of ~300A.. making the avoidance of emission and absorption lines in the measurement more difficult.," The measurements had coarse resolution of $\sim
300$, making the avoidance of emission and absorption lines in the measurement more difficult."
 In the numbers tabulated by Neugebauer et al..," In the numbers tabulated by Neugebauer et al.,"
 interpolation is needed. for most of the quasars. to obtain the flux at rest wavelength. 5100A.," interpolation is needed, for most of the quasars, to obtain the flux at rest wavelength 5100."
. As opposed to the measurements of Kaspi et al. (, As opposed to the measurements of Kaspi et al. (
2000). only one epoch per object exists. inducing scatter due to the time variability of quasars.,"2000), only one epoch per object exists, inducing scatter due to the time variability of quasars."
 More importantly. the Neugebauer et al. (," More importantly, the Neugebauer et al. ("
1987) measurements were obtained about 20 years before the Kaspi et al. (,1987) measurements were obtained about 20 years before the Kaspi et al. (
2000) size measurements. and it has been shown in Seyfert galaxies that BLR size and emission-line width change with time in individual objects (Peterson Wandel 2000).,"2000) size measurements, and it has been shown in Seyfert galaxies that BLR size and emission-line width change with time in individual objects (Peterson Wandel 2000)."
 Figure | compares the observer-frame quasar flux densities measured at wavelengths corresponding to rest-wavelength, Figure 1 compares the observer-frame quasar flux densities measured at wavelengths corresponding to rest-wavelength
nebular continuum contribution to the integrated spectrum.,nebular continuum contribution to the integrated spectrum.
 At present. we consider a basic mocdel as discussed. above and parameters could be adjusted to gain a better fit. for this star-forming region.," At present, we consider a basic model as discussed above and parameters could be adjusted to gain a better fit for this star-forming region."
 Secondly. it may be possible to include the elfects of material external to the svstem in addition to gas ancl dust local to the star-forming region which is modelled: usingCLOUDY.," Secondly, it may be possible to include the effects of material external to the system in addition to gas and dust local to the star-forming region which is modelled using."
.. Thirdlv. we only use a sparse grid of O-star models rather than a dense grid. of atmospheres and it may. be possible to improve upon this in future 2004).," Thirdly, we only use a sparse grid of O-star models rather than a dense grid of atmospheres and it may be possible to improve upon this in future ."
. Fourthlv. we have assumed an instantaneous burst.," Fourthly, we have assumed an instantaneous burst."
 A better fit might. be. achieved by allowing multiple star formation episodes with slightly dillerent ages., A better fit might be achieved by allowing multiple star formation episodes with slightly different ages.
 Fifthlv we have not considered. the effects of stellar rotation on the evolution or the spectra., Fifthly we have not considered the effects of stellar rotation on the evolution or the spectra.
 This will have two elfects on our results., This will have two effects on our results.
 It would. extend the niain-sequence life-time and would also rotationally broaden rw CIV. line. further. particularly if the stars rotate at Peloceities above 200kmis.1.," It would extend the main-sequence life-time and would also rotationally broaden the CIV line further, particularly if the stars rotate at velocities above $200 {\rm km \, s^{-1}}$."
 This would potentially improve 10 fit. between models ancl observed data., This would potentially improve the fit between models and observed data.
 We speculate ju study. of the CIV. profile could. potentially eive a way o evaluate the importance of rotation versus binarity in gaellar populations., We speculate that study of the CIV profile could potentially give a way to evaluate the importance of rotation versus binarity in stellar populations.
 Finally. some WNL stars have CIV in vmission.," Finally, some WNL stars have CIV in emission."
 We have not included any such CIV. emission in Ίο WNL stars that are not represented. by. the Potsdam model atmospheres., We have not included any such CIV emission in the WNL stars that are not represented by the Potsdam model atmospheres.
 Including an empirical correction for such emission in these models may. broaden the emission component of the CIV. line ancl simultaneously. decrease the absorption component., Including an empirical correction for such emission in these models may broaden the emission component of the CIV line and simultaneously decrease the absorption component.
 However we do not choose to introduce arbitrary line emission here., However we do not choose to introduce arbitrary line emission here.
 In this paper we set out to describe the construction of model stellar populations incorporating massive stars and massive stellar binaries. and then the svnthesis of spectra for this population.," In this paper we set out to describe the construction of model stellar populations incorporating massive stars and massive stellar binaries, and then the synthesis of spectra for this population."
 We have compared our mocel spectra to observations of comparable unresolved: stellar populations and found agreement is fair over a large range of wavelengths. from the UV to the near infra-red.," We have compared our model spectra to observations of comparable unresolved stellar populations and found agreement is fair over a large range of wavelengths, from the UV to the near infra-red."
 In general our population svnthesis including binary models predict that such svstems have less emission at long wavelengths around the L band., In general our population synthesis including binary models predict that such systems have less emission at long wavelengths around the I band.
 This is because. binary interactions remove the hydrogen envelope. of some red supergiants to form Wolf-Ravet stars., This is because binary interactions remove the hydrogen envelope of some red supergiants to form Wolf-Rayet stars.
 These WIR stars then leac to more blue colours in D-V and V-I broad-band colours and a larger. UV. flux., These WR stars then lead to more blue colours in B-V and V-I broad-band colours and a larger UV flux.
 The latter increases the timespan over which nebula emission is important to the evolution of stellar populations from 6 Myrs to 20 Alves., The latter increases the timespan over which nebula emission is important to the evolution of stellar populations from 6 Myrs to 20 Myrs.
 Another binary elfect is that it is more likely that strong WR. emission lines are observed since binary interactions tend to spread. WI emission features over a longer timespan., Another binary effect is that it is more likely that strong WR emission lines are observed since binary interactions tend to spread WR emission features over a longer timespan.
 For single stars the WI stars all exist over a short timespan so WI features are only present for a short. period. of time., For single stars the WR stars all exist over a short timespan so WR features are only present for a short period of time.
 This suggests that ages derived from Woll-Rayvet features to date may have been systematically uncerestimated., This suggests that ages derived from Wolf-Rayet features to date may have been systematically underestimated.
 Lt may be somewhat surprising that bv including binary evolution we find no great dillerence from. predictions fromsharbursit).. a code based on single-star evolution alone.," It may be somewhat surprising that by including binary evolution we find no great difference from predictions from, a code based on single-star evolution alone."
 However the stellar evolution models used in the model present in Section 3.1. are nearly. two decades: old001992)., However the stellar evolution models used in the model present in Section \ref{sec:WR_in_M31} are nearly two decades old.
. The important dillerence. between these models and those used. here is that the mass-Ioss rates have substantially decreased. for OB-stars ancl WR stars., The important difference between these models and those used here is that the mass-loss rates have substantially decreased for OB-stars and WR stars.
 Therefore. with incorrect. single-star niass-loss rates the older models reproduce the observed stellar population., Therefore with incorrect single-star mass-loss rates the older models reproduce the observed stellar population.
 Our model binary population therefore should broadly agree with this older code. but now we model the same magnituce of mass-loss as a combination of lower single-star mass-loss rates enhanced for some stars by binary interactions which is a physically reasonable model.," Our model binary population therefore should broadly agree with this older code, but now we model the same magnitude of mass-loss as a combination of lower single-star mass-loss rates enhanced for some stars by binary interactions which is a physically reasonable model."
 We demonstrate that our code produces a good fit to the observational data of local stellar populations in which massive stars are important., We demonstrate that our code produces a good fit to the observational data of local stellar populations in which massive stars are important.
 However. further refinements are possible and additional verification data on local star-forming regions would be welcome.," However, further refinements are possible and additional verification data on local star-forming regions would be welcome."
 Nonetheless. the stellar models and. synthesis code presented. here may now be used as a tool to study stellar populations in a range of different) observational domains and to derive their physical parameters. as demonstrated by their application to extragalactic star-forming regions Tol A and D. The authors would. like to thank the referee Jarle Brinchmann for his very helpful ancl constructive comments.," Nonetheless, the stellar models and synthesis code presented here may now be used as a tool to study stellar populations in a range of different observational domains and to derive their physical parameters, as demonstrated by their application to extragalactic star-forming regions Tol A and B. The authors would like to thank the referee Jarle Brinchmann for his very helpful and constructive comments."
 The authors also thank Malcolm: Bremer. Max Pettini. raul Crowther. Stephen Smartt. Nate Bastian and eorbert Langer for useful discussions.," The authors also thank Malcolm Bremer, Max Pettini, Paul Crowther, Stephen Smartt, Nate Bastian and Norbert Langer for useful discussions."
 ERS acknowledges xostedoctoral research. support. from the Ulx Science. and ‘Technology Facilities Council (SPEC)., ERS acknowledges postdoctoral research support from the UK Science and Technology Facilities Council (STFC).
" JJ began this work when he was supported by the award. ""Understanding the ives of massive stars from birth. to. supernovae” made under the European Leads of Research Councils and European Science Foundation EURYL Awards scheme which is supported by funds from the Participating Organisations of EURBRXI and the EC Sixth Framework Programme.", JJE began this work when he was supported by the award “Understanding the lives of massive stars from birth to supernovae” made under the European Heads of Research Councils and European Science Foundation EURYI Awards scheme which is supported by funds from the Participating Organisations of EURYI and the EC Sixth Framework Programme.
 JJ also acknowledges support [rom the Ul Science and Technology. Facilities. Council (SPEC) under the rolling theory grant for the Institute of Astronomy., JJE also acknowledges support from the UK Science and Technology Facilities Council (STFC) under the rolling theory grant for the Institute of Astronomy.
 rearest protoplanetary disks are located. with a few rotable exceptions such as the TW Ίνα association and a few Ierbig Ac stars. at distances in excess of 100 pc.," nearest protoplanetary disks are located, with a few notable exceptions such as the TW Hya association and a few Herbig Ae stars, at distances in excess of 100 pc."
 Resolved images of typical PFZs ust © obtained at a spatial resolution better than 0711., Resolved images of typical PFZs must be obtained at a spatial resolution better than 1.
 Further compounding the problems is that many of he main tracers of PFZs are found iu the infrared wavelength rauge μια). because the relevant gas cluperatures are in excess of LOOIWIS makine erounud-sedo observations challenging., Further compounding the problem is that many of the main tracers of PFZs are found in the infrared wavelength range $\mu$ m) because the relevant gas temperatures are in excess of $100$ K – making ground-based observations challenging.
 Cenerally. specialized lustrmuentation is needed to obtain the requisite spatial resolution.," Generally, specialized instrumentation is needed to obtain the requisite spatial resolution."
 In recent vears. eround-breaking progress lias )oen mde in infrared iaterferometry of the iuueriuost regions of protoplauctary disks from 1-2jan and near LOgan(?).. However. interferometric observations of gas have thus far been lianited to single baud or relatively low spectral resolving power aud to voung A aud B stars.," In recent years, ground-breaking progress has been made in infrared interferometry of the innermost regions of protoplanetary disks from $\mu$ m and near $\mu$ m. However, interferometric observations of gas have thus far been limited to single band or relatively low spectral resolving power and to young A and B stars."
 This has. for the most part. limited interferometry to dust. eas continuun and bydrogcu recombination lines(???7?).," This has, for the most part, limited interferometry to dust, gas continuum and hydrogen recombination lines."
. Our understanding of the structure aud dvnandes of inner protoplanetary disks aud PFZs remains hnuited., Our understanding of the structure and dynamics of inner protoplanetary disks and PFZs remains limited.
 While there appears to be a conuuon oeud result of disk evolution the formation of a planetary svsteni. coluplete with =sas elauts and perhaps simaller rocky planets. as well as. prestmably. plauctesimals. comets and zodiacal dust — the pathway is unclear.," While there appears to be a common end result of disk evolution – the formation of a planetary system, complete with gas giants and perhaps smaller rocky planets, as well as, presumably, planetesimals, comets and zodiacal dust – the pathway is unclear."
 It is known that disks exist and that they carry most of the angular momentum of a voung stellay svsteui as evidenced by resolved sub-uullimeter Πασάς of their outer regions(277)... and matching the momentum distribution of the solar syste aud other plauctary svstenis.," It is known that disks exist and that they carry most of the angular momentum of a young stellar system, as evidenced by resolved sub-millimeter imaging of their outer regions, and matching the momentum distribution of the solar system and other planetary systems."
 Their spectral, Their spectral
" RS Cauuu Venaticorum (RS CVn) variables ave described by Hall (1972: see also Biermann Tall 1976). as close binaries comprising a Ce or K-tvpoe subegiaut aie a F- or C-type star of huninosity class TVV. with the following special property (see also Zeilik 11979): ""Their light curves are characterized by long waves with amplitudes up to 0.2 mae which. if the nary js eclipsing. as a rule move towards smaller phase of the orbital ligl: curve."," RS Canum Venaticorum (RS CVn) variables are described by Hall (1972; see also Biermann Hall 1976), as close binaries comprising a G- or K-type subgiant and a F- or G-type star of luminosity class IV–V with the following special property (see also Zeilik 1979): Their light curves are characterized by long waves with amplitudes up to 0.2 mag which, if the binary is eclipsing, as a rule move towards smaller phase of the orbital light curve."
" This is typically interpreted as the effect of huge star spots on one hemisphere of the cool ""Olniponieli when rotating with a speed sliehtlv different from) the svuchronous rotation (e.g. Catalano 1983).", This is typically interpreted as the effect of large star spots on one hemisphere of the cool component when rotating with a speed slightly different from the synchronous rotation (e.g. Catalano 1983).
 Thus. these waves are assumed to be the beat between the orbital period aud the shehth out-of-plase differential rotation of the spotted star.," Thus, these waves are assumed to be the beat between the orbital period and the slightly out-of-phase differential rotation of the spotted star."
 The assumption of substantial chromospheric activity is supported bv additional observational features. such as stroug enudssion lines. ively flaring activity iu the optical aud other spectral regions and variable X-ray cussion (e.g. Walter 11980. Charles 1983).," The assumption of substantial chromospheric activity is supported by additional observational features, such as strong emission lines, lively flaring activity in the optical and other spectral regions and variable X-ray emission (e.g. Walter 1980, Charles 1983)."
 As part of our proegramune ofinvestigating the optical loug-term beliaviour of selected ROSAT sources (c.g. Richter 11995) we discovered a new variable star as the likely optical counterpart of RN J2000.8|1557., As part of our programme of investigating the optical long-term behaviour of selected ROSAT sources (e.g. Richter 1995) we discovered a new variable star as the likely optical counterpart of RX J2009.8+1557.
 This object which we called S 10917 Αα. varies in the B baud between Lb and E888.," This object which we called S 10947 Aql, varies in the B band between 4 and 8."
 It is identical to CSC 0161801655., It is identical to GSC 0161801655.
 The optical coordinates are: (equinox 2000.0). consistent with our measurement on the digitized Palomar Sky Survey plate (see marked object iu Fie.," The optical coordinates are: (equinox 2000.0), consistent with our measurement on the digitized Palomar Sky Survey plate (see marked object in Fig."
 1)., 1).
 RN J2009.8|1557 was detected during the ROSAT all-sky survey in 1990 at a count rate of O.01940.005 cts/s with a likelibood of detection of 12.1 (correspouding to about lo confidence)., RX J2009.8+1557 was detected during the ROSAT all-sky survey in 1990 at a count rate of $\pm$ 0.005 cts/s with a likelihood of detection of 12.4 (corresponding to about $\sigma$ confidence).
 With a tota of ouly 9 counts collected during the total exposure time of 195 sec this source is below the brightness limit of he 1 RNS catalog (Voges 11999)., With a total of only 9 counts collected during the total exposure time of 495 sec this source is below the brightness limit of the 1 RXS catalog (Voges 1999).
" Nonetheless. these few photous supply both a well-defined position of RA (2000.0) = 20509151709, Decl. ("," Nonetheless, these few photons supply both a well-defined position of RA (2000.0) = 9, Decl. ("
2000.0) = wwith an error radius of as well as an iucication for an absorbed. hard X- spectrum (iun ROSAT standards).,"2000.0) = with an error radius of as well as an indication for an absorbed, hard X-ray spectrum (in ROSAT standards)."
" Theharcucss ratios are URL = (Noovor Ni nC)n d Nap vor) = 0.702:0.32.. and IIR2 = (Noy290 Ns)ου) σου200 = O.52+0.31. where Nz7, denotes the nuuber of couuts in ROSAT's position sensitive proportional counter between channel a and channel b).Adopting a Ravineud-Sinith spectrum with a 1 keV temperature aud an absorbing column of ZrNj «10734n2 cu.27? (correspouding: to —504 of] the total cohuuu of Nj21.95 «10?! ci.2 in this direction: Dickey Lockiiau 1990) we derive an unabsorbed N-ray intensity of LOL—2.1]keV)=L8«10D ore ? +."," Thehardness ratios are HR1 = $_{\rm 52-201}$ – $_{\rm 11-41}$ $_{\rm 11-41}$ + $_{\rm 52-201}$ ) = $\pm$ 0.32, and HR2 = $_{\rm 91-200}$ – $_{\rm 50-90}$ $_{\rm 50-200}$ = $\pm$ 0.34, where $_{\rm a-b}$ denotes the number of counts in ROSAT's position sensitive proportional counter between channel a and channel b).Adopting a Raymond-Smith spectrum with a 1 keV temperature and an absorbing column of $N_{\rm H}$ $\times$ $^{21}$ $^{-2}$ (corresponding to $\sim$ of the total column of $N_{\rm H}$ $\times$ $^{21}$ $^{-2}$ in this direction; Dickey Lockman 1990) we derive an unabsorbed X-ray intensity of $L(0.1-2.4\ {\rm keV}) = 1.8 \times 10^{-13}$ erg $^{-2}$ $^{-1}$ ."
 This corresponds to au euission micasure of 107? (D/100 pc)? 7. No ROSAT poiuting exists for this sky area which could allow to derive better constraints., This corresponds to an emission measure of $\times$ $^{52}$ (D/100 $^2$ $^{-5}$ No ROSAT pointing exists for this sky area which could allow to derive better constraints.
in the amplitude of the rrelation by a factor of approximately 3 tracks closely the increase in the growth of bulges from cisrupted disks.,in the amplitude of the relation by a factor of approximately $3$ tracks closely the increase in the growth of bulges from disrupted disks.
 This demonstrates the simple idea we set out in Section 1.. that if mergers are the primary mechanism. which shape both black hole and bulee erowth in the galaxy population. then a larger fractional contribution to the bulge from disrupted disks should result in an inevitable evolution of the rrelation.," This demonstrates the simple idea we set out in Section \ref{intro}, that if mergers are the primary mechanism which shape both black hole and bulge growth in the galaxy population, then a larger fractional contribution to the bulge from disrupted disks should result in an inevitable evolution of the relation."
 This holds true even when no explicit evolution the the growth modes of black holes and bulges is assumed., This holds true even when no explicit evolution the the growth modes of black holes and bulges is assumed.
 ln Fig., In Fig.
 3. we redo the analysis of Fig 2. now using the dynamical model described in Section 2.1.2.., \ref{fig3} we redo the analysis of Fig \ref{fig2} now using the dynamical model described in Section \ref{dynamic}.
 This allows us to explore the sensitivity of cevolution when some evolution —in the black hole growth rate has been assumed., This allows us to explore the sensitivity of evolution when some evolution in the black hole growth rate has been assumed.
 The top panel of Fig., The top panel of Fig.
 3. clearly, \ref{fig3} clearly
we neglect the structure linking the screens. the resulting transit can be considered as the sum of (wo identical lishteurves. slightly phase-shifted. of compact convex bodies.,"we neglect the structure linking the screens, the resulting transit can be considered as the sum of two identical lightcurves, slightly phase-shifted, of compact convex bodies."
 Clearly a sphere would only poorly fit the resulting lighteurve., Clearly a sphere would only poorly fit the resulting lightcurve.
 This is verifiel on Fie.9 showing residuals three times larger peak-to-peak than lor the triangle. although both transiting objects have (he same cross-section.," This is verified on \ref{residual_transit_louver2} showing residuals three times larger peak-to-peak than for the triangle, although both transiting objects have the same cross-section."
 For a given impact parameter. the 2-screen object ingress ancl egress duration is longer than for a compact object of similar cross-section.," For a given impact parameter, the 2-screen object ingress and egress duration is longer than for a compact object of similar cross-section."
 Therefore the fit bv a sphere converges toward a larger impact parameter al which ingress and egress lengthen. and. in order to maintain (he overall transit duration aud depth. planet and star radii increased by 2230%.," Therefore the fit by a sphere converges toward a larger impact parameter at which ingress and egress lengthen, and, in order to maintain the overall transit duration and depth, planet and star radii increased by $\approx30\%$."
 Note that here again. the residuals Hehcurve shows an ambiguity with a ringed planet (Barnes&Fortney2004).," Note that here again, the residuals lighcurve shows an ambiguity with a ringed planet \citep{baf04}."
. The last example is a louver-like object. an elongated structure composed of six screens (Fig.10)).," The last example is a louver-like object, an elongated structure composed of six screens \ref{image_transit_louver6}) )."
 The screens produce undulating structures in (he ingress ancl egress transit Leltcurve which are visible in the residuals alter best-fit sphere subtraction (Fie.1l1))., The screens produce undulating structures in the ingress and egress transit lightcurve which are visible in the residuals after best-fit sphere subtraction \ref{residual_transit_louver6}) ).
 Here again. the object elongated shape induces a longer ineress ancl egress (han for the previously considered objects: The fit converges Coward a larger impact parameter and planet and star radii increased by 2280% and z6054 respectively.," Here again, the object elongated shape induces a longer ingress and egress than for the previously considered objects: The fit converges toward a larger impact parameter and planet and star radii increased by $\approx80\%$ and $\approx60\%$ respectively."
 The louver practically produces multiple (rausits., The louver practically produces multiple transits.
 Each screen can indeed be considered as a single objects transiting one alter the other. like (he two screens of the previous object.," Each screen can indeed be considered as a single objects transiting one after the other, like the two screens of the previous object."
 The detectability of the louver is twice better than for the triaugle if we consider the peak-to-peak residuals. (he cross-section remaining the same.," The detectability of the louver is twice better than for the triangle if we consider the peak-to-peak residuals, the cross-section remaining the same."
 The oscillations curing ingress and egress could be considered as an altention-getng signal from a communicative civilization. alühough the signal is in the 10.| range (in our examples) and requires a sampling time of the order of 5iin.," The oscillations during ingress and egress could be considered as an attention-getting signal from a communicative civilization, although the signal is in the $10^{-4}$ range (in our examples) and requires a sampling time of the order of $5\ min$."
 Detectability can even be higher as we will show in the next Section with mulüple (ransits., Detectability can even be higher as we will show in the next Section with multiple transits.
 For the record. note also that non-planetary transit could be created by Dysons sunllowers (Dyson2003).. a hypothetical live [orm spread into space to efficientlv collect the energy of a distant star.," For the record, note also that non-planetary transit could be created by Dyson's sunflowers \citep{dyson2003}, a hypothetical live form spread into space to efficiently collect the energy of a distant star."
 These structures. if they exist. could be almost circular and (quite compact. but also could look like parse or dendritic arrangements.," These structures, if they exist, could be almost circular and quite compact, but also could look like parse or dendritic arrangements."
 Nevertheless. large sunflowers are fay from the star and transit probability consequently is very low.," Nevertheless, large sunflowers are far from the star and transit probability consequently is very low."
 Let us now consider a formation of several objects spatially arranged in eroups to ingress the star according to a remarkable manner. such as a series of prime munbers. or powers of two.," Let us now consider a formation of several objects spatially arranged in groups to ingress the star according to a remarkable manner, such as a series of prime numbers, or powers of two."
 Even more remarkable would be a sudden swap between these (wo fIving Iormations alter several orbits., Even more remarkable would be a sudden swap between these two flying formations after several orbits.
 We consider (hat such multiple transits would clearly be attention-getting signals and the will of communication would be obvious., We consider that such multiple transits would clearly be attention-getting signals and the will of communication would be obvious.
 The photometric accuracy required, The photometric accuracy required
 , 
Figure 5. shows the interior temperature as a function of density for models of mass 0.8 Mo and core mass M;=0.49Mo.,Figure \ref{fig:str} shows the interior temperature as a function of density for models of mass 0.8 $\msun$ and core mass $M_{1} = 0.49 \msun$.
 A constant temperature “plateau” spreads inward from the base of the hydrogen-burning shell., A constant temperature “plateau” spreads inward from the base of the hydrogen-burning shell.
 This implies that the thermal structure of the helium core results mainly from the temperature in the hydrogen burning shell and neutrino energy losses in central regions (Fujimotoetal.1984)., This implies that the thermal structure of the helium core results mainly from the temperature in the hydrogen burning shell and neutrino energy losses in central regions \citep{fuj84}.
" Along with a decrease in the initial abundance of CNO elements, the hydrogen-burning rate decreases, and, in order to compensate for this, the temperature in the hydrogen-burning shell increases, which makes the contribution of compressional heating smaller."," Along with a decrease in the initial abundance of CNO elements, the hydrogen-burning rate decreases, and, in order to compensate for this, the temperature in the hydrogen-burning shell increases, which makes the contribution of compressional heating smaller."
" Accordingly, an isothermal plateau develops in the region where the radiative heat transport dominates over electron conduction for the metallicity of [Fe/H]S—5 (seeFujimotoetal.1995)."," Accordingly, an isothermal plateau develops in the region where the radiative heat transport dominates over electron conduction for the metallicity of $[{\rm Fe}/{\rm H}] \lesssim -5$ \citep[see ][]{fuj95}."
". In our models, the compressional heating plays a small part to raise the maximum temperature in the helium zone slightly higher than the temperature in the hydrogen burning shell, as seen from the lowest mass model; note that for the massive models, the helium burning already contributes appreciably to increase the temperatures in the right shoulder of structure lines in this figure."," In our models, the compressional heating plays a small part to raise the maximum temperature in the helium zone slightly higher than the temperature in the hydrogen burning shell, as seen from the lowest mass model; note that for the massive models, the helium burning already contributes appreciably to increase the temperatures in the right shoulder of structure lines in this figure."
 This behavior contrasts with that of model stars of younger population in which the compressional heating play the dominant part in determining the maximum temperature in the helium core and makes it much larger than the temperature in the hydrogen-burning shell., This behavior contrasts with that of model stars of younger population in which the compressional heating play the dominant part in determining the maximum temperature in the helium core and makes it much larger than the temperature in the hydrogen-burning shell.
" Because of high temperature in the hydrogen burning shell, therefore, Z=ϐ model stars experience central ignition at smaller initial masses than do model stars of younger populations for which the minimum initial mass for central helium ignition is ~2.5Mo."," Because of high temperature in the hydrogen burning shell, therefore, $Z=0$ model stars experience central ignition at smaller initial masses than do model stars of younger populations for which the minimum initial mass for central helium ignition is $\sim 2.5 \msun$."
" While the density and temperature of the hydrogen burning shell are in local maximum just before the helium ignition, hot CNO-cycle is not still effective."," While the density and temperature of the hydrogen burning shell are in local maximum just before the helium ignition, hot CNO-cycle is not still effective."
" At this stage, the nuclear timescale of !?N against proton capture reaction (~10000 sec) is much larger than that against 3-decay reaction (863 sec) and is negligible in the outcome of neither the nucleosynthesis nor the nuclear energy output."," At this stage, the nuclear timescale of $\nucm{13}{N}$ against proton capture reaction $\sim 10000$ sec) is much larger than that against $\beta$ -decay reaction (863 sec) and is negligible in the outcome of neither the nucleosynthesis nor the nuclear energy output."
" If neutrino energy losses are sufficiently effective before helium is ignited, the central region cools, and helium is ignited off-center as is the case in low-mass stars of younger populations."," If neutrino energy losses are sufficiently effective before helium is ignited, the central region cools, and helium is ignited off-center as is the case in low-mass stars of younger populations."
" In the central region of the model, plasma neutrinos are the dominant neutrino energy-loss mechanism and produce a steep gradient in the rate of released energy."," In the central region of the model, plasma neutrinos are the dominant neutrino energy-loss mechanism and produce a steep gradient in the rate of released energy."
" On the other hand, the conduction, which is important in the electron-degenerate core, transport the energy towards the center where the neutrino loss works."," On the other hand, the conduction, which is important in the electron-degenerate core, transport the energy towards the center where the neutrino loss works."
" Consequently, both conductivity and neutrino energy-loss rates promote cooling of the core and, thus, delay the off-center ignition of helium until a larger core mass gives rise to a larger hydrogen-shell burning rate and a larger temperature in the burning shell, as seen from Table 2.."," Consequently, both conductivity and neutrino energy-loss rates promote cooling of the core and, thus, delay the off-center ignition of helium until a larger core mass gives rise to a larger hydrogen-shell burning rate and a larger temperature in the hydrogen-burning shell, as seen from Table \ref{tab:he-flash}."
" In all of the models which experience off-center ignition, convection driven by helium burning extends into the upper hydrogen-rich layers during the decay phase of the core helium flash, because of smaller entropy in the hydrogen burning shell, as shown by Fujimotoetal.(1990,1995)."," In all of the models which experience off-center ignition, convection driven by helium burning extends into the upper hydrogen-rich layers during the decay phase of the core helium flash, because of smaller entropy in the hydrogen burning shell, as shown by \cite{fuj90,fuj95}."
. The ingestion of hydrogen into the helium convective zone begins a sequence of events that leads to the enrichment of the surface with carbon and nitrogen (Hollowelletal.1990;FujimotoWeissetal. 2004).," The ingestion of hydrogen into the helium convective zone begins a sequence of events that leads to the enrichment of the surface with carbon and nitrogen \citep{hol90,fuj00,sch02,pic04,wei04}."
. Characteristics of hydrogen mixing are also given in Table 2.., Characteristics of hydrogen mixing are also given in Table \ref{tab:he-flash}.
 The ingestion of hydrogen occurs within a matter of days after the helium-burning luminosity reaches its peak., The ingestion of hydrogen occurs within a matter of days after the helium-burning luminosity reaches its peak.
" In the models undergoing central helium burning, we do not find à hydrogen-mixing event; The core helium flash is rather weak (see [μοι in Table 2)), which makes it difficult for the outer edge of the convective region driven by helium burning to reach the hydrogen-containing layer (Fujimoto1977)."," In the models undergoing central helium burning, we do not find a hydrogen-mixing event; The core helium flash is rather weak (see $L_{\rm He}^{\rm max}$ in Table \ref{tab:he-flash}) ), which makes it difficult for the outer edge of the convective region driven by helium burning to reach the hydrogen-containing layer \citep{fuj77}."
". For the models of M=1.2Mo, we follow the evolution through the thermally pulsing AGB phase to find that the He-FDDM is triggered during the helium shell flashes."," For the models of $M=1.2 \msun$, we follow the evolution through the thermally pulsing AGB phase to find that the He-FDDM is triggered during the helium shell flashes."
" In this section, we compare the models in the literature with our models adopting similar input physics to see which cause the dominant effect on the differences in the evolution and to confirm the correctness of numerical computations."," In this section, we compare the models in the literature with our models adopting similar input physics to see which cause the dominant effect on the differences in the evolution and to confirm the correctness of numerical computations."
" In addition, we compare the models without non- effect for a-capture reactions to see the influence of uncertainty in"," In addition, we compare the models without non-resonant effect for $\alpha$ -capture reactions to see the influence of uncertainty in"
It can be shown that photons of energy £. traveling trough (he vacuum resonance region. experience resonant mode conversion with probability 2=1—exp|-x(EZ/E44)/*/2].,"It can be shown that photons of energy $E$, traveling through the vacuum resonance region, experience resonant mode conversion with probability $P_C = 1-\exp\left[-\pi(E/E_{\rm ad})^3/2\right]$."
" The adiabatie energv. £44. is defined: where 05 is the angle between (he photon propagation direction and the magnetic field. Lf, is (he atmosphere scale height. and Ej;©0.63(Z/A)D4,; keV is the ion evclotron energy. with atomic and mass numbers Z and A. respectfullv."," The adiabatic energy, $E_{\rm ad}$, is defined: where $\theta_B$ is the angle between the photon propagation direction and the magnetic field, $H_{\rho}$ is the atmosphere scale height, and $E_{Bi}\approx 0.63 (Z/A) B_{14}$ keV is the ion cyclotron energy, with atomic and mass numbers $Z$ and $A$, respectfully."
 The parameter fy; is a slowly varying [unetion of D. whose magnitude is of order munity.," The parameter $f_B$ is a slowly varying function of $B$ , whose magnitude is of order unity."
 At photon energies F<Fg. mode conversion is ineffective. while lor B=1.44. mode conversion is essentially complete.," At photon energies $E\ll E_{\rm ad}$, mode conversion is ineffective, while for $E\ga 1.4 E_{\rm ad}$, mode conversion is essentially complete."
 The vacuum resonance density occurs at: The vacuum resonance phenomenon leaves a unique signature on (he emission of polarized radiation [rom the NS (?).., The vacuum resonance density occurs at: The vacuum resonance phenomenon leaves a unique signature on the emission of polarized radiation from the NS \citep[][]{LaiHo03a}.
 For magnetic field strengths D.«7xLOM G. photons of both modes encounter (he vacuum resonance after decoupling from the NS atmosphere. while [or 7x107G«B510! G. the vacuum resonance occurs al an alimospheric depth between (he X and. O mode photospheres (?)..," For magnetic field strengths $B \ll 7\times 10^{13}$ G, photons of both modes encounter the vacuum resonance after decoupling from the NS atmosphere, while for $7\times 10^{13}\mbox{ G}<B< 5\times 10^{16}$ G, the vacuum resonance occurs at an atmospheric depth between the X and O mode photospheres \citep[][]{LaiHo03a}."
 In (he former case. photons encounter the vacuum resonance while propagating into the magnetosphere.," In the former case, photons encounter the vacuum resonance while propagating into the magnetosphere."
 For low enerev photons with E« Ej.mode conversion is ineffective. and the vacuum resonance has no effect on the emerging radiation.," For low energy photons with $E\ll E_{\rm ad}$, mode conversion is ineffective, and the vacuum resonance has no effect on the emerging radiation."
 For high energv photons. mode conversion is ellective. and results in conversion between the X and O polarization states.," For high energy photons, mode conversion is effective, and results in conversion between the X and O polarization states."
 Thus. the vacuum polarization effect causes (he plane of linear polarization to rotate 90° between low ancl high energies.," Thus, the vacuum polarization effect causes the plane of linear polarization to rotate $90^{\circ}$ between low and high energies."
 This can occur in models with magnetic fields as high as B=7xLOM G. though there is a dependence on the NS geometry (see re[sect:Results)).," This can occur in models with magnetic fields as high as $B=7\times 10^{13}$ G, though there is a dependence on the NS geometry (see \\ref{sect:Results}) )."
 After emerging from the atmosphere. photons propagate through the NS magnetosphere before reaching the observer.," After emerging from the atmosphere, photons propagate through the NS magnetosphere before reaching the observer."
 As discussed above. we ignore scattering and only consider magnetic field effects on the polarization mode evolution.," As discussed above, we ignore scattering and only consider magnetic field effects on the polarization mode evolution."
 The solutionto Maxwells equations in magnetized vacuum vields (he unit polarization, The solutionto Maxwell's equations in magnetized vacuum yields the unit polarization
The present investigation was based on the assumption that the circulation luminosity in the secondary’s interior is positive and large enough to enable thermal equilibrium.,The present investigation was based on the assumption that the circulation luminosity in the secondary's interior is positive and large enough to enable thermal equilibrium.
 The first part of this assumption can be abandoned., The first part of this assumption can be abandoned.
 As shown in the preceding section the circulation luminosity is certainly positive in. the radiative interior. not only in the secondary but also in the primary.," As shown in the preceding section the circulation luminosity is certainly positive in the radiative interior, not only in the secondary but also in the primary."
 In systems with intermediate or small mass ratios the effects in the primary are small while the effects in the secondary are important since the core luminosity is small., In systems with intermediate or small mass ratios the effects in the primary are small while the effects in the secondary are important since the core luminosity is small.
 In systems with large mass ratios (e.g. OO Αα) the neglect of the circulation luminosity in the primary introduces a slight error., In systems with large mass ratios (e.g. OO Aql) the neglect of the circulation luminosity in the primary introduces a slight error.
 The resulting treatment of contact binaries is. simple., The resulting treatment of contact binaries is simple.
 Sources of the circulation luminosity deep in the secondary's interior. are necessary. but details are unimportant.," Sources of the circulation luminosity deep in the secondary's interior are necessary, but details are unimportant."
 Free parameters are the temperature difference between the components AT)... bound to be positive and small. and the efficiency fg of the energy transfer between the components.," Free parameters are the temperature difference between the components $\Delta T_{\rm e}$, bound to be positive and small, and the efficiency $f_{\rm E}$ of the energy transfer between the components."
 The fractional extent of radiative regions is larger in the secondary than m the primary., The fractional extent of radiative regions is larger in the secondary than in the primary.
 In the course of evolution the period increases and the mass ratio decreases., In the course of evolution the period increases and the mass ratio decreases.
 Adopting standard values for the parameters. a survey of unevolved and evolved contact configuration has been obtained.," Adopting standard values for the parameters, a survey of unevolved and evolved contact configuration has been obtained."
 The results are successful in explaining the basic observational facts. in particular the preference for shallow contact. the minimum period and the minimum temperature. the existence of stable configurations with mass ratios very close to unity. and the period-colour relation.," The results are successful in explaining the basic observational facts, in particular the preference for shallow contact, the minimum period and the minimum temperature, the existence of stable configurations with mass ratios very close to unity, and the period-colour relation."
 Since stability considerations are essential in these observational tests we consider the results as strong arguments not only in favour of thermal equilibrium but also in favour of the treatment of the stability problem., Since stability considerations are essential in these observational tests we consider the results as strong arguments not only in favour of thermal equilibrium but also in favour of the treatment of the stability problem.
 Assuming thermal equilibrium. models of individual observed systems with reliable data are well-determined apart from some freedom in AT)...," Assuming thermal equilibrium, models of individual observed systems with reliable data are well-determined apart from some freedom in $\Delta T_{\rm e}$."
 When stable they can be used to to calibrate the efficiency and to determine metallicity and age., When stable they can be used to to calibrate the efficiency and to determine metallicity and age.
 The models obtained so far are stable. which lends again support to the assumption of thermal equilibrium.," The models obtained so far are stable, which lends again support to the assumption of thermal equilibrium."
 Evolutionary effects are important., Evolutionary effects are important.
" The efficiency is very small (fg=10%... 100%),", The efficiency is very small $f_{\rm E}=10^{-3}\ldots 10^{-5}$ ).
 Concerning the hydrodynamic problem. arguments have been presented that the velocity field in the common envelope has a reversing layer with motions from the secondary to the primary in the layers just above the critical surface and from the primary to the secondary in the surface layers.," Concerning the hydrodynamic problem, arguments have been presented that the velocity field in the common envelope has a reversing layer with motions from the secondary to the primary in the layers just above the critical surface and from the primary to the secondary in the surface layers."
 Details of the circulation. in. particular the circulation function. can be determined. only in à three-dimensional discussion.," Details of the circulation, in particular the circulation function, can be determined only in a three-dimensional discussion."
 In comparison with rotating stars the problem ts more difficult since the geometry 1s more complex., In comparison with rotating stars the problem is more difficult since the geometry is more complex.
 In other respects the problem is simpler since rotation law and geometry are known in a close approximation., In other respects the problem is simpler since rotation law and geometry are known in a close approximation.
 This problem will be investigated in a forthcoming paper. again assuming thermal equilibrium.," This problem will be investigated in a forthcoming paper, again assuming thermal equilibrium."
 Although we have strong arguments in favour of thermal equilibrium we have no proof. and it is uncertain whether a proof is possible.," Although we have strong arguments in favour of thermal equilibrium we have no proof, and it is uncertain whether a proof is possible."
We (hank Dragan Huterer For initial guidance in the project and many useful suggestions.,We thank Dragan Huterer for initial guidance in the project and many useful suggestions.
 We thank Glenn Starkman for introducing the idea of C(O.0).," We thank Glenn Starkman for introducing the idea of $C(\hat{\Omega},\theta)$."
 We thank IHiranva Peiris for a useful communication., We thank Hiranya Peiris for a useful communication.
 We acknowledge use of the IIEALPix package (Gorski2005) and the Legacy. Archive for Microwave Dackground Data Analvsis (LAMDDA)., We acknowledge use of the HEALPix package \citep{healpix} and the Legacy Archive for Microwave Background Data Analysis (LAMBDA).
 This work was partly supported by NSF under contract AST-0807564. NASA under contract NNNOO9ACSOG. and the University of Michigan.," This work was partly supported by NSF under contract AST-0807564, NASA under contract NNX09AC89G, and the University of Michigan."
the data imply that the magnetic field. once again. becomes weaker on larger scales.,"the data imply that the magnetic field, once again, becomes weaker on larger scales."
 GLAST is located at a distance of  2 kpe (Arcquilla&Ciold-smith 1954)., GL437 is located at a distance of $\sim$ 2 kpc \citep{arquilla}.
. It is composed of a compact. cluster of voung stars (including B stars). and a rellection nebulosity which is centered on WIX34 (Weintraub&Wastner1996).," It is composed of a compact cluster of young stars (including B stars), and a reflection nebulosity which is centered on WK34 \citep{weintraubkastner96}."
.. Phere is also a broad. molecularbipolar outflow that extends ~ l pc and is orientated roughly north-south (Gomezetal. 1992).. with the south lobe being blueshifted.," There is also a broad molecularbipolar outflow that extends $\sim$ 1 pc and is orientated roughly north-south \citep{gomez92}, with the south lobe being blueshifted."
 The SCUBA data (Fig., The SCUBA data (Fig.
 If) reveal that this source is an elongated core orientated in the northwest-southeast direction., \ref{fig1}f f) reveal that this source is an elongated core orientated in the northwest-southeast direction.
 The polarimetry is aligned. such that the magnetic field appears to be projected. perpendicular to the ridge of gas and dust forming the elongated source. roughly parallel to the outflow axis.," The polarimetry is aligned such that the magnetic field appears to be projected perpendicular to the ridge of gas and dust forming the elongated source, roughly parallel to the outflow axis."
 There is à polarisation null to the southeast of the core. and the polarisation percentage changes seemingly randomly across the core.," There is a polarisation null to the southeast of the core, and the polarisation percentage changes seemingly randomly across the core."
 IRAS 20126|4104 is located in the Cyvenus-X region. at a distance of 1.7 kpe (Wilkingetal.1989).," IRAS 20126+4104 is located in the Cygnus-X region, at a distance of 1.7 kpc \citep{wilking}."
.. Lt is. classified as a high mass protostellar object (LIMLPO:Sridharanetal.2002) although at a later stage of evolution than those discussed in. Curranetal.(2004) it is the hieh mass equivalent of a class O/L protostar., It is classified as a high mass protostellar object \citep[HMPO;][]{sridharan02} — although at a later stage of evolution than those discussed in \citet{hmpos} — it is the high mass equivalent of a class 0/I protostar.
 Lt has a well known molecular outllow and an ionised jet., It has a well known molecular outflow and an ionised jet.
 Phe CO outtlow is roughly north-south in orientation (position angle of 171)) with redshifted gas in the south and blueshiftecl eas in the north., The CO outflow is roughly north-south in orientation (position angle of ) with redshifted gas in the south and blueshifted gas in the north.
 The jet. detected by emission knots of H» and Su] has a position angle of (Shepherdetal.," The jet, detected by emission knots of $_{2}$ and ] has a position angle of \citep{shepherd00}."
2000).. Shepherdetal.(2000). conclude that the most. likely interpretation for this ijs for the collimated jet to be precessing through an angle of ~45°, \citet{shepherd00} conclude that the most likely interpretation for this is for the collimated jet to be precessing through an angle of $\sim$.
 Alore recently. Lebronetal.(2006). studied the kinematics of this region and. coneluded that whilst a precessing jet is supported by the data. multiple lows driven by independent sources cannot be ruled out.," More recently, \citet{lebron06}
 studied the kinematics of this region and concluded that whilst a precessing jet is supported by the data, multiple flows driven by independent sources cannot be ruled out."
 The polarimetry data. (Fig., The polarimetry data (Fig.
 lee) reveal that the projected magnetic field. has. in general a northeast-southwest orientation.," \ref{fig1}g g) reveal that the projected magnetic field has, in general a northeast-southwest orientation."
 This is roughly at an angle of ~ to the ionisecl jet. ancl almost perpendicular to the CO outllow. which mav be evidence of a helical field.," This is roughly at an angle of $\sim$ to the ionised jet, and almost perpendicular to the CO outflow, which may be evidence of a helical field."
 There are several polarisation nulls to the north of the core. away from the intensity peak.," There are several polarisation nulls to the north of the core, away from the intensity peak."
 The plane of skv component of the magnetic field is small only 0.1 mC. IRAS 20188|3928 is located in the Cvgnus region at an uncertain distance between 0.44 kpe (Littleetal. 1955)., The plane of sky component of the magnetic field is small – only 0.1 mG. IRAS 20188+3928 is located in the Cygnus region at an uncertain distance between 0.4–4 kpc \citep{little}. .
 All further caleulations in this paper assume a maximum distance of 4 Κρο and so all values quoted: are upper limits.," All further calculations in this paper assume a maximum distance of 4 kpc, and so all values quoted are upper limits."
 Littleetal.(19585) observed. this region in CO and and found a CO bipolar outllow. which has a north-southwestO direction.," \citet{little} observed this region in CO and $^{+}$, and found a CO bipolar outflow, which has a north-southwest direction."
 The redshiftecl eas is to the north. and the blueshifted eas is in the southwest.," The redshifted gas is to the north, and the blueshifted gas is in the southwest."
 The SCUBA data (Fig., The SCUBA data (Fig.
 Ihh) shows a curved ridge of gas and dust extending from the north of the core. curling around to the west.," \ref{fig1}h h) shows a curved ridge of gas and dust extending from the north of the core, curling around to the west."
 The polarimetry (in general) indicate that the magnetic field. lines are cast-west in orientation. projected onto the plane of the sky.," The polarimetry (in general) indicate that the magnetic field lines are east-west in orientation, projected onto the plane of the sky."
 Towards the southwest. of the core. the polarimetry vectors begin to deviate from the cast-west orientation towards a northeast-southwest orientation. aligning with the outflow in the southwest.," Towards the southwest of the core, the polarimetry vectors begin to deviate from the east-west orientation towards a northeast-southwest orientation, aligning with the outflow in the southwest."
 The degree of polarisation is measured. to be (on average) lower across the ridge than across the brighter core., The degree of polarisation is measured to be (on average) lower across the ridge than across the brighter core.
 A plane of sky field strength component of —0.2 mG has been calculated., A plane of sky field strength component of $\sim$ 0.2 mG has been calculated.
 LI2NT is à dark cloud at a distance of 850. pe (Yangοἱ , L1287 is a dark cloud at a distance of 850 pc \citep{yang}. .
There is an energetic οιμον associated with the cloud. orientated in a northeast-soutwest. direction.," There is an energetic outflow associated with the cloud, orientated in a northeast-southwest direction."
 At the centre of the outllow there is a very cold. LRAS source IRAS 00338]6312 (Yangetal.1991)., At the centre of the outflow there is a very cold IRAS source IRAS $+$ 6312 \citep{yang}.
. Phe submillimetre data (Fig., The submillimetre data (Fig.
 lii) reveal a marginally resolved object., \ref{fig1}i i) reveal a marginally resolved object.
 There is little polarimetry (2 vectors) as the integration time for this object was only 0.85 hrs. and so the errors in. polarisation are Large.," There is little polarimetry (2 vectors) as the integration time for this object was only 0.85 hrs, and so the errors in polarisation are large."
 Alonoceros 12 is a site of ongoing star formation. with a compact region (Wood&Churchwell1989). and. 120 and OLL masers (Downesetal.1975:Knapp&Brown1976).," Monoceros R2 is a site of ongoing star formation, with a compact region \citep{wood89} and $_{2}$ O and OH masers \citep{downes,knapp}."
. The outllow in this region is one of the largest known. with a total extent of 6.8 pe (assumingadistanceof950.pe:Racine&vandeBereh 1970).," The outflow in this region is one of the largest known, with a total extent of 6.8 pc \citep[assuming a distance of 950
pc;][]{racine}."
. Phis outllow extends to the north. east and. south-west of the region.," This outflow extends to the north, east and south-west of the region."
 Giannakopoulou.Mitchell&LHasegawa(1907) assume that the large scale north-southwest outflow is old. and that the source of this outllow is now inactive. as the lobes are not well collimated.," \citet{giannakopoulou} assume that the large scale north-southwest outflow is old, and that the source of this outflow is now inactive, as the lobes are not well collimated."
 The easterly outflow is thought to be a relatively voung. small-scale outflow. unresolved. from another source within the region.," The easterly outflow is thought to be a relatively young, small-scale outflow, unresolved, from another source within the region."
 The polarimetry of this region (Fig., The polarimetry of this region (Fig.
 1jj) indicate that the magnetic field aliens with the outllows — the vectors are oriented. north-south in the north of the region. east-west in the cast of the region and. northeast-southwest in the southwest of the region.," \ref{fig1}j j) indicate that the magnetic field aligns with the outflows – the vectors are oriented north-south in the north of the region, east-west in the east of the region and northeast-southwest in the southwest of the region."
 The morphology of the magnetic field is complex. with a plane of skv magnetic field strength of ~ 0.2 mG. NGC6334A4 is a giant. region/molecular cloud. complex ancl is located. close to the galactic plane at a distance of L7 kpc (Neckel1978).," The morphology of the magnetic field is complex, with a plane of sky magnetic field strength of $\sim$ 0.2 mG. NGC6334A is a giant region/molecular cloud complex and is located close to the galactic plane at a distance of 1.7 kpc \citep{neckel}."
. NGCG6334-submim has two cores. one in the east. ancl a brighter core (probably multiple. but not resolved within the JCAIT beam) in the west. of the image.," NGC6334-submm has two cores, one in the east, and a brighter core (probably multiple, but not resolved within the JCMT beam) in the west of the image."
 There is a CO bipolar outllow associated. with NGC6334A. which lies almost in the plane of the sky etal.2000) in a north-south direction. (DePree 1995)..," There is a CO bipolar outflow associated with NGC6334A, which lies almost in the plane of the sky \citep{sarma00} in a north-south direction \citep{depree}. ."
 Phere is a ridge of material between the two cores in which the polarimetry (fig., There is a ridge of material between the two cores in which the polarimetry (fig.
 ΤΚ) suggests that the magnetic Ποιά.isparallel to the ridge., \ref{fig1}k k) suggests that the magnetic fieldisparallel to the ridge.
 The magnetic field. across the faint core has a southeast-northwest direction. which curls round. to east-west across the ridge.," The magnetic field across the faint core has a southeast-northwest direction, which curls round to east-west across the ridge."
 The magnetic, The magnetic
"Indeed, the far-IR emission provides the most important and necessary constraints for the modeling of the dust SED of star forming region in collision debris.","Indeed, the far-IR emission provides the most important and necessary constraints for the modeling of the dust SED of star forming region in collision debris."
 The far-IR emission by dust is dominated by grains in thermal equilibrium with the radiation field., The far-IR emission by dust is dominated by grains in thermal equilibrium with the radiation field.
" Therefore, for a given set of grain properties, the far-IR SED gives direct access to the opticaldistribution of equilibrium temperatures of grains located mainly in dense PDR (photodissociation regions) and in the diffuse ISM (interstellar medium)."," Therefore, for a given set of grain optical properties, the far-IR SED gives direct access to the distribution of equilibrium temperatures of grains located mainly in dense PDR (photodissociation regions) and in the diffuse ISM (interstellar medium)."
" Consequently, it is the most reliable tracer of the conditions that the bulk of the dust is exposed to."," Consequently, it is the most reliable tracer of the physical conditions that the bulk of the dust is exposed to."
" On thephysical contrary, the mid-IR emission is dominated by grains stochastically heated."," On the contrary, the mid-IR emission is dominated by grains stochastically heated."
 The extent of their temperature fluctuations is unrelated to the energy density of the general radiation field., The extent of their temperature fluctuations is unrelated to the energy density of the general radiation field.
 Fig., Fig.
 2 illustrates this phenomenon., \ref{fig:sed_Zubko} illustrates this phenomenon.
 It shows the dust SED of the diffuse Galactic ISM exposed to ISRF of different intensities., It shows the dust SED of the diffuse Galactic ISM exposed to ISRF of different intensities.
" The dust is made of carbon and silicate grains of different sizes (3.5 tto 0.25 um) responsible for most of the continuum emission, and of PAH containing 10 to 1000 C atoms, mainly responsible for the mid-IR emission features."," The dust is made of carbon and silicate grains of different sizes (3.5 to 0.25 $\mu$ m) responsible for most of the continuum emission, and of PAH containing 10 to 1000 C atoms, mainly responsible for the mid-IR emission features."
" The integrated intensity of the ISRF is quantified by the standard Go parameter, defined by: where J, is the monochromatic mean intensity of the ISRF which is determined by the fit of the photospheric+nebular model from FUV to near-infrared."," The integrated intensity of the ISRF is quantified by the standard $G_0$ parameter, defined by: where $J_\lambda$ is the monochromatic mean intensity of the ISRF which is determined by the fit of the photospheric+nebular model from FUV to near-infrared."
" Using this ISRF, we compute the temperatures of our dust composition following the method of?."," Using this ISRF, we compute the temperatures of our dust composition following the method of."
". Increasing the intensity of the ISRF makes the equilibrium grain SED shift to shorter wavelengths (because of hotter temperatures), while the spectrum of PAH and very small grains remains almost identical, since they are excited by single photon events."," Increasing the intensity of the ISRF makes the equilibrium grain SED shift to shorter wavelengths (because of hotter temperatures), while the spectrum of PAH and very small grains remains almost identical, since they are excited by single photon events."
" From Fig. 2,,"," From Fig. \ref{fig:sed_Zubko},"
" it is obvious that in the absence of far-IR observations, the mean intensity of the radiation field is unconstrained."," it is obvious that in the absence of far-IR observations, the mean intensity of the radiation field is unconstrained."
" Unfortunately, we do not have any constraints in the far-infrared."," Unfortunately, we do not have any constraints in the far-infrared."
 Consequently we do not pretend to constrain the distribution of equilibrium grain temperatures within the galaxy., Consequently we do not pretend to constrain the distribution of equilibrium grain temperatures within the galaxy.
" In addition, without further information, the diagnostics that we can extract from the mid-IR are very limited."," In addition, without further information, the diagnostics that we can extract from the mid-IR are very limited."
" However, we can put some constraints on the dust properties by relying on the information provided by the attenuation estimate (section 3.1.0))."," However, we can put some constraints on the dust properties by relying on the information provided by the attenuation estimate (section \ref{sssec:extinction}) )."
" Indeed, by assuming that the global emission from the star forming region is isotropic, the energy absorbed in the UV and in the optical -- calculated from the attenuation determined by the photospheric+nebular fit — should be re-radiated in the IR."," Indeed, by assuming that the global emission from the star forming region is isotropic, the energy absorbed in the UV and in the optical – calculated from the attenuation determined by the photospheric+nebular fit – should be re-radiated in the IR."
 Another important point is that the ISRF illuminate the PAHs., Another important point is that the ISRF illuminate the PAHs.
" Indeed, for a typical cloud, the radiation field becomes redder, when the optical depth increases."," Indeed, for a typical cloud, the radiation field becomes redder, when the optical depth increases."
" The shape of this radiation field will have an effect on the emissivity of the PAHs, which are stochastically heated, but will not have any effect on the equilibrium grains which are only sensitive to the total absorbed power."," The shape of this radiation field will have an effect on the emissivity of the PAHs, which are stochastically heated, but will not have any effect on the equilibrium grains which are only sensitive to the total absorbed power."
" In our case, we consider the unattenuated ISRF (optically thin case), since PAH emission is dominated by photodissociation regions and cirrus clouds, which are both optically thin."," In our case, we consider the unattenuated ISRF (optically thin case), since PAH emission is dominated by photodissociation regions and cirrus clouds, which are both optically thin."
 This assumption introduces an additional on the determination of the PAH-to-dust mass ratiouncertainty that we will derive in the next paragraph., This assumption introduces an additional uncertainty on the determination of the PAH–to–dust mass ratio that we will derive in the next paragraph.
" However, this error should be small compared to the variations we probe here as we will see in Sec. 5.2.0.."," However, this error should be small compared to the variations we probe here as we will see in Sec. \ref{sec:pahvsZ}."
" Thus, we can perform a simple SED fit, where the total power emitted by the dust is determined by the energy balance, and where the PAH-to—dust mass ratio is constrained by the mid-IR fluxes."," Thus, we can perform a simple SED fit, where the total power emitted by the dust is determined by the energy balance, and where the PAH–to–dust mass ratio is constrained by the mid-IR fluxes."
" Indeed, in first approximation, LgoMpanX Go."," Indeed, in first approximation, $L_8\propto M_{PAH}\times G_0$ ."
 The fact that PAH are stochastically heated guarantees that they emit at the same wavelengths., The fact that PAH are stochastically heated guarantees that they emit at the same wavelengths.
" Hence, only one point is needed to scale their emission."," Hence, only one point is needed to scale their emission."
 The total infrared power follows a similar relation: L;jgοMaus:X Go., The total infrared power follows a similar relation: $L_{IR}\propto M_{dust}\times G_0$ .
 Therefore the 8 um-to-IR luminosity is proportional to the PAH-to-dust mass ratio: Lg/LigοMpag/Muust., Therefore the 8 $\mu$ m-to-IR luminosity is proportional to the PAH-to-dust mass ratio: $L_8/L_{IR}\propto M_{PAH}/M_{dust}$.
" As shown in Fig. 2,,"," As shown in Fig. \ref{fig:sed_Zubko},"
" the PAH-to-dust mass ratio, fpAn, is approximately independent of Go."," the PAH-to-dust mass ratio, $f_{\mbox{\scriptsize PAH}}$, is approximately independent of $G_0$."
" Such an estimate of fpay is strongly dependent on the accuracy of the attenuation estimate, since we use the attenuation to guess the level of the far-IR emission."," Such an estimate of $f_{\mbox{\scriptsize PAH}}$ is strongly dependent on the accuracy of the attenuation estimate, since we use the attenuation to guess the level of the far–IR emission."
" It is therefore not extremely reliable, in itself."," It is therefore not extremely reliable, in itself."
" However, the comparison of this derived PAH-to-dust mass ratio to its expected value from chemical evolution considerations is more interesting, since it provides an insight on the amount of embedded IR emission."," However, the comparison of this derived PAH--to--dust mass ratio to its expected value from chemical evolution considerations is more interesting, since it provides an insight on the amount of embedded IR emission."
" In other words, the attenuation estimate from the visible recombination lines is biased towards evolved star forming regions, where the molecular cloud covering factor is low enough for the radiation to "," In other words, the attenuation estimate from the visible recombination lines is biased towards evolved star forming regions, where the molecular cloud covering factor is low enough for the radiation to escape."
"On the contrary, the attenuation towards very young escape.embedded regions will not be traced by the Balmer decrement."," On the contrary, the attenuation towards very young embedded regions will not be traced by the Balmer decrement."
" In consequence, comparing the derived PAH-to-dust mass ratio to its expected value is a way to estimate the fraction of embedded infrared power."," In consequence, comparing the derived PAH-to-dust mass ratio to its expected value is a way to estimate the fraction of embedded infrared power."
 Another way to probe the presence of embedded star-forming regions would be to compare the estimate of the dust masses with the attenuation., Another way to probe the presence of embedded star-forming regions would be to compare the estimate of the dust masses with the attenuation.
" Unfortunately, this would require to make assumptions on the geometry of the dust and the stellar populations, which is beyond the scope of this paper."," Unfortunately, this would require to make assumptions on the geometry of the dust and the stellar populations, which is beyond the scope of this paper."
 This also prevents us from using the 24 um emission to trace embedded regions as it on many unconstrained physical star-formingparameters such as the dependscompacity of star-forming regions and the composition of the grains., This also prevents us from using the 24 $\mu$ m emission to trace embedded star-forming regions as it depends on many unconstrained physical parameters such as the compacity of star-forming regions and the composition of the grains.
 The PAHs are a better tracer as they do not exhibit systematic temperature variations and finally we do not have 24 um observastions for all systems., The PAHs are a better tracer as they do not exhibit systematic temperature variations and finally we do not have 24 $\mu$ m observastions for all systems.
are strongly self gravitating and have a large vertical extent.,are strongly self gravitating and have a large vertical extent.
 As an example let us consider the central regions of halo 4 which bears the greatest resemblance to a classical disk out of the set of simulations., As an example let us consider the central regions of halo 4 which bears the greatest resemblance to a classical disk out of the set of simulations.
" Taking a density cut of material greater than 101? the disk-like region had a radius which varies between 200 to 300? AU, and has a vertical extent of 150 AU measured from the mid plane."," Taking a density cut of material greater than $10^{10}$ the disk-like region had a radius which varies between 200 to 300 AU, and has a vertical extent of 150 AU measured from the mid plane."
" The thin disk approximation requires that the vertical extent of the disk, Z, is much smaller than the radial extent of the disk, R, and is only appropriate for disks where H/Re0.1 or lower (?).."," The thin disk approximation requires that the vertical extent of the disk, $H$, is much smaller than the radial extent of the disk, $R$, and is only appropriate for disks where $H/R \approx 0.1$ or lower \citep{Lodato08}."
 ? propose that to avoid artificial fragmentation in disks the scale height must be resolved by 4 SPH particle smoothing lengths per scale height., \citet{Nelson06} propose that to avoid artificial fragmentation in disks the scale height must be resolved by 4 SPH particle smoothing lengths per scale height.
" In the inner 100 AU of the disk in halo 4, the average particle smoothing length is 14.3 AU."," In the inner 100 AU of the disk in halo 4, the average particle smoothing length is 14.3 AU."
" However, it is unclear how best to define a scale height for this irregular puffy disk."," However, it is unclear how best to define a scale height for this irregular puffy disk."
" Given that even our best-case scenario cannot be considered a classical disk, and that as the simulations proceed fragmentation increasingly takes place in large spiral arm features or filaments (as seen in the later panels of 3)), a better resolution criteria to use is the Jeans mass."," Given that even our best-case scenario cannot be considered a classical disk, and that as the simulations proceed fragmentation increasingly takes place in large spiral arm features or filaments (as seen in the later panels of \ref{column}) ), a better resolution criteria to use is the Jeans mass."
" At our sink creation density of 105cm-?,, the Jeans mass is 2x107?,, at least two times more massive than our resolution limit."," At our sink creation density of $10^{15}$, the Jeans mass is $2\E^{-2}$, at least two times more massive than our resolution limit."
" Further to this, as shown in 2,, even when our resolution was increased by a factor of ten, the same number of fragments were formed."," Further to this, as shown in \ref{resolution}, even when our resolution was increased by a factor of ten, the same number of fragments were formed."
" In truth, the greatest limitation in our treatment of the disks is that the large size of our sinks means that we do not fully capture the behaviour of the inner regions of the disk."," In truth, the greatest limitation in our treatment of the disks is that the large size of our sinks means that we do not fully capture the behaviour of the inner regions of the disk."
" As such, there would probably be more fragmentation in the disks in reality than we find in our simulation."," As such, there would probably be more fragmentation in the disks in reality than we find in our simulation."
" For a fuller discussion of the physics of disk fragmentation around population III protostars see (?,b) where this issue is considered in great detail.", For a fuller discussion of the physics of disk fragmentation around population III protostars see \citep[][b]{Clark11b} where this issue is considered in great detail.
" The fragmentation here differs slightly from that shown in ?,, as in that case the sink radius was similar to the stellar radius, whereas here it is 20 AU."," The fragmentation here differs slightly from that shown in \citet{Greif11}, as in that case the sink radius was similar to the stellar radius, whereas here it is $20$ AU."
 The simulations of ? used the AREPO method which has less artificial viscosity than Gadget., The simulations of \citet{Greif11} used the AREPO method which has less artificial viscosity than Gadget.
" In order to test whether the results here were reproducible with AREPO, we re-simulated halo 4 (which showed the maximum difference between the feedback and no feedback cases) using AREPO with no feedback and the same sink sizes used in this study."," In order to test whether the results here were reproducible with AREPO, we re-simulated halo 4 (which showed the maximum difference between the feedback and no feedback cases) using AREPO with no feedback and the same sink sizes used in this study."
"and. when Lj,«LU. we correct for incompleteness using the expression: In the above formulae Lj, is in ss!. T. is the cluster age in yr for clusters older than 3 Myr and is T=10° yyr for younger clusters forwhich the time evolution of Li, 1s negligible.","and, when $L_{bol} < L_{bol }^{crit}$, we correct for incompleteness using the expression: In the above formulae $L_{bol}$ is in $^{-1}$, $T$ is the cluster age in yr for clusters older than 3 Myr and is $T=10^6$ yr for younger clusters forwhich the time evolution of $L_{bol}$ is negligible."
" For Ly>LU the SED models show that the cluster mass is well approximated by the relation logLy,—44.3+L2logT .", For $L_{bol} > L_{bol}^{crit}$ the SED models show that the cluster mass is well approximated by the relation $log M/\msun = log L_{bol}-44.3+1.2 log T$ .
 Our main conclusion of the analysis of YSCs in M33 is that a stochastically sampled IMF with a maximum mass of 100Μ.: well describes many properties of their integrated light distribution throughout the whole disk.," Our main conclusion of the analysis of YSCs in M33 is that a stochastically sampled IMF with a maximum mass of $100\,\msun$ well describes many properties of their integrated light distribution throughout the whole disk."
 We have corrected the mass of the less luminous clusters inferred by the SED model fits for incompleteness of the IMF according to the results of the previous Section., We have corrected the mass of the less luminous clusters inferred by the SED model fits for incompleteness of the IMF according to the results of the previous Section.
 The average mass of low luminosity clusters is underestimated if no correction ts applied., The average mass of low luminosity clusters is underestimated if no correction is applied.
 The mass-radius relation of the YSCs of our sample is illustrated in the right panel of Fig. 12.., The mass-radius relation of the YSCs of our sample is illustrated in the right panel of Fig. \ref{mr}. .
 Here. the size is the," Here, the size is the"
whom find evidence for «wicutation effects among the properties that they investigated.,whom find evidence for orientation effects among the properties that they investigated.
 None of these studies. however. have produced a genera technique for deterniuπιο the orientation of a eiver object. or for identifving subsets of analogous objects differiic onlv by oricutation.," None of these studies, however, have produced a general technique for determining the orientation of a given object, or for identifying subsets of analogous objects differing only by orientation."
 Despite the lack of an ΠΡΓΕ orientation indicator. a xeture of the ecometiic structure of active nuclei has αμασος (Cary&Paclovani1995:Caskel2000:Crenshawetal. 2010).," Despite the lack of an unambiguous orientation indicator, a picture of the geometric structure of active nuclei has emerged \citep{urry95,gaskell09,crenshaw10}."
. The region hat cuits the ower konüzation broad lines ds hought to be somewhat cüsklike., The region that emits the lower ionization broad lines is thought to be somewhat disklike.
" Material iu this cisk nay be flowing Πατ», but there may. also ve rotation and turbulence."," Material in this disk may be flowing inwards, but there may also be rotation and turbulence."
 The optical contiuuua nav also be cuutted from a disk. interior to the ITOAC lime region (BLR).," The optical continuum may also be emitted from a disk, interior to the broad line region (BLR)."
 These disks are assed o be coplanar with the «lusty torus that obscures our view of the continuuu and BLR in Type 2 objects., These disks are assumed to be coplanar with the dusty torus that obscures our view of the continuum and BLR in Type 2 objects.
 The higher ionization marrow lines are jlieved. to arise du a biconical reeiou hat is verpendicular to the disk aud flowing outward., The higher ionization narrow lines are believed to arise in a biconical region that is perpendicular to the disk and flowing outward.
 Some fraction of he higher ionization broad lines nay also participate in this polar outflow., Some fraction of the higher ionization broad lines may also participate in this polar outflow.
 The ower ionization narrow lines probably arise over a very large region doniüuated by the poteutial of 16 host galaxy rather than the central black hole., The lower ionization narrow lines probably arise over a very large region dominated by the potential of the host galaxy rather than the central black hole.
" Iu section 2. a diagramC» is constructed showiusC» je relationship between the motions of the ""meril that cits bro Fe II aud that which cinits narrow |O ΤΠ."," In section 2, a diagram is constructed showing the relationship between the motions of the material that emits broad Fe II and that which emits narrow [O III]."
" This diagram is most obviously interpreted as ixdicating the oricutation of two eroups of quasai* one pole-on and one edee-on. within the ecreral picture described ADOVE,"," This diagram is most obviously interpreted as indicating the orientation of two groups of quasars, one pole-on and one edge-on, within the general picture described above."
 Exploration of imuermediate objects show tha another factor. m acdlition to orientation. is at work: this is identified as Lia Lev ox Eddinetou ratio (ER).," Exploration of intermediate objects show that another factor, in addition to orientation, is at work; this is identified as $_{bol}$ $_{Edd}$ or Eddington ratio (ER)."
 These internediate objects can be separated into subsets that match the pole-ou aud edge-on subsets in ER. mt differ iu orientation.," These intermediate objects can be separated into subsets that match the pole-on and edge-on subsets in ER, but differ in orientation."
 Section 3 preseuts the claracteristic properties of each of these four subses. allowing a separation of the effects of the two factors.," Section 3 presents the characteristic properties of each of these four subsets, allowing a separation of the effects of the two factors."
 Finally. section l discusses the miplicatlous of the oricutation-driven effects for deterninations of black hole mass (Mp) aud ER.," Finally, section 4 discusses the implications of the orientation-driven effects for determinations of black hole mass $_{BH}$ ) and ER."
 Section [concludes with sole questions that coukl be answered by further development of this techique., Section 4 concludes with some questions that could be answered by further development of this technique.
 The study by Huetal.(2008) was alnied at understanding the properties of Fe II emissiou iu quasars., The study by \citet{hu08} was aimed at understanding the properties of Fe II emission in quasars.
 They decomposed the spectra of 1037 Z-U.8 quasars iuto coutiuuun aud multiple line contributions., They decomposed the spectra of 4037 $<$ 0.8 quasars into continuum and multiple line contributions.
" These authors measure Dmuuinosity. width. and velocity shift parameters for I+ (broad and narrow separately). ο TM) A5007. Mg IT A2800. and |O II], A3727."," These authors measure luminosity, width, and velocity shift parameters for $\beta$ (broad and narrow separately), [O III] $\lambda$ 5007, Mg II $\lambda$ 2800, and [O II] $\lambda$ 3727."
 Because of the range of redshift covered. the Mg IT aud [O IH] lines are only measured iu a subset of the full s;uuple.," Because of the range of redshift covered, the Mg II and [O II] lines are only measured in a subset of the full sample."
 They also 1ieasure these parameters for the eusemible of Fe II lines in the range AA L13I-1681., They also measure these parameters for the ensemble of Fe II lines in the range $\lambda\lambda$ 4434-4684.
 They devote considerable effort to testing and demonstrating the accuracy of their technique - in particular. the velocity shifts.," They devote considerable effort to testing and demonstrating the accuracy of their technique - in particular, the velocity shifts."
 The suuple is drawn from the SDSS DRS (Adchnuan-AMcCarthyetal.2007) quasar catalog (Schneideretal.2007) with the following additional restrictions., Their sample is drawn from the SDSS DR5 \citep{adelman07} quasar catalog \citep{schneider07} with the following additional restrictions.
 They attempt to ft onlv objects having signal-to-noise ratio (S/N) > 10 iu the range AA LI30-5550., They attempt to fit only objects having signal-to-noise ratio (S/N) $>$ 10 in the range $\lambda\lambda$ 4430-5550.
" They thendiscard objects with poor fits to their model or with equivaleut width of the Fe II cmussion (EWp,47) <«", They thendiscard objects with poor fits to their model or with equivalent width of the Fe II emission $_{Fe II}$ ) $<$.
 Finally. they discard objects with large errors πι either the width of the broad IL? or the velocity shift of |O III] A5007.," Finally, they discard objects with large errors in either the width of the broad $\beta$ or the velocity shift of [O III] $\lambda$ 5007."
" These restrictions. yarticularly the niunninuuini ENrz,4,. should be kept in miud for statistical conclusions craw fromm heir sample."," These restrictions, particularly the minimum $_{Fe II}$, should be kept in mind for statistical conclusions drawn from their sample."
 IIuetal.(2008). adopt the [ο ΤΠ A5007 line o define the systemic redshift because they cau neasure its position in all objects iu their sample., \citet{hu08} adopt the [O III] $\lambda$ 5007 line to define the systemic redshift because they can measure its position in all objects in their sample.
 Ilowever. A5SOOF is known to © blIneshüfted or to lave OCXCOSS Olülssjon oni its due Wine in many objects (Zamanovctal.2002:Boroson2005).," However, $\lambda$ 5007 is known to be blueshifted or to have excess emission on its blue wing in many objects \citep{zamanov02,boroson05}."
 Shifts as huge as several huudred km 1 are occasionally seeu., Shifts as large as several hundred km $^{-1}$ are occasionally seen.
" For this study. all the velocities have been renormalized to the |O II] A3727 line, which is measured in 2265 of the objects."," For this study, all the velocities have been renormalized to the [O II] $\lambda$ 3727 line, which is measured in 2265 of the objects."
 Objects without a measurement of velocity shift for [O II] were discarded., Objects without a measurement of velocity shift for [O II] were discarded.
 Figure 1 shows the |O TH] velocity shift plotted against the Fe II velocity shift for the 2265 objects., Figure 1 shows the [O III] velocity shift plotted against the Fe II velocity shift for the 2265 objects.
 The raudoni errors on the Fe II velocity shifts are about 100 kin 3 for small shifts and rise to about 300 kn) s| for the shifts of 3000 lau +., The random errors on the Fe II velocity shifts are about 100 km $^{-1}$ for small shifts and rise to about 300 km $^{-1}$ for the shifts of 3000 km $^{-1}$ .
 The errors on the |O III] velocity shifts are about 25 kins | for small shifts aud rise ο about, The errors on the [O III] velocity shifts are about 25 km $^{-1}$ for small shifts and rise to about
The work described in (his paper is inspired by the success achieved by (he OGLE (Udalski 2003) ancl MOA (Bond et 2001) teams in discovering and monitoring events ol short curation.,The work described in this paper is inspired by the success achieved by the OGLE (Udalski 2003) and MOA (Bond et 2001) teams in discovering and monitoring events of short duration.
 Short-duration events now constitute a significant fraction of all event candidates., Short-duration events now constitute a significant fraction of all event candidates.
 While not all of the candidate events of short duration are likely to be lensing events. many of the light curves appear to be well fit by lensine models.," While not all of the candidate events of short duration are likely to be lensing s, many of the light curves appear to be well fit by lensing models."
 The prospects for increasing the rate οἱ discovery of short-duration events are good., The prospects for increasing the rate of discovery of short-duration events are good.
 Ongoing programs are continuing to implement improvenienis. while new projects are starting or being planned.," Ongoing programs are continuing to implement improvements, while new projects are starting or being planned."
 The value of 7j; is an estimate of the time during which the gravitational magnification would exceed. 1.34 for close approaches between the source and lens., The value of $\tau_E$ is an estimate of the time during which the gravitational magnification would exceed $1.34$ for close approaches between the source and lens.
 Current. observing programs are sensitive enough to detect ongoing events when (he magnification is just a [ew percent. increasing the effective event duration bv a [actor that can be as large as 3.5.! These programs are therefore able to call alerts on ongoing events with small values of Tj.," Current observing programs are sensitive enough to detect ongoing events when the magnification is just a few percent, increasing the effective event duration by a factor that can be as large as $3.5$ These programs are therefore able to call alerts on ongoing events with small values of $\tau_E$."
" Just as alerts on events deemed likely to produce caustic crossings inspire world- ""follow-up"" with more frequent observations (see Griest Salizadeh 1993). alerts on short-duration events can be accorded high prioritv."," Just as alerts on events deemed likely to produce caustic crossings inspire world-wide “follow-up” with more frequent observations (see Griest Safizadeh 1998), alerts on short-duration events can be accorded high priority."
 The goal is to collect enough data to perimit detailed moclel fits., The goal is to collect enough data to permit detailed model fits.
and all the strong lines detected in are seen in6.,and all the strong lines detected in are seen in.
. In Fig. 9..," In Fig. \ref{linecomp},"
 we compare ihe intensity ratios of the lines detected in both objects., we compare the intensity ratios of the lines detected in both objects.
 The results show that the line intensilv ratios in the (wo objects are in agreement within one order of magnitude with an average value of 0.08 with a standard deviation of 0.05., The results show that the line intensity ratios in the two objects are in agreement within one order of magnitude with an average value of 0.08 with a standard deviation of 0.05.
 The intensity ratios. however. might be affected by beam dilution.," The intensity ratios, however, might be affected by beam dilution."
 is a more compact object aud (hus suffers from a larger bean dilution ellect., is a more compact object and thus suffers from a larger beam dilution effect.
" If we correct for the beam dilution effect. the Z(CIT //ORC+ 10216) ratios should increase bv a factor of £. where Assuming θες=20"" and Viequay=30"" (Fukasakuοἱal.1994) and using 9;=40” and 30"" lor the ARO 12m and the SMT respectively. we obtain the € values of 1.8 and 1.6 for the 12mm and SAIT data. respectively."," If we correct for the beam dilution effect, the $I$ $I$ (IRC+10216) ratios should increase by a factor of $\xi$ , where Assuming $\theta_{\rm CIT\,6}=20''$ and $\theta_{\rm IRC+10216}=30''$ \citep{fukasaku94} and using $\theta_b=40''$ and $''$ for the ARO 12m and the SMT respectively, we obtain the $\xi$ values of 1.8 and 1.6 for the m and SMT data, respectively."
 While recognizing Chat eq., While recognizing that eq.
 9 assumes uniform brightness temperature which is probably not realistic. (he result does suggest that there is no large correction [actor difference between the A 2 mm and 3 mm bands.," 9 assumes uniform brightness temperature which is probably not realistic, the result does suggest that there is no large correction factor difference between the $\lambda$ 2 mm and 3 mm bands."
 Our conclusion that for most of the species the Z(CET (11210216) line ratios are in good agreement therefore stands., Our conclusion that for most of the species the $I$ $I$ (IRC+10216) line ratios are in good agreement therefore stands.
 Consequently. we conclude that is indeed likely representative of C-rich envelopes although comparisons wilh a larger sample of objects will be needed {ο make a stronger statement.," Consequently, we conclude that is indeed likely representative of C-rich envelopes although comparisons with a larger sample of objects will be needed to make a stronger statement."
 The richness of molecular lines in the spectra of is mainlv due to its relatively nearby distance. and not due (o any special circumstances.," The richness of molecular lines in the spectra of is mainly due to its relatively nearby distance, and not due to any special circumstances."
 Fi, Fig.
e.9 also shows that (here are a few molecular species lor which the Z(CIET Z IRC+10216) ratios depart from (he average value., \ref{linecomp} also shows that there are a few molecular species for which the $I$ $I$ (IRC+10216) ratios depart from the average value.
 For the. CN and Που lines. (he ratios are higher. whereas for ICN. SiS. and CIL. the ratios are lower.," For the CN and $_3$ N lines, the ratios are higher, whereas for HCN, SiS, and $_4$ H, the ratios are lower."
 As discussed in Sect. 4.1.2...," As discussed in Sect. \ref{carbon},"
 this partly reflects the chemical evolution in the cireumstellar envelope around (he more evolved AGB star6., this partly reflects the chemical evolution in the circumstellar envelope around the more evolved AGB star.
. However. the cause of the abnormally strong C;II emission in remains unknown.," However, the cause of the abnormally strong $_4$ H emission in remains unknown."
 Far-IR spectra of and have been obtained by the ISO Lone Waveleneth Spectrometer (LAWS) (Schoieretal.2002:Cernicharo 1996)..," Far-IR spectra of and have been obtained by the ISO Long Wavelength Spectrometer (LWS) \citep{schoier02,cernicharo96}. ."
 The rotational transitions revealed by the ISO spectra can trace the inner regions of the circumstellar envelopes 2002)., The rotational transitions revealed by the ISO spectra can trace the inner regions of the circumstellar envelopes \citep{herpin02}.
. In Fig 10.. we compare the far-IR spectra of the two objects.," In Fig \ref{iso_con}, we compare the far-IR spectra of the two objects."
 The spectra were retrieved from (he ISO archive., The spectra were retrieved from the ISO archive.
 Inspection of the figure shows (hat their far-IR. spectra are dominated by thermal continuum emission from the dust with some superimposed molecular lines., Inspection of the figure shows that their far-IR spectra are dominated by thermal continuum emission from the dust with some superimposed molecular lines.
 Lindqvistetal.(2000) fitted the dust continuum of IRC+10216 with a single blackbody ol 510 Ix and the continuum of CIT 6 with two blackbodies of 1000 IX and 510 Ix. In the long wavelength region. the dust temperatures of the two objects are therefore almost identical.," \citet{lin00} fitted the dust continuum of IRC+10216 with a single blackbody of 510 K and the continuum of CIT 6 with two blackbodies of 1000 K and 510 K. In the long wavelength region, the dust temperatures of the two objects are therefore almost identical."
 Fig., Fig.
 11. gives the continuum-subtracted ISO LWS spectra of the two C-rich envelopes., \ref{isocomp} gives the continuum-subtracted ISO LWS spectra of the two C-rich envelopes.
 A number of lines from CO. HCN. HECN. andvibrationally excitedIICN have been identified," A number of lines from CO, HCN, $^{13}$ CN, andvibrationally excitedHCN have been identified"
under the hypothesis that their inclination is the same.,under the hypothesis that their inclination is the same.
 It is immediately clear from the Figure that our couple of low inclination orbits show an excess of counter-rotating stars in the outer halo., It is immediately clear from the Figure that our couple of low inclination orbits show an excess of counter-rotating stars in the outer halo.
" Table 2 reports the number of satellite star particles in the positive and negative peak in all of our cases, and the number of particles having rotation velocity Όνοι«—10 and Όνοι>10 It also reports the same numberskm/s obtained from our km/s.set B, in which the DM halo has no spin."," Table \ref{tab:npart} reports the number of satellite star particles in the positive and negative peak in all of our cases, and the number of particles having rotation velocity $v_{\rm rot}< -10$ km/s and $v_{\rm rot}>10$ km/s. It also reports the same numbers obtained from our set B, in which the DM halo has no spin."
 We give such numbers for the final time of our simulations., We give such numbers for the final time of our simulations.
" The counter-rotating excess signal, defined as the fraction of the number of counter- to co- rotating satellite stellar particles, is 3.39 at peaks and 1.98 overall for the low-inclination orbits, while it is 1.3 for the high-inclination ones."," The counter-rotating excess signal, defined as the fraction of the number of counter- to co- rotating satellite stellar particles, is 3.39 at peaks and 1.98 overall for the low-inclination orbits, while it is 1.3 (at peaks) for the high-inclination ones."
" Note that, for high-inclination(atpeaks) orbits, the retrograde case shows a peak at a rotation velocity near to Όνοι=0."," Note that, for high-inclination orbits, the retrograde case shows a peak at a rotation velocity near to $v_{\rm
 rot}=0$."
 This is because high inclination orbits have a larger impact parameters., This is because high inclination orbits have a larger impact parameters.
" The disk responds by tilting more than in the low-inclination case, and as a result, rotation velocities in the disk reference frame are shifted towards more positive values."," The disk responds by tilting more than in the low-inclination case, and as a result, rotation velocities in the disk reference frame are shifted towards more positive values."
" This effect is also present in our low-inclination runs, but it is much smaller because the tilt of the disk is smaller."," This effect is also present in our low-inclination runs, but it is much smaller because the tilt of the disk is smaller."
" It remains to be determined if the counter-rotating signal is caused by the interaction with the stellar disk of the main halo, or by the DM of the halo itself."," It remains to be determined if the counter-rotating signal is caused by the interaction with the stellar disk of the main halo, or by the DM of the halo itself."
" The upper row of Figure 2 shows the results of the same analysis performed on our set A, but in the case in which no halo spin is present (set B)."," The upper row of Figure \ref{fig:spin0} shows the results of the same analysis performed on our set A, but in the case in which no halo spin is present (set B)."
" Here we only show the results for our last snapshot; as in Figure l, averaging over five snapshots makes no appreciable difference."," Here we only show the results for our last snapshot; as in Figure \ref{fig:spin1}, averaging over five snapshots makes no appreciable difference."
" Also in our set B, low inclination retrograde orbits produce more counter-rotating star particle than co-rotating stars produced by prograde orbits."," Also in our set B, low inclination retrograde orbits produce more counter-rotating star particle than co-rotating stars produced by prograde orbits."
" Again, this is not the case for high inclination orbits."," Again, this is not the case for high inclination orbits."
" From Figure 2,, it is clear that the excess of counter-rotating star is clearly smaller: and this is due to the fact thatmore co-rotating stars are produced by prograde orbit in our no spin case than in spin 1 case."," From Figure \ref{fig:spin0}, it is clear that the excess of counter-rotating star is clearly smaller: and this is due to the fact that co-rotating stars are produced by prograde orbit in our no spin case than in spin 1 case."
" From Table 2,, the counter-rotating excess signal is 1.60 at peaks and 1.59 overall in the low inclination case, and 1.17 at peaks and 1.02 overall in the high inclination one."," From Table \ref{tab:npart}, the counter-rotating excess signal is 1.60 at peaks and 1.59 overall in the low inclination case, and 1.17 at peaks and 1.02 overall in the high inclination one."
" Therefore, both disk rotation and halo spin contribute to the slowing-down of prograde orbits and to the consequent smaller amount of high-energy star particles stripped from satellites that can reach the outer halo."," Therefore, both disk rotation and halo spin contribute to the slowing-down of prograde orbits and to the consequent smaller amount of high-energy star particles stripped from satellites that can reach the outer halo."
 'The lower row of Figure 2 shows that our result does not depend on the mass resolution of our satellite halo., The lower row of Figure \ref{fig:spin0} shows that our result does not depend on the mass resolution of our satellite halo.
" Even using 10 times more particles in the secondary, only our low inclination couple of mergers shows an excess of counter-rotating stars in the outer halo."," Even using 10 times more particles in the secondary, only our low inclination couple of mergers shows an excess of counter-rotating stars in the outer halo."
 'Two effects are acting on the satellite halo in these kind of mergers: dynamical friction and tidal disruption., Two effects are acting on the satellite halo in these kind of mergers: dynamical friction and tidal disruption.
 The first one is exerted both by the main halo DM particles and by the disk star particles., The first one is exerted both by the main halo DM particles and by the disk star particles.
" The second is most important near the center of the main halo, where the gravitational potential is stronger."," The second is most important near the center of the main halo, where the gravitational potential is stronger."
" But, dynamical friction depends upon the details of the satellite orbits."," But, dynamical friction depends upon the details of the satellite orbits."
 It is already known (Quinn&Goodman1986;Walkeretal.Villalobos&Helmi2008) that prograde orbits tends to decay faster than retrograde ones.," It is already known \citep{Quinn86,Walker96,Huang97,Velazquez99,Alvaro08} that prograde orbits tends to decay faster than retrograde ones."
" The dynamical frictionforce goes as Faynx1/v2 (Binney&Tremaine 2008), where v, is the velocity of the satellite relative to the field particles; retrograde satellites have higher v; with respect to prograde one, since in the first case the rotation velocity of the satellite is opposite to that of the disk and the main DM halo particles."," The dynamical friction goes as $F_{\rm dyn} \propto 1/v_s^2$ \citep{BinneyTremaine}, where $v_s$ is the velocity of the satellite relative to the field particles; retrograde satellites have higher $v_s$ with respect to prograde one, since in the first case the rotation velocity of the satellite is opposite to that of the disk and the main DM halo particles."
 Particles stripped from the satellites will remain on the orbit on which, Particles stripped from the satellites will remain on the orbit on which
and the velocity interval is in km s!.,and the velocity interval is in km $^{-1}$.
 No HI emission is detected from the smaller spiral galaxy (galaxy-2) for which the 30 upper limit of H mass assuming a similar distance and linewidth is , No H emission is detected from the smaller spiral galaxy (galaxy-2) for which the $3\sigma$ upper limit of H mass assuming a similar distance and linewidth is $2.4\times10^9 {\rm M}_\odot$ .
"Since galaxy-2 is about factor of two smaller in size than galaxy-1,Ms. for a similar H surface density, the total mass is expected to be ~1.9x109Ms, consistent with the derived upper limit."," Since galaxy-2 is about factor of two smaller in size than galaxy-1, for a similar H surface density, the total mass is expected to be $\sim1.9\times10^9 {\rm M}_\odot$, consistent with the derived upper limit."
" Figure 3 and 4 show the SDSS DR6 r-band image of the field overlaid with the integrated H emission contours for synthesized beam size of 16.6""x13.5"" and 6.9""x5.1"" respectively."," Figure \ref{fig:gp3} and \ref{fig:gp4} show the SDSS DR6 $r$ -band image of the field overlaid with the integrated H emission contours for synthesized beam size of $16.6\arcsec 
\times 13.5\arcsec$ and $6.9\arcsec \times 5.1\arcsec$ respectively."
 Both these H maps show the emission to be extended and significantly offset from the optical position of the galaxy., Both these H maps show the emission to be extended and significantly offset from the optical position of the galaxy.
 Figure 5 shows the H emission spectrum for galaxy-7 with a spectral resolution of ~45 km s! and a >4c peak flux density of ~2.8 mJy.," Figure \ref{fig:gp5} shows the H emission spectrum for galaxy-7 with a spectral resolution of $\sim 
45$ km $^{-1}$ and a $> 4\sigma$ peak flux density of $\sim 2.8$ mJy."
 This marginally significant emission feature at the expected systemic velocity of the galaxy coinciding with the optical position makes it most likely to be H emission from this galaxy., This marginally significant emission feature at the expected systemic velocity of the galaxy coinciding with the optical position makes it most likely to be H emission from this galaxy.
 The emission profile shows weak indication of asymmetry and the total estimated H mass is 5.3+1.3x10?Mg for a luminosity distance of 299 Mpc quoted in NED.," The emission profile shows weak indication of asymmetry and the total estimated H mass is $5.3\pm1.3 
\times 10^9 {\rm M}_\odot$ for a luminosity distance of $299$ Mpc quoted in NED."
" Figure 6 shows the SDSS DR6 r-band image of the field overlaid with the H1 emission contours for a synthesized beam size of 25.2""x23.0"".", Figure \ref{fig:gp6} shows the SDSS DR6 $r$ -band image of the field overlaid with the H emission contours for a synthesized beam size of $25.2\arcsec \times 23.0\arcsec$.
" Apart from the disturbed optical morphology, for this galaxy also there is some hint that the H emission is extended and offset from the optical position."," Apart from the disturbed optical morphology, for this galaxy also there is some hint that the H emission is extended and offset from the optical position."
 But this offset is smaller than the beam size and hence needs to be confirmed., But this offset is smaller than the beam size and hence needs to be confirmed.
" This low surface brightness emission is resolved in the high resolution map, making it difficult to draw any definitive conclusion on interaction between this pair of galaxies."," This low surface brightness emission is resolved in the high resolution map, making it difficult to draw any definitive conclusion on interaction between this pair of galaxies."
 One intriguing thing here is the striking similarity in the morphology of the central object of our study (galaxy-1) and that of the other spiral galaxy (galaxy-2) in the field., One intriguing thing here is the striking similarity in the morphology of the central object of our study (galaxy-1) and that of the other spiral galaxy (galaxy-2) in the field.
 The small spiral galaxy-2 looks like a scaled down image of galaxy-1 with similar tidal-tail-like extended features., The small spiral galaxy-2 looks like a scaled down image of galaxy-1 with similar tidal-tail-like extended features.
 The total H1 mass estimated for galaxy-1 from the GMRT observations is 7.7x 10° Μο., The total H mass estimated for galaxy-1 from the GMRT observations is $\times$ $^9$ $_{\odot}$.
" As has been mentioned earlier, the H mass is mostly concentrated on galaxy-1 and seems to have an offset with respect to the optical position which can be due to tidal interaction between the two neighbours."," As has been mentioned earlier, the H mass is mostly concentrated on galaxy-1 and seems to have an offset with respect to the optical position which can be due to a tidal interaction between the two neighbours."
" From the optical image,a the galaxy diameter quoted in NED is ~ 22”."," From the optical image, the galaxy diameter quoted in NED is $\sim 22\arcsec$ ."
" Assuming the entire H is only from galaxy-1, the H surface density of this galaxy turns out to be log(Mir/d7)6.86 Mo kpe~?, where d; is the linear diameter of the galaxy."," Assuming the entire H is only from galaxy-1, the H surface density of this galaxy turns out to be $\log(M_{\rm HI}/d_l^2) = 6.86$ $_{\odot}$ $^{-2}$, where $d_l$ is the linear diameter of the galaxy."
 The quantity log(My1/d?) is known to be a good diagnostic of the H content of galaxies., The quantity $\log(M_{\rm HI}/d_l^2)$ is known to be a good diagnostic of the H content of galaxies.
" Comparing the derived H surface density for this galaxy with the expected surface density for spiral galaxies of similar size (Haynes&Giovanelli1984),, we find no H deficiency in this particular case."," Comparing the derived H surface density for this galaxy with the expected surface density for spiral galaxies of similar size \citep{hg84}, we find no H deficiency in this particular case."
" Thus even if we detect a reasonably disturbed and displaced H1 disk in the main galaxy of the system possibly due to an ongoing interaction, we do not see any signs of gas loss from the system."," Thus even if we detect a reasonably disturbed and displaced H disk in the main galaxy of the system possibly due to an ongoing interaction, we do not see any signs of gas loss from the system."
 This conclusion is based on the assumption that the entire H detected belongs to galaxy-1 and there is no contribution from the other galaxy., This conclusion is based on the assumption that the entire H detected belongs to galaxy-1 and there is no contribution from the other galaxy.
" But, its external origin can not be ruled out completely."," But, its external origin can not be ruled out completely."
" A possible explanation of the disturbed and displaced H1 disk of galaxy-1, presumably used to be a nicely rotating spiral, is tidal interaction with galaxy-2."," A possible explanation of the disturbed and displaced H disk of galaxy-1, presumably used to be a nicely rotating spiral, is tidal interaction with galaxy-2."
" Based on the absence of m—2 symmetry andthe fact thatits center being too off from the optical center, it is less likelyto be caused only through tidal interaction."," Based on the absence of $m = 2$ symmetry andthe fact thatits center being too off from the optical center, it is less likelyto be caused only through tidal interaction."
 Gas accretion from satellites, Gas accretion from satellites
Fe structures.,Fe structures.
 Very few of the papers listed in Table |. report that the detections were made on the basis of a systematic and uniform search over a wide spectral range (a notable exception is ?))., Very few of the papers listed in Table \ref{table} report that the detections were made on the basis of a systematic and uniform search over a wide spectral range (a notable exception is \citet{nandra07}) ).
 Either or both these effects may enhance the number of false detections made in the 5—6 keV range., Either or both these effects may enhance the number of false detections made in the $5-6$ keV range.
 Of the 36 lines listed in Table |. (ignoring the BALs) with strength and uncertainty estimates. only five have strengths that are greater than three times their uncertainty (half width of the 90 per cent confidence interval).," Of the $36$ lines listed in Table \ref{table} (ignoring the BALs) with strength and uncertainty estimates, only five have strengths that are greater than three times their uncertainty (half width of the 90 per cent confidence interval)."
 These are: emission at 5.57 keV in NGC 3816 (EW523eV:ο). absorption at 6.2 keV in EI8214643 (EW=—54eV: ο). absorption at 7.7 keV in MCG-5-23-16 (EWx—33eV:2). 1.63 keV and 2.94 keV absorption in PG [2114143 CEWx-—I4+eVand—36eV:?2)..," These are: emission at $5.57$ keV in NGC 3516 \citep[$EW \approx 23$~eV;][]{turner02}, absorption at $6.2$ keV in E1821+643 \citep[$EW \approx -54$~eV;][]{yaqoob05}, , absorption at $7.7$ keV in MCG-5-23-16 \citep[$EW \approx -33$~eV;][]{braito07}, $1.63$ keV and $2.94$ keV absorption in PG 1211+143 \citep[$EW \approx -14$~eV and $-36$~eV;][]{pounds03a,pounds05}."
 These are considered in turn below., These are considered in turn below.
 The emission line in NGC 3516 was found in HHETOS data. but could not be detected in the partly simultaneous ddata. those data gave an upper limit on the line flux an order of magnitude smaller than the ddetection (2).," The emission line in NGC 3516 was found in HETGS data, but could not be detected in the partly simultaneous data, those data gave an upper limit on the line flux an order of magnitude smaller than the detection \citep{turner02}."
. If real. the line flux must have decreased by at least one order of magnitude between the observations.," If real, the line flux must have decreased by at least one order of magnitude between the observations."
 Furthermore it should be noted that the line energy was held fixed during the evaluation of the confidence interval on the flux (seethefootnotetoTable|of ?).., Furthermore it should be noted that the line energy was held fixed during the evaluation of the confidence interval on the flux \citep[see the footnote to Table~1 of][]{turner02}.
 As the line energy was not predicted but obtained from fitting the data. it should have remained a free parameter throughout the caleulation of confidence intervals. otherwise the confidence interval may be artificially reduced.," As the line energy was not predicted but obtained from fitting the data, it should have remained a free parameter throughout the calculation of confidence intervals, otherwise the confidence interval may be artificially reduced."
 In any event. the procedure described by ? differs from the standard procedure adopted in the other cases and so is perhaps best considered a lower limit on the size of the contidence region of that particular line (as indicated in Figure |).," In any event, the procedure described by \citet{turner02} differs from the standard procedure adopted in the other cases and so is perhaps best considered a lower limit on the size of the confidence region of that particular line (as indicated in Figure \ref{fig}) )."
 The absorption line found in the oobservation of EI8214642. was seen in the HEG spectrum but could not be confirmed in the lower signal-to-noise MEG data., The absorption line found in the observation of E1821+642 was seen in the HEG spectrum but could not be confirmed in the lower signal-to-noise MEG data.
 In the case of PG 12114143 the simultaneous detection of multiple lines. and their identification at similar blueshifts. in both CCD (EPIC) and grating (RGS) data from citeppoundsO3a would seem to put this case on firmer ground.," In the case of PG 1211+143 the simultaneous detection of multiple lines, and their identification at similar blueshifts, in both CCD (EPIC) and grating (RGS) data from \\citep{pounds03a} would seem to put this case on firmer ground."
 For completeness. it should be noted that the identification of the lines in terms of highly blueshifted features has been debated. e.g. 22??..," For completeness, it should be noted that the identification of the lines in terms of highly blueshifted features has been debated, e.g. \citet{mckernan05,kaspi06,pounds07,reeves08}."
 ? presented gyrating data of the same object and again reportedabsorption lines. except redshifted not blueshifted.," \citet{reeves05} presented grating data of the same object and again reportedabsorption lines, except redshifted not blueshifted."
 The 7.6 keV absorption line found in the ddata was not detected in more recent ddata (withanupperlimit~4timessmallerthantheoriginalXMM-Newtondete, The $7.6$ keV absorption line found in the data was not detected in more recent data \citep[with an upper limit $\sim 4$ times smaller than the original \xmm\ detection;][]{reeves08}.
ction: ?).. ? and ? studied MCG-S-23-16 using simultaneousNewton.Chandra... aand oobservations.," \citet{braito07} and \citet{reeves07} studied MCG-5-23-16 using simultaneous, and observations."
 The detection of 7.7 keV absorption is based on the EPIC pn spectrum fromXMAM-Newton., The detection of $7.7$ keV absorption is based on the EPIC pn spectrum from.
. The EPIC MOS spectrum is consistent with the pn spectrum but is unable to confirm the presence of the line due to the smaller photon sample size: the ddata show a possible absorption feature. but it was poorly constrained compared to the EPIC pn data: and the ddata were unable to contirm the line detection.," The EPIC MOS spectrum is consistent with the pn spectrum but is unable to confirm the presence of the line due to the smaller photon sample size; the data show a possible absorption feature, but it was poorly constrained compared to the EPIC pn data; and the data were unable to confirm the line detection."
 The lines that gave the largest improvement in the 7 fit statistic were both absorption lines: at 5.9 keV in NGC 3516 (with ASCA)) and at 7.45 keV in NGC 4151 (with XMM-Newtonn., The lines that gave the largest improvement in the $\chi^2$ fit statistic were both absorption lines: at $5.9$ keV in NGC 3516 (with ) and at $7.45$ keV in NGC 4151 (with ).
 The former has no published EW. while the latter has a surprisingly low |JEW|/error ratio given its apparent ettect on the fit.," The former has no published EW, while the latter has a surprisingly low $|EW|/error$ ratio given its apparent effect on the fit."
 However. as ? noted. there is some ambiguity over whether the NGC 4151 feature should be identified with a line or an edge (see their section 8.8.4).," However, as \citet{nandra07} noted, there is some ambiguity over whether the NGC 4151 feature should be identified with a line or an edge (see their section 8.8.4)."
 The Seyfert galaxy NGC 3516 provides one more interesting example., The Seyfert galaxy NGC 3516 provides one more interesting example.
 ? described a 6.08 keV emission line in an EEPIC pn spectrum of NGC 3516 taken in 2001 April (seealso?) and suggested the feature is varied in EW and/or energy., \citet{dovciak04} described a $6.08$ keV emission line in an EPIC pn spectrum of NGC 3516 taken in 2001 April \citep[see also][]{bianchi04} and suggested the feature is varied in $EW$ and/or energy.
 The significance of the feature was assessed using an Z-test (see section 4.13) but no uncertainties were given on the EW and so the observation is not represented in Figure |.., The significance of the feature was assessed using an $F$ -test (see section \ref{sect:detect}) ) but no uncertainties were given on the $EW$ and so the observation is not represented in Figure \ref{fig}.
 The same observation was analysed by ?.. Who claimed a periodic modulation in the spectral shape of a broad 5.6—6.5 keV line based on ~3 'eyeles'.," The same observation was analysed by \cite{iwasawa04}, who claimed a periodic modulation in the spectral shape of a broad $5.6-6.5$ keV line based on $\sim 3$ `cycles'."
 Although this claim is intriguing it is not an independent confirmation of the significance of the ~6.1 keV line: the analysis is an attempt to better understand and model the feature reported by assuming its reality and using the same data. not an independent assessment of it.," Although this claim is intriguing it is not an independent confirmation of the significance of the $\sim 6.1$ keV line; the analysis is an attempt to better understand and model the feature reported by \citet{dovciak04}, assuming its reality and using the same data, not an independent assessment of it."
 This leaves the possibility that many or most of the line detections are false detections caused by random sampling fluctuations., This leaves the possibility that many or most of the line detections are false detections caused by random sampling fluctuations.
 The number of false detections may at first sight appear large. but one must remember that each spectrum fromXAMM-Newton.. eetc.," The number of false detections may at first sight appear large, but one must remember that each spectrum from, etc."
. contains >—50 resolution. elements. and there are many hundreds of spectra (especially considering that. longer observations are routinely split into multiple spectra corresponding to ditterent time intervals. flux levels. ete.).," contains $\gsim 50$ resolution elements, and there are many hundreds of spectra (especially considering that longer observations are routinely split into multiple spectra corresponding to different time intervals, flux levels, etc.)."
 The non-detections from each resolution element of each spectrum contributes an (unpublished) point inside (or just above) the shaded region of Fig. l.., The non-detections from each resolution element of each spectrum contributes an (unpublished) point inside (or just above) the shaded region of Fig. \ref{fig}.
 The human analyst. or an automated search algorithm. has a tendency to focus on the largest fluctuations — this is a selection bias.," The human analyst, or an automated search algorithm, has a tendency to focus on the largest fluctuations – this is a selection bias."
 These may then be subjected to an hypothesis test. and those satisfying some conventional criterion (p.«a with e.g. a= 0.05) may be chosen for publication.," These may then be subjected to an hypothesis test, and those satisfying some conventional criterion $p < \alpha$ with e.g. $\alpha = 0.05$ ) may be chosen for publication."
 The more ‘significant’ the result (i.e. the smaller p is). the more likely it is to be chosen for publication: publication bias.," The more `significant' the result (i.e. the smaller $p$ is), the more likely it is to be chosen for publication: publication bias."
 These two biases act in the same direction but at ditferent stages in the process. and are examples of what Francis Bacon. in hisOreganun.. described as the tendency to “notice the events where they are fulfilled. but where they fail. though this happens much more often. neglect and pass them by” (2).," These two biases act in the same direction but at different stages in the process, and are examples of what Francis Bacon, in his, described as the tendency to “notice the events where they are fulfilled, but where they fail, though this happens much more often, neglect and pass them by” \citep{bacon1620}."
 The result is that the vastly greater number of null results (whether with strong or weak limits) will go largely unpublished. making it dithcult to estimate the global signiticance of any individual detection.," The result is that the vastly greater number of null results (whether with strong or weak limits) will go largely unpublished, making it difficult to estimate the global significance of any individual detection."
 The question that needs to be addressed is whether any given excess or deficit in a spectrum is unlikely to be a sampling fluctuation given a large number of spectra each with many resolution elements., The question that needs to be addressed is whether any given excess or deficit in a spectrum is unlikely to be a sampling fluctuation given a large number of spectra each with many resolution elements.
 The only systematic attempts to address this specitic problem are those of ?. and ?.. both of which describe surveys of narrow. shifted lines in samples of observations.," The only systematic attempts to address this specific problem are those of \citet{nandra07} and \citet{longi06}, both of which describe surveys of narrow, shifted lines in samples of observations."
 One can make an order of magnitude estimate of the number of unpublished non-detections using the followingsimple argument., One can make an order of magnitude estimate of the number of unpublished non-detections using the followingsimple argument.
 Let us assume that >500 spectra have been examined in the last few years and each has >50 resolution elements. the number of independent spectral resolution elements that have been examined must be >2.5x 10+.," Let us assume that $\gsim 500$ spectra have been examined in the last few years and each has $\gsim 50$ resolution elements, the number of independent spectral resolution elements that have been examined must be $\gsim 2.5 \times 10^4$ ."
 If the residuals (after fitting a suitable continuum model) in each of these is approximately, If the residuals (after fitting a suitable continuum model) in each of these is approximately
Skinner Brown found T Tau to have a falling spectra and circular polarization. suggestive of a maguetic origin for the cussion.,"Skinner Brown found T Tau S to have a falling spectrum and circular polarization, suggestive of a magnetic origin for the emission."
 They also observed a flux imereasc. accompanied by a reversal m polarization. from left- to right-handed.," They also observed a flux increase, accompanied by a reversal in polarization, from left- to right-handed."
 Rav ct al. (, Ray et al. (
"1997) claim to have detected polarized ΙσΙΟ associated with T Tau ο, with reeions of left and right handed circular polarization on cither side of the star offset from one another bv some 0.157 along a northlwest-southeast axis.","1997) claim to have detected polarized emission associated with T Tau S, with regions of left and right handed circular polarization on either side of the star offset from one another by some $''$ along a northwest-southeast axis."
 This was interpreted as enisson from a maenetically-threaded outflow. with oppositelv-directed field lines leading to opposite observed circular polarization.," This was interpreted as emission from a magnetically-threaded outflow, with oppositely-directed field lines leading to opposite observed circular polarization."
 Ou a slightly larger scale. White (2000) reported high resolution VLA observations showing extended structure to the northeast of the southern coniponoent.," On a slightly larger scale, White (2000) reported high resolution VLA observations showing extended structure to the northeast of the southern component."
 It has been suggested that T Tu N is seen roughly »ole on to the line of sieht. based on a comparison of the photometric period. the stellar radius aud the rotational xoadenime of photospheric lines.," It has been suggested that T Tau N is seen roughly pole on to the line of sight, based on a comparison of the photometric period, the stellar radius and the rotational broadening of photospheric lines."
" This sugecstion has con used to explain the discrepant extinction between T Tau N aud ο, with the northern component viewed hrough a cavity iu the envelope cleared bv the outflow (Moose et al.."," This suggestion has been used to explain the discrepant extinction between T Tau N and S, with the northern component viewed through a cavity in the envelope cleared by the outflow (Momose et al.,"
 1996)., 1996).
 The jet from the T. Tau S system. iowever. appears to lie nearly in the plane of the sky (vau Langevelde et al.," The jet from the T Tau S system, however, appears to lie nearly in the plane of the sky (van Langevelde et al."
 1991). a view reinforced by modelling of he scattered optical and NIR leh surrounding the stars (Wood et al..," 1994), a view reinforced by modelling of the scattered optical and NIR light surrounding the stars (Wood et al.,"
 2001)., 2001).
 If this is so. he T Tau compoucuts uust have strongly unisalicned rotation axes.," If this is so, the T Tau components must have strongly misaligned rotation axes."
 Tn this paper. we present VLBI observations of the T Tau system.," In this paper, we present VLBI observations of the T Tau S system."
 Our main intention was to probe tle iuner reeions of the system in order to study the collimating nechanisin of the jet at higher resolution than achieved w Rav ct al., Our main intention was to probe the inner regions of the system in order to study the collimating mechanism of the jet at higher resolution than achieved by Ray et al.
 We serendipitously observed the T Tau s svsteni undergo two strong flux increases. acconrpained * changes in circular polarization.," We serendipitously observed the T Tau S system undergo two strong flux increases, accompained by changes in circular polarization."
 The extra fiux iu he secoud of these brighteuiugs appears to be cireularlv polarized. something which to our kuowledec las never been seen previously iu an accreting preanain sequence system.," The extra flux in the second of these brightenings appears to be circularly polarized, something which to our knowledge has never been seen previously in an accreting pre-main sequence system."
" The spatial resolution offered by the VLBI interferometry network should enable us to separate he two components of T Tau ο, but the maps reveal ouly one source."," The spatial resolution offered by the VLBI interferometry network should enable us to separate the two components of T Tau S, but the maps reveal only one source."
 We identity our detection as being T Tau Sh. he M star of the southern binary svstem.," We identify our detection as being T Tau Sb, the M star of the southern binary system."
 We also fail o detect T Tau N at high resolution. despite detecting it clearly with the VLA.," We also fail to detect T Tau N at high resolution, despite detecting it clearly with the VLA."
 In the further discussions. we will refer to the cutire T Tau S system when the single conrponeuts are not explicitly named.," In the further discussions, we will refer to the entire T Tau S system when the single components are not explicitly named."
 The observations were conducted on December 15th. 1999.," The observations were conducted on December 15th, 1999."
 All times referred to are measured from OUT. on this date., All times referred to are measured from $0^{h}$ UT on this date.
 Ten VLBA stations were used. together with the Effelshere 100m telescope and the phased VLA iu ID configuration.," Ten VLBA stations were used, together with the Effelsberg 100m telescope and the phased VLA in `B' configuration."
 The observing frequency was 8.IGITz (Όσσα. N band).," The observing frequency was 8.4GHz (3.6cm, X band)."
 The bright quasars J0131]|1731 aud J0357|2319. Iviug at 3° and 77 from T. Tau respectively. were used as phase calibrators.," The bright quasars J0431+1731 and J0357+2319, lying at $^{\circ}$ and $^{\circ}$ from T Tau respectively, were used as phase calibrators."
 Since astromietry was not he primary goal of the experiment. phase calibrators with positions dotermüned from eeocdectic observations were not required.," Since astrometry was not the primary goal of the experiment, phase calibrators with positions determined from geodectic observations were not required."
 The two phase calibrators used had yOSITIOUS determined from VLBA survey observations and were disposed along a line with T Tau between hei., The two phase calibrators used had positions determined from VLBA survey observations and were disposed along a line with T Tau between them.
 The phase solutious ou both of thei should have allowed us to get precise phase solutions at the tarect ocation., The phase solutions on both of them should have allowed us to get precise phase solutions at the target location.
 Unfortunately the observing schedule coutaimed oo few scans on the secondary calibrator making the signal to noise ratio on J0357|2319 too bad to ect reliable phase solutions., Unfortunately the observing schedule contained too few scans on the secondary calibrator making the signal to noise ratio on J0357+2319 too bad to get reliable phase solutions.
 J0357|2319 was therefore used to check he quality of the phase referencing., J0357+2319 was therefore used to check the quality of the phase referencing.
" Ibvbrid maps of JO3572319 were produced showing au offset from the phase center of approximately πας,", Hybrid maps of J0357+2319 were produced showing an offset from the phase center of approximately mas.
 Since T Tau 5 ies between the two calibrators. we can be sure that the astrometric accuracy at the T Tau S position is at least ΗΝ. aud probably better iu proportion to the relative xoxiuitv of T Tau S to the primary pliase calibrator.," Since T Tau S lies between the two calibrators, we can be sure that the astrometric accuracy at the T Tau S position is at least mas, and probably better in proportion to the relative proximity of T Tau S to the primary phase calibrator."
 This will be discussed further iu Section 3.1., This will be discussed further in Section \ref{sourcepositions}.
 The T Ti S observations were scheduled in 9 blocks consisting of three alteruatious between T Tau S (3 minutes) anc JOIST|1731 (2 minutes). with one scan ou JO357|2319 (2 minutes) at the ους aud the cud of cach block.," The T Tau S observations were scheduled in 9 blocks consisting of three alternations between T Tau S (3 minutes) and J0431+1731 (2 minutes), with one scan on J0357+2319 (2 minutes) at the beginning and the end of each block."
 The total observation time on T Tau s was about 1.5 hours., The total observation time on T Tau S was about 1.5 hours.
 The blocks were separated by gaps of approximately one hour. during which we observed another two targets which are not discussed in the preseut paper.," The blocks were separated by gaps of approximately one hour, during which we observed another two targets which are not discussed in the present paper."
 The VLBI imap of the T Tau field is plotted in Figure I.., The VLBI map of the T Tau S field is plotted in Figure \ref{posnfig}.
" A single poiutlike source can be seen at position a=|21"" 595.1263. à=19?32/5"".730 (equinox 2000)."," A single pointlike source can be seen at position $\alpha=4^{h}~21^{m}~59^{s}.4263$ , $\delta=19^{\circ}~32'~5''.730$ (equinox 2000)."
" This position is epoch 1909005, and has heen corrected for parallax using the Iipparcos value of 7I66 mas."," This position is epoch 1999.958, and has been corrected for parallax using the Hipparcos value of 5.66 mas."
 The astrometric accuracy was estimated from the position of the secondary calibrator JOB572319 in the hybrid map conipared to its expected xositiou from the VLBA survey., The astrometric accuracy was estimated from the position of the secondary calibrator J0357+2319 in the hybrid map compared to its expected position from the VLBA survey.
 The detected source was found to be offset 23 uas to the cast and 17 mas to the south of the catalogue position., The detected source was found to be offset 23 mas to the east and 17 mas to the south of the catalogue position.
" Assunius that this offse arises froni a combination of position errors for the primary and secondary calibration sources, we adopt an uncertainty at the T Tau ο position wil be V2. 20 mas. or Ll mas."," Assuming that this offset arises from a combination of position errors for the primary and secondary calibration sources, we adopt an uncertainty at the T Tau S position will be $\sqrt{2}\times$ 20 mas, or 14 mas."
 This uncertaity is indicated with a circle arom the source in Figure I.., This uncertainty is indicated with a circle around the source in Figure \ref{posnfig}.
 We can compare the position of the VEDI source to the expected positious of T Tau Sa aud Sb in the infrared from the recent literature., We can compare the position of the VLBI source to the expected positions of T Tau Sa and Sb in the infrared from the recent literature.
 These positions are available as offsets from the IR position of T Tau N. We take as the fundamental reference position the Iipparcos position of T Tau N. which is determined in the V-baud.," These positions are available as offsets from the IR position of T Tau N. We take as the fundamental reference position the Hipparcos position of T Tau N, which is determined in the V-band."
 The, The
"lline right)), overlaid on the silicate-based extinction Av(9.85um) map.","line ), overlaid on the silicate-based extinction $A_{\rm V}(9.85~\mu \rm m)$ map."
 The molecular hydrogen emission avoids the obscured nucleus and peaks between one and three pixels away from the highest obscuration., The molecular hydrogen emission avoids the obscured nucleus and peaks between one and three pixels away from the highest obscuration.
 Although the Hy S(2) 12.3 ((Fig. 6)), Although the $_2$ S(2) 12.3 (Fig. \ref{fig:SH-maps2}) )
" has a similar distribution as the Hz S(1) line, its peak emission lies closer (~2.3"", one pixel) to the peak obscuration and to the H.O mega maser than the other Hp lines."," has a similar distribution as the $_2$ S(1) line, its peak emission lies closer $\sim2.3''$, one pixel) to the peak obscuration and to the $_2$ O mega maser than the other $_2$ lines."
 Since the distribution of the silicate-based extinction Av(9.85um) is similar to that of the stellar light-based extinction Ay(H—K) we think that the mid-IR derived extinction towards the starburst ring may be accurate to correct the fluxes of the starburst tracers., Since the distribution of the silicate-based extinction $A_{\rm V}(9.85~\mu \rm m)$ is similar to that of the stellar light-based extinction $A_{\rm V}(H-K)$ we think that the mid-IR derived extinction towards the starburst ring may be accurate to correct the fluxes of the starburst tracers.
" That is, the elongated shape of the silicate map indicates that the obscuration is mostly associated with the starburst ring rather than with the line of sight to the AGN."," That is, the elongated shape of the silicate map indicates that the obscuration is mostly associated with the starburst ring rather than with the line of sight to the AGN."
" Hence, the Ay(9.85um) map can be used to perform an extinction correction for species whose emission emerge from the disk."," Hence, the $A_{\rm V}(9.85~\mu \rm m)$ map can be used to perform an extinction correction for species whose emission emerge from the disk."
 Species associated with the AGN (BLR and NLR gas) may be suffering more extinction than the disk does., Species associated with the AGN (BLR and NLR gas) may be suffering more extinction than the disk does.
" Therefore, the corrected flux derived for the 14.32 ccorresponds only to the best lower limit we can derive from the current data."," Therefore, the corrected flux derived for the 14.32 corresponds only to the best lower limit we can derive from the current data."
 Figures 10 and 11 show the same set of maps as in Figs., Figures \ref{fig:SH-maps1-extcorr} and \ref{fig:SH-maps2-extcorr} show the same set of maps as in Figs.
" 5 and 6,, but corrected for extinction using the estimated Ay(9.85um) map."," \ref{fig:SH-maps1} and \ref{fig:SH-maps2}, but corrected for extinction using the estimated $A_{\rm V}(9.85~\mu \rm m)$ map."
" Note that after correcting for extinction, all the fine-structure emission lines, the average 14.5-15.0 ccontinuum, and the PAH features peak at about the same position, that of the H2O mega maser."," Note that after correcting for extinction, all the fine-structure emission lines, the average 14.5–15.0 continuum, and the PAH features peak at about the same position, that of the $_2$ O mega maser."
 The corrected for extinction H; lines (Fig. 11)), The corrected for extinction $_2$ lines (Fig. \ref{fig:SH-maps2-extcorr}) )
 show an offset of ~2.3” (one pixel) with respect to the water maser., show an offset of $\sim2.3''$ (one pixel) with respect to the water maser.
" However, there is a difference of only «1"" between the centroids obtained from a two-dimensional Gaussian profile fit of all the lines."," However, there is a difference of only $<1''$ between the centroids obtained from a two-dimensional Gaussian profile fit of all the lines."
 The integrated flux densities of the co-added spectra from the whole 10x10 aperture of the IRS/SH map (described in Sec. ??)), The integrated flux densities of the co-added spectra from the whole $\times$ 10 aperture of the IRS/SH map (described in Sec. \ref{sec:results}) )
 are summarized in Table 1.., are summarized in Table \ref{tab-c4:fluxes-fov}.
 In thiswork we include only the fluxes of the most prominent emission lines., In thiswork we include only the fluxes of the most prominent emission lines.
" These fluxes are larger than those reported by ?,, because our 10x10 aperture is larger than the SH staring aperture."," These fluxes are larger than those reported by \citet{bernards09}, , because our $\times$ 10 aperture is larger than the SH staring aperture."
" In Table 1 we also include the fluxes corrected with the silicate-based extinction Ay(9.85 um), andthe line widths (FWHM) as estimated from the Gaussian fit."," In Table \ref{tab-c4:fluxes-fov} we also include the fluxes corrected with the silicate-based extinction $A_{\rm V}(9.85~\mu \rm m)$ , andthe line widths (FWHM) as estimated from the Gaussian fit."
Turbulence plavs au dnmuportaut role in a variety of astrophysical plenomena. incliding amplification of imaenetic fields in planetary aud stellar interiors. magnetization aud angular momentum transport iu accretion disks. seatteriug of cosmic raves. formation of simallscale ceusitv structures im the interstellar medi. heating of the solar corona aud the solar wind.,"Turbulence plays an important role in a variety of astrophysical phenomena, including amplification of magnetic fields in planetary and stellar interiors, magnetization and angular momentum transport in accretion disks, scattering of cosmic rays, formation of small-scale density structures in the interstellar medium, heating of the solar corona and the solar wind."
 Theoretical description of magnetized. plasma turbulence is a complicated task., Theoretical description of magnetized plasma turbulence is a complicated task.
 A valuable euidance for phenomenological modeling is provided by ligh-resolution numerical simulations (c.¢..Braudeubure&Nordland2011:Ivitsilsetal. 2011).," A valuable guidance for phenomenological modeling is provided by high-resolution numerical simulations \citep[e.g.,][]{brandenburg_n11,kritsuk_etal11}."
. On the analytical side. a significant progress is made in the Πιτ of weak turbulence. that is. turbulence consisting of weakly interacting linear waves (e...Newelletal.2001:Galtieretal2000. 2002). which provides a /testbed for fundamental ideas iu the theory of turbulence. such as scale invariance. locality of interactions. euergy cascacles. and anisotropy.," On the analytical side, a significant progress is made in the limit of weak turbulence, that is, turbulence consisting of weakly interacting linear waves \citep[e.g.,][]{newell_nb01,galtier_nnp00,galtier_nnp02}, which provides a testbed for fundamental ideas in the theory of turbulence, such as scale invariance, locality of interactions, energy cascades, and anisotropy."
 Weak MIID turbulence is also interesting on its own as it may plav a role in the interstellar medi. iu the solar corona and the solu wind. iu planetary aud stellar iuagnetosphleres (o...Bhattachar- 2010)..," Weak MHD turbulence is also interesting on its own as it may play a role in the interstellar medium, in the solar corona and the solar wind, in planetary and stellar magnetospheres \citep[e.g.,][]{bhattacharjee_n01,saur_etal02,melrose06,rappazzo_etal2008,chandran10,chandran_etal10}."
 Iu the present Letter we address the problem of nüsiateh between kinetic and magnetic enereies recently reported iu observatious of the solu wind turbulence., In the present Letter we address the problem of mismatch between kinetic and magnetic energies recently reported in observations of the solar wind turbulence.
 Our interest is also motivated bv a similar nüsmateh found iu recent numerical simulations of maenetolvdrodvnamic (ATID) turbulence. suggesting that the phenomenon ax. m fact. have a fundamental nature rather than reflect a possible non-uuiversalitv of solar wiud turbulence.," Our interest is also motivated by a similar mismatch found in recent numerical simulations of magnetohydrodynamic (MHD) turbulence, suggesting that the phenomenon may, in fact, have a fundamental nature rather than reflect a possible non-universality of solar wind turbulence."
 By eiiploviug the framework of weak ΑΠΟ turbulence we present the first direct analytical derivation of residual energv generation in AIIID turbulence., By employing the framework of weak MHD turbulence we present the first direct analytical derivation of residual energy generation in MHD turbulence.
 We demonstrate that kiueticauagnetic equipartitiou ects spontaneouslv broken bv imteractiug random Δανό waves. even if it is present initially.," We demonstrate that kinetic-magnetic equipartition gets spontaneously broken by interacting random Alfvénn waves, even if it is present initially."
 Our results indicate that this effect plavs a fundamental role in the energy cascade. aud it should be taken iuto account in phenomenological 1iodeliug of MIID turbulence.," Our results indicate that this effect plays a fundamental role in the energy cascade, and it should be taken into account in phenomenological modeling of MHD turbulence."
 For our analysis we write the incompressible MIID equations m ters of the Elsasser variables: where the Elsasser variables are defined as z=vdb.," For our analysis we write the incompressible MHD equations in terms of the Elsasser variables: where the Elsasser variables are defined as $\vec z^\pm=\vec
v\pm\vec b$."
 vis the fluctuating plasma velocity. b is the fuctuating inaeuetie field normalized byJ/Izpy. v4=Buffit is the Alfvénn velocity corresponding to the uniformpo magnetic field By. P=(p/py|57/2) includes the plasma pressure p and the magnetic pressure. py ds the constant mass deusitv. and we neglected the terms describing viscous aud resistive dissipation.," $\vec v$ is the fluctuating plasma velocity, $\vec b$ is the fluctuating magnetic field normalized by$\sqrt{4 \pi \rho_0}$ , ${\bf v}_A={\bf B}_0/\sqrt{4\pi \rho_0}$ is the Alfvénn velocity corresponding to the uniform magnetic field ${\bf B}_0$, $P=(p/\rho_0+b^2/2)$ includes the plasma pressure $p$ and the magnetic pressure, $\rho_0$ is the constant mass density, and we neglected the terms describing viscous and resistive dissipation."
 One observes that when z(x.f£)=0. au arbitrary fiction zZ'xti)-F(xcvaÉ) isan exact solution of (13). which represeuts a non-dispersive Alfvén wave propagating parallel or anti-parallel to By with the Alfvéóun speed.," One observes that when $\vec z^\mp(\vec x,t)\equiv 0$, an arbitrary function $\vec z^\pm(\vec x,t)=F^\pm(\vec x\pm\vec v_At)$ is an exact solution of \ref{mhd-elsasser}) ), which represents a non-dispersive Alfvénn wave propagating parallel or anti-parallel to ${\bf B}_0$ with the Alfvénn speed."
 Noulincar interactions are the result of collisions between counter-propagating Alfvén wave packets., Nonlinear interactions are the result of collisions between counter-propagating Alfvénn wave packets.
 This qualitative picture plays a central vole in phenomenological wodels of MIID turbulence: comprehensive reviews of recent analytical aud numerical results can be found in. e.g. (Caltier2009:Sridhar2010:Miniuni2011:Newell&Riuupt2011:Tobiasetal. 2011).," This qualitative picture plays a central role in phenomenological models of MHD turbulence; comprehensive reviews of recent analytical and numerical results can be found in, e.g., \citep{galtier09,sridhar10,mininni11,newell_r11,tobias_cb11}."
. The ideal ΑΠΟ system (1)) conserves the two independent Elsasser energies. ET=1 related to the total enerev aud cross-lelicity. Eat?)=E!/| aud IHE!E .vespoctivolv.," The ideal MHD system \ref{mhd-elsasser}) ) conserves the two independent Elsasser energies, $E^\pm=\langle |\vec z^\pm|^2\rangle /4$, related to the total energy and cross-helicity, $E=E^++E^-$ and $H_c=E^+-E^-$ , respectively."
In a turbulent state. when cherey is supplied to the svsteni at huge scales. both energies E cascade toward sinall scales where they are,"In a turbulent state, when energy is supplied to the system at large scales, both energies $E^\pm$ cascade toward small scales where they are"
object revealed a huge mass of molecular gas (1015.7;2Al. (Brown&vandenBout1991:Solomon.Downesford1992))). a Sevlert emission. spectrum (Ilstonctal. 1994).. high optical polarisation. (Lawrencectal.1993).. and evidence for lensing with a magnification of about 10 at infrared wavelengths (Graham&Liu1995:BroadhurstRobinson 19960).,"object revealed a huge mass of molecular gas $10^{11}h_{50}^{-2}M_{\sun}$ \cite{br,so3}) ), a Seyfert emission spectrum \cite{el}, high optical polarisation \cite{la}, and evidence for lensing with a magnification of about 10 at infrared wavelengths \cite{gra,bro,ei,gr2}."
 These objects appeared. to. presage an entirely new class of infrared. galaxy., These objects appeared to presage an entirely new class of infrared galaxy.
 The source and trigger of the Ht emission in ΗΕτις is currently the subject. of considerable debate., The source and trigger of the IR emission in HLIRGs is currently the subject of considerable debate.
 LILURCs may simply be the high luminosity tail of the ULIRG population. where mergers between evolved galaxics trigecr dust. shroucecl starburst and. AGN activity (Sanders.&Alirabel 19906).," HLIRGs may simply be the high luminosity tail of the ULIRG population, where mergers between evolved galaxies trigger dust shrouded starburst and AGN activity \cite{sa1}."
.. There is evidence from4481 observations (Farrahetal.2002) that at least some HILIBC:s are merging galaxies., There is evidence from observations \cite{far1} that at least some HLIRGs are merging galaxies.
 ;X second. possibility is that ΕΙΠας may be very voung. or ‘primeval’ galaxies.," A second possibility is that HLIRGs may be very young, or 'primeval' galaxies."
 Rowan-Robinson (2000) argues that the majority of the emission at rest-wavelengths > 505m in LLIRGs is due to starburst activity. implying star formation rates 100034;yrt.," Rowan-Robinson \shortcite{rr2} argues that the majority of the emission at rest-wavelengths $>50\mu$ m in HLIRGs is due to starburst activity, implying star formation rates $>1000M_{\sun}yr^{-1}$."
 UW the. rest-frame zw infrared and sub-mm emission [from LILIRGs is due to gaar formation. then the star formation rates would be the jighest for any objects in the Universe.," If the rest-frame far infrared and sub-mm emission from HLIRGs is due to star formation, then the star formation rates would be the highest for any objects in the Universe."
 This would strongly suggest these galaxies are going through their. maximal star formation periods. implving that they are galaxies in 10 first stages of formation.," This would strongly suggest these galaxies are going through their maximal star formation periods, implying that they are galaxies in the first stages of formation."
 | final possibility. is that 1e LR emission arises via some other mechanism (e.g. a ransient I. luminous phase in QSO evolution not trigeered »v interactions). LILIRGs would then be an entirely different ‘lass of object.," A final possibility is that the IR emission arises via some other mechanism (e.g. a transient IR luminous phase in QSO evolution not triggered by interactions), HLIRGs would then be an entirely different class of object."
 In this paper we study the infrared. emission. [rom LILIRGs using data from the optical to the sub-mm., In this paper we study the infrared emission from HLIRGs using data from the optical to the sub-mm.
 We present new sub-nim data for a sample of 11 LILERGs. and use radiative transfer. mocels for starbursts and GN. in conjunction with previously published. HI. photometry. to examine the power source behind the Li emission.," We present new sub-mm data for a sample of 11 HLIRGs, and use radiative transfer models for starbursts and AGN, in conjunction with previously published IR photometry, to examine the power source behind the IR emission."
 Sample selection. observations ancl data analysis are described. in £2.," Sample selection, observations and data analysis are described in 2."
 The radiative transfer models used to evaluate the Li emission. are described in 83., The radiative transfer models used to evaluate the IR emission are described in 3.
 Results are. presented. in 4. and notes on individual sources are given in 85.," Results are presented in 4, and notes on individual sources are given in 5."
 Discussion is presented in 86 and conclusions are summarized in 87., Discussion is presented in 6 and conclusions are summarized in 7.
 We have taken df)=65 km + |. Q=I0 and Q4—0.0.," We have taken $H_{0}=65$ km $^{-1}$ $^{-1}$, $\Omega=1.0$ and $\Omega_{\Lambda}=0.0$."
 The objects in our sample are taken from the FUR selected sample presented by Rowan-Robinson (2000).., The objects in our sample are taken from the FIR selected sample presented by Rowan-Robinson \shortcite{rr2}.
 Unlike nearly all previous stuclies of HLIBCs. the sample in our study. is selected in a manner independent of obscuration. inclination or AGN content.," Unlike nearly all previous studies of HLIRGs, the sample in our study is selected in a manner independent of obscuration, inclination or AGN content."
 Together with the statistical homogeneity and completeness of the parent samples. our saniple is therefore entirely free from ACN bias and suitable for drawing global conclusions about the LLIRG population.," Together with the statistical homogeneity and completeness of the parent samples, our sample is therefore entirely free from AGN bias and suitable for drawing global conclusions about the HLIRG population."
 Observations were mace using the Submillimetre Common User Bolometer Array (SCUBA. Holland οἱ al.," Observations were made using the Submillimetre Common User Bolometer Array (SCUBA, Holland et al."
 1999) on the James Clerk Maxwell Telescope. (JCAL) on October 11-12. 2000 and on January 8S-17 2001., 1999) on the James Clerk Maxwell Telescope (JCMT) on October 11-12 2000 and on January 8-17 2001.
 SCUBA contains two bolometer arravs. one containing 37) pixels and optimized. for observations at S5bOjum. and the other containing 91 pixels and optimized. for observations at 4507/m. In most circumstances. both arrays are operated simultaneously using a cdichroic.," SCUBA contains two bolometer arrays, one containing 37 pixels and optimized for observations at $850\mu$ m, and the other containing 91 pixels and optimized for observations at $450\mu$ m. In most circumstances, both arrays are operated simultaneously using a dichroic."
 Observations were performed using SCULDA's photometry mode. in which data is taken using only a single pixel on cach array.," Observations were performed using SCUBA's photometry mode, in which data is taken using only a single pixel on each array."
 For each integration the secondary mirror was Jigeled so that the selected: bolometer in each array sampled a3. 3 grid with spacing between grid. positions. centred. on the source.," For each integration the secondary mirror was jiggled so that the selected bolometer in each array sampled a $3\times3$ grid with spacing between grid positions, centred on the source."
 During cach integration the secondary mirror is chopped bv 45° in azimuth with a frequency of τς in order to remove sky variations., During each integration the secondary mirror is chopped by 45” in azimuth with a frequency of 7Hz in order to remove sky variations.
" After cach integration. the telescope is then nodded to a reference position 45"" away in azimuth to remove hotspots in the internal SCUBA optics."," After each integration, the telescope is then nodded to a reference position 45” away in azimuth to remove hotspots in the internal SCUBA optics."
 Each object was observed. for approximately 40 minutes. depending on weather conditions.," Each object was observed for approximately 40 minutes, depending on weather conditions."
 Sky opacities during the observations were of moderate quality. with measured opacities at 225Chz from the Caltech Submillimetre Observatory (CSO) in the range 0.075<τους0.15.," Sky opacities during the observations were of moderate quality, with measured opacities at 225Ghz from the Caltech Submillimetre Observatory (CSO) in the range $0.075 < \tau_{225} < 0.15$."
 Calibration observations were mace of Mars or Uranus. or of a secondary. calibrator if no primary calibrator was available.," Calibration observations were made of Mars or Uranus, or of a secondary calibrator if no primary calibrator was available."
 Skyelips were taken before and after cach object and calibrator observation., Skydips were taken before and after each object and calibrator observation.
 The SCUBA User Recluction Facility (SURE) software was used. to reduce the data for all objects., The SCUBA User Reduction Facility (SURF) software was used to reduce the data for all objects.
 The data were first Hatficlelecl ancl despiked., The data were first flatfielded and despiked.
 Atmospheric extinction corrections at 45070. andl 8507420. were derived. by extrapolating from the CSO τους extinction. values following the prescription of Archibald. Wage Jenness (2000)...," Atmospheric extinction corrections at $450\mu$ m and $850\mu$ m were derived by extrapolating from the CSO $\tau_{225}$ extinction values following the prescription of Archibald, Wagg Jenness \shortcite{awj}."
 These extrapolated values were checked. for. consistency against. the observed 450jimn ancl S5bO0f/ extinctions from the skvedips., These extrapolated values were checked for consistency against the observed $450\mu$ m and $850\mu$ m extinctions from the skydips.
 Resiclual sky gradients not removed by nodding ancl chopping were removed. by averaging over all the bolometers in. each array. ane subtracting this value from the measured source lux.," Residual sky gradients not removed by nodding and chopping were removed by averaging over all the bolometers in each array, and subtracting this value from the measured source flux."
 Individual integrations for each source were then concatenated into a single. exposure., Individual integrations for each source were then concatenated into a single exposure.
 Each concatenated dataset was checked. for internal consistency using a Ixolmogorov-Smürnov (Ix-8) test., Each concatenated dataset was checked for internal consistency using a Kolmogorov-Smirnov (K-S) test.
 Finally. Hux calibration for each source was carried. out using the FLUXES package together with the calibrators listed in Table 1..," Finally, flux calibration for each source was carried out using the FLUXES package together with the calibrators listed in Table \ref{scubaobs}."
 In those cases where the calibrator has a larger spatial extent than the 9 point photometry mode jigelemap the Lux of the calibrator in a single photometry jigglemap was used to calibrate the sources. rather than the total flux of the calibrator.," In those cases where the calibrator has a larger spatial extent than the 9 point photometry mode jigglemap the flux of the calibrator in a single photometry jigglemap was used to calibrate the sources, rather than the total flux of the calibrator."
 To model the LR emission. due to starburst activity we, To model the IR emission due to starburst activity we
We next relate the temporal change of μμ to that of c.,We next relate the temporal change of $J_{\rm bin}$ to that of $\omega$.
 We first note that the binarys moment of inertia may be written in terms of the separation πω aucl mass ratio q: Thus. if we again use equation (11) to eliminate πω. the augular momentum may be written as Duriug contraction of the binary. therefore. Taking the time derivative of equation (16) aud applying equation (19) now gives As claimed earlier. the energy aud angular momentum of the binary change at the same relative rates as tliese same quantities in the outgoing wave.," We first note that the binary's moment of inertia may be written in terms of the separation $a_{\rm tot}$ and mass ratio $q$: Thus, if we again use equation (44) to eliminate $a_{\rm tot}$, the angular momentum may be written as During contraction of the binary, therefore, Taking the time derivative of equation (46) and applying equation (49) now gives As claimed earlier, the energy and angular momentum of the binary change at the same relative rates as these same quantities in the outgoing wave."
 The rate of energy trausport by the wave cau be recast in another way that provides a check on our derivation., The rate of energy transport by the wave can be recast in another way that provides a check on our derivation.
 We first uote. after applying equation (17) to the negative of equation (39). tliat [If we then use equation (L1) to eliminate CAM; in favor of πω aucl w. we obtain Thus the energy emission rate cau be written as For fixed ei. the binary components. relative speed scales with c.," We first note, after applying equation (47) to the negative of equation (39), that If we then use equation (44) to eliminate $G\,M_{\rm tot}$ in favor of $a_{\rm tot}$ and $\omega$, we obtain Thus the energy emission rate can be written as For fixed $a_{\rm tot}$, the binary components' relative speed scales with $\omega$."
 This last expressiou thus reproduces the fact that the acoustic enerey radiated by a quadrupolarsource increases as the eighth power of the Mach number (Lighthill 1952)., This last expression thus reproduces the fact that the acoustic energy radiated by a quadrupolarsource increases as the eighth power of the Mach number \citep{L52}. .
Calculations are performed using the AIC method for particle transport ina thermal background described by Longo Diomede (2009) aud by Panarese et al.,Calculations are performed using the MC method for particle transport in a thermal background described by Longo Diomede (2009) and by Panarese et al.
 This method has been receutly validated by comparing the calculated values of binary diffusion cocfiicicuts in different eases with calculations based on the Chapman-Enskoe development extended to high orders (Panareseetal.2011)., This method has been recently validated by comparing the calculated values of binary diffusion coefficients in different gases with calculations based on the Chapman-Enskog development extended to high orders \cite{panarese2011}.
. Although the inethod is described in the above references. here a sclfcousistent short description is provided.," Although the method is described in the above references, here a self-consistent short description is provided."
 The starting poiut is the expression of the real collision frequency for a IT particle moving with velocity v. given by where f is the velocity distribution function of tarect particles and à is the collision pair frequency defined as σ isthe total cross section. g=|vwl is the relative speed of the collision pair aud 0; is the target particle clensity.," The starting point is the expression of the real collision frequency for a H particle moving with velocity v, given by where $f$ is the velocity distribution function of target particles and $\alpha$ is the collision pair frequency defined as $\sigma$ isthe total cross section, $g=\left|{\bf{v}}-{\bf w}\right|$ is the relative speed of the collision pair and $n_b$ is the target particle density."
 The method is based ou the preliminary selection of a nmnaxiuun value for the product go(g) denoted bv (οσο) ρω., The method is based on the preliminary selection of a maximum value for the product $g\sigma(g)$ denoted by $(g\sigma(g))_{max}$ .
" By replaciug οσο) with (galas. iu the integral expression (13) this last can be rewritten iuto the form VOW)=αμ Where This replacement implies a potentially non-plivsical increase of the collision frequency. which can he conrpeusated by using the coucept of null-collision. i.c. the inchision of artificial scattering eveuts which accounts for the difference o,0 but has uo effect on the motion of IT ato1us."," By replacing $g\sigma(g)$ with $(g\sigma(g))_{max}$ in the integral expression \ref{eqnuno}) ) this last can be rewritten into the form $\nu({\bf v}) = \alpha_{max}$ where This replacement implies a potentially non-physical increase of the collision frequency, which can be compensated by using the concept of null-collision, i.e. the inclusion of artificial scattering events which accounts for the difference $\alpha_{max}- \alpha$ but has no effect on the motion of H atoms."
 This solution allows au exact siuple treatiment of collisions ina Test Particle Monte Carlo (TPAIC) iiocdel., This solution allows an exact simple treatment of collisions in a Test Particle Monte Carlo (TPMC) model.
 This umumerical method describes the motion of test particles diluted iu à bulk mediun of target particles., This numerical method describes the motion of test particles diluted in a bulk medium of target particles.
" Iu our case. the svsteii is constituted by test particles of II moving in a IH uniform bulk. im equilibrium at temperature T, aud pressure p."," In our case, the system is constituted by test particles of H moving in a $_{2}$ uniform bulk, in equilibrium at temperature $T_{g}$ and pressure $p$."
 Tuitially test particles are put iu the origin of a threc-dimensional space aud are let to diffuse across the bulk., Initially test particles are put in the origin of a three-dimensional space and are let to diffuse across the bulk.
 Test particles are initialized with the same enerey aud interact with bulk particles by means of binary collisions., Test particles are initialized with the same energy and interact with bulk particles by means of binary collisions.
 For cach collision. the bulls particle velocity is selected according to the Maxwell-DBoltzuian distribution at the temperature 2.ye using a direct method of sampling.," For each collision, the bulk particle velocity is selected according to the Maxwell-Boltzmann distribution at the temperature $T_{g}$, using a direct method of sampling."
 For this purpose. setting (0;=rsind as the velocity conrponeut along the direction. a pair of values of roand J is sampled from ¢@=227i and ro=ΝΤΠινHina) using Ωμ and go unifoniulv distributed between 0 aud 1.," For this purpose, setting $v_{i}=r \sin \vartheta$ as the velocity component along the i-direction, a pair of values of $r$ and $\vartheta$ is sampled from $\vartheta=2\pi \eta_{1}$ and $r=(-2kT_{g}\ln \eta_{2}/m_{bulk})^{1/2}$ , using two random numbers $\eta_{1}$ and $\eta_{2}$ uniformly distributed between $0$ and $1$."
 Finally the value of the i-coniponeut of the thermal velocity iu the equilibrium bulk is sampled as 0;=rcos8., Finally the value of the i-component of the thermal velocity in the equilibrium bulk is sampled as $v_{i}=r \cos \vartheta$.
" Tn order to remove the extra collision events used to equalize the collision frequency to 0,,4,. a further random number jj is compared to the fraction of real collisions given by α jy."," In order to remove the extra collision events used to equalize the collision frequency to $\alpha_{max}$, a further random number $\eta_{3}$ is compared to the fraction of real collisions given by $\alpha/\alpha_{max}$ ."
 Tf ay3 is sinaller than this quantity. the collision is effective.," If $\eta_{3}$ is smaller than this quantity, the collision is effective."
 After an effective collision. the relative velocity vector must be rotated according to two polar aueles. namely J. the scattering angle. aud y. the azimuthal angle.," After an effective collision, the relative velocity vector must be rotated according to two polar angles, namely $\vartheta$, the scattering angle, and $\varphi$, the azimuthal angle."
 This last is uniformly sampled in the interval [0.27]. while the selection of 0 depends on the interaction model.," This last is uniformly sampled in the interval $\pi$ ], while the selection of $\vartheta$ depends on the interaction model."
 Ouce the scattering angle is known. the scattering is treated taking into account the correlation with the old particle velocity using Euler angles: the relative velocity vector after the collision. g. is calculated as where B=UP|ge)? aud g—(qu.du.d.) as above is the relative velocity before the collision.," Once the scattering angle is known, the scattering is treated taking into account the correlation with the old particle velocity using Euler angles: the relative velocity vector after the collision, $g^{*}$, is calculated as where $B=\left(g^{2}_{y}+g^{2}_{z}\right)^{1/2}$ and $\textbf{g}=\left(g_{x},g_{y},g_{z}\right)$ as above is the relative velocity before the collision."
" υ is determined from a quadrature of the iuteractiou potential o(r) based ou the known value of the impact parameter b where py, is the distance of closest approach.", $\vartheta$ is determined from a quadrature of the interaction potential $\phi(r)$ based on the known value of the impact parameter $b$ where $r_{m}$ is the distance of closest approach.
 The value of b is obtained from b—Oya1]., The value of $b$ is obtained from $b = b_{max} \sqrt{\eta_{3}}$.
 o ds siluply given by 0o=25., $\phi$ is simply given by $\phi = 2\pi \eta_{4}$.
 In case of isotropic clastic scattering. cos() 2p.," In case of isotropic elastic scattering, $\cos(\vartheta) = 1 - 2\eta_{5}$ ."
" The motion of the colliding particle of niass 15, relative to the bulk target particle of mass ay is equivalent to the motion of a particle of mass p—ig(n.|orig) relative to a centre of force.", The motion of the colliding particle of mass $m_{c}$ relative to the bulk target particle of mass $m_{t}$ is equivalent to the motion of a particle of mass $\mu=m_{c}m_{t}/\left(m_{c}+m_{t}\right)$ relative to a centre of force.
 The collision energv is calculated as a function of the relative speed g of the interacting pair bv Γιο=ng?2., The collision energy is calculated as a function of the relative speed $g$ of the interacting pair by $E_{coll}=\mu g^{2}/2$.
 As ina binary interaction the centre of mass velocity is a coustant. the velocity of the colliding particle after the collision is given by The tine difference between oue collisiou Gucludiug uull collisions) aud the next one is eiven by the forma where yg is again a random number from a uniformi distribution. O<yg1l.," As in a binary interaction the centre of mass velocity is a constant, the velocity of the colliding particle after the collision is given by The time difference between one collision (including null collisions) and the next one is given by the formula where $\eta_{6}$ is again a random number from a uniform distribution, $0<\eta_{6}<1$."
 This time is inversely proportional to the gas density., This time is inversely proportional to the gas density.
 An appropriate paraueter to measure the degree of thermalization of a IL atom in the bulk medium is therefore provided by the product v7. where r is the average lifetime of a IT atomi before its chemical or diffusion loss.," An appropriate parameter to measure the degree of thermalization of a H atom in the bulk medium is therefore provided by the product $n\tau$, where $\tau$ is the average lifetime of a H atom before its chemical or diffusion loss."
 Calewlatious performed using different values of à and 7. but leaving the value of the product v7 unchanged. produce the same results.," Calculations performed using different values of $n$ and $\tau$ , but leaving the value of the product $n\tau$ unchanged, produce the same results."
 The finite lifetime 7 is obtained iu the simulation bv removing the injectedparticles with a probability 1 before cach collision., The finite lifetime $\tau$ is obtained in the simulation by removing the injectedparticles with a probability $1-exp(-\Delta t/\tau)$ before each collision.
 The removed particle is labelled iu such a way that the computer simulation is uot accountiue for itauvinore., The removed particle is labelled in such a way that the computer simulation is not accounting for itanymore.
 The simulation proceeds until all particles are removed., The simulation proceeds until all particles are removed.
"In the spectrum of wwe detected two emission lines, located at photospheric rest wavelengths 11 1136.5, reffig,pd; ; radialvelocity+35 kkmss~!; Werner 1996).","In the spectrum of we detected two emission lines, located at photospheric rest wavelengths $\lambda\lambda$ 1136.5, \\ref{fig_kpd}; ; radial velocity $+$ $^{-1}$; Werner 1996)."
" We identify these lines as due the 4p?S ?P° transition in the ion reffig,rotrian)).", We identify these lines as due the $^2$ $^2$ $^{\rm o}$ transition in the ion \\ref{fig_grotrian}) ).
"ComparedtotheRitzwavelengthsintheN database, both observed lines are located at wavelengths shorter byAA."," Compared to the Ritz wavelengths in the database, both observed lines are located at wavelengths shorter by."
". Their NIST gf-values are logg;fa=0.23 and 0.46, respectively."," Their NIST $gf$ -values are $\log g_i f_{ik}=0.23$ and 0.46, respectively."
" The third line component of this transition is located at 4 aaccording to NIST, so that in reality we expect it to be found at AAA.."," The third line component of this transition is located at $\lambda$ according to NIST, so that in reality we expect it to be found at $\lambda$."
" Its g f-value, however, is much smaller (logg;fix= —0.47) explaining the fact that we cannot detect it in the observation. ("," Its $gf$ -value, however, is much smaller $\log g_i f_{ik}=-0.47$ ) explaining the fact that we cannot detect it in the observation. ("
This is confirmed by our line-formation calculations.),This is confirmed by our line-formation calculations.)
" We searched for the twoCaX lines in other hot DO white dwarfs and PG1159 stars, but to no avail."," We searched for the two lines in other hot DO white dwarfs and PG1159 stars, but to no avail."
" As we will demonstrate below refresults)), this is a consequence of the extremely high oof ccombined with a relatively low surface gravity."," As we will demonstrate below \\ref{results}) ), this is a consequence of the extremely high of combined with a relatively low surface gravity."
" However, these lines are seen in the very hot, early-type Wolf-Rayet central star 22371 reffig,gc2371))."," However, these lines are seen in the very hot, early-type Wolf-Rayet central star 2371 \\ref{fig_ngc2371}) )."
Forthisob jectwedidnotattemptto fitthe selineswithou¥t aR daeSPIRIRSs BE SAS fectedbyt ," For this object we did not attempt to fit these lines with our (static) model atmospheres, because the profiles might be affected by the stellar wind."
In the course of our model calculations we found that further lines might be detectable in other wavelength regions., In the course of our model calculations we found that further lines might be detectable in other wavelength regions.
" The 4s—4p transition gives rise to a line doublet at AA 1461.8,AA."," The 4s–4p transition gives rise to a line doublet at $\lambda\lambda$ 1461.8,."
". Our models predict only marginal emission features for0005+5106,, which cannot be detected in archival spectra taken with the aboard and high-resolution spectra from theExplorer."," Our models predict only marginal emission features for, which cannot be detected in archival spectra taken with the aboard and high-resolution spectra from the."
". However, for NGC2246 our models predict absorption lines that are possibly present in a spectrum reffig,gc246))."," However, for 246 our models predict absorption lines that are possibly present in a spectrum \\ref{fig_ngc246}) )."
"T hepositionsof thetentativelyidenti f iedabsorption featur: 0.7and-0.2, respectively."," The positions of the tentatively identified absorption features in 246 differ from the NIST wavelengths by $-0.7$ and $-0.2$, respectively."
" CENREMS f'NQrRTEORBEBI, lines of yet another transition of (4d—4f) to be located in the optical UV at 11 3478,AA,, respectively."," From NIST level energies one expects the two strongest lines of yet another transition of (4d–4f) to be located in the optical UV at $\lambda\lambda$ 3478, respectively."
 A high-resolution spectrum of ttaken with ESO's and the UVES spectrograph as part of the SPY survey (Napiwotzki 2003) reveals no line features there., A high-resolution spectrum of taken with ESO's and the UVES spectrograph as part of the SPY survey (Napiwotzki 2003) reveals no line features there.
 Our model for 2246 predicts absorption lines with a depth of only relative to the continuum., Our model for 246 predicts absorption lines with a depth of only relative to the continuum.
" The relatively fast rotation (vssini ss!) smears theline features considerably, and they remain hidden in the noise."," The relatively fast rotation $v$ $i$ $^{-1}$ ) smears theline features considerably, and they remain hidden in the noise."
" Similar weak absorption profiles for these lines are predicted for 0005--5106,, but no appropriate observations are available."," Similar weak absorption profiles for these lines are predicted for , but no appropriate observations are available."
parameter cog)=4+.,parameter $c_{200} = 4$.
 We further assume the eas density profile follows a distribution typical of observed: “cool core? clusters (see VOG: MeCarthy et 22007)., We further assume the gas density profile follows a distribution typical of observed “cool core” clusters (see V06; McCarthy et 2007).
 Finally. we try different hidden cold mass distributions.," Finally, we try different hidden cold mass distributions."
 The predicted ICM temperature distributions are computed by placing the gas in LSE., The predicted ICM temperature distributions are computed by placing the gas in HSE.
 In the case where there is no hidden component. we find the resulting temperature. profiles. quantitatively match the observed. ones.," In the case where there is no hidden component, we find the resulting temperature profiles quantitatively match the observed ones."
 This is not unexpected. given the results of AleCarthy et ((2007) (sec 83.1 of tha study).," This is not unexpected, given the results of McCarthy et (2007) (see 3.1 of that study)."
 However. the predicted. temperature. profiles are quite sensitive to the addition of a centrallv-concentratec cold. component.," However, the predicted temperature profiles are quite sensitive to the addition of a centrally-concentrated cold component."
 In particular. if we require that. the predicted temperature profiles are within the observationa uncertainties (typically. binned. temperature. profiles have a measurement uncertainty). we can rule out. a cole component weighing >=3«101234. within the centra 50 kpe or >zm5«1072A/.. within the central LOO kpe. (," In particular, if we require that the predicted temperature profiles are within the observational uncertainties (typically, binned temperature profiles have a measurement uncertainty), we can rule out a cold component weighing $> \ \approx 3\times10^{12} M_{\odot}$ within the central 50 kpc or $> \ \approx 5\times10^{12}M_{\odot}$ within the central 100 kpc. ("
Note that our results at a fixed radius are not sensitive o the distribution of the cold component. only to the tota enclosed mass within that radius.),"Note that our results at a fixed radius are not sensitive to the distribution of the cold component, only to the total enclosed mass within that radius.)"
 This is small compared to he 5οAL. required to resolve the cluster vs. CIB cliscrepaney for a massive. LOAL. cluster., This is small compared to the $\sim 5\times10^{13} M_\odot$ required to resolve the cluster vs. CMB discrepancy for a massive $10^{15} M_\odot$ cluster.
 The only way o reconcile the cluster vs. CALB ciserepaney ancl maintain consistency with the observed temperature profiles is if the jiididen component is very spatially-extended., The only way to reconcile the cluster vs. CMB discrepancy and maintain consistency with the observed temperature profiles is if the hidden component is very spatially-extended.
 Llowever. such a configuration seems physically contrived.," However, such a configuration seems physically contrived."
 lt therefore appears that the. stellar/cool barvon contribution to the total cluster barvon budget too small to account for the cluster vs. CMD. discrepancy., It therefore appears that the stellar/cool baryon contribution to the total cluster baryon budget too small to account for the cluster vs. CMB discrepancy.
 la particular. the resolved stellar/cold barvon component is a approximately factor of 3 smaller than required.," In particular, the resolved stellar/cold baryon component is a approximately factor of 3 smaller than required."
 Moreover. we have demonstrated that any hidden component. cannot be very massive. otherwise this would. violate the observed temperature profiles of relaxed. clusters.," Moreover, we have demonstrated that any hidden component cannot be very massive, otherwise this would violate the observed temperature profiles of relaxed clusters."
 As discussed. above. the small [fraction of cooled. barvons observed. relative. to. the. predictions ο cooling-only simulations. implies that some form. of non-gravitational heating (7[eedback) is at work in the LOCAL," As discussed above, the small fraction of cooled baryons observed, relative to the predictions of cooling-only simulations, implies that some form of non-gravitational heating (“feedback”) is at work in the ICM."
" Without significant feedback to prevent this overcooling. theoretical moclels are unable to account for the observed. eut-oll of the galaxy. luminosity Function at the bright end or the fact that BC'Gs are. by and large. ""red and dead” (e... Benson et 22003: Bower et 22006). nor would it be possible to explain the lack of N-ray. emission from intracluster gas with temperatures below about 1 keV (e... Peterson et 22003)."," Without significant feedback to prevent this overcooling, theoretical models are unable to account for the observed cut-off of the galaxy luminosity function at the bright end or the fact that BCGs are, by and large, “red and dead” (e.g., Benson et 2003; Bower et 2006), nor would it be possible to explain the lack of X-ray emission from intracluster gas with temperatures below about 1 keV (e.g., Peterson et 2003)."
 Non-gravitational heating also appears to. be necessary to account for the X-ray sealing properties of clusters (e.g. Waiser 1901: Evrard LHenry 1991: Babul et 22002: Voit et 22002: AleCarthy ct 22004).," Non-gravitational heating also appears to be necessary to account for the X-ray scaling properties of clusters (e.g., Kaiser 1991; Evrard Henry 1991; Babul et 2002; Voit et 2002; McCarthy et 2004)."
 Injecting thermal energv into the ICM will cause it to expand (while leaving the dark matter in-tact) and therefore will reduce the gas mass fraction within a given radius., Injecting thermal energy into the ICM will cause it to expand (while leaving the dark matter in-tact) and therefore will reduce the gas mass fraction within a given radius.
 Can this heating explain the cliscrepaney between cluster and CAIB measurements?, Can this heating explain the discrepancy between cluster and CMB measurements?
 To test this. we have computed the bulk energy required to transform clusters with the universal barvon fraction (all in hot gas) into the observed svstems.," To test this, we have computed the bulk energy required to transform clusters with the universal baryon fraction (all in hot gas) into the observed systems."
" For ourclusters. we assume the eas traces the total matter at all raclii and. therefore. within any radius the integrated eas mass fraction is always the best-fitWALA? 3-vear value O,/0,,=0.176."," For our, we assume the gas traces the total matter at all radii and, therefore, within any radius the integrated gas mass fraction is always the best-fit 3-year value $\Omega_b/\Omega_m = 
0.176$."
 Phe total mass profiles (which are dominated by dark matter) are assumed to be the same as those measured by VOG and. PAP. thus we construct a vaseline model cluster for cach of the observed. svstenis.," The total mass profiles (which are dominated by dark matter) are assumed to be the same as those measured by V06 and PAP, thus we construct a baseline model cluster for each of the observed systems."
 The temperature profiles are. determined. by placing the eas in LSE., The temperature profiles are determined by placing the gas in HSE.
 Calculating the total enerey of the eas in hese svstenis (ic.. the summation of the total internal and »otential energies) is then straightforward.," Calculating the total energy of the gas in these systems (i.e., the summation of the total internal and potential energies) is then straightforward."
 For theobserced systems. we extrapolate the eas ancl total density. profiles xvond the maximum radius to which they can be observed. until the integrated gas mass fraction is the universal ratio. (," For the systems, we extrapolate the gas and total density profiles beyond the maximum radius to which they can be observed, until the integrated gas mass fraction is the universal ratio. ("
Note that we also extrapolate the baseline models to this radius to ensure that both the baseline and observed systems iive the same integrated eas and total masses.).,Note that we also extrapolate the baseline models to this radius to ensure that both the baseline and observed systems have the same integrated gas and total masses.).
 We assume he total density profiles continue to follow the NEW form it by VOG anc DAP., We assume the total density profiles continue to follow the NFW form fit by V06 and PAP.
 We try various dillerent. powerlaw extrapolations for the gas density profiles with plausible indices ranging from Oto 2.5., We try various different powerlaw extrapolations for the gas density profiles with plausible indices ranging from $0$ to $-2.5$.
 We find that the minimum energy required to. convert. the baseline models into the observed. svstems occurs when the gas density is constant with radius outside the maximum observable radius., We find that the minimum energy required to convert the baseline models into the observed systems occurs when the gas density is constant with radius outside the maximum observable radius.
 This configuration is perhaps unlikely. but it does provide a useful to the amount of heating required.," This configuration is perhaps unlikely, but it does provide a useful to the amount of heating required."
 Like the baseline models. we place the gas in the observed systems in LISE.," Like the baseline models, we place the gas in the observed systems in HSE."
 Reassuringly. we verify that the resulting temperature profiles are in good. agreement with the observed. profiles.," Reassuringly, we verify that the resulting temperature profiles are in good agreement with the observed profiles."
 The total (minimum) energy. required. to heat. the ICM is just the dillerence of the total energy. of the observed. aud baseline model systems., The total (minimum) energy required to heat the ICM is just the difference of the total energy of the observed and baseline model systems.
 For hot (Zi1 keV) clusters. we find that a substantial amount of energy. is required. ranging between G645.107 eres with a mean of z2.2«107 eres (assuming fuas iS Constant outside the maximunr observable radius - Le. the minimum required energy).," For hot $T_{\rm spec} > 4$ keV) clusters, we find that a substantial amount of energy is required, ranging between $6-45\times10^{62}$ ergs with a mean of $\approx 2.2\times10^{63}$ ergs (assuming $\rho_{gas}$ is constant outside the maximum observable radius - i.e., the minimum required energy)."
 Lt is interesting to compare this minimum energy estimate with the energy that can potentially be deposited. by AGN. the most. powerful source of non-gravitational heating we know of in clusters of ealaxies. (," It is interesting to compare this minimum energy estimate with the energy that can potentially be deposited by AGN, the most powerful source of non-gravitational heating we know of in clusters of galaxies. ("
We note that the energy available from ACN far exceeds that available from stellar feedback.),We note that the energy available from AGN far exceeds that available from stellar feedback.)
 We estimate the amount of AGN energy available to be tapped as follows., We estimate the amount of AGN energy available to be tapped as follows.
 First. we convert. Lin ct (2003)s observed. relationship between stellar mass fraction and total mass within σου (seo equation LO of that. studyi)) into a stellartotal mass relation d.e. Αμ(πο). Λου) assuming a gas mass [fraction of 0.12 within σου (sce 11).," First, we convert Lin et (2003)'s observed relationship between stellar mass fraction and total mass within $r_{500}$ (see equation 10 of that ) into a stellar—total mass relation [i.e., $M_{\rm 
star}(r_{500})-M_{500}$ ] assuming a gas mass fraction of 0.12 within $r_{500}$ (see 1)."
 This relation can be converted into a relationship between the total mass in black holes within roy by Coptimistically) assuming that the entire stellar mass is contained in bulges ancl adopting the black holebulge mass (AMpg—Muy ϱ) relation of Warring Rix (2004)., This relation can be converted into a relationship between the total mass in black holes within $r_{500}$ by (optimistically) assuming that the entire stellar mass is contained in bulges and adopting the black hole—bulge mass $M_{\rm BH}-M_{\rm bulge}$ ) relation of Härring Rix (2004).
 This leads to the following Αμ(rsa)Alsou relation:, This leads to the following $M_{\rm BH}(r_{500})-M_{500}$ relation:
Using the high angular resolution of the and higher sensitivity of the.. combined with optical imaging and spectroscopy and near IR imaging. we studied the most massive and hot fossil galaxy system. which also hosts the the most luminous giant elliptical galaxy observed in fossils.,"Using the high angular resolution of the and higher sensitivity of the, combined with optical imaging and spectroscopy and near IR imaging, we studied the most massive and hot fossil galaxy system, which also hosts the the most luminous giant elliptical galaxy observed in fossils."
" We confirm the result of an earlier ROSAT study that RX J1416.44-2315 meets the requirements of being a ""fossil"" discussed in Jonesetal.(2003).", We confirm the result of an earlier ROSAT study that RX J1416.4+2315 meets the requirements of being a ”fossil” discussed in \citet{jones03}.
. These are the minimum X-ray luminosityof 1077 erg | and a minimum of two magnitude difference in R-band luminosity between the two brightest galaxy belonging to the system (within 0.5)., These are the minimum X-ray luminosityof $10^{42}$ erg $^{-1}$ and a minimum of two magnitude difference in R-band luminosity between the two brightest galaxy belonging to the system (within $r_{vir}$ ).
 The examination was necessary as the gas temperature in the earlier study was underestimated by ~40% resulting in a smaller virial radius within half of which we had to examine the above optical criteria., The examination was necessary as the gas temperature in the earlier study was underestimated by $\sim 40\%$ resulting in a smaller virial radius within half of which we had to examine the above optical criteria.
 The global X-ray distribution is relaxed. except a small region at the core. indicating an absence of any recent group-group merger and major galaxy-galaxy merger indicative of an early formation epoch.," The global X-ray distribution is relaxed, except a small region at the core, indicating an absence of any recent group-group merger and major galaxy-galaxy merger indicative of an early formation epoch."
 The central dominant galaxy is a giant elliptical galaxy aligned with the X-ray emission and hence the dark matter distribution., The central dominant galaxy is a giant elliptical galaxy aligned with the X-ray emission and hence the dark matter distribution.
 The alignment which is also seen in NGC 6482 (Khosroshahi.Jones&Ponman2004). and ESO 3060170 (Sunetal.2004) is an interesting feature which is known to be the consequence of an anisotropie collapse. such as along a filament.," The alignment which is also seen in NGC 6482 \citep{kjp04} and ESO 3060170 \citep{sun04} is an interesting feature which is known to be the consequence of an anisotropic collapse, such as along a filament."
 The alignment effect only works for the central giant elliptical galaxy and is seen in the N-body simulations (Dubinski1998)., The alignment effect only works for the central giant elliptical galaxy and is seen in the N-body simulations \citep{dubinski98}.
 Recent hydrodynamical cosmological simulations. to investigate the effect of baryonic dissipation on halo shapes (Kazantzidis.Zentner&Nagai 2005).. show that there is also a strong correlation between dark matter halo shape and stellar remnant morphology in collisionless mergers.," Recent hydrodynamical cosmological simulations, to investigate the effect of baryonic dissipation on halo shapes \citep{kzd05}, show that there is also a strong correlation between dark matter halo shape and stellar remnant morphology in collisionless mergers."
 Despite the high degree of relaxation and symmetry. there are signs of asymmetry at the core of the system.," Despite the high degree of relaxation and symmetry, there are signs of asymmetry at the core of the system."
 These are seen in the X-ray soft emission within the core., These are seen in the X-ray soft emission within the core.
 This could be a result of AGN activity., This could be a result of AGN activity.
 VLA FIRST catalog gives the integrated flux of the central AGN in the giant elliptical galaxy as 3.39 mJy at 1.4 GHz. with a spatial size of4 arcsec.," VLA FIRST catalog gives the integrated flux of the central AGN in the giant elliptical galaxy as 3.39 mJy at 1.4 GHz, with a spatial size of 4 arcsec."
 Radio observations of the system are currently being performed to investigate the detailed morphology of the radio emission and the status of the AGN duty cycle., Radio observations of the system are currently being performed to investigate the detailed morphology of the radio emission and the status of the AGN duty cycle.
 If fossil groups are old systems. as argued. they should provide ideal environments for the formation of cool cores. due to the absence of any recent major merger which can remove the cool core.," If fossil groups are old systems, as argued, they should provide ideal environments for the formation of cool cores, due to the absence of any recent major merger which can remove the cool core."
 The temperature profile shows no significantly cooler core. suggesting the existence of a heating mechanism.," The temperature profile shows no significantly cooler core, suggesting the existence of a heating mechanism."
 Given the limited statistics and the point spread function. the presence of a small €50 Kpe) cooler region cannot be ruled out.," Given the limited statistics and the point spread function, the presence of a small $<50$ kpc) cooler region cannot be ruled out."
 However the absence of cool core is also found in the case of the nearest observed fossil. NGC 6482 (Khosroshahi.Jones&Ponman where the core is resolved with | kpe resolution.," However the absence of cool core is also found in the case of the nearest observed fossil, NGC 6482 \citep{kjp04} where the core is resolved with 1 kpc resolution."
 In addition ESO 3060170 (Sunetal.2004). shows a cool core much smaller than the expected cooling region for a system of few 10 AZ.., In addition ESO 3060170 \citep{sun04} shows a cool core much smaller than the expected cooling region for a system of few $10^{14}$ $M_\odot$.
 As mentioned above there is evidence of AGN activity in the core of RX 1141642315. similarly to ESO 3060170.," As mentioned above there is evidence of AGN activity in the core of RX J1416+2315, similarly to ESO 3060170."
 This study shows that the X-ray temperature from a previous ROSAT study was underestimated by about half its value., This study shows that the X-ray temperature from a previous ROSAT study was underestimated by about half its value.
 It is therefore important to study the effects of these changes on the sealing relations involving X-ray quantities., It is therefore important to study the effects of these changes on the scaling relations involving X-ray quantities.
 This is the subject of a separate study concentrating on the X-ray scaling relations of fossi galaxy groups., This is the subject of a separate study concentrating on the X-ray scaling relations of fossil galaxy groups.
 The initial argument for the early formation of fossils relied on the relaxed X-ray morphology. absence of recent mergers and the large gap (2 mag) in the optical luminosity function of their constituent galaxies.," The initial argument for the early formation of fossils relied on the relaxed X-ray morphology, absence of recent mergers and the large gap (2 mag) in the optical luminosity function of their constituent galaxies."
 Their mass distribution. and in particular their mass concentration parameter. can provide more solid and independent evidence for their early formation.," Their mass distribution, and in particular their mass concentration parameter, can provide more solid and independent evidence for their early formation."
 Wechsler(2002) show that halos which have not had a major merger since z = 2 ure more concentrated in comparison to systems which had experienced a major merger since then., \citet{wechs02} show that halos which have not had a major merger since z = 2 are more concentrated in comparison to systems which had experienced a major merger since then.
 We measure copo and find that the mass concentration for this system is relatively high compared to clusters of similar mass (Pratt&Arnaud2005)., We measure $c_{200}$ and find that the mass concentration for this system is relatively high compared to clusters of similar mass \citep{pratt05}.
.. This is excluding the stellar mass contribution to the total gravitationa mass from the central giant elliptical galaxy., This is excluding the stellar mass contribution to the total gravitational mass from the central giant elliptical galaxy.
 We have already reported a very high mass concentration parameter for NGC 6482 (Khosroshahi.Jones&Ponman2004)., We have already reported a very high mass concentration parameter for NGC 6482 \citep{kjp04}.
 In addition. Sunetal.(2004) find a high mass concentration (~ 8.9) when they fit a NFW profile for ESO 3060170. another massive fossil.," In addition, \citet{sun04} find a high mass concentration $\sim 8.9$ ) when they fit a NFW profile for ESO 3060170, another massive fossil."
 If fossils are old systems. their core must have formed in an early stages when the density of matter was higher. resulting in formation of a dense and compact core whilst later accretion of mass builds an extended halo.," If fossils are old systems, their core must have formed in an early stages when the density of matter was higher, resulting in formation of a dense and compact core whilst later accretion of mass builds an extended halo."
 Relatively higher mass concentration of low mass systems compared to more massive halos are seen in cosmological simulations (Bullocketal.2001., Relatively higher mass concentration of low mass systems compared to more massive halos are seen in cosmological simulations \citep{bullock01}.
 The central galaxy in this cluster has disey isophotes despite the fact that majority of giant elliptical galaxies at the cores of clusters and groups show boxy isophotes (Kormendy&Bender 1996)., The central galaxy in this cluster has discy isophotes despite the fact that majority of giant elliptical galaxies at the cores of clusters and groups show boxy isophotes \citep{kb96}.
. There are other examples of discy isophotes in fossils? central galaxies., There are other examples of discy isophotes in fossils' central galaxies.
 NGC 6482 also shows highly disey isophotes., NGC 6482 also shows highly discy isophotes.
 The original fossil. RX 3134044017 has pure elliptical isophotes.," The original fossil, RX J1340+4017 has pure elliptical isophotes."
 We are currently studying the possibility that this. non-boxy isophotes. is a generic property of giant ellipticals at the core of fossil systems.," We are currently studying the possibility that this, non-boxy isophotes, is a generic property of giant ellipticals at the core of fossil systems."
 A more accurate galaxy luminosity function. compared to the first published for a fossil (Jones etal 2001). is presented in this study.," A more accurate galaxy luminosity function, compared to the first published for a fossil (Jones etal 2001), is presented in this study."
 The contirmed members form a LF which is steeper than most of the galaxy luminosity functions presented for the groups., The confirmed members form a LF which is steeper than most of the galaxy luminosity functions presented for the groups.
" We finde=L28340.28 and Mj,=—2040+ 0.22. by fitting the Schecter function to the LF of the spectroscopically confirmed member galaxies."," We find $\alpha= -1.23\pm0.28$ and $M^\star_R= -20.40\pm0.22$ , by fitting the Schecter function to the LF of the spectroscopically confirmed member galaxies."
 This agrees well with the results of an independent study of this cluster by (2006)., This agrees well with the results of an independent study of this cluster by \citet{mendes06}. .
"velocity such that if O exceeds Q,,.;¢. then the magnetic switch also should be triggered aud a fast jet of the type found w MECPL should cusue.","velocity such that if $\Omega$ exceeds $\Omega_{crit}$, then the magnetic switch also should be triggered and a fast jet of the type found by MEGPL should ensue."
" The author tested this hwpothesis by performing munerical ATID simulations simular to MEGPL. but ποιαπο B,y fixed aud varving the disk rotation speed instead."," The author tested this hypothesis by performing numerical MHD simulations similar to MEGPL, but holding $B_{p0}$ fixed and varying the disk rotation speed instead."
 The left panel of Figure 3) shows the resulting jet speed as a function of O., The left panel of Figure \ref{sim_fig} shows the resulting jet speed as a function of $\Omega$ .
 For angular velocities below Oi the terminal jet/wind speed is approximately the escape speed at the injection point., For angular velocities below $\Omega_{crit}$ the terminal jet/wind speed is approximately the escape speed at the injection point.
 When ©>µε the naterial is unbound and quickly accelerated (n a wellcollimated jet. similar to MEGPL) to a speed such that a laree fraction of the power Lajup is in the form of sinctic energy (see the right paucl of Figure 3)).," When $\Omega > \Omega_{crit}$, the material is unbound and quickly accelerated (in a well-collimated jet, similar to MEGPL) to a speed such that a large fraction of the power $L_{MHD}$ is in the form of kinetic energy (see the right panel of Figure \ref{sim_fig}) )."
" It is this strong difference iu the character of the outflows below and above O,,5; that is so simular to the FR I/II division: while the total power varies suoothlv as λε exceeds Log. the character of the outflow varies dramatically frou a trausonic. maguetic-cucrey-dominated jet to a highly Supersonic. kinetic-cacrev-domunated onc."," It is this strong difference in the character of the outflows below and above $\Omega_{crit}$ that is so similar to the FR I/II division: while the total power varies smoothly as $L_{MHD}$ exceeds $L_{crit}$, the character of the outflow varies dramatically from a transonic, magnetic-energy-dominated jet to a highly supersonic, kinetic-energy-dominated one."
 The details of the behavior of the curves in Fieure 2. can be understood qualitatively as follows., The details of the behavior of the curves in Figure \ref{sim_fig} can be understood qualitatively as follows.
 When O<O4. the injected. plasma is no louger supported fully by rotation.," When $\Omega < \Omega_{K}$, the injected plasma is no longer supported fully by rotation."
 Therefore. as soon as it emerges from the disk boundary. it dynamically contracts m cvlindical radius. seckiug a iw Ieplerian equilibrium.," Therefore, as soon as it emerges from the disk boundary, it dynamically contracts in cylindrical radius, seeking a new Keplerian equilibrium."
" The jet then is ejected. frou his new. smaller. radius Ri with a speed similar to the Keplerian velocity there and a power slightly larger than hat given by equation (1)). because of compressional increase in the axial maguetic field componcut (BY,xRi 7j TE"," The jet then is ejected from this new, smaller, radius $R_{0} '$ with a speed similar to the Keplerian velocity there and a power slightly than that given by equation \ref{nonrel_mhd_power_eq}) ), because of compressional increase in the axial magnetic field component $B_{Z0}' \propto {R_{0} '}^{-2}$ )."
 Q de decreased. further. the plasuia contracts nore. increasing the eventual jet velocity slielhtlv.," If $\Omega$ is decreased further, the plasma contracts more, increasing the eventual jet velocity slightly."
 Ou the other haud. when ©Or. the reverse occurs. with he plasma expanding to a new injection radius.," On the other hand, when $\Omega > \Omega_{K}$, the reverse occurs, with the plasma expanding to a new injection radius."
" As ο increases, the value of O.,;; calculated from equation (6)) changes as well. chiefly because only the radial coumponcu of the maenetic field is iuportaut iu the expaucded stat (BixRy! y, "," As $\Omega$ increases, the value of $\Omega_{crit}$ calculated from equation \ref{omega_crit_eq}) ) changes as well, chiefly because only the radial component of the magnetic field is important in the expanded state $B_{R0}' \propto {R_{0} '}^{-1}$ )."
"This increases the critical value to Y=QLVLT""uar (where Vig is the radial component of the Alfvénn velocity at the new injection radius). which is ereater than the value in equation (6))."," This increases the critical value to ${\Omega_{crit} '} = {\Omega_{K} '} {V_{esc} '}^{2} 
/ {V_{AR} '}^{2}$ (where $V_{AR} '$ is the component of the Alfvénn velocity at the new injection radius), which is greater than the value in equation \ref{omega_crit_eq}) )."
" Finallv. wel above the magnetic switch (παν Q7 O,.,4) the well jet flow gives wav to a less-collimated radia wind. and the iuerease in outflow speed with rotation speed begins to break down."," Finally, well above the magnetic switch (when $\Omega >> \Omega_{crit}$ ) the well-collimated jet flow gives way to a less-collimated radial wind, and the increase in outflow speed with rotation speed begins to break down."
 It is not clear whether such a decrease at high 9 will occur in the relativistic case. where helt cvlinder effects will be important.," It is not clear whether such a decrease at high $\Omega$ will occur in the relativistic case, where light cylinder effects will be important."
 Relativistic MIID simulations are necded to investigate this behavior., Relativistic MHD simulations are needed to investigate this behavior.
 Some brief conuueuts should be mace about boundary conditions., Some brief comments should be made about boundary conditions.
" Iu these simulations. the kinematic auc magnetic properties of the inflowing matter at the equator (the ""disk corona’) are fixed. throughout the calculation."," In these simulations, the kinematic and magnetic properties of the inflowing matter at the equator (the “disk corona”) are fixed throughout the calculation."
 Uowever. some authors allow these paramcters to be determined by the solution. particularly the azimuthal aud radial components of the field at the equator and the inflow velocity there (sec.e.g... Ustyugovactal. 1999)).," However, some authors allow these parameters to be determined by the solution, particularly the azimuthal and radial components of the field at the equator and the inflow velocity there (see, \cite{ust99}) )."
 Both of these sets of conditions (fixed or seclt-adjusting) are easv to apply uuuericallv. aud cach is appropriate under different circumstances.," Both of these sets of conditions (fixed or self-adjusting) are easy to apply numerically, and each is appropriate under different circumstances."
 The latter is useful for gencratiug steady state solutious to compare with Blandford&Pavue(1982) s solf-3inilar. semiu-analytic models.," The latter is useful for generating steady state solutions to compare with \cite{bp82}' 's self-similar, semi-analytic models."
 Iu this case. conditions at the base of the outflow re-adjust themselves. subject only to the streneth of the vertical magnetic field. the density of inflow. and the speed aud shear of the disk rotation.," In this case, conditions at the base of the outflow re-adjust themselves, subject only to the strength of the vertical magnetic field, the density of inflow, and the speed and shear of the disk rotation."
 The svstem finally reaches a steady state with parameters line ou (or near) the Blaudford-Pavue solution surface im (i.A.£o)-space (sce DP). with the flow varving s200thlv from the outer regions all the way back o the equatorial boundary.," The system finally reaches a steady state with parameters lying on (or near) the Blandford-Payne solution surface in $(\kappa, \lambda, \xi_{0}')$ -space (see BP), with the flow varying smoothly from the outer regions all the way back to the equatorial boundary."
 Ou the other haud. fixing the inflow couditions on he equator. as we have done here. is perhaps more appropriate for studving the astrophysical situation where he magnetic field is anchored in an accretion disk (or other type of rotator).," On the other hand, fixing the inflow conditions on the equator, as we have done here, is perhaps more appropriate for studying the astrophysical situation where the magnetic field is anchored in an accretion disk (or other type of rotator)."
 In that case the local magnetic Ποια is deteriuued by processes mterual to the rotator. rot bv the outflow that originates above it.," In that case the local magnetic field is determined by processes internal to the rotator, not by the outflow that originates above it."
 When the )ndary conditions are fixed. a steady flow still develops above the disk. aud the azimuthal aud radial fields aud How velocity in the corona also re-adjust.," When the boundary conditions are fixed, a steady flow still develops above the disk, and the azimuthal and radial fields and flow velocity in the corona also re-adjust."
 However. because hat resulting outflow docs not necessarily match outo the fixed disk boundary couditious. there is a transition region vetween the disk aud the base of the outflow region. m which the maguetic field and velocity dramatically chauge Indmagnitude and direction.," However, because that resulting outflow does not necessarily match onto the fixed disk boundary conditions, there is a transition region between the disk and the base of the outflow region, in which the magnetic field and velocity dramatically change in magnitude and direction."
 It is not vet fully clear whether Us transition region merely provides a wav to connect an outflow solution that would have occurred anyway to ie disk conditions. or whether it plavs an active role iu determining what outflow solution is chosen.," It is not yet fully clear whether this transition region merely provides a way to connect an outflow solution that would have occurred anyway to the disk conditions, or whether it plays an active role in determining what outflow solution is chosen."
 What is clear is that the values of the parameters chosen on je disk. the transition regiondo matter: the flow Llauges dramatically with the streneth of the poloidalnd azithal field) specified aud. to a lesser extent. with ie angle of the poloidal feld (Meier.Payne.&Lind 1996: MEGPL).," What is clear is that the values of the parameters chosen on the disk the transition region matter: the flow changes dramatically with the strength of the poloidal azimuthal field specified and, to a lesser extent, with the angle of the poloidal field \cite{mpl96}; MEGPL)."
 Furthermore. jet simulations in which a substantialequatorial field (azimuthal plus radial) is specified seein to display more collimatiou. presuniablv because of the strong azimuthal field that is initially prescut or that develops after a few turus of the disk.," Furthermore, jet simulations in which a substantial field (azimuthal plus radial) is specified seem to display more collimation, presumably because of the strong azimuthal field that is initially present or that develops after a few turns of the disk."
 The proposed model of jet production from the accreting. rotating black hole system is esscutially a marriage between the disk acceleration model of Dlaudtord&Pavue(1982). and the rotating black hole acceleration model of Blandford&Zuajek(977) aud MacDonald&Thorne (1982).," The proposed model of jet production from the accreting, rotating black hole system is essentially a marriage between the disk acceleration model of \cite{bp82} and the rotating black hole acceleration model of \cite{bz77} and \cite{mt82}."
. The former authors considered the case when the maguetic field threads ouly the IKepleriau portion of the disk. while the latter examined the situation when it threads both hole aud disk.," The former authors considered the case when the magnetic field threads only the Keplerian portion of the disk, while the latter examined the situation when it threads both hole and disk."
 Iu this paper we cousider the intermediate case when the maguetic field threads both a region of the svstem outside the static lamit that docs not experience appreciable frame drageine and a region closer to the black hole that does the ereosphere., In this paper we consider the intermediate case when the magnetic field threads both a region of the system outside the static limit that does not experience appreciable frame dragging and a region closer to the black hole that does — the ergosphere.
 The main features of this model are the following: (1) Uulike the BZ model. the maeuetic field does not have to thread the horizon itself in order to extract the black hole rotational energv. which avoids the problems discussed bv Puuslv&Coroniti (1990a).. aud the primary source of material inthe jet can be a corona created by the accretion disk itself. not necessarily particles created in spark gaps near the horizon. (," The main features of this model are the following: (1) Unlike the BZ model, the magnetic field does not have to thread the horizon itself in order to extract the black hole rotational energy, which avoids the problems discussed by \cite{pc90a}, and the primary source of material inthe jet can be a corona created by the accretion disk itself, not necessarily particles created in spark gaps near the horizon. ("
2) It is unlike theBP modelbecause the disk aud feld there are frame-drageedao by the hole relative to,2) It is unlike theBP modelbecause the disk and field there are frame-dragged by the hole relative to
1G.,0.6.
 The other raw measurements for differcut multiplicitv eroups are preseuted in the correspouding graphs aud ables. however such is the inherent biasing iu these lower uultiplicitv data that little can be eleaned from their analysis alone.," The other raw measurements for different multiplicity groups are presented in the corresponding graphs and tables, however such is the inherent biasing in these lower multiplicity data that little can be gleaned from their analysis alone."
 Iustead. our last use of the raw data is to determine the cousistency between the real universe aud he ACDAL dependent mock universe.," Instead, our last use of the raw data is to determine the consistency between the real universe and the $\Lambda$ CDM dependent mock universe."
 It is completely justifiable to make a comparison between the raw data and thei multiplicity associated mock data. since the shape biasing (being a sampling artifact) will be the sane.," It is completely justifiable to make a comparison between the raw data and their multiplicity associated mock data, since the shape biasing (being a sampling artifact) will be the same."
 By directly comparing the minimal 47 fits and considering the cross catalog ISS tests we will have au iusieht iuto the reliability of both the erouping algoritliu and the cosmological model used., By directly comparing the minimal $\chi^2$ fits and considering the cross catalog KS tests we will have an insight into the reliability of both the grouping algorithm and the cosmological model used.
 When considering table 1. and table we can see remarkably good agreenient between the real aud mock data (c.f.Pazetal. 2006)., When considering table \ref{tableraw} and table \ref{kstruemock} we can see remarkably good agreement between the real and mock data \citep[c.f.][]{paz06}.
. Cousidering the fits. we fd near perfect agreement for 2 between the real and mock data. the major disagreenieuts beige for multiplicities of 5-9 (for prolate distributions the real data las)=0.31 compared to 3=0.30 for the mock data) and for 10-19 (for oblate istributious the real data has σ=0.0L compared to ao=0.1 for the mock data).," Considering the fits, we find near perfect agreement for $\bar{\beta}$ between the real and mock data, the major disagreements being for multiplicities of 5-9 (for prolate distributions the real data has $\bar{\beta}=0.34$ compared to $\bar{\beta}=0.30$ for the mock data) and for 10-19 (for oblate distributions the real data has $\sigma=0.04$ compared to $\sigma=0.1$ for the mock data)."
 We fiud for the most part extremely similar 4? between the two data sets; another mvwdication that they are truly describing the same aποιο distributions.," We find for the most part extremely similar $\chi^2$ between the two data sets, another indication that they are truly describing the same underlying distributions."
 Furthermore. when looking at —he KS test results we find that all direct commparisous ‘van be drawn from the same underlying distribution to within lo expectations.," Furthermore, when looking at the KS test results we find that all direct comparisons can be drawn from the same underlying distribution to within $1\sigma$ expectations."
 The self similaritics between the ata for different uniltiplicities arc not quite as οσους tables 1. and 5)). however. the real aud mock data th. reject shared self similar distributious to a high ceree of significance.," The self similarities between the data for different multiplicities are not quite as good (tables \ref{kstrue} and \ref{ksmock}) ), however, the real and mock data both reject shared self similar distributions to a high degree of significance."
 As would be expected from the fitting results. the most similar distributions are those for multiplicitics of 20| aud 10-20 for both the real aud mock data (2.9% and 0.115€ respectivelv).," As would be expected from the fitting results, the most similar distributions are those for multiplicities of 20+ and 10-20 for both the real and mock data $2.9\%$ and $0.14\%$ respectively)."
 It can be seen foi the 4? contour data that these uniltiplicities ouly agree upon paramcters when cousidering the 26 error contours. which is cousistent with this KS test statistic.," It can be seen from the $\chi^2$ contour data that these multiplicities only agree upon parameters when considering the $2\sigma$ error contours, which is consistent with this KS test statistic."
 We now consider the convolution corrected data with weighted histogram bins., We now consider the convolution corrected data with weighted histogram bins.
 The results for the real and mock forms of the 2PIGG catalog can be seen iu figures 10.. 11.. 12 and 13.. and in table 2..," The results for the real and mock forms of the 2PIGG catalog can be seen in figures \ref{fitschi5+}, \ref{fitschi510}, \ref{fitschi1020} and \ref{fitschi20+}, and in table \ref{tablecorrect}."
 Siguificautly. we now fud that the distributions are not necessarily prolate.," Significantly, we now find that the distributions are not necessarily prolate."
 The complete data aud the nultiplicitv cut for 5-9 is better fit by an oblate distribution. the cut-off for 19 is equally well fit by either. aud only the 20) cut is better ft bv a prolate distribution aud not with a definite rejection of an oblate distribution.," The complete data and the multiplicity cut for 5-9 is better fit by an oblate distribution, the cut-off for 10-19 is equally well fit by either, and only the 20+ cut is better fit by a prolate distribution and not with a definite rejection of an oblate distribution."
 If oblate distributions are accepted for the corrected data then we see a strong divide hetween the low iuultiplicitv cut (5-9) aud the other two larger multiplicity cuts., If oblate distributions are accepted for the corrected data then we see a strong divide between the low multiplicity cut (5-9) and the other two larger multiplicity cuts.
 The former has a distribution mean ο=0.2 whilst the latter two have a nean 3~0.3: in all cases 64=0.1., The former has a distribution mean $\bar{\beta}=0.2$ whilst the latter two have a mean $\bar{\beta}\sim0.3$: in all cases $\sigma_{\beta}=0.1$.
" Tf we still assume prolate distributions. despite evidence of strong rejection in our 5-9 imultipliitv cut. then we still see a small variation iu distribution mean: ο=0.12+0.01 for 5-9, ο=OLLE0.0L for 10-19 and 3=OIG40.02 for 20|."," If we still assume prolate distributions, despite evidence of strong rejection in our 5-9 multiplicity cut, then we still see a small variation in distribution mean: $\bar{\beta}=0.42\pm0.01$ for 5-9, $\bar{\beta}=0.44\pm0.01$ for 10-19 and $\bar{\beta}=0.46\pm0.02$ for 20+."
" The latter two have σι=GL whilst the 5-9 cut has 0,=0.2."," The latter two have $\sigma_{\beta}=0.14$, whilst the 5-9 cut has $\sigma_{\beta}=0.2$."
" Within the lo error contours the two higher multiplicity cuts share the same distribution. §~LL and o,=0.11. whilst the 5-9 cut is strongly indicated to be a truly different distribution. Wing outside the Jo error contours."," Within the $1\sigma$ error contours the two higher multiplicity cuts share the same distribution, $\bar{\beta}\sim0.44$ and $\sigma_{\beta}=0.14$, whilst the 5-9 cut is strongly indicated to be a truly different distribution, lying outside the $3\sigma$ error contours."
 This fuding is cousistant with the previous INS tests which showed the two high multiplicity cuts to be the most simular: once the distributions are corrected they become identical., This finding is consistant with the previous KS tests which showed the two high multiplicity cuts to be the most similar; once the distributions are corrected they become identical.
" Evidence of oblate shapes Las been preseuted before in literature: Fasanoetal.(1993) and Orlov.Petrova&Tarautaey(2001). being notable examples. though with a smaller sample size aud different inultiplicitv. limits respectively,"," Evidence of oblate shapes has been presented before in literature: \citet{fasano93} and \citet{orlov01} being notable examples, though with a smaller sample size and different multiplicity limits respectively."
 These authors also found a non-rejectable oblate distribution. although Fasanoetal.(1993). favoured the prolate fit.," These authors also found a non-rejectable oblate distribution, although \cite{fasano93} favoured the prolate fit."
 Orlov.Petrova&Tarautaey(2001). found au axial ratio mean of 3:1 for their sample of small galaxy groups (1nultiplicities of 3-8). a figure that can be considered consistent witli our data (0=0.2 for group imultiplicities 5-9) eiven that the statistics are significautly poorer (2703 calaxics in [85 eroups).," \cite{orlov01} found an axial ratio mean of 3:1 for their sample of small galaxy groups (multiplicities of 3-8), a figure that can be considered consistent with our data $\bar{\beta}=0.2$ for group multiplicities 5-9) given that the statistics are significantly poorer (2703 galaxies in 485 groups)."
 Prolate fits have been found to be the best ft for huge cluster samples (Cooray2000) which would be consistant with our 20) multiplicity cut (these systems being the most comparable to cluster environments)., Prolate fits have been found to be the best fit for large cluster samples \citep{cooray00} which would be consistant with our 20+ multiplicity cut (these systems being the most comparable to cluster environments).
 Reearding comparisous to the mock catalog we once again see cousistent fittime between the real aud mock data. accounting for the corresponding error ellipses the fitting between the real and mock data is actually better than for the raw data.," Regarding comparisons to the mock catalog we once again see consistent fitting between the real and mock data, accounting for the corresponding error ellipses the fitting between the real and mock data is actually better than for the raw data."
 However. the mock catalogs do uot find such similar distributions for the 10-19 aud 20| cuts. either for oblate (3=0.28. 3=0.36 respectively} or prolate (3=QLL 9j=0.18 respectively) solutions.," However, the mock catalogs do not find such similar distributions for the 10-19 and 20+ cuts, either for oblate $\bar{\beta}=0.28$, $\bar{\beta}=0.36$ respectively) or prolate $\bar{\beta}=0.44$, $\bar{\beta}=0.48$ respectively) solutions."
 It should be noted that a KS test cannot be done iu this situation since it is onlv applicable to data with nuweielited histogram frequency poiuts., It should be noted that a KS test cannot be done in this situation since it is only applicable to data with unweighted histogram frequency points.
 Perhaps the most interesting aspect of this analysis is the tentative evidence it shows for generally more spherical groups for larger inultiplicities. not just as a ποΊσα. artifact.," Perhaps the most interesting aspect of this analysis is the tentative evidence it shows for generally more spherical groups for larger multiplicities, not just as a numerical artifact."
 This would be cousistaut with arge inultipliitv eroups being amore virialised aud dynamically more evolved. consistant with hierarchical ormnation of πα. groups collapsing alone fibuneuts and larger eroups formine at nodes.," This would be consistant with large multiplicity groups being more virialised and dynamically more evolved, consistant with hierarchical formation of small groups collapsing along filaments and larger groups forming at nodes."
 The features that account for the ligh cllipticity sul sample iu low uultiplicitv eroups are a low ./ for au oblate solution. and a large σ for a prolate solution.," The features that account for the high ellipticity sub sample in low multiplicity groups are a low $\bar{\beta}$ for an oblate solution, and a large $\sigma$ for a prolate solution."
 Either way there does appear to be a low axial ratio population that must vc accounted for. and that is not present m the higher uultiplicity data.," Either way there does appear to be a low axial ratio population that must be accounted for, and that is not present in the higher multiplicity data."
 It is worth remembering that at low nagnitude cuts (corresponding to higher multiplicitiex) he Local Group demonstrates oblate characteristics. πο his result is not necessarily suprising.," It is worth remembering that at low magnitude cuts (corresponding to higher multiplicities) the Local Group demonstrates oblate characteristics, so this result is not necessarily suprising."
 This result is even nore reliable occiring as it does in the region of the PDF that is least corrected in convolution., This result is even more reliable occurring as it does in the region of the PDF that is least corrected in convolution.
 In fact iu lis region highlv clliptical results for low unltiplicity eroups are reduced in weighting. so the fact the signal is still so strong indicates a real population difference.," In fact in this region highly elliptical results for low multiplicity groups are reduced in weighting, so the fact the signal is still so strong indicates a real population difference."
 This trend would appear to be iu the opposite seuse to what is expected in simulations (sec (Allgoodetal.2006) and references therein). but it should be noted that these results are much closer to agreement after correction.," This trend would appear to be in the opposite sense to what is expected in simulations (see \citep{allgood06} and references therein), but it should be noted that these results are much closer to agreement after correction."
" We fouud good agreement between our results for raw data with a simple 20] multiplicitv cut aud fiidiugs bv other authors. indicating/ a strouely prolate distribution with a mean οὐ=(0.16 aud a,= 0.11."," We found good agreement between our results for raw data with a simple 20+ multiplicity cut and findings by other authors, indicating a strongly prolate distribution with a mean $\bar{\beta}=0.46$ and $\sigma_{\beta}=0.14$ ."
We. present à survey to detect extragalactic. racio. sources at 432 Cllzο with. repeated. independent.. pointings.. of⋅ the D-configuration. .VLA in. snapshot mode.,We present a survey to detect extragalactic radio sources at 43 GHz with repeated independent pointings of the D-configuration VLA in snapshot mode.
 ..The survey is. centered on ELAIS-N2 (RA 16h 36m 48s. Dec 01 45° J2000). and covers about 0.5 down to 7 mv.," The survey is centered on ELAIS-N2 (RA 16h 36m 48s, Dec 01' 45”, J2000), and covers about 0.5 down to 7 mJy."
 The rimary objective was to establish the 10Ον source surface densitym at mJy levels. an issue. of 2.importance in. assessing. cdiscrete-source foreground.: contamination⋅ of⋠⋠ the⋅ CMD objects. signal.," The primary objective was to establish the 43-GHz source surface density at mJy levels, an issue of importance in assessing the discrete-source foreground contamination of the CMB signal."
 Lt is only recently that we have gained insight into what »opulates the radio sky at survey frequencies above 5 Cllz., It is only recently that we have gained insight into what populates the radio sky at survey frequencies above 5 GHz.
" Dv this frequency. some 60. per cent of the sources at Jv evels have ‘Lat’ spectra (a>0.5. where Lux density S is related to frequency £j via Sx 7"") over some region of he radio domain."," By this frequency some 60 per cent of the sources at Jy levels have `flat' spectra $\alpha > 
-0.5$, where flux density $S$ is related to frequency $\nu$ via $S \propto \nu^\alpha$ ) over some region of the radio domain."
 This is the signature of optically-thick emission- and so-called “Lat-spectrum synchrotronsources. (," This is the signature of optically-thick synchrotron emission and compact structure, so-called `flat-spectrum' sources. ("
Surveys at compactfrequencies siructure.above 150 Giz do find invertec-spectrum sources presumably dominated by [rec-free emission. probably dust emitters related to or part of the sub-nim galaxy population.,"Surveys at frequencies above 150 GHz do find inverted-spectrum sources presumably dominated by free-free emission, probably dust emitters related to or part of the sub-mm galaxy population."
 These do not appear at frequencies as low as 43 κ που Viera et al., These do not appear at frequencies as low as 43 GHz – see Viera et al.
 2009). -Flat- ⊳∖↓≻∢⊾≼∼, 2009). `Flat-spectrum'
↿↓⋅⊔⊔↓↼∪∣⋡≯↕⋖⋅≼∙↿≱∖∠⇂∪⊔∪↥⋏∙≟∢⊾⊔⋖⋅↓⋅⋜↧∐∙∖⇁↓⋯∖⇁⋖⋅⋂∶∪∪∖⇁∢⊾↓⋅⋜↧↓⋜⊔⋅⋏∙≟⋖⊾ frequeney range: rather their spectra show inversions. points of inllection. and. sometimes single peaks (c.g. Cigahertz-Peaked Spectra or GPS sources).," objects do not generally have $\alpha = 0$ over a large frequency range; rather their spectra show inversions, points of inflection, and sometimes single peaks (e.g. Gigahertz-Peaked Spectra or GPS sources)."
 The optical counterparts of the great of these sources are stellar in appearance and are Flat-Speetrummajority adio QSOs (ESItQs) or LLac he, The optical counterparts of the great majority of these sources are stellar in appearance and are Flat-Spectrum Radio QSOs (FSRQs) or Lac objects.
 Llow rapidly does this trend to such objects and away from conventional stccp-spectrum generally opticallv-thin spectrum radio galaxies continue with increasing frequency above 5 Gllz?, How rapidly does this trend to such objects and away from conventional steep-spectrum – generally optically-thin spectrum radio galaxies – continue with increasing frequency above 5 GHz?
 When and at what IHux-density level do we find new populations of flat or inverteda-aspectrum objects?, When and at what flux-density level do we find new populations of flat or inverted-spectrum objects?
 The pioneering survey of Brandie( Brielle (1974) covered 0.8 sterad at S Gilz and found 55 sources. of which “more than 2/27 were stellar objects. ie. QSOs or LLac objects.," The pioneering survey of Brandie Bridle (1974) covered 0.8 sterad at 8 GHz and found 55 sources, of which “more than 2/3” were stellar objects, i.e. QSOs or Lac objects."
 Subsequent deep surveys by Winclhorst et al. (, Subsequent deep surveys by Windhorst et al. (
1993) enabled definition of an S-ECllz source. count down to 10 pJy.,1993) enabled definition of an 8.4-GHz source count down to 10 $\mu$ Jy.
 Beandic and. Beielle found that at their relatively high Hux densities. 70 per cent of the sources have Dat or inverted spectra.," Brandie and Bridle found that at their relatively high flux densities, 70 per cent of the sources have flat or inverted spectra."
 Some spectrallv-extreme sources certainly exist ]5dge et al. (, Some spectrally-extreme sources certainly exist – Edge et al. (
1998) found two with spectral peaks,1998) found two with spectral peaks
In conclusion. only a few resonances. namely the most external of the first or second order resonances (4:3. 3:2. 5:3 and 2:1 resonance) capture most of the planetesimals.,"In conclusion, only a few resonances, namely the most external of the first or second order resonances $4$ $3$ , $3$ $2$, $5$ $3$ and $2$ $1$ resonance) capture most of the planetesimals."
 However. depending on the planet mass. the planetesimals that are not trapped in resonances can. or not. change the shape of the disk.," However, depending on the planet mass, the planetesimals that are not trapped in resonances can, or not, change the shape of the disk."
 In our simulations. planets with a mass above | Jupiter mass eject almost all the non-resonant planetesimals: the MMR structures therefore appear clearly.," In our simulations, planets with a mass above $1$ Jupiter mass eject almost all the non-resonant planetesimals: the MMR structures therefore appear clearly."
 Below 1 Jupiter mass. planets cannot eject all the non-resonant planetesimals. which can then partly hide the MMR structures.," Below $1$ Jupiter mass, planets cannot eject all the non-resonant planetesimals, which can then partly hide the MMR structures."
 In the circular orbit case (Fig., In the circular orbit case (Fig.
 2 and 3)). differences appear between thesimulations. depending on the planet mass:," \ref{figure_e0} and \ref{figure_e0_ae}) ), differences appear between thesimulations, depending on the planet mass:"
the individual particles.,the individual particles.
" However. near appropriately strong molecular absorption bands the particles’ optical properties become strongly. waveleneth-dependent. ancl there can even be wavelengths where the particles’ index of refraction has a real part n, close to that of the background medium while (he imaginary part 1; is still well less than unitv."," However, near appropriately strong molecular absorption bands the particles' optical properties become strongly wavelength-dependent, and there can even be wavelengths where the particles' index of refraction has a real part $n_r$ close to that of the background medium while the imaginary part $n_i$ is still well less than unity."
 At these so-called Christiansen wavelengths. surface scattering and diffraction are strongly suppressed. and provided (he particles are sulliciently small. this produces a dip in the overall opacity (Llapke1993:ElachiVahidiniaefal. 2011).," At these so-called Christiansen wavelengths, surface scattering and diffraction are strongly suppressed, and provided the particles are sufficiently small, this produces a dip in the overall opacity \citep{Hapke93, EvZ06, Vahidinia11}."
". since (he primary constituent of Saturns rings is ervstalline water ice 2009).. ancl the ring particles are dispersed in [ree space. the E-rings Christiansen wavelengths should occur wherever ervstalline water ice has both n,ο1 and »;<<1."," Since the primary constituent of Saturn's rings is crystalline water ice \citep{Cuzzi09}, and the ring particles are dispersed in free space, the F-ring's Christiansen wavelengths should occur wherever crystalline water ice has both $n_r\simeq1$ and $n_i<<1$."
" Examining the optical constants for cervstalline water ice plotted in Figure 2.. we find the indices of relraction vary dramatically in the vicinity of the strong 3.1 jan absorption band. with », falling below unity between 2.37 and 3.10 jan. In (his range of wavelengths. ice only has n,21 and m;<<lI at 2.87 jan. Thus the deepest part of the observed opacity dip falls verv close to the Christiansen wavelength for crvstalline. water ice7.. providing strong evidence that the Christiansen Effect is indeed responsible for this feature."," Examining the optical constants for crystalline water ice plotted in Figure \ref{specplot}, we find the indices of refraction vary dramatically in the vicinity of the strong 3.1 $\mu$ m absorption band, with $n_r$ falling below unity between 2.87 and 3.10 $\mu$ m. In this range of wavelengths, ice only has $n_r\simeq1$ and $n_i<<1$ at 2.87 $\mu$ m. Thus the deepest part of the observed opacity dip falls very close to the Christiansen wavelength for crystalline water ice, providing strong evidence that the Christiansen Effect is indeed responsible for this feature."
 The opacity dip at the Christiansen wavelength is only observed in the F ring and other similarly vcustw” ring features (see below)., The opacity dip at the Christiansen wavelength is only observed in the F ring and other similarly “dusty” ring features (see below).
 It is not seen in the i¢v main rings because (he particles in these rings are all millimeters in size or larger (Cuzzi 2009).. and the Christiansen Effect is only important for waveleneth-sized particles.," It is not seen in the icy main rings because the particles in these rings are all millimeters in size or larger \citep{Cuzzi09}, and the Christiansen Effect is only important for wavelength-sized particles."
 This can be most easily understood. by considering the exünction coefficient μι. which is delined as the ratio of a particles optical cross section (including loses due to both scattering and absorption) to its physical cross section.," This can be most easily understood by considering the extinction coefficient $Q_{ext}$, which is defined as the ratio of a particle's optical cross section (including loses due to both scattering and absorption) to its physical cross section."
 In general. the relationship between this parameter and the measured opacity of a ring is rather complex and depends on the width of the ring particles’ forwarel scattering lobe. as well as the apparent size of the observed ring feature and (he instrument5 fielcl of view (Cuzzi1985:FrenchandNicholson2000).," In general, the relationship between this parameter and the measured opacity of a ring is rather complex and depends on the width of the ring particles' forward scattering lobe, as well as the apparent size of the observed ring feature and the instrument's field of view \citep{Cuzzi85, FN00}."
". However. for rings composed primarily of sub-millimeter particles (which have broad. forwardescattering diffraction lobes) the measured opacity is simplv proportional to the appropriately weighted average ol (Qu.,, over all the particles in the ring."," However, for rings composed primarily of sub-millimeter particles (which have broad forward-scattering diffraction lobes) the measured opacity is simply proportional to the appropriately weighted average of $Q_{ext}$ over all the particles in the ring."
" For dielectric spheres (Q.,; can be computed analvticallv (vandeIIulst. 1957):", For dielectric spheres $Q_{ext}$ can be computed analytically \citep{vandeHulst}: :
simple svuchrotvon model. aud it would seem Chat Chis fact is forcing our code to choose hieher break lrequency models. despite their poor fit in (he ultraviolet.,"simple synchrotron model, and it would seem that this fact is forcing our code to choose higher break frequency models, despite their poor fit in the ultraviolet."
 By refitüng the models without the X-rav. flix. we can generate models that are much better fits to the radio through ultraviolet emission.," By refitting the models without the X-ray flux, we can generate models that are much better fits to the radio through ultraviolet emission."
 The best fitting parameters for these new fits are given in table 3.. and the new svnchrotron fits are shown in the rieht panels of Figure 1..," The best fitting parameters for these new fits are given in table \ref{tab:modelX}, and the new synchrotron fits are shown in the right panels of Figure \ref{fig:All}."
 These fits provide much lower break lrequences for the emission. aud for the three knots with obvious ultraviolet turnovers (knots D. A. and D). these fits have much lower 4? values compared io those generated using the N-rav measurement.," These fits provide much lower break frequencies for the emission, and for the three knots with obvious ultraviolet turnovers (knots D, A, and B), these fits have much lower $\chi^2$ values compared to those generated using the X-ray measurement."
 The remaining two knots (knots E and F). have minimal difference in the fit equality. which suggests that we still do not have direct measurements around the break frecuency.," The remaining two knots (knots E and F), have minimal difference in the fit quality, which suggests that we still do not have direct measurements around the break frequency."
 Ifthe radio through [ar UV emission is generated bv a single burst svnchrotron model. then the excess of X-ray flax compared to these models must be accounted. for.," If the radio through far UV emission is generated by a single burst synchrotron model, then the excess of X-ray flux compared to these models must be accounted for."
 Inverse Compton is a tempting option. as it has been used (o explain X-ray enission in other jets.," Inverse Compton is a tempting option, as it has been used to explain X-ray emission in other jets."
 llowever. previous analvsis has shown that the geometry of the jet required [or inverse Compton to contribute significantly to the X-ray [lux [rom (he knots is inconsistent with the eeomelry indicated by proper motion studies.," However, previous analysis has shown that the geometry of the jet required for inverse Compton to contribute significantly to the X-ray flux from the knots is inconsistent with the geometry indicated by proper motion studies."
 Specilicallv. the beaming model presented by lluris&Ixrawczvnski(2002) predicts that in order (ο explain the observed X-ray Πας. the angle of the jet to the line of sight must be nearly zero. in contradiction to the 20° value derived fom proper motions (Direttaetal.1999).," Specifically, the beaming model presented by \citet{HK}
 predicts that in order to explain the observed X-ray flux, the angle of the jet to the line of sight must be nearly zero, in contradiction to the $20^\circ$ value derived from proper motions \citep{BSM}."
. It is clear then that there must be some reinjection or re-excitation of electrons in the emitting regions., It is clear then that there must be some reinjection or re-excitation of electrons in the emitting regions.
 This scenario has some direct evidence to support it., This scenario has some direct evidence to support it.
 The kinematic ages of the knots in the jet ave much longer than the calculated synchrotron ages., The kinematic ages of the knots in the jet are much longer than the calculated synchrotron ages.
 Therefore. even to explain the optical emission of the jet. some amount of fresh high energv material [rom shocks along the jet must be added.," Therefore, even to explain the optical emission of the jet, some amount of fresh high energy material from shocks along the jet must be added."
invalidate the idea that ασ realistic stable distribution of simall 1nass objects can account for the observed potential.,invalidate the idea that any realistic stable distribution of small mass objects can account for the observed potential.
 Iu the future. it may be of interest to carry out a calculation simular to that reported here. but for a muore coudeused dark matter distribution consistent with this latest observation.," In the future, it may be of interest to carry out a calculation similar to that reported here, but for a more condensed dark matter distribution consistent with this latest observation."
 As we shall see. however. the present hydrodvunamical calculation strougly suggests that a cluster potential cannot easily reproduce Ser A*'« spectrum. aud this next generation of calculations way therefore be unnecessiuv.," As we shall see, however, the present hydrodynamical calculation strongly suggests that a cluster potential cannot easily reproduce Sgr A*'s spectrum, and this next generation of calculations may therefore be unnecessary."
 Iu a complex flow. eeucrated by mamy wind sources. the wind velocity aud density are uot uniforni. so the accretion radius may not be incepeudent of augle.," In a complex flow, generated by many wind sources, the wind velocity and density are not uniform, so the accretion radius may not be independent of angle."
" To set the length scale for the simulations. we shall therefore adopt the value Ry=LSpc (for which 1"" = 2.3 R4) as a reasonable mean representation of this quantity."," To set the length scale for the simulations, we shall therefore adopt the value $R_A = .018 \pc$ (for which $^{\prime\prime}$ = 2.3 $R_A$ ) as a reasonable mean representation of this quantity."
 We use a iodified version of the mmucerical aleorithii ZEUS. a general purpose code for MIID fiuids developed at NCSA (Stone&Norman1992: Norman 1991).," We use a modified version of the numerical algorithm ZEUS, a general purpose code for MHD fluids developed at NCSA \cite{SN92}; \cite{No94}) )."
 The code uses Eulerian finite differencing with the following relevant characteristics: fully explicit iu time: operator aud directional splitting of the lydrodvuamical variables: fully stageered grid: second-order (van Leer) upwiuded. monotonic interpolation for advection: consistent advection to evolve internal energy aud momenta: aud explicit solution of 1uterual energy.," The code uses Eulerian finite differencing with the following relevant characteristics: fully explicit in time; operator and directional splitting of the hydrodynamical variables; fully staggered grid; second-order (van Leer) upwinded, monotonic interpolation for advection; consistent advection to evolve internal energy and momenta; and explicit solution of internal energy."
 More details can be found in the references., More details can be found in the references.
 The code was yun ou the massively parallel Cray T3E at NASA's Cioddard Space Flight Center under the Wiel Performance Computing Challenge program., The code was run on the massively parallel Cray T3E at NASA's Goddard Space Flight Center under the High Performance Computing Challenge program.
 The production ruu for the calculation preseuted here used 125 T3E processors aud ran at more than 8 Cops., The production run for the calculation presented here used 125 T3E processors and ran at more than 8 Gflops.
" The code assigus cach processor to a “tile” aud cach tile consists of 60"" zones.", The code assigns each processor to a “tile” and each tile consists of $60^3$ zones.
" Zone sizes are ecometrically scaled by a factor of 1.02 so that the central zoues are 20 times smaller than the outermost zones. ninickius the ""1uultiple nested erids arrangement used by other researchers (c.e.. Ruffert&Melia199 1))."," Zone sizes are geometrically scaled by a factor of 1.02 so that the central zones are $\sim 20$ times smaller than the outermost zones, mimicking the “multiply nested grids” arrangement used by other researchers (e.g., \cite{RM94}) )."
 This allows for maximal resolution of the central regionOo (within the computer memory luuits available} while sufficieutlv resolving the wind sources aud miuimuizimg zouc-to-zone boundary effects., This allows for maximal resolution of the central region (within the computer memory limits available) while sufficiently resolving the wind sources and minimizing zone-to-zone boundary effects.
 The total volume is (LORAY? or ~(0.7pc)? with the center of the dark clister distribution being located at the origin., The total volume is $(40 R_A)^3$ or $\sim(0.7 \pc)^3$ with the center of the dark cluster distribution being located at the origin.
 The density of the 8eas initially filius5 the vohune of solution is set to a small value aud the velocity is ↴∖↴↸∖↑↑∪∑↸∖↥, The density of the gas initially filling the volume of solution is set to a small value and the velocity is set to zero.
⋅∪⋅↕∐∪↥⋅≼∐∖↥⋅↑∪↥⋅↸∖⋯⊳∐↸∖≺∣∏∐∏⋝↥⋅↕∏⋯⋯∪↥⋅↸∖≺∏∏↸⊳↨↘↽↕⋅↖⇁∙↑∐↸∖↕∐↑↸∖↥⋅∐⋜↧↕↸∖∐↸∖↥⋅∶↴∙⊾⋅↖↽≼∐∖∐↴∖↴↕↑⋅↖⇁↕↴∖↴↸⊳∐∪↴∖↴↸∖∐↴∖↴⋯⊳∐↑∐⋜↧↑↑∐↸∖ initialBeye temperature isB —+107Di Ik.- Free outflow couditiousm are miposedB on the outermost zones aud cach time," In order to reach equilibrium more quickly, the internal energy density is chosen such that the initial temperature is $\sim10^2$ K. Free outflow conditions are imposed on the outermost zones and each time"
⋅∪⋅↕∐∪↥⋅≼∐∖↥⋅↑∪↥⋅↸∖⋯⊳∐↸∖≺∣∏∐∏⋝↥⋅↕∏⋯⋯∪↥⋅↸∖≺∏∏↸⊳↨↘↽↕⋅↖⇁∙↑∐↸∖↕∐↑↸∖↥⋅∐⋜↧↕↸∖∐↸∖↥⋅∶↴∙⊾⋅↖↽≼∐∖∐↴∖↴↕↑⋅↖⇁↕↴∖↴↸⊳∐∪↴∖↴↸∖∐↴∖↴⋯⊳∐↑∐⋜↧↑↑∐↸∖ initialBeye temperature isB —+107Di Ik.- Free outflow couditiousm are miposedB on the outermost zones aud cach timeB," In order to reach equilibrium more quickly, the internal energy density is chosen such that the initial temperature is $\sim10^2$ K. Free outflow conditions are imposed on the outermost zones and each time"
"and smaller photospheric radit 7,.",and smaller photospheric radii $\Rph$.
 The low-energy emission is increased with increasing 2p., The low-energy emission is increased with increasing $\eB$.
 An important effect of magnetization is the softening of the spectral slope below £j..., An important effect of magnetization is the softening of the spectral slope below $\Epeak$.
 This is the result of synchrotron emission from nonthermal pairs., This is the result of synchrotron emission from nonthermal pairs.
 At energies below ~30 keV aclear soft excess is predicted. even for small 2j~10°," At energies below $\sim 30$ keV a clear soft excess is predicted, even for small $\eB\sim 10^{-3}$."
 Our model also predicts that the spectrum below a few keV should have the photon index —1., Our model also predicts that the spectrum below a few keV should have the photon index $-1$.
 This spectrum extends down to the optical band., This spectrum extends down to the optical band.
 Note that photon spectra in this paper are not corrected for the cosmological redshift :., Note that photon spectra in this paper are not corrected for the cosmological redshift $z$.
 The actually observed energy of all photons is smaller by the factor of (1|:)ἂν, The actually observed energy of all photons is smaller by the factor of $(1+z)^{-1}$.
 In particular. the soft excess is redshifted to ~10 keV for a typical : of GRBs.," In particular, the soft excess is redshifted to $\sim 10$ keV for a typical $z$ of GRBs."
 Thompson(1994) proposed that GRB photons are Comptonized in the subphotospheric region by turbulent bulk motions of the plasma in the jet., \citet{Thompson94} proposed that GRB photons are Comptonized in the subphotospheric region by turbulent bulk motions of the plasma in the jet.
 His model assumec that the jet is dominated by magnetic field and picturec Alfvénnic turbulence generated by some instabilities (e.g. by reconnection)., His model assumed that the jet is dominated by magnetic field and pictured Alfvénnic turbulence generated by some instabilities (e.g. by reconnection).
 Building a detailed theory of turbulence from first principles would be a formidable task. and Thompsor(1994) made estimates assuming that the plasma motions are limited by radiation drag.," Building a detailed theory of turbulence from first principles would be a formidable task, and \citet{Thompson94} made estimates assuming that the plasma motions are limited by radiation drag."
 Then the bulk Comptonizatioi occurs in the unsaturated regime and controlled by the efficiency of turbulence generation against the drag., Then the bulk Comptonization occurs in the unsaturated regime and controlled by the efficiency of turbulence generation against the drag.
 The predicted GRB. spectrum is. sensitive to. this unknowr efficiency. which may vary with radius.," The predicted GRB spectrum is sensitive to this unknown efficiency, which may vary with radius."
 The mechanism is very different from collisional heating., The mechanism is very different from collisional heating.
 Direct comparison of the model predictions is difficult. because Thompson(1994) used the theory of Comptonization in static sources. which is not valid for GRBs (see discussion in Section 3)).," Direct comparison of the model predictions is difficult, because \citet{Thompson94} used the theory of Comptonization in static sources, which is not valid for GRBs (see discussion in Section \ref{sec:method}) )."
 Like in the expanding Universe. the number of scatterings in a GRB jet scales as ~7r (not as Li) and is controlled by the decrease in density. not the escape of photons from the plasma.," Like in the expanding Universe, the number of scatterings in a GRB jet scales as $\sim\taut$ (not as $\taut^2$ ) and is controlled by the decrease in density, not the escape of photons from the plasma."
 The expected spectrum can only be found by solving radiative transfer in the expanding jet. and we emphasize the importance of accurate transfer calculations for GRB Mészáros&Rees(2000b) and Rees&Meszaros(2005) discussed subphotospheric Comptonization in general terms and highlighted its importance for GRBs.," The expected spectrum can only be found by solving radiative transfer in the expanding jet, and we emphasize the importance of accurate transfer calculations for GRB \citet{MR00b} and \citet{ReesMeszaros05} discussed subphotospheric Comptonization in general terms and highlighted its importance for GRBs."
 They argued that Comptonization helps explain observations. although they did not consider any concrete model.," They argued that Comptonization helps explain observations, although they did not consider any concrete model."
 More detailed calculations were performed by Pe'eretal.(2006)., More detailed calculations were performed by \citet{Peer06}.
. They still did not focus on any concrete physical scenario and. experimented with various phenomenological models. using free parameters to describe the heating process: the amount of dissipated energy and its fraction that is given to electrons. the shape of the electron distribution. and the optical depth at which the dissipation occurs.," They still did not focus on any concrete physical scenario and experimented with various phenomenological models, using free parameters to describe the heating process: the amount of dissipated energy and its fraction that is given to electrons, the shape of the electron distribution, and the optical depth at which the dissipation occurs."
 Various combinations of these parameters led to various spectra., Various combinations of these parameters led to various spectra.
 In contrast. we investigated the concrete physical model of collisional dissipation. which predicts the thermal and nonthermal heating rates.both scaling as r+.," In contrast, we investigated the concrete physical model of collisional dissipation, which predicts the thermal and nonthermal heating rates,both scaling as $r^{-1}$."
 It is therefore hard to compare their results with ours., It is therefore hard to compare their results with ours.
 We can only compare the technical tools., We can only compare the technical tools.
 Pe'eretal.(2006) used the kinetic code developed by Pe'er&Waxman(2005)., \citet{Peer06} used the kinetic code developed by \citet{PW05}.
. In many respects. their code is similar to ours.," In many respects, their code is similar to ours."
 Their code assumes. however. that radiation is isotropic in the comoving frame of the jet.," Their code assumes, however, that radiation is isotropic in the comoving frame of the jet."
 We gave up this approximation because it is violated even below the photosphere (see Section. 3)) and instead solved the radiative transfer equation., We gave up this approximation because it is violated even below the photosphere (see Section \ref{sec:method}) ) and instead solved the radiative transfer equation.
 Notealso that Pe'eretal.(2006) argued that GRB jets never become dominated by c* pairs., Notealso that \citet{Peer06} argued that GRB jets never become dominated by $e^\pm$ pairs.
 In contrast. we find that jets with magnetization sp<1 have a large pair loading factor fi~200. which changes the photospherie radius AR. by the factor of 20-40.," In contrast, we find that jets with magnetization $\eB\ll 1$ have a large pair loading factor $\npoel/n\sim 20-40$, which changes the photospheric radius $\Rph$ by the factor of 20–40."
 Giannios(2006) investigated Comptonization in jets that cury alternating magnetic fields and are heated by magnetic reconnection (Drenkhahn&Spruit2002).., \citet{Giannios06} investigated Comptonization in jets that carry alternating magnetic fields and are heated by magnetic reconnection \citep{DS02}. .
 This model has two phenomenological parameters — the reconnection rate and the fraction of dissipated energy that is deposited into electrons., This model has two phenomenological parameters — the reconnection rate and the fraction of dissipated energy that is deposited into electrons.
 The effect of magnetic dissipation on the electron distribution is unknown., The effect of magnetic dissipation on the electron distribution is unknown.
 Giannios(2006) assumed a thermal electron distribution and studied the radiation produced by this model., \citet{Giannios06} assumed a thermal electron distribution and studied the radiation produced by this model.
 He used a Monte-Carlo code and an iterative technique to simulate the radiative transfer in the We note that the jet model in Giannios(2006) 15 significantly different from Β10 and our work., He used a Monte-Carlo code and an iterative technique to simulate the radiative transfer in the We note that the jet model in \citet{Giannios06} is significantly different from B10 and our work.
 In particular. the jet continues to accelerate as D.(7/7 in the heated subphotospheric region while we focus on matter-dominated jets with D.zcost.," In particular, the jet continues to accelerate as $\Gamma\propto r^{1/3}$ in the heated subphotospheric region while we focus on matter-dominated jets with $\Gamma\approx const$."
 The electron temperature scales with radius approximately as Tj;X°°? while in our model we find that 7;. varies only by a factor of —2 for three decades in radius (Figure 4))., The electron temperature scales with radius approximately as $\Te\propto r^{5/3}$ while in our model we find that $\Te$ varies only by a factor of $\sim 2$ for three decades in radius (Figure \ref{fig:ele_temp}) ).
 Note also that Giannios(2006). assumed pure thermal heating. leading to negligible pair creation.," Note also that \citet{Giannios06} assumed pure thermal heating, leading to negligible pair creation."
 In contrast. our work and B10 predict a nonthermal e cascade. which has significant effects on the observed spectrum at high energies (above 10 MeV) and at low energies where synchrotron emission from the cascade becomes dominant.," In contrast, our work and B10 predict a nonthermal $e^\pm$ cascade, which has significant effects on the observed spectrum at high energies (above 10 MeV) and at low energies where synchrotron emission from the cascade becomes dominant."
 Lazzati&Begelman(2010) considered a simplified model for Comptonization following an impulsive heating event below the photosphere., \citet{Lazzati10} considered a simplified model for Comptonization following an impulsive heating event below the photosphere.
 In their model. all electrons in a slab of optical depth 7r~2 are suddenly given a Lorentz factor ~ and then cooled by blackbody radiation.," In their model, all electrons in a slab of optical depth $\taut\sim 2$ are suddenly given a Lorentz factor $\gamma$ and then cooled by blackbody radiation."
" This setup can produce a variety of photon spectra depending on electron ratio »-/i,. rr. and 5."," This setup can produce a variety of photon spectra depending on photon-to-electron ratio $n_\gamma/n_e$, $\taut$, and $\gamma$."
 It has. however. a problem.," It has, however, a problem."
 The sudden heating of the slab (faster than Compton cooling) can hardly deseribe any physicaldissipation mechanism. as different parts of the slab are causally disconnected.," The sudden heating of the slab (faster than Compton cooling) can hardly describe any physicaldissipation mechanism, as different parts of the slab are causally disconnected."
 The causal contact dight-crossing) time for a region with 7p~2 is orders of magnitude longer than Compton cooling time., The causal contact (light-crossing) time for a region with $\taut\sim 2$ is orders of magnitude longer than Compton cooling time.
 The main part of observed GRB radiation is emitted with a Band-type spectrum — a smoothly broken power law that peaks near 0.3 MeV. The spectral slopes below and above the peak (a and J) vary (Preeceetal.2000)., The main part of observed GRB radiation is emitted with a Band-type spectrum — a smoothly broken power law that peaks near $0.3$ MeV. The spectral slopes below and above the peak $\alpha$ and $\beta$ ) vary \citep{Preece00}.
. The most frequently measured low-energy slope à is close to 1. but in some cases it reaches 1 (Rayleigh-Jeans slope of the Planck spectrum).," The most frequently measured low-energy slope $\alpha$ is close to $-1$ , but in some cases it reaches $1$ (Rayleigh-Jeans slope of the Planck spectrum)."
 The most frequently measured high-energy slope ./ is near 2.5; it also significantly varies from burst to burst. roughly between 2and 2.," The most frequently measured high-energy slope $\beta$ is near $-2.5$ ; it also significantly varies from burst to burst, roughly between $-2$ and $-3$ ."
 The spectra predicted by our model appear to agree with observations., The spectra predicted by our model appear to agree with observations.
The predicted spectral peak is close to MeV/(1| :).,The predicted spectral peak is close to $(1+z)$ .
 The average observed valuesof a~—1 and jej2.5 are consistent with the collisionally heated jet with, The average observed valuesof $\alpha\sim -1$ and $\beta\sim -2.5$ are consistent with the collisionally heated jet with
"Filling the positions listed in Table 1 (o a second-order polynomial 0.5¢,A7) allows an estimate of the acceleration.",Fitting the positions listed in Table 1 to a second-order polynomial $x = x_0 + v_x \Delta t + 0.5 a_x \Delta t^2$ ) allows an estimate of the acceleration.
 We do not detect anv significant acceleration., We do not detect any significant acceleration.
" Acceleration uncertainties (1o) in equatorial components are σι,=0.024. 0.035 mas 7 and in Galactic components are 2;=0.035. σι=0.026 mas 7. A value of 0.03 mas v7? corresponds to zz| km tv ! at the distance of the Galactic center."," Acceleration uncertainties $1\sigma$ ) in equatorial components are $\sigma_\alpha=0.024,$ $\sigma_\delta=0.035$ mas $^{-2}$ and in Galactic components are $\sigma_l=0.035,$ $\sigma_b=0.026$ mas $^{-2}.$ A value of 0.03 mas $^{-2}$ corresponds to $\approx 1$ km $^{-1}$ $^{-1}$ at the distance of the Galactic center."
 discuss the implications of upper limits to the acceleration of for the nature of and Galactic structure constants., \citet{GR98} discuss the implications of upper limits to the acceleration of for the nature of and Galactic structure constants.
 The orbits of stars as close as LOO AU [rom require a central mass of about 4xLO°.., 	The orbits of stars as close as 100 AU from require a central mass of about $4\times10^6$.
 Our limits on the proper motion of can be used to indicate how much of this mass must be cirectly associated withA*., Our limits on the proper motion of can be used to indicate how much of this mass must be directly associated with.
. Since the motion of in the plane of the Galaxy. has a much greater uncertainty (LO15. owing primarily (o uncertainties in and )) than the motion out of the plane of the Galaxy (0.9 )). we focus mainlv on the out-of-plane component.," Since the motion of in the plane of the Galaxy has a much greater uncertainty (10–15, owing primarily to uncertainties in and ) than the motion out of the plane of the Galaxy (0.9 ), we focus mainly on the out-of-plane component."
" We use our strong upper limit [or ss peculiar motion out of the plane of the Galaxy. in conjunction with 1) the constraint (hat z4x10"" ool dark matter is contained in the central LOO AU region (Schódel 2003).. 2) the observation that is contained in this region (Menten 2003).. and 3) the existence of ~LO®to10* stars orbiting the galactic center within the central lew parsecs (Genzelοἱal.2003) to place firma lower limits to the mass directly associated withA*."," We use our strong upper limit for s peculiar motion out of the plane of the Galaxy, in conjunction with 1) the constraint that $\approx4\times10^6$ of dark matter is contained in the central 100 AU region \citep{S03,Ghez03}, 2) the observation that is contained in this region \citep{MREG97,R03}, and 3) the existence of $\sim10^6~{\rm to}~10^7$ stars orbiting the galactic center within the central few parsecs \citep{Genzel03} to place firm lower limits to the mass directly associated with."
. In the next sub-sections. we consider the possibility that ss mass does not necessarily dominate (he region and compare (he expected positions. velocity ancl acceleration ol that might result with our observational limits.," In the next sub-sections, we consider the possibility that s mass does not necessarily dominate the region and compare the expected positions, velocity and acceleration of that might result with our observational limits."
 We will first consider various possibilities [or configurations of the zz4x105 of material known to occupy the central 100 AU region., We will first consider various possibilities for configurations of the $\approx4\times10^6$ of material known to occupy the central 100 AU region.
 The discussion divides along observational lines by a possible offset. of [rom the dynamical center of a hypothetical dark matter distribution., The discussion divides along observational lines by a possible offset of from the dynamical center of a hypothetical dark matter distribution.
 Since. is observed to be within 100 AU of the center of this region (Schódeletοἱal. 2003).. we do not consider offsets >LOO AU.," Since, is observed to be within 100 AU of the center of this region \citep{S03,Ghez03,R03}, we do not consider offsets $>100$ AU."
 In $44.1. we consider offsets of [rom the center of mass of the svstem between 4«r100 AU. where we would easily detect changes in position. but not necessarily be able to measure a velocity or acceleration. because the orbital period would be short compared to the time span of our observations used to get accurate motions.," In 4.1, we consider offsets of from the center of mass of the system between $4 < r < 100$ AU, where we would easily detect changes in position, but not necessarily be able to measure a velocity or acceleration, because the orbital period would be short compared to the time span of our observations used to get accurate motions."
 In 844.2. we consider offsets of <4 AU. where we could detect neither a positional offset nor A*’ss velocity or acceleration even though the latter two," In 4.2, we consider offsets of $<4$ AU, where we could detect neither a positional offset nor s velocity or acceleration even though the latter two"
where 0<axI alone with a higher dimensional isotropic pressure py existing in the extra space such the the field equations reduce to where &=dim?(d—1)+601dim). FFor positive∙− energy density.⋅  must be greater (han zero which⋅ gives⋅ mj«——————addidAd(d1)2) or. nb»adOM21,"where $0\leq \alpha \leq 1$ along with a higher dimensional isotropic pressure $p_{d}$ existing in the extra space such the the field equations reduce to where $k = dm^{2}(d-1) + 6(1-dm)$ For positive energy density, $k$ must be greater than zero which gives $m < \frac{3d-\sqrt{3d(d+2)}}{d(d-1)}$ or, $m > \frac{3d+\sqrt{3d(d+2)}}{d(d-1)}$."
 AAgain Bianchi identity will be in our case using equations (36). (37) & we gel Λ' equation (41) we get lil AM and e is the integration constant.," Again Bianchi identity will be in our case Now using equations (36), (37) (40) we get Solving equation (41) we get where and $c$ is the integration constant."
" From physical considerations we can finally put the restriction on as m.<apAdQ/Ad(d2) , ford@j andan<1. lor. d=lI. The relation (43) needs some interpretation."," From physical considerations we can finally put the restriction on $m$ as $m <
\frac{3d-\sqrt{3d(d+2)}}{d(d-1)}$ for $ d \neq 1$ and $m <1$ for $
d = 1$ The relation (43) needs some interpretation."
 The last term is a (vpical higher dimensional effect. absent in 4D (d= 0).," The last term is a typical higher dimensional effect, absent in 4D $d=0$ )."
 Now we see that several possibilities for different. values of m present G)) m=—L: Here A=HR and we have to sacrifice (he desirable feature of dimensional reduction to get an isotropic expansion in all dimensions., Now we see that several possibilities for different values of $m$ present ) $m = - 1$: Here $ A=R$ and we have to sacrifice the desirable feature of dimensional reduction to get an isotropic expansion in all dimensions.
 Hlowever it not exactly the, However it not exactly the
models.,models.
 From these sizes we estimate the radial extension of the emitting CO rings., From these sizes we estimate the radial extension of the emitting CO rings.
 We compute the outer radius of the ring as a function of its inner radius., We compute the outer radius of the ring as a function of its inner radius.
 The results for both stars are depicted by the solid lines in reffig:rout. indicating that for increasing inner radii the emission area converges towards infinitesimally narrow rings., The results for both stars are depicted by the solid lines in \\ref{fig:rout} indicating that for increasing inner radii the emission area converges towards infinitesimally narrow rings.
" With the stellar luminosities given in reftab:parameters and an assumed dust sublimation temperature of Yauc1600 KK (ez.Lodders2003) we find individual sublimation radii of fou,224/2. and cSOR. for S773 and 112. respectively."," With the stellar luminosities given in \\ref{tab:parameters} and an assumed dust sublimation temperature of $T_{\rm sub}\simeq 1600$ K \citep[e.g.,][]{Lodders} we find individual sublimation radii of $R_{\rm
sub} \simeq 24\,R_{*}$ and $\simeq 86\,R_{*}$ for 73 and 12, respectively."
 These maximum outer edges of the CO rings are shown in reffig:rout ας dashed lines., These maximum outer edges of the CO rings are shown in \\ref{fig:rout} as dashed lines.
 An estimate for the inner edge of the CO ring is less precise., An estimate for the inner edge of the CO ring is less precise.
 The CO gas has to have a temperature of about KK. only slightly higher than the dust sublimation temperature but definitely much lower than the effective temperatures of the stars.," The CO gas has to have a temperature of about K, only slightly higher than the dust sublimation temperature but definitely much lower than the effective temperatures of the stars."
 A reasonable estimate could thus be Zi<LOL. for S773 and (much) larger for 1102.," A reasonable estimate could thus be $R_{\rm in} \ga 10\,R_{*}$ for 73 and (much) larger for 12."
 For the latter we plotted the relation between inner and outer radius for two different values of disk inclination., For the latter we plotted the relation between inner and outer radius for two different values of disk inclination.
" However. independent of the exact value of Z, Gand ὁ as in the case of 112) we can conclude that the CO emitting area is concentrated within a rather narrow ring at short distance from the central star."," However, independent of the exact value of $R_{\rm in}$ (and $i$ as in the case of 12) we can conclude that the CO emitting area is concentrated within a rather narrow ring at short distance from the central star."
 Using high-quality A -band spectra obtained with SINFONI. we detected. for the first time. a 13C enhancement in two B[e]SGs in the LMC. thereby confirming their evolved nature.," Using high-quality $K$ -band spectra obtained with SINFONI, we detected, for the first time, a 13C enhancement in two B[e]SGs in the LMC, thereby confirming their evolved nature."
 The enrichment of their disk-forming wind material with C and the resulting strong CO band emission. as was suggested by Kraus(2009).. is thus shown to be a robust method for the identification of evolved stars.," The enrichment of their disk-forming wind material with $^{13}$ C and the resulting strong $^{13}$ CO band emission, as was suggested by \citet{Kraus09}, is thus shown to be a robust method for the identification of evolved stars."
 This method can now be applied to the Galactic B[e]SG candidates to disentangle the real Galactic B[e]SGs from the more abundant Herbig Ae/Be stars., This method can now be applied to the Galactic B[e]SG candidates to disentangle the real Galactic B[e]SGs from the more abundant Herbig Ae/Be stars.
 In our analysis. we also found a deficiency of hot CO gas in he disks of both stars.," In our analysis, we also found a deficiency of hot CO gas in the disks of both stars."
 This can be explained in two ways., This can be explained in two ways.
 Either both stars rotate at their critical limit so that equatorial surface emperatures (and hence. maximum disk temperatures) drop below ~3000 KK and no hot CO gas is expected. or the CO emitting region is located in a high-density. detached ring of material rather han in an equatorially enhanced. disk-forming wind.," Either both stars rotate at their critical limit so that equatorial surface temperatures (and hence, maximum disk temperatures) drop below $\sim 3000$ K and no hot CO gas is expected, or the CO emitting region is located in a high-density, detached ring of material rather than in an equatorially enhanced, disk-forming wind."
 While the ormer Scenario cannot be excluded. we consider the latter as the more realistic one since it also agrees with the recently found detached gas ring around the SMC B[e]SG star 665 by Krausetal.(2010).," While the former scenario cannot be excluded, we consider the latter as the more realistic one since it also agrees with the recently found detached gas ring around the SMC B[e]SG star 65 by \citet{Kraus2010}."
. A.L. is supported by the Max-Planck-Institut fürr Radioastronomie. M.K. acknowledges financial support from GÀ CRR grant number 205/08/0003.," A.L. is supported by the Max-Planck-Institut fürr Radioastronomie, M.K. acknowledges financial support from GA ČRR grant number 205/08/0003."
 Financial support was provided to O.S. by the Science and Technology Facilities Council (UK)., Financial support was provided to O.S. by the Science and Technology Facilities Council (UK).
 M.B.F. acknowledges Conselho Nacional de Desenvolvimento Cientiftico e Tecnoldggico (CNPq-Brazil) for the post-doctoral grant., M.B.F. acknowledges Conselho Nacional de Desenvolvimento Científfico e Tecnológgico (CNPq-Brazil) for the post-doctoral grant.
wide-ficld imagers at meciumesizec telescopes. as long as enough observing time is available to provide a large number of datapoints and thus a high level of redundancy.,"wide-field imagers at medium-sized telescopes, as long as enough observing time is available to provide a large number of datapoints and thus a high level of redundancy."
 We show that main-sequence open clusters like Pracsepe are ideal environments to probe the long-term angular momentuni evolution and the underlying regulation mechanisms., We show that main-sequence open clusters like Praesepe are ideal environments to probe the long-term angular momentum evolution and the underlying regulation mechanisms.
 The five periods range from 5hh to almost ddays and have been confirmed by various. independent: period search procedures as well as plausibility checks., The five periods range from h to almost days and have been confirmed by various independent period search procedures as well as plausibility checks.
 We attribute hese photometric variations to cool. magnetically induced spots. co-rotating with the objects thus the periods ikelv correspond. to the rotation periods.," We attribute these photometric variations to cool, magnetically induced spots, co-rotating with the objects – thus the periods likely correspond to the rotation periods."
 Comparing the small sample of periods in Praesepe with sin? data for he coeval Livades and with periods measured in vounger clusters. we conclude that they give a reasonable first elance on the very low mass (VLM) period. distribution in praesepe.," Comparing the small sample of periods in Praesepe with $v\sin i$ data for the coeval Hyades and with periods measured in younger clusters, we conclude that they give a reasonable first glance on the very low mass (VLM) period distribution in Praesepe."
 Particularly. it is unlikely that Praesepe harbours a large. undetected population of VLM objects with periods significantly longer than ddavs.," Particularly, it is unlikely that Praesepe harbours a large, undetected population of VLM objects with periods significantly longer than days."
 Still. the big caveat in the discussion of the periods is the small number of datapoints.," Still, the big caveat in the discussion of the periods is the small number of datapoints."
 As long as we do not have a more substantial dataset. we refrain from carrying out a vigorous statistical analysis ancl rely instead on more qualitative assessments.," As long as we do not have a more substantial dataset, we refrain from carrying out a vigorous statistical analysis and rely instead on more qualitative assessments."
 We find a trend. of decreasing period with decreasing mass in the VLM regime. confirming previous claims in vounger clusters.," We find a trend of decreasing period with decreasing mass in the VLM regime, confirming previous claims in younger clusters."
 This trend is roughly linear in our small sample with P=240(M/M.)h.," This trend is roughly linear in our small sample with $P \approx 240\,(M/M_{\odot})\,h$."
 To probe the periomass relation over a broad range of stellar masses. we combine our dataset with periods in the coeval IEvades from the literature.," To probe the period-mass relation over a broad range of stellar masses, we combine our dataset with periods in the coeval Hyades from the literature."
 We find a dramatic change in the periods at spectral ἵνρο INS-M2.. corresponding to masses of Ανν Without clear transition regime. the periods dro from ~ lO0cdcdays in mid WKestars to « 4eldavs in carly A," We find a dramatic change in the periods at spectral type K8-M2, corresponding to masses of $\,M_{\odot}$: Without clear transition regime, the periods drop from $\sim 10$ days in mid K-stars to $<4$ days in early M-stars."
 Even considering that the VEM period sample may he incomplete. the mere existence of the fast rotating Meciwarfs in Praesepe indicates a significant change in the rotational regulation.," Even considering that the VLM period sample may be incomplete, the mere existence of the fast rotating M-dwarfs in Praesepe indicates a significant change in the rotational regulation."
 We note that this change coincides with the onset of chromospheric activity occuring at spectral tvpe MI in Pracsepe (?).., We note that this change coincides with the onset of chromospheric activity occuring at spectral type M1 in Praesepe \citep{2006AJ....132.1517K}.
 Comparing the periods in Praesepe with rotational data in vounger clusters (ages of 100-200MM we find. that some angular momentum loss is likely to occuryr). in the VLM range.," Comparing the periods in Praesepe with rotational data in younger clusters (ages of Myr), we find that some angular momentum loss is likely to occur in the VLM range."
 The exponential timescale of spin-down due to stellar winds is clearly longer than MMyr. most likely between 400 and MMSr. and thus 10-100 times longer than in solar-mass stars.," The exponential timescale of spin-down due to stellar winds is clearly longer than Myr, most likely between 400 and Myr, and thus 10-100 times longer than in solar-mass stars."
 Hence. rotational braking on the main-sequence is clearly less etlicient in the VEM regime. resulting in fast rotators.," Hence, rotational braking on the main-sequence is clearly less efficient in the VLM regime, resulting in fast rotators."
 Thus. the periods in Praesepe provide evidence for two observational manifestations of a fundamental change in the rotational regulation: à clear drop of the rotation. periods at AL. and very. long spin-down timescales in the VLM regime.," Thus, the periods in Praesepe provide evidence for two observational manifestations of a fundamental change in the rotational regulation: a clear drop of the rotation periods at $\,M_{\odot}$ and very long spin-down timescales in the VLM regime."
 This is in line with previous findings: ? mace an attempt to model the rotational evolution of VLM objects and conclude that it is impossible to simultaneously reproduce the observed stellar spin-down in the 0.6.1.1AL. range and for stars between 0.1 and 0.6AZ. ὃν implying that the solar-tvpe rotational models are not applicable at. very low masses.," This is in line with previous findings: \citet{2000ApJ...534..335S} made an attempt to model the rotational evolution of VLM objects and conclude that it is impossible to 'simultaneously reproduce the observed stellar spin-down in the $0.6-1.1\,M_{\odot}$ range and for stars between 0.1 and $\,M_{\odot}$ ', implying that the solar-type rotational models are not applicable at very low masses."
 In the evolutionary scheme proposed by 2.. the break between “L and C-sequence’ (which is identical to the break in the period-mass relation reported in this work) can be seen as another way to describe a transition. between wo different rotational regulation regimes.," In the evolutionary scheme proposed by \citet{2003ApJ...586..464B}, the break between 'I- and C-sequence' (which is identical to the break in the period-mass relation reported in this work) can be seen as another way to describe a transition between two different rotational regulation regimes."
 ln summary. he solar-tvpe rotational evolution. picture of a transition αναπιο driving à magnetic wind. which leads to Skumanich-ype angular momentum losses. does not appear to be applicable in the VLM regime.," In summary, the solar-type rotational evolution picture of a transition dynamo driving a magnetic wind, which leads to Skumanich-type angular momentum losses, does not appear to be applicable in the VLM regime."
 The underlving physical reason for this fundamental eakcdown of the solar-type rotational braking is not clear., The underlying physical reason for this fundamental breakdown of the solar-type rotational braking is not clear.
 We argue that it might be possible to explain it with 1¢ disappearance of the radiative core and the resulting inability to operate a solar-tvpe interlace cvnameo., We argue that it might be possible to explain it with the disappearance of the radiative core and the resulting inability to operate a solar-type interface dynamo.
 However. lis transition is thought to occur at masses 0.3-0.4A7. in solar-metallicitv stars. ie at somewhat lower masses than 10 drop in rotation periods.," However, this transition is thought to occur at masses $\,M_{\odot}$ in solar-metallicity stars, i.e at somewhat lower masses than the drop in rotation periods."
 The metallicity dilference tween the Sun and. voung clusters like Pracsepe might xplain parts of this discrepancy., The metallicity difference between the Sun and young clusters like Praesepe might explain parts of this discrepancy.
 Fundamental changes in 1e magnetic field structure. the wind properties. or dynamo action (possibly independent. of interior structure) have o be considered as alternative ways to understand: the observational findings.," Fundamental changes in the magnetic field structure, the wind properties, or dynamo action (possibly independent of interior structure) have to be considered as alternative ways to understand the observational findings."
 The obvious signature of a fundamental change seen in he rotation periods holds great potential to use rotation as a probe for magnetic properties and/or interior structure., The obvious signature of a fundamental change seen in the rotation periods holds great potential to use rotation as a probe for magnetic properties and/or interior structure.
 Follow-up studies aimed. a) to enlarge the rotational database for VLM objects on the main-sequence ancl b) to provide complementary activitv-related: data in. this mass regime are thus highly desirable to clarify the open questions., Follow-up studies aimed a) to enlarge the rotational database for VLM objects on the main-sequence and b) to provide complementary activity-related data in this mass regime are thus highly desirable to clarify the open questions.
ELGs is detected in the hard band (2-8 keV). and their stacked-average X-ray spectral slope of Γz2.0 suggests a low fraction of obscured AGN activity within the X-ray detected ELG population.,"ELGs is detected in the hard band (2–8 keV), and their stacked-average X-ray spectral slope of $\Gamma\approx2.0$ suggests a low fraction of obscured AGN activity within the X-ray detected ELG population."
 The range of rest-frame soft-band luminosities for the [1 z— 0.4-1.3 X-ray detected ELGs suggests they are probably X-ray detected starburst galaxies similar to M82 and NGC 3256 (see Figure 3: GGriffiths 2000: Lira 2002)., The range of rest-frame soft-band luminosities for the 11 $z=$ 0.4–1.3 X-ray detected ELGs suggests they are probably X-ray detected starburst galaxies similar to M82 and NGC 3256 (see Figure 3; Griffiths 2000; Lira 2002).
 The most X-ray luminous sources are up to ~ 3 times more luminous than NGC 3256. the most X-ray luminous nearby starburst galaxy.," The most X-ray luminous sources are up to $\approx$ 3 times more luminous than NGC 3256, the most X-ray luminous nearby starburst galaxy."
" The range of vL,s,,, luminosities (2«1019 to zz2«IO depending on the K-correction used: EElbaz 2002) are also consistent with M82. and NGC 3256-type. Moran. Lehnert. Helfand (1999) predicted that X-ray detected starbursts would contribute ~ of the X-ray background (XRB)."," The range of $\nu$$L_{15\mu m}$ luminosities $\approx2\times10^{10}$ to $\approx2\times10^{12}$ depending on the $K$ -correction used; Elbaz 2002) are also consistent with M82 and NGC 3256-type Moran, Lehnert, Helfand (1999) predicted that X-ray detected starbursts would contribute $\approx$ of the X-ray background (XRB)."
 However. the total soft and hard band fluxes of the X-ray detected starbursts in this study (7.7«10715 erg em sc! and 1.5«107% erg em s. respectively) only account for ~ of the XRB (assuming the XRB flux densities given in Garmire 2002).," However, the total soft and hard band fluxes of the X-ray detected starbursts in this study $7.7\times10^{-16}$ erg $^{-2}$ $^{-1}$ and $1.5\times10^{-15}$ erg $^{-2}$ $^{-1}$, respectively) only account for $\approx$ of the XRB (assuming the XRB flux densities given in Garmire 2002)."
 This fraction is below the lower limit given by Moran (1999). although we note that these X-ray observations have probably not fully resolved the XRB CCowie 2002) and we might expect these sources to dominate the faint X-ray source counts.," This fraction is below the lower limit given by Moran (1999), although we note that these X-ray observations have probably not fully resolved the XRB Cowie 2002) and we might expect these sources to dominate the faint X-ray source counts."
 The large fraction of X-ray detected ELGs with 15 um counterparts suggests an association between the production of X-ray and IR emission., The large fraction of X-ray detected ELGs with 15 $\mu$ m counterparts suggests an association between the production of X-ray and IR emission.
 Since the 15 gm emission is considered an unobscured indicator of star-formation rate RRowan-Robinson 1997). this further suggests that the X-ray emission can provide an indication of star-formation rate.," Since the 15 $\mu$ m emission is considered an unobscured indicator of star-formation rate Rowan-Robinson 1997), this further suggests that the X-ray emission can provide an indication of star-formation rate."
 However. we urge caution since the 17 15 um detected ELGs undetected at X-ray energies have /-band magnitude. 7)K color. V|7 color. 15 um flux density. and redshift distributions indistinguishable from the I] X-ray detected starbursts (a Kolmogorov-Smirnov test gives probabilities of98%..97%.48%.. and 55%.. These results suggest a broad range of X-ray luminosities for starbursts of similar IR luminosities and implies that there may be another factor that effects the production of the X-ray emission.," However, we urge caution since the 17 15 $\mu$ m detected ELGs undetected at X-ray energies have $I$ -band magnitude, $I-K$ color, $V-I$ color, 15 $\mu$ m flux density, and redshift distributions indistinguishable from the 11 X-ray detected starbursts (a Kolmogorov-Smirnov test gives probabilities of, and , These results suggest a broad range of X-ray luminosities for starbursts of similar IR luminosities and implies that there may be another factor that effects the production of the X-ray emission."
 The lack of a statistical difference between the V.7 and 7.K colors suggest that galaxy type and dust extinction. effects are not important., The lack of a statistical difference between the $V-I$ and $I-K$ colors suggest that galaxy type and dust extinction effects are not important.
 However. we note that a much larger fraction of the X-ray detected starbursts have radio counterparts; see Bauer (2002) for further constraints.," However, we note that a much larger fraction of the X-ray detected starbursts have radio counterparts; see Bauer (2002) for further constraints."
 The four z«0.2 sources have X-ray luminosities typical of normal galaxies 107 erg s: see Figure 3: Hornschemeier 2002. hereafter Paper VIID.," The four $z<0.2$ sources have X-ray luminosities typical of normal galaxies $\simlt10^{40}$ erg $^{-1}$; see Figure 3; Hornschemeier 2002, hereafter Paper VIII)."
" Their average VLs, luminosity of =6«10° lis also typical of normal galaxies.", Their average $\nu$$L_{15\mu m}$ luminosity of $\approx6\times10^{8}$ is also typical of normal galaxies.
 These sources arise at smaller. X-ray-to-optical flux ratios than the X-ray detected starbursts log(>)L22], These sources arise at smaller X-ray-to-optical flux ratios than the X-ray detected starbursts $\log{({{f_{\rm X}}\over{f_{\rm I}}})}\simlt-2$ ].
 Thus to detect higher redshift normal galaxies will require either deeper oobservations or the application of X-ray stacking techniques PPaper VILL., Thus to detect higher redshift normal galaxies will require either deeper observations or the application of X-ray stacking techniques Paper VIII).
 Deeper IR observations GGOODS?)) are planned withSIRTF over z: 160 aremin? of the most sensitive region of the CDF-N: these should reach a depth in the 24 4m MIPS band to an equivalent 15 um flux density of z 10 gy., Deeper IR observations ) are planned with over $\approx$ 160 $^2$ of the most sensitive region of the CDF-N: these should reach a depth in the 24 $\mu$ m MIPS band to an equivalent 15 $\mu$ m flux density of $\approx$ 10 $\mu$ Jy.
 In addition to providing a sample of ~ 100 X-ray detected ELGs with 24 um counterparts. these observations should be deep enough to detect counterparts for the majority of the X-ray detected ALGs and AGN-dominated sources.," In addition to providing a sample of $\approx$ 100 X-ray detected ELGs with 24 $\mu$ m counterparts, these observations should be deep enough to detect counterparts for the majority of the X-ray detected ALGs and AGN-dominated sources."
 This work would not have been possible without the support of the entire aand ACIS teams., This work would not have been possible without the support of the entire and ACIS teams.
 We acknowledge the financial support of NASA grant NAS 8-38252 (GPG. ΡΕ. NSF CAREER award AST-9983783 (DMA. FEB. WNB. CV). NASA GSRP grant NGT5-50247 and the Pennsylvania Space Grant Consortium (AEH). and NSF grant AST-9900703 (DPS).," We acknowledge the financial support of NASA grant NAS 8-38252 (GPG, PI), NSF CAREER award AST-9983783 (DMA, FEB, WNB, CV), NASA GSRP grant NGT5-50247 and the Pennsylvania Space Grant Consortium (AEH), and NSF grant AST-9900703 (DPS)."
SZ effect observations.,SZ effect observations.
 Fortunately. substantial upgrades are planned for currently operational arrays [for example. the current aud.OVRO arrays are being upgraded and iuersed into the Combined Array for Research in \Gillimeter-wave ÁAstrononiv CCARALAY| and a nunbor of new interferometers are planned as well (e... 71.AABDA. and AME Holder et al.," Fortunately, substantial upgrades are planned for currently operational arrays [for example, the current and arrays are being upgraded and merged into the Combined Array for Research in Millimeter-wave Astronomy $^6$ ] and a number of new interferometers are planned as well (e.g., $^7$, and ; Holder et al."
 2000: Lo et al., 2000; Lo et al.
 2000: Ikuceissl et al., 2000; Kneissl et al.
 2001)., 2001).
 These experiments will substautiallv iuprove the qualitv of SZ effect observatious., These experiments will substantially improve the quality of SZ effect observations.
 Below. we compare the expected performance of the aud the (soon to be) upgradedOVRO array with the currently operationalBIVA/OVRO arrays.," Below, we compare the expected performance of the and the (soon to be) upgraded array with the currently operational arrays."
" We show that the degeneracy in the 30 plane is greatlv reduced with data from the or the uperadedOVRO array aud this will allow for accurate determinations of gg aud S$, aud without the need for any N-rav results.", We show that the degeneracy in the $\beta-\theta_c$ plane is greatly reduced with data from the or the upgraded array and this will allow for accurate determinations of $y_0$ and $S_{\nu}$ and without the need for any X-ray results.
 To compare the arravs. we first generate “mock” observations of model galaxy clusters.," To compare the arrays, we first generate “mock” observations of model galaxy clusters."
 We follow a method simular to that outlined by Iolder et al. (, We follow a method similar to that outlined by Holder et al. (
2000).,2000).
" First. we create a Compton parameter imap. g(0,.0,). for cach model cluster."," First, we create a Compton parameter map, $y(\theta_x,\theta_y)$, for each model cluster."
 These maps are converted iuto SZ effect intensity decrement maps., These maps are converted into SZ effect intensity decrement maps.
 Before we cau do this. however. we lust asstme an observing frequency for cach of the interferometers.," Before we can do this, however, we must assume an observing frequency for each of the interferometers."
 The current audOVRO arrays use auplificrs which operate over the range 2636 GUIIz (Reese et al., The current and arrays use amplifiers which operate over the range $26-36$ GHz (Reese et al.
 2000)., 2000).
 We assume a frequency centered ou 30 GIIz for these arravs., We assume a frequency centered on 30 GHz for these arrays.
 The upgradedOVRO array aud the are expected to have amplifiers which operate at 26—36 CGIIz andat ου115 GIIz., The upgraded array and the are expected to have amplifiers which operate at $26-36$ GHz and at $85-115$ GHz.
 We asstune frequencies centered on 30 and 90 GIlz for these arravs., We assume frequencies centered on 30 and 90 GHz for these arrays.
 Thus. we generate six different decrement maps for cach model cluster: 30 CdIz naps for the current. arrays and 30 and 90 CGIIz maps for the aud upgraded.OVRO array.," Thus, we generate six different decrement maps for each model cluster: 30 GHz maps for the current arrays and 30 and 90 GHz maps for the and upgraded array."
 Tuterferometers do not image the sky, Interferometers do not image the sky.
 Rather. they measure the Fourier transform of the SZ effect mteusitv decrement multiplied by the primary beam.," Rather, they measure the Fourier transform of the SZ effect intensity decrement multiplied by the primary beam."
" A particular value of this observable is referred to as a “visibility” aud 1s given by where s and c are the conjugate variables. 0, aud 06, arc the projected sky coordinates. aud ACO).0,) is the primary beam"," A particular value of this observable is referred to as a “visibility” and is given by where $u$ and $v$ are the conjugate variables, $\theta_x$ and $\theta_y$ are the projected sky coordinates, and $A(\theta_x,\theta_y)$ is the primary beam."
" At 30 GIIz. the scusitivity patterns for the primary bemus of the aud arrays (both cureut aud uperaded) and the are. or will be. well-vepreseuted by Caussiaus with PFWHAD = 6.6% L2'. aud 10—12% respectively,"," At 30 GHz, the sensitivity patterns for the primary beams of the and arrays (both current and upgraded) and the are, or will be, well-represented by Gaussians with FWHM $\approx$ $6.6\arcmin$, $4.2\arcmin$, and $10-12\arcmin$, respectively."
" We axstume a EWIIM. = 10.8"" for the SZA.", We assume a FWHM = $10.8\arcmin$ for the SZA.
 For more ou the above formalisi see White et al. (, For more on the above formalism see White et al. (
1999).,1999).
 We anultiply the SZ effect maps by the Caussiau primary beams described above and Fourier trausforii the result., We multiply the SZ effect maps by the Gaussian primary beams described above and Fourier transform the result.
 We discard all visibilitics for baseBues (Ry.—vu?| 02) that are not probed by these interferometers.," We discard all visibilities for baselines $R_{u,v} \equiv 
\sqrt{u^2+v^2}$ ) that are not probed by these interferometers."
" At 30CIIz. the andOVRO arrays probe uultipole moments of f=27R,. 21000 ancl 7000. respectively."," At 30GHz, the and arrays probe multipole moments of $\ell = 2 \pi R_{u,v}$ $\gtrsim 4000$ and $7000$, respectively."
 We assume uniform coverage over the ranges 1200=(z20000 and 1200&(τς20000 for these arravs. respectively.," We assume uniform coverage over the ranges $4200 \lesssim \ell 
\lesssim 20000$ and $7200 \lesssim \ell \lesssim 20000$ for these arrays, respectively."
 The smaller 3.5 ia dishes of the will be able to probe shorter baselines (larger aneular scales) aud are expected ο sample (6z2000., The smaller 3.5 m dishes of the will be able to probe shorter baselines (larger angular scales) and are expected to sample $\ell \gtrsim 2000$.
 We assume uniform coverage over the range 2100ες20000 (also at 30 CIIz)., We assume uniform coverage over the range $2400 \lesssim \ell \lesssim 20000$ (also at 30 GHz).
" Finally, we acd random Caussian noise to each of he mock observations."," Finally, we add random Gaussian noise to each of the mock observations."
 At 30 CILz. we assume syste cluperatures (scaled to the atinosphliere) of 15 Is for the aud arrays aud 30 IS for theSZA.," At 30 GHz, we assume system temperatures (scaled to the atmosphere) of 45 K for the and arrays and 30 K for the."
 At 90 Cz. we assiunae svstem teiiperatures of LOO I& for the ix or the upgradedOVRO array.," At 90 GHz, we assume system temperatures of 100 K for the and for the upgraded array."
 The curveutDIMA/OVRO arravs have effective noise bandwidthns of approximately 550/1000 ΑΠΕ. while both the upgraded.OVRO array aix he SZA are expected to have bandwidths of 8 1).," The current arrays have effective noise bandwidths of approximately 550/1000 MHz, while both the upgraded array and the SZA are expected to have bandwidths of 8 GHz."
 A otal iutegration tine of LO hours is assumed for cach of he axrays at both 30 aud 90 GIIz., A total integration time of 40 hours is assumed for each of the arrays at both 30 and 90 GHz.
 For specificity. the mock observations for the current and arrays have rum noise levels of approximately 150 pJy | an 70 µ. beam+. respectively. while the uperadedOVRO array has noise levels of 25 444 + and 50 44x + at 30 and 90 ΟΠΣ (respectively) aud the lias noise levels of 10 pJy beam1 aud 110 jy | at 30 and 90 11. respectively.," For specificity, the mock observations for the current and arrays have rms noise levels of approximately 150 $\mu$ Jy $^{-1}$ and 70 $\mu$ Jy $^{-1}$ , respectively, while the upgraded array has noise levels of 25 $\mu$ Jy $^{-1}$ and 50 $\mu$ Jy $^{-1}$ at 30 and 90 GHz (respectively) and the has noise levels of 40 $\mu$ Jy $^{-1}$ and 140 $\mu$ Jy $^{-1}$ at 30 and 90 GHz, respectively."
 To analyse the mock observations. we follow the method described in a number of observational papers based on eenuimeBIALA/OVRO data (e.g... Carlstrom ct al.," To analyse the mock observations, we follow the method described in a number of observational papers based on genuine data (e.g., Carlstrom et al."
 1996: Grego et al., 1996; Grego et al.
 2000: 2001: Reese et al., 2000; 2001; Reese et al.
 2000: 2002: Jov ct al., 2000; 2002; Joy et al.
 2001: LaRoque et al., 2001; LaRoque et al.
 2003)., 2003).
 We model the observations with the isothermal § models®., We model the observations with the isothermal $\beta$ $^8$.
 We create model SZ effect laps for for various choices of the three free paraiaeters (yy. O-. aud 2). multiply the maps by the appropriate primary beams. Fourier transform tle maps. aud compare the results to the mock observatious in §LL1 via a \? statistic.," We create model SZ effect maps for for various choices of the three free parameters $y_0$, $\theta_c$, and $\beta$ ), multiply the maps by the appropriate primary beams, Fourier transform the maps, and compare the results to the mock observations in 4.1 via a $\chi^2$ statistic."
 The best-fit parameters are those which result vthe mininuun value of 47., The best-fit parameters are those which result in the minimum value of $\chi^2$.
 Note that the commparisous oe.)eteen the > models aud the mock observations are done iu the Fourier domain (also referred to as the à0 planc). where the noise characteristics aud spatial filtering of the interferometers are well understood.," Note that the comparisons between the $\beta$ models and the mock observations are done in the Fourier domain (also referred to as the $u-v$ plane), where the noise characteristics and spatial filtering of the interferometers are well understood."
 Like the observers. we o not “decouvolve” the mock observations for analysis.," Like the observers, we do not “deconvolve” the mock observations for analysis."
" To demonstrate the quality of the datato be produced bv the aud the uperadedOVRO array. we plot iu Figure 9 confidence contours for the > and 0, parameters from “observations of a Ly=6.7 keV (Mj;z5.6XN 101437.) cluster at 2=1."," To demonstrate the quality of the datato be produced by the and the upgraded array, we plot in Figure 9 confidence contours for the $\beta$ and $\theta_c$ parameters from “observations” of a $T_X = 6.7$ keV $M_{tot} \approx 5.6 \times 10^{14} 
M_{\odot}$ ) cluster at $z = 1$."
" The third free parameter. gy. is allowed to asstme its best-fit value at cach pair of Jj and 06, (ie. the plot is a projection of the coutidence ""volume onto the >0, plane)."," The third free parameter, $y_0$, is allowed to assume its best-fit value at each pair of $\beta$ and $\theta_c$ (i.e., the plot is a projection of the confidence “volume” onto the $\beta-\theta_c$ plane)."
 The contours correspoucd to a Ay? and the filled squares indicate. the best-fit models., The contours correspond to a $\Delta \chi^2$ and the filled squares indicate the best-fit models.
 The projection of Ay?=1 line outo the axes eives the confidence interval for cach of the two paralmcters., The projection of $\Delta \chi^2 = 1$ line onto the axes gives the confidence interval for each of the two parameters.
 The contour labeled “current” is the result of fitting the SZ effect surface brightuess model to the mock obscervatious with the currently operational /OVRO arrays (both arravs lave similar scusitivity patterus)., The contour labeled “current” is the result of fitting the SZ effect surface brightness model to the mock observations with the currently operational arrays (both arrays have similar sensitivity patterns).
" Thecontours labeled “SZA™ and ""OVROare the results of sinultaucouslv fitting the surface brightuess model to the mock 30 aud 90 GIIz observations aud the mock 30", Thecontours labeled “SZA” and “OVRO”are the results of simultaneously fitting the surface brightness model to the mock 30 and 90 GHz observations and the mock 30
We find that the equivalent width (EW) of the he Fe line shows a variation through the beat evele. with no trend being visible.,"We find that the equivalent width (EW) of the $\alpha$ Fe line shows a variation through the beat cycle, with no trend being visible."
 his is likely to be due to a varving continuum. since the Ίνα Fe line normalisation is highest during the main-on state while it remains low ancl almost constant at other 35 day. phases.," This is likely to be due to a varying continuum, since the $\alpha$ Fe line normalisation is highest during the main-on state while it remains low and almost constant at other 35 day phases."
 Similarly. the line width is ercatest during the main-on state.," Similarly, the line width is greatest during the main-on state."
 The energy of the second component remains approximately constant at ~7.0 keV. except during the short-on state where it is significantly lower (76.85 keV).," The energy of the second component remains approximately constant at $\sim$ 7.0 keV, except during the short-on state where it is significantly lower $\sim$ 6.85 keV)."
 The feature is strongest (larger. LW) during the states of low X-ray intensity. although i0..is undetectable in some of the exposures in this phase range.," The feature is strongest (larger EW) during the states of low X-ray intensity, although it is undetectable in some of the exposures in this phase range."
 On the other hand. its normalisation is strongest. curing the short-on.," On the other hand, its normalisation is strongest during the short-on."
 In some of the observations. we do not find evidence for a significant Fe XXVI line.," In some of the observations, we do not find evidence for a significant Fe XXVI line."
 In order to determine upper limits to EW and. line normalisation. we added a Gaussian component to the best fit spectral model. by fixing he central energy at 7 keV and the width at 0.1. keV. The upper limits obtained for the datasets taken. during he main-on are consistent with the line being undetected xeause of the strong continuum emission.," In order to determine upper limits to EW and line normalisation, we added a Gaussian component to the best fit spectral model, by fixing the central energy at 7 keV and the width at 0.1 keV. The upper limits obtained for the datasets taken during the main-on are consistent with the line being undetected because of the strong continuum emission."
 In order to show the significance of both the Fe features. we show in Figure 7 the X-ray spectrum between 5-9 keV or O44;—0.02 ancl &35=0.79.," In order to show the significance of both the Fe features, we show in Figure \ref{fe_zoom} the X-ray spectrum between 5-9 keV for $\Phi_{35}$ =0.02 and $\Phi_{35}$ =0.79."
 The solid line is the best fit spectral model with the normalisations of the two Gaussian components set to zero., The solid line is the best fit spectral model with the normalisations of the two Gaussian components set to zero.
 We find no evidence for the presence of a Compton downscattered shoulder (expected near £oLufame|ES) where fy is the centroid of the line) of the kind recently observed in LUETCGS spectra of the high mass X-ray binary GX 301-2 (Watanabe ct al., We find no evidence for the presence of a Compton downscattered shoulder (expected near $E_0 - E_0/(m_ec^2 + E_0)$ where $E_0$ is the centroid of the line) of the kind recently observed in HETGS spectra of the high mass X-ray binary GX 301-2 (Watanabe et al.
 2003), 2003).
 The variation of the Fe Whe line over the neutron star spin period and its correlation with the thermal or non-thermal spectral Components was studied. by ROR., The variation of the Fe $\alpha$ line over the neutron star spin period and its correlation with the thermal or non-thermal spectral components was studied by R02.
 We found. no evidence for a significant variation in the line parameters, We found no evidence for a significant variation in the line parameters
disk.,disk.
 The radius of disk. raja. is determined following Iirashita&Ferrara(2002).," The radius of disk, $r_{\rm disk}$, is determined following \citet{Hir02}."
". We estimate the typical scale height. £7. from lvdrostatic οπα (Shakura&Suuvaev1988) for ου otherwise. assume LPrai= 0.1. where e;=(kp""ME is the isothermal sound velocity."," We estimate the typical scale height, $H$, from hydrostatic equilibrium \citep{Sha88}
 for $H/r_{\rm disk}\le0.1$ , otherwise, assume $H/r_{\rm disk}=0.1$ , where $v_{s}=(k_{B}T/\mu m_{\rm H})^{\frac{1}{2}}$ is the isothermal sound velocity."
" Therefore. initial hydrogen density of disk. ο. is estimated as The initial mass ofhydrogen iu the galaxy. My. js written as where Ay. is the mitial mass of helimm in the ealaxv. Ma,=Mg|Adie. ye isthe mass of a heliuni atom aud yy. is the helinuun abuudauce."," Therefore, initial hydrogen density of disk, $n_{\rm H}$, is estimated as The initial mass ofhydrogen in the galaxy, $M_{\rm H}$, is written as where $M_{\rm He}$ is the initial mass of helium in the galaxy, $M_{\rm gas}=M_{\rm H}+M_{\rm He}$, $m_{\rm He}$ isthe mass of a helium atom and $y_{\rm He}$ is the helium abundance."
 We assmue gp.=0.0972 (COnukai 2000).., We assume $y_{\rm He}=0.0972$ \citep{Omu00}. .
 Note that since massive star jonizes surrounding eas and forms an expanding II IT region (e.g.Whalenetal.2001:INitaviunaetal. 200L).. we adopt eas density around SN progenitor. (ax. as beime different from the hvdrosen gas density iu our one-zoue galaxy model. og.," Note that since massive star ionizes surrounding gas and forms an expanding H II region \citep[e.g.][]{Wha04, Kit04}, we adopt gas density around SN progenitor, $n_{\rm SN}$, as being different from the hydrogen gas density in our one-zone galaxy model, $n_{\rm H}$."
 This is because the eas density around SN progenitor. max. ds closely related with dust destruction processby reverse shocks driveu by SNe. as shown in Section 2.2..," This is because the gas density around SN progenitor, $n_{\rm SN}$, is closely related with dust destruction processby reverse shocks driven by SNe, as shown in Section \ref{subsec:dustprod}."
" We expect that the SFR. Wit). is roughly proportional to f,d. since a represeutative timescale of the dvuauiics of the eas disk is f4,."," We expect that the SFR, $\Psi(t)$ , is roughly proportional to $t_{\rm cir}^{-1}$, since a representative timescale of the dynamics of the gas disk is $t_{\rm cir}$."
" We assume that where fi, is the mass fraction of molecular hydrogen to the total eas.", We assume that where $f_{\rm H_{2}}$ is the mass fraction of molecular hydrogen to the total gas.
 We should note that observationally Dieieletal.(2008) find that IIo is converted iuto stars af a constant eficicucy ii nearby spirals and Cmnedinetal.(2009) show that the star formation recipe iu ealaxy formation simulation in which star formation occurs onlv in the molecular gas can reproduce the observational correlations between SER aud the total gas density., We should note that observationally \citet{Big08} find that ${\rm H}_{2}$ is converted into stars at a constant efficiency in nearby spirals and \citet{Gne09} show that the star formation recipe in galaxy formation simulation in which star formation occurs only in the molecular gas can reproduce the observational correlations between SFR and the total gas density.
" We calculate the time evolutions of the masses of hydrogen and helium gases; Ma. stars. Mi. and metal of clement ἐν AMj,;. in the galaxy bx using the following equatious where ποΕΣπε Teas|7 aud mua. and Wye are the gas mass of So)hydrogen aud helium aud the moetal mass of clement ;/ iu SN ejecta. respectively."," We calculate the time evolutions of the masses of hydrogen and helium gases, $M_{\rm gas}$, stars, $M_{\rm star}$, and metal of element $i$, $M_{{\rm m},i}$, in the galaxy by using the following equations where $M_{\rm ISM}(t) = M_{\rm gas}(t)+\sum_{i}M_{{\rm m},i}(t)$, $\overline{m_{\rm ejecta}} = \overline{m_{\rm gas}}+\sum_{i}\overline{m_{{\rm m},i}}$ and $\overline{m_{\rm gas}}$ and $\overline{m_{{\rm m},i}}$ are the gas mass of hydrogen and helium and the metal mass of element $i$ in SN ejecta, respectively."
 The mass retuunug to the ISML ino(n). μαςΗΕ) aud My) through a SN with progenitor mass. my. are taken from Uineda&Nomoto(2002) iu the case of i=13. 20. 25 and 30 AL... ma.," The mass returning to the ISM, $m_{\rm ejecta}(m)$, $m_{\rm gas}(m)$ and $m_{{\rm m},i}(m)$ through a SN with progenitor mass, $m$, are taken from \citet{Ume02} in the case of $m=13$, $20$, $25$ and $30$ $M_{\odot}$. $\overline{m_{\rm gas}}$,"
" Hidecra Ald Wy, ἀπὸ IME-averaged μας11}. HhgeeraA] aud Ayia). respectively,"," $\overline{m_{\rm ejecta}}$ and $\overline{m_{m,i}}$ are IMF-averaged $m_{\rm gas}(m)$, $m_{\rm ejecta}(m)$ and $m_{\rm m,i}(m)$, respectively."
 Note that metals consist of not oulv heavy clemeuts iu gas phase but also those in dust eraius., Note that metals consist of not only heavy elements in gas phase but also those in dust grains.
 We follow the time evolution of molecular lnassfraction. fj. ionization degree. c. aud gas temperature. 2.," We follow the time evolution of molecular massfraction, $f_{{\rm H}_{2}}$ , ionization degree, $x$ , and gas temperature, $T$ ."
" We definethe molecular fraction of hydrogen as where vy, aud my are the nuuber cdensitics of molecular Lydrogen and hwdrogen nuclei. respectively."," We definethe molecular fraction of hydrogen as where $n_{{\rm H}_{2}}$ and $n_{\rm H}$ are the number densities of molecular hydrogen and hydrogen nuclei, respectively."
 The molecular fraction is very miportant in our models. because it determines the final cooling rate of low-inetallicity gas.," The molecular fraction is very important in our models, because it determines the final cooling rate of low-metallicity gas."
 Themetal-free, Themetal-free
where (?)..,where \citep{1967ApJ...150..131R}.
" The Fermi chemical potential prj. which includes the rest-mass energy. is connected with the numberdifference. electrons minus positrons. through where 0,(T.V.με) denotes the partial derivative OO.V.μι)Ὁμι."," The Fermi chemical potential $\mu_F$, which includes the rest-mass energy, is connected with the number, electrons minus positrons, through where $\Omega_{\mu_F}(T,V,\mu_F)$ denotes the partial derivative $\partial\Omega(T,V,\mu_F)/\partial\mu_F$."
 Clearly the positron contribution. which is due to the second term of DE. becomes insignificant whenever © is large (sav ©Z 15).," Clearly the positron contribution, which is due to the second term of $D_\pm$, becomes insignificant whenever $\phi$ is large (say $\phi\geq15$ )."
 The last equation determines µε.VN) as a function of T. V and NV (actually the Fermi parameter ej=(pj—mc)/kT in terms of T and N/V).," The last equation determines $\mu_F(T,V,N)$ as a function of $T$ , $V$ and $N$ (actually the Fermi parameter $\epsilon_F=(\mu_F-mc^2)/kT$ in terms of $T$ and $N/V$ )."
" The number difference V must satisfy the equation of charge neutrality If, in the expression for the free energy. electrons appeared only in the second line. (hen it would follow that where (he subscript V denotes the partial derivative wilh respect to AN. at constant f£. and V."," The number difference $N$ must satisfy the equation of charge neutrality If, in the expression for the free energy, electrons appeared only in the second line, then it would follow that where the subscript $N$ denotes the partial derivative with respect to $N$, at constant $T$ and $V$."
 Thus py would indeed be the electron chemical potential jp9Fy+ine?., Thus $\mu_F$ would indeed be the electron chemical potential $\mu=F_N+mc^2$.
 We maintain the distinction (between se and pry) because other parts of the free energyfor example the pressure ionization term Fp; loo. depend on the electron number density.," We maintain the distinction (between $\mu$ and $\mu_F$ ) because other parts of the free energy—for example the pressure ionization term $F_{PI}$ —too, depend on the electron number density."
 The pressure p=—£Fi. the entropy 5=—Fy and their derivatives require derivatives of OQ. will respect (ο 9 or ©. up to the second order.," The pressure $p=-F_V$, the entropy $S=-F_T$ and their derivatives require derivatives of $\Omega$ , with respect to $\beta$ or $\phi$ up to the second order."
" This leads to five additional Fermi-Dirac integrals. in which P(e/) is replaced by D'(e/3). (e/2)D""(e/9). E""(ef/2). (ef⋡∕2)""(ef) or ie/3)T""Ced)."," This leads to five additional Fermi-Dirac integrals, in which $\Gamma(\epsilon/\beta)$ is replaced by $\Gamma'(\epsilon/\beta)$, $(\epsilon/\beta)\Gamma'(\epsilon/\beta)$, $\Gamma''(\epsilon/\beta)$, $(\epsilon/\beta)\Gamma''(\epsilon/\beta)$ or $(\epsilon/\beta)^2\Gamma''(\epsilon/\beta)$."
 In the degenerate case. when ej=ὡ-ο>5. O is caleulated by Sommerleld's nethod. and then differentiated.," In the degenerate case, when $\epsilon_F=\phi-\beta>5$, $\Omega$ is calculated by Sommerfeld's method, and then differentiated."
 Otherwise the six integrals involving the first. electronic. part of De are calculated in one swoop. using Gaussian quadrature.," Otherwise the six integrals involving the first, electronic, part of $D_\pm$ are calculated in one swoop, using Gaussian quadrature."
 The nodes aud weights or this equadrature are calculatedat the beginning of the run. and their number can be chosen by the user (the code's default is 12 nodes).," The nodes and weights for this quadrature are calculatedat the beginning of the run, and their number can be chosen by the user (the code's default is 12 nodes)."
 Thepositronic contribution.which is," Thepositronic contribution,which is"
vicinity of the ionization front.,vicinity of the ionization front.
 It would be prohibitively expensive computationally and unnecessary to force the hydrodynamics to evolve at the same time scale as the radiative transfer., It would be prohibitively expensive computationally and unnecessary to force the hydrodynamics to evolve at the same time scale as the radiative transfer.
" Hence, in our implementation the radiative transfer time scale is allowed to drop below that of the hydrodynamics."," Hence, in our implementation the radiative transfer time scale is allowed to drop below that of the hydrodynamics."
 The energy equation is solved through operator splitting by computing the ionization state of the gas and the amount of photoionization heating and atomic cooling between the hydrodynamical time steps., The energy equation is solved through operator splitting by computing the ionization state of the gas and the amount of photoionization heating and atomic cooling between the hydrodynamical time steps.
" During the ionization timesteps, all parameters of the gas are taken to be those at the end of the immediately preceding hydrodynamical timestep with the exception of the fluid density and temperature."," During the ionization timesteps, all parameters of the gas are taken to be those at the end of the immediately preceding hydrodynamical timestep with the exception of the fluid density and temperature."
" The density is varied linearly from the previous hydrodynamical timestep to that of the current ionization timestep while the radiative transfer processes along with the gas temperature are evolved on their own, shorter, adaptive time scale."," The density is varied linearly from the previous hydrodynamical timestep to that of the current ionization timestep while the radiative transfer processes along with the gas temperature are evolved on their own, shorter, adaptive time scale."
 The simulations were carried out in a box 25h~'Mpc (comoving) on a side.," The simulations were carried out in a box $25h^{-1}\,{\rm Mpc}$ (comoving) on a side."
" This corresponds to the order of the expected size of a iionization front at the helium reionization epoch, when the characteristic diameter of the zzone produced by an individual QSO is comparable to the mean distance between QSOs, which is typically ~50 Mpc (comoving) at zc3 (?7)."," This corresponds to the order of the expected size of a ionization front at the helium reionization epoch, when the characteristic diameter of the zone produced by an individual QSO is comparable to the mean distance between QSOs, which is typically $\sim50$ Mpc (comoving) at $z\simeq3$ \citep{2003A&A...408..499W, Meiksin05}."
" Then simulations were performedwith 512? gravitating particles and 256? grid zones for the fluid component, corresponding to a resolution of 98 kpc (comoving), sufficient to resolve the Jeans length."," Then simulations were performedwith $512^3$ gravitating particles and $256^3$ grid zones for the fluid component, corresponding to a resolution of 98 kpc (comoving), sufficient to resolve the Jeans length."
" The incident spectrum corresponds to a hybrid starburst-QSO model as in ?:: aan initial starburst spectrum is turned on at z—8 and transformed into a power-law spectrum f,«v??, turning on the power-law spectrum initially at z—4 and ramping it up to its full value at z—3, while the starburst spectrum is ramped down during the same interval."," The incident spectrum corresponds to a hybrid starburst–QSO model as in \citet{2007MNRAS.380.1369T}: an initial starburst spectrum is turned on at $z=8$ and transformed into a power-law spectrum $f_\nu\propto\nu^{-0.5}$, turning on the power-law spectrum initially at $z=4$ and ramping it up to its full value at $z=3$, while the starburst spectrum is ramped down during the same interval."
" The intensity of the starburst spectrum was adjusted to ensure the hydrogen ionization front swept across the simulation volume by z6, although some shadowed regions persist along the edge of the box until z~5.5."," The intensity of the starburst spectrum was adjusted to ensure the hydrogen ionization front swept across the simulation volume by $z\simeq6$, although some shadowed regions persist along the edge of the box until $z\simeq5.5$."
" The intensity of the power-law spectrum was adjusted so that the Herr--ionization front swept across the box by z~3.2, again with some shadowing persisting to z<3."," The intensity of the power-law spectrum was adjusted so that the -ionization front swept across the box by $z\simeq3.2$, again with some shadowing persisting to $z\lsim3$."
 This corresponds to a hydrogen ionization rate from the QSO at full intensity half that of the starburst at its full intensity., This corresponds to a hydrogen ionization rate from the QSO at full intensity half that of the starburst at its full intensity.
 The heating rate per ionization of aas the iionization front passes through the box at z—3.3 is shown in Fig. 2., The heating rate per ionization of as the ionization front passes through the box at $z=3.3$ is shown in Fig. \ref{fig:epsHeII}.
 The filtering effect of the IGM in hardening the ionizing radiation field is clearly discernible in the pre-ionized gas ahead of the ionization front., The filtering effect of the IGM in hardening the ionizing radiation field is clearly discernible in the pre-ionized gas ahead of the ionization front.
" The heating rate per ionization is substantially enhanced, reaching values of up to 100—400 eV. By comparison, once the front has passed, the energy per iionization is ~22 eV. 'The resulting temperature structure is shown in Fig."," The heating rate per ionization is substantially enhanced, reaching values of up to $100-400$ eV. By comparison, once the front has passed, the energy per ionization is $\sim22$ eV. The resulting temperature structure is shown in Fig."
" as the IL-front passes through the box at z=3.3 and at z=2.0, after iionization has completed."," \ref{fig:Temp_map} as the I-front passes through the box at $z=3.3$ and at $z=2.0$, after ionization has completed."
 The effect of shadowing produced by density inhomogeneities in the IGM are apparent as the advance of the ionization front is delayed beyond the densest regions., The effect of shadowing produced by density inhomogeneities in the IGM are apparent as the advance of the ionization front is delayed beyond the densest regions.
 This may best be seen by comparing the temperature map at z—3.3 in Fig., This may best be seen by comparing the temperature map at $z=3.3$ in Fig.
 3 with the density map in Fig. 4.. (, \ref{fig:Temp_map} with the density map in Fig. \ref{fig:Gas_density_map}. (
The principal density structures change little between z—3.3 and z—2.0 in comoving coordinates.),The principal density structures change little between $z=3.3$ and $z=2.0$ in comoving coordinates.)
" Large temperature boosts of 3—4x10* K occur downstream from dense clumps, consistent with the striations of high heating rates per ionization in Fig."," Large temperature boosts of $3-4\times10^4$ K occur downstream from dense clumps, consistent with the striations of high heating rates per ionization in Fig."
" 2 as expected for photoionization by a radiation field hardened by intervening optically thick clumps with optical depths of several tens, as discussed in Sec. 2.."," \ref{fig:epsHeII} as expected for photoionization by a radiation field hardened by intervening optically thick clumps with optical depths of several tens, as discussed in Sec. \ref{sec:photoheating}."
" By z—2.0, most of the high temperature striations have been replaced by gas temperatures that trace the matter distribution now that rreionization has completed."," By $z=2.0$, most of the high temperature striations have been replaced by gas temperatures that trace the matter distribution now that reionization has completed."
" Pronounced hot regions, however, persist, particularly in underdense regions unable to achieve thermal equilibrium."," Pronounced hot regions, however, persist, particularly in underdense regions unable to achieve thermal equilibrium."
 These lay behind dense structures while the gas was reionized., These lay behind dense structures while the gas was reionized.
" The filtering of the radiation field by the dense structures hardened the spectrum of the incident radiation passing downstream to the underdense gas, as shown in Fig. 2, "," The filtering of the radiation field by the dense structures hardened the spectrum of the incident radiation passing downstream to the underdense gas, as shown in Fig. \ref{fig:epsHeII}, ,"
producing the high temperatures., producing the high temperatures.
 This effect is especially prevalent in the centre-right portion of the figure., This effect is especially prevalent in the centre-right portion of the figure.
" Including radiative transfer in the ionization of ssubstantially boosts the temperature of the gas, increasing it typically by 15—20x10? K over the optically thin reionization limit, as shown in Fig. 5.."," Including radiative transfer in the ionization of substantially boosts the temperature of the gas, increasing it typically by $15-20\times10^3$ K over the optically thin reionization limit, as shown in Fig. \ref{fig:T_hist}."
 The temperature-density relation is shown in Fig. 6.., The temperature-density relation is shown in Fig. \ref{fig:dT_rho_contours}.
" The plateau with a gas temperature of ~20—30x10? K at the mean baryon density agrees well with the estimate of ? of To~20—30x10? K at z— 3, with only a weakly increasing temperature with overdensity, as inferred from the Doppler widths measured in high resolution spectraof the fforest."," The plateau with a gas temperature of $\sim20-30\times10^3$ K at the mean baryon density agrees well with the estimate of \citet{Schaye00} of $T_0\sim20-30\times10^3$ K at $z=3$ , with only a weakly increasing temperature with overdensity, as inferred from the Doppler widths measured in high resolution spectraof the forest."
" The temperature found in the simulation, however, is much larger than the temperature at the mean density"," The temperature found in the simulation, however, is much larger than the temperature at the mean density"
(Biges1992)..,\citep{big92a}.
 Furthermore. it aided. the identification. of nulline pulsars whose olf times are more than a few minutes. sources that would often be missed by surveys with observation times of this order.," Furthermore, it aided the identification of nulling pulsars whose off times are more than a few minutes, sources that would often be missed by surveys with observation times of this order."
" Phe extreme example of this is the discovery of ""rotating radio transients” or RILATs (MeLaughlinetal.2006).. pulsars in which just single pulses are observed at intervals of minutes to hours."," The extreme example of this is the discovery of “rotating radio transients” or RRATs \citep{mll+06}, pulsars in which just single pulses are observed at intervals of minutes to hours."
 As can be seen from Table 1.. seven of the 23 pulsars in our sample have NEs of or more (within the uncertainties). with two having NES of order95'4.. the largest. known (apart. [rom the ΑΕ).," As can be seen from Table \ref{tb:obs}, seven of the 23 pulsars in our sample have NFs of or more (within the uncertainties), with two having NFs of order, the largest known (apart from the RRATs)."
 Table 2. lists previously known nulling pulsars and their NE., Table \ref{tb:null} lists previously known nulling pulsars and their NF.
 While some of these values may be misleading because of low signal-to-noise ratio in the observations. the table gives the best available indication of the distribution of nulling in known pulsars.," While some of these values may be misleading because of low signal-to-noise ratio in the observations, the table gives the best available indication of the distribution of nulling in known pulsars."
 Fig., Fig.
 7 shows that the present work has substantially increased the number of pulsars with high NE. doubling the number known with NE z50%.," \ref{fg:null_hist} shows that the present work has substantially increased the number of pulsars with high NF, doubling the number known with NF $\ge 50$."
. Although our observations are generally not very sensitive to short nulls (because of limited. signal/noise on these relatively weak pulsars). it is clear that pulsars with low NE are much more common than those with high NE.," Although our observations are generally not very sensitive to short nulls (because of limited signal/noise on these relatively weak pulsars), it is clear that pulsars with low NF are much more common than those with high NF."
 Fig., Fig.
 S shows the distribution of known nulling pulsars on the P P nocitemhthO05— This diagram suggests that the degree. of nulling is related more to age than period., \ref{fg:ppdot} shows the distribution of known nulling pulsars on the $P$ $\dot P$ \\nocite{mhth05} This diagram suggests that the degree of nulling is related more to age than period.
 Although some »ulsars with large 7. have small NE. all of the pulsars with aree NE. have z 1. Myr ancl most have 7.c 5 Mvwr.," Although some pulsars with large $\tau_c$ have small NF, all of the pulsars with large NF have $\tau_c >$ 1 Myr and most have $\tau_c >$ 5 Myr."
 In contrast. even excluding the recycled. pulsar 3922 with a period of 0.172 s. PSR 1820. 0509. with period of 3437 s and age of 5.7 Myr has a NE of. whereas PSR 0525|21. with a period of 3.75 s and an age of 1.5 Myr vas à NE of25'4.," In contrast, even excluding the recycled pulsar $-$ 3922 with a period of 0.172 s, PSR $-$ 0509, with period of 0.337 s and age of 5.7 Myr has a NF of, whereas PSR B0525+21, with a period of 3.75 s and an age of 1.5 Myr has a NF of."
. Phe longest-period pulsar in our sample. SR 1223 (2.86 8). has a NE of justΕκ.," The longest-period pulsar in our sample, PSR $-$ 1223 (2.86 s), has a NF of just."
.. There is no strong correlation of NE. with. profile morphology., There is no strong correlation of NF with profile morphology.
 As Figs 2 6 shows. nulls are observed. in oulsars. with both narrow single-component pulse. profiles and wider multi-component profiles.," As Figs \ref{fg:nulls} – \ref{fg:1401} shows, nulls are observed in pulsars with both narrow single-component pulse profiles and wider multi-component profiles."
" As a fraction of the population. nulling tends to be more comnmion in multi-component profiles (Rankin1986:Biges1992): such pulsars tend to have larger 7 so this is just the same correlation as that shown in Fig. ὃν,"," As a fraction of the population, nulling tends to be more common in multi-component profiles \citep{ran86,big92a}; such pulsars tend to have larger $\tau_c$ so this is just the same correlation as that shown in Fig. \ref{fg:ppdot}."
 Unfortunately. polarisation data exist for only a few of these pulsars. so no clear relationships with impact parameter or other polarisation characteristics can be crawn.," Unfortunately, polarisation data exist for only a few of these pulsars, so no clear relationships with impact parameter or other polarisation characteristics can be drawn."
 There is no obvious correlation of null evele time with pulse period or characteristic age., There is no obvious correlation of null cycle time with pulse period or characteristic age.
 Do RILCVXES represent an extreme form of pulse nulling or are they an intrinsically cifferent phenomenon?, Do RRATs represent an extreme form of pulse nulling or are they an intrinsically different phenomenon?
 Based on a study of the pulse modulation in PSR. 0056|14. Weltevrecdeetal.(2006). have suggested. that IUIS are highly modulated: pulsars which occasionally emit. very," Based on a study of the pulse modulation in PSR B0656+14, \citet{wsrw06} have suggested that RRATs are highly modulated pulsars which occasionally emit very"
In this section we discuss in more detail the properties of the X-ray sources with radio identification ancl available spectroscopic information.,In this section we discuss in more detail the properties of the X-ray sources with radio identification and available spectroscopic information.
 In addition to these sources we also discuss the properties of the candidate star (source 3£225) identified in our sample., In addition to these sources we also discuss the properties of the candidate star (source 25) identified in our sample.
 his source has a narrow emission. line optical spectrum at z=(0.276 (see Figure. A4))., This source has a narrow emission line optical spectrum at $z=0.276$ (see Figure \ref{fig_spec}) ).
 The relative strengths of the 112 and OLI]5007A emission lines as ve as the presence of high excitation lines (c.g NeLEE]386is Ne V]3426X). suggest. obscured AGN activity.," The relative strengths of the $\rm H\beta$ and $\rm [O\,III]\,5007\AA$ emission lines as well as the presence of high excitation lines (e.g $\rm
[Ne\,III]\,3869$, $\rm [Ne\,V]\,3426\,\AA$ ) suggest obscured AGN activity."
 This consistent with the flat N-rayv spectrum suggesting a rest frame absorbing column of z41077em?E (p= 1.7).," This is consistent with the flat X-ray spectrum suggesting a rest frame absorbing column of $\rm
\approx 4\times10^{22}\,cm^{-2}$ $\Gamma=1.7$ )."
 The optical L-band image with the X-ray and the radio positions overlaid are shown in Figure AS., The optical $R$ -band image with the X-ray and the radio positions overlaid are shown in Figure \ref{fig_image}.
 Visual inspection of the L-band image shows two components likely to be in the process of interaction/merging., Visual inspection of the $R$ -band image shows two components likely to be in the process of interaction/merging.
We follow the evolution of the MBH population resulting from the seed formation process delineated above in a ACDM Universe.,We follow the evolution of the MBH population resulting from the seed formation process delineated above in a$\Lambda$ CDM Universe.
" Our approach is similar to the one described in Volonteri, Haardt Madau (2003)."," Our approach is similar to the one described in Volonteri, Haardt Madau (2003)."
" We simulate the merger history of present-day halos with masses in the range 101«M starting from z=20, via a Monte Carlo algorithm based on the extended Press-Schechter formalism."," We simulate the merger history of present-day halos with masses in the range $10^{11}<M<10^{15}\,\msun$ starting from $z=20$, via a Monte Carlo algorithm based on the extended Press-Schechter formalism."
 Every halo entering the merger tree is assigned a spin parameter according to eqn. 1.., Every halo entering the merger tree is assigned a spin parameter according to eqn. \ref{plambda}.
 Recent work on the fate of halo spins during mergers in cosmological simulations has led to conflicting results: Vitvitska et al. (, Recent work on the fate of halo spins during mergers in cosmological simulations has led to conflicting results: Vitvitska et al. (
"2002) suggest that the spin parameter of a halo increases after a major merger, and the angular momentum decreases after a long series of minor mergers; D'Onghia Navarro (2007) find instead no significant correlation between spin and merger history.","2002) suggest that the spin parameter of a halo increases after a major merger, and the angular momentum decreases after a long series of minor mergers; D'Onghia Navarro (2007) find instead no significant correlation between spin and merger history."
" Given the unsettled nature of this matter, we adopt Occam's razor to guide us, and assume that the spin parameter of a halo is not modified by its merger history."," Given the unsettled nature of this matter, we adopt Occam's razor to guide us, and assume that the spin parameter of a halo is not modified by its merger history."
 When a halo enters the merger tree we assign seed MBHs by determining if the halo meets all the requirements described in Section 2 for the formation of a central mass concentration., When a halo enters the merger tree we assign seed MBHs by determining if the halo meets all the requirements described in Section 2 for the formation of a central mass concentration.
" As we do not self-consistently trace the metal enrichment of the intergalactic medium, we consider here a sharp transition threshold, and assume that the MBH formation scenario suggested by Lodato Natarajan ceases at z&15 (seealso??).."," As we do not self-consistently trace the metal enrichment of the intergalactic medium, we consider here a sharp transition threshold, and assume that the MBH formation scenario suggested by Lodato Natarajan ceases at $z\approx 15$ \citep[see also][]{sesana07, volonteriprep2007}."
" At z>15, therefore, whenever a new halo appears in the merger tree (because its mass is larger than the mass resolution), or a pre-existing halo modifies its mass by a merger, we evaluate if the gaseous component meets the conditions for efficient transport of angular momentum to create a large inflow of gas which can either form a MBH seed, or feed one if already present."," At $z>15$, therefore, whenever a new halo appears in the merger tree (because its mass is larger than the mass resolution), or a pre-existing halo modifies its mass by a merger, we evaluate if the gaseous component meets the conditions for efficient transport of angular momentum to create a large inflow of gas which can either form a MBH seed, or feed one if already present."
" The efficiency of MBH formation is strongly dependent on critical value of the Toomre parameter Q., which sets the frequency of formation, and consequently the number density of MBH seeds."," The efficiency of MBH formation is strongly dependent on critical value of the Toomre parameter $Q_{\rm c}$, which sets the frequency of formation, and consequently the number density of MBH seeds."
 We investigate the influence of this parameter in the determination of the global evolution of the MBH population., We investigate the influence of this parameter in the determination of the global evolution of the MBH population.
 Fig., Fig.
" 2 shows the number density of seeds formed in three models, with Φε=1.5 (low efficiency model A), Φε= (intermediate efficiency model B), and Φε=3 (high efficiency2 model C)."," \ref{fig2} shows the number density of seeds formed in three models, with $Q_{\rm c}=1.5$ (low efficiency model A), $Q_{\rm c}=2$ (intermediate efficiency model B), and $Q_{\rm c}=3$ (high efficiency model C)."
" The solid histograms show the total mass function of seeds formed by z—15 when this formation channel ceases, while the dashed histograms refer to seeds formed at a specific redshift slice at z—18."," The solid histograms show the total mass function of seeds formed by $z=15$ when this formation channel ceases, while the dashed histograms refer to seeds formed at a specific redshift slice at $z=18$."
" The number of seeds changes by about one order of magnitude from the least efficient to the most efficient model, consistent with the probabilities shown in Fig."," The number of seeds changes by about one order of magnitude from the least efficient to the most efficient model, consistent with the probabilities shown in Fig."
 1., 1.
" We assume that, after seed formation ceases, the z«15 population of MBHs evolves according to a “merger driven scenario"", as described in ?.."," We assume that, after seed formation ceases, the $z<15$ population of MBHs evolves according to a “merger driven scenario"", as described in \citet{volonteri06}."
 We assume that during major mergers MBHs accrete gas mass that scales with the fifth power of the circular velocity (or equivalently the velocity dispersion σε) of the host halo (?).., We assume that during major mergers MBHs accrete gas mass that scales with the fifth power of the circular velocity (or equivalently the velocity dispersion $\sigma_c$ ) of the host halo \citep{ferrarese02}.
" We thus set the final mass of the MBH at the end of the accretion episode to of the mass predicted by the Mpu--σε correlation, assuming that the scaling does not evolve with redshift."," We thus set the final mass of the MBH at the end of the accretion episode to of the mass predicted by the $M_{\rm BH}-\sigma_c$ correlation, assuming that the scaling does not evolve with redshift."
 Major mergers are defined as mergers between two dark matter halos with mass ratio between 1 and 10., Major mergers are defined as mergers between two dark matter halos with mass ratio between 1 and 10.
 BH mergers contribute to the mass increase of the remaining 10%.., BH mergers contribute to the mass increase of the remaining .
 In order to calculate the accreting black hole luminosity function and to follow the black hole mass growth during, In order to calculate the accreting black hole luminosity function and to follow the black hole mass growth during
"halo’s mass, but it will add a considerable amount of angular momentum.","halo's mass, but it will add a considerable amount of angular momentum."
" Thus, a measurement of any particular halo’s angular momentum or spin parameter will be dependent on the halo finding algorithm used in the study (Trentietal. 2010)."," Thus, a measurement of any particular halo's angular momentum or spin parameter will be dependent on the halo finding algorithm used in the study \citep{Trenti10}."
" In addition to the boundary issue, particle discreteness and resolution is also problematic when measuring halo properties such as spin and angular momentum."," In addition to the boundary issue, particle discreteness and resolution is also problematic when measuring halo properties such as spin and angular momentum."
" As seen in Table 2,, the dispersion decreases with increasing number of particles across all quantities."," As seen in Table \ref{Comp_stats}, the dispersion decreases with increasing number of particles across all quantities."
 This implies that a large source of the scatter is due to particle resolution., This implies that a large source of the scatter is due to particle resolution.
 The ensemble properties of the haloes’ spin parameter reaches convergence at the medium resolution., The ensemble properties of the haloes' spin parameter reaches convergence at the medium resolution.
 Figure 4 shows the log-normal fit to the spin distribution for the three runs., Figure \ref{CT_ensemble} shows the log-normal fit to the spin distribution for the three runs.
" The three runs have a mean λ of 0.039 (LoRes), 0.040 (MedRes), and 0.038 (HiRes)."," The three runs have a mean $\lambda$ of $0.039$ (LoRes), $0.040$ (MedRes), and $0.038$ (HiRes)."
" In the further analysis reported below, we use the results from the WM5 run, which has the same mass resolution as the MedRes run."," In the further analysis reported below, we use the results from the WM5 run, which has the same mass resolution as the MedRes run."
" In our previous work (Davis&Natarajan2009),, we found that at high redshift, higher spin haloes are more clustered than lower spin haloes at a given mass, using the halo-halo correlation function as a proxy for the environment."," In our previous work \citep{Davis09}, we found that at high redshift, higher spin haloes are more clustered than lower spin haloes at a given mass, using the halo-halo correlation function as a proxy for the environment."
" As in our previous work, we again define distinct spin bins from the lognormal distribution, each with roughly one third of the haloes, and denote them as high, median, and low spin bins."," As in our previous work, we again define distinct spin bins from the lognormal distribution, each with roughly one third of the haloes, and denote them as high, median, and low spin bins."
" Using several different ways to characterize the environment, we report our findings of correlations between measures of the environment with the spin parameter for haloes in the mass range 10°Mo to 10°Mo."," Using several different ways to characterize the environment, we report our findings of correlations between measures of the environment with the spin parameter for haloes in the mass range $10^6 \Msun$ to $10^8 \Msun$."
 We first use the distance to close neighbors as a measure of environment., We first use the distance to close neighbors as a measure of environment.
 This distance is calculated as the separation between each halo’s centre (i.e. the densest particle)., This distance is calculated as the separation between each halo's centre (i.e. the densest particle).
" We measure both the distance to the 3' (Ds) and τί"" (Dz) nearest neighbors.", We measure both the distance to the $3^{rd}$ $D_3$ ) and $7^{th}$ $D_7$ ) nearest neighbors.
 Both distances show the same result: haloes with closer neighbors have a small tendency to have larger spins., Both distances show the same result: haloes with closer neighbors have a small tendency to have larger spins.
 The second method that we use to measure environment is the over-density within ten times the virial radius., The second method that we use to measure environment is the over-density within ten times the virial radius.
 We use as a proxy for the virial radius Ri7zg: the radius of a spherical region which has an average density of 178pcrit., We use as a proxy for the virial radius $R_{178}$: the radius of a spherical region which has an average density of $178 \rho_{crit}$ .
 The factor 178 is the overdensity criterion A(z) evaluated at z=6., The factor $178$ is the overdensity criterion $\Delta(z)$ evaluated at $z=6$.
" Bryan&Norman(1998) give a fitting formula in a LCDM universe for A(z), but at high redshifts, A(z) is approximately the same as it would be in an Einstein-de Sitter universe, such that A(z)=18x?e178."," \citet{Bryan98} give a fitting formula in a LCDM universe for $\Delta(z)$, but at high redshifts, $\Delta(z)$ is approximately the same as it would be in an Einstein-de Sitter universe, such that $\Delta(z) = 18\pi^2 \approx 178$."
" To calculate this radius, we use all particles surrounding the centre of the potential for the given halo, regardless of whether or not the particles are part of the halo."," To calculate this radius, we use all particles surrounding the centre of the potential for the given halo, regardless of whether or not the particles are part of the halo."
 This is to account for the fact that not all haloes extended out far enough to reach the required density threshold., This is to account for the fact that not all haloes extended out far enough to reach the required density threshold.
" In other words, some haloes have Ri7g outside the most distant particle assigned to the group."," In other words, some haloes have $R_{178}$ outside the most distant particle assigned to the group."
" We then calculate the over-density within 10R;7s, and findthat haloes in over-dense regions have a slight tendency to have higher spins."," We then calculate the over-density within $10 R_{178}$, and findthat haloes in over-dense regions have a slight tendency to have higher spins."
" Figure 5 shows the trends with spin for both measures of environment in two mass bins: 1095?Mo and 10759?Mo, and histograms for the lowest mass bin after separating our sample by spin."," Figure \ref{environ} shows the trends with spin for both measures of environment in two mass bins: $10^{6\pm0.2} \Msun$ and $10^{7\pm0.2} \Msun$, and histograms for the lowest mass bin after separating our sample by spin."
 The histograms for the 10°Mo and 108Mo bins show qualitatively the same trends., The histograms for the $10^7 \Msun$ and $10^8 \Msun$ bins show qualitatively the same trends.
" We find no strong correlation between environment and spin — it is possible to find low spin haloes in very dense environments, and high spin haloes in sparse environments."," We find no strong correlation between environment and spin – it is possible to find low spin haloes in very dense environments, and high spin haloes in sparse environments."
" However, we do find a small excess probability that haloes with lower spins are in less dense environments, as evidenced by the histograms in Figure 5.."," However, we do find a small excess probability that haloes with lower spins are in less dense environments, as evidenced by the histograms in Figure \ref{environ}."
" This is in contrast to Avila-etal. (2005),, who report that galaxy scale haloes (masses «5x1011 Mo/h) in clusters have lower spin than isolated haloes in the field."," This is in contrast to \citet{Avila05}, who report that galaxy scale haloes (masses $< 5 \times 10^{11} \Msunh$ ) in clusters have lower spin than isolated haloes in the field."
" Interestingly, they find that the trend they report also disappears as they go back to z=1."," Interestingly, they find that the trend they report also disappears as they go back to $z=1$."
" This may imply a trend with redshift where at high redshift, high spin haloes are in denser environments, and at very low redshift, high spin haloes are in less dense environments."," This may imply a trend with redshift where at high redshift, high spin haloes are in denser environments, and at very low redshift, high spin haloes are in less dense environments."
" However, a strict comparison of these two numerical studies cannot be made as we have significantly more haloes, probe significantly higher redshift slices and significantly lower halo masses."," However, a strict comparison of these two numerical studies cannot be made as we have significantly more haloes, probe significantly higher redshift slices and significantly lower halo masses."
" Avila-Reeseetal.(2005) only use the over-density criterion to quantify the environment, whereas our findings hold across three different measures of environment."," \citet{Avila05} only use the over-density criterion to quantify the environment, whereas our findings hold across three different measures of environment."
" We note that the trends reported here have shallow slopes, similar to the Maccióetal.(2007) work."," We note that the trends reported here have shallow slopes, similar to the \citet{Maccio07} work."
 We conclude that the dependence of spin on environment is weak for the halo masses and epochs studied here., We conclude that the dependence of spin on environment is weak for the halo masses and epochs studied here.
" In determining the effect of halo spin on halo structural properties, we first turn to the mass distribution within the halo."," In determining the effect of halo spin on halo structural properties, we first turn to the mass distribution within the halo."
" Wefit the NFW profile, (Navarro,Frenk,&White1996, 1997),, to each halo through a least squares fit to the radially averaged density profile calculated using all particles assigned to the halo by"," Wefit the NFW profile, \citep{NFW96, NFW97}, , to each halo through a least squares fit to the radially averaged density profile calculated using all particles assigned to the halo by"
were all applied.,were all applied.
 The data were then filtered for bad event grades. good time intervals and flaring.," The data were then filtered for bad event grades, good time intervals and flaring."
 The final usable exposures Were 9.6 ks and 37.5 ks for the 2000 and 2005 data. respectively.," The final usable exposures were 9.6 ks and 37.5 ks for the 2000 and 2005 data, respectively."
 An image from the 2005 data in the energy range 0.3-10 keV is shown in Figure 4.4((a)., An image from the 2005 data in the energy range 0.3--10 keV is shown in Figure \ref{fig_main}( (a).
 A smoothed version. with radio contours from O2001 overlaid. is presented in Figure 4.4((b).," A smoothed version, with radio contours from O2001 overlaid, is presented in Figure \ref{fig_main}( (b)."
 The resulting image ts ~6 times deeper than that from the 2000 data of 02001., The resulting image is $\sim6$ times deeper than that from the 2000 data of O2001.
 For spectroscopy. we combined the two data-sets.," For spectroscopy, we combined the two data-sets."
 We used identical source extraction regions for each observation (see Fig. 4.4[[, We used identical source extraction regions for each observation (see Fig. \ref{fig_main}[ [
b]). and then created corresponding weighted response and effective area files using the CIAO script,"b]), and then created corresponding weighted response and effective area files using the CIAO script."
 were grouped so that there was a minimum of 25 counts per channel (with the exception of the 2000-epoch spectrum of thecentral source CXOU J061705.34+222127 [see $ refsecjmages ||. whichwasgroupedtolScountsperchannelbecauseo fit," Spectra were grouped so that there was a minimum of 25 counts per channel (with the exception of the 2000-epoch spectrum ofthe central source CXOU J061705.3+222127 [see \\ref{sec_images}] ], which was grouped to 15 counts per channel because of the low number of counts)."
destavtneabeiesi n.," For background spectra, we used annular regions immediately surrounding each of the regions shown in Figure \ref{fig_main}( (b)."
 Tiereermi, Spectra were then modeled using XSPEC v11.3.1.
lpsaicu atebsgtecesdrnmeeia udur regronsst," In all spectral fitting, we jointly fit the two observations from 2000 and 2005, with all model parameters locked between the two data-sets."
 We identify à number of notable features from the X-ray and radio images shown in Figure 4.4:: We have extracted spectra for the three regions shown in Figure 4.4((b). corresponding to J0617-222]1.. the tongue. and the tail.," We identify a number of notable features from the X-ray and radio images shown in Figure \ref{fig_main}: : We have extracted spectra for the three regions shown in Figure \ref{fig_main}( (b), corresponding to , the tongue, and the tail."
 To accurately estimate the foreground absorbing column. Ny. we fit simultaneous absorbed power laws to the large number of counts obtained for the tongue and the tail. constrained to have the same Nj. but with independent photon indices.," To accurately estimate the foreground absorbing column, $N_H$, we fit simultaneous absorbed power laws to the large number of counts obtained for the tongue and the tail, constrained to have the same $N_H$, but with independent photon indices."
 The results are given in Table 1.. showing that both regions are good fits to power law spectra. with a foreground column Ny2(7.20.6)«107! em™ (this and all subsequer= errors are at confidence).," The results are given in Table \ref{tab_spec}, , showing that both regions are good fits to power law spectra, with a foreground column $N_H =
(7.2\pm0.6)\times10^{21}$ $^{-2}$ (this and all subsequent errors are at confidence)."
 The photon indices for the two regions are the same within the errors., The photon indices for the two regions are the same within the errors.
 We then fit to the spectrum ofJO617+2221.. excluding ond el surrounding it by using a circular extraction region of radius 173.," We then fit to the spectrum of, excluding most of the emission from the compact structures immediately surrounding it by using a circular extraction region of radius $1\farcs4$."
 If we constrain Nj; to lie in the range inferred above. we find that both a power law and a blackbody are poor fits to the data. both models not accounting for a hard excess seen above 2 keV. A combined fit of a blackbody and a power law resolves this problem and produces an excellent fit. as listed in Table | and shown in Figure 4..," If we constrain $N_H$ to lie in the range inferred above, we find that both a power law and a blackbody are poor fits to the data, both models not accounting for a hard excess seen above 2 keV. A combined fit of a blackbody and a power law resolves this problem and produces an excellent fit, as listed in Table \ref{tab_spec} and shown in Figure \ref{fig_spec}."
 We have also tried to fit spectra to the compact knots seen on either side of the pulsar., We have also tried to fit spectra to the compact knots seen on either side of the pulsar.
 However. since each knot only contains 40 counts in the combined data-set. we are unable to derive useful spectral parameters for these regions.," However, since each knot only contains $\sim40$ counts in the combined data-set, we are unable to derive useful spectral parameters for these regions."
 Simple hardness ratios indicate that both these knots are significantly harder thanJ0617-2221.. suggesting that their emission is probably non-thermal.," Simple hardness ratios indicate that both these knots are significantly harder than, suggesting that their emission is probably non-thermal."
 Our new data clearly confirm the main result of O2001: cconsists of a compact source at the apex of a striking cometary nebula. which presumably represent an energetic pulsar and its wind-driven bow shock. respectively.," Our new data clearly confirm the main result of O2001: consists of a compact source at the apex of a striking cometary nebula, which presumably represent an energetic pulsar and its wind-driven bow shock, respectively."
 Here we focus on the additional information provided by the higher sensitivity and improved angular resolution of our new data., Here we focus on the additional information provided by the higher sensitivity and improved angular resolution of our new data.
 We first consider the X-ray emission from JO617+2221., We first consider the X-ray emission from .
. 02001 assumed that the X-rays from wwere thermal., O2001 assumed that the X-rays from were thermal.
" They adopted a low absorbing column Ny=1.3«107! em. and consequently found a very high temperature (~0.7 keV) and a small emitting area (270.025d,4 km?)"," They adopted a low absorbing column $N_H = 1.3\times10^{21}$ $^{-2}$, and consequently found a very high temperature $\approx 0.7$ keV) and a small emitting area $\approx0.025d_{1.5}$ $^2$ )."
 With better statistics but limited spatial resolution. Boechino&Bykov(2001) were unable to directly detect X-rays fromJO617+2221.. but found that a blackbody with a temperature 0.13 keV and radius 3.3 km was consistent with their data.," With better statistics but limited spatial resolution, \cite{bb01b} were unable to directly detect X-rays from, but found that a blackbody with a temperature 0.13 keV and radius 3.3 km was consistent with their data."
 Here we directly detect thermal emission fromJ06174-222].. confirming tts identification as à neutron star.," Here we directly detect thermal emission from, confirming its identification as a neutron star."
 The blackbody component of the fit to sshown in Figure 4. implies an emitting radius (as viewed at infinity) of s km., The blackbody component of the fit to shown in Figure \ref{fig_spec} implies an emitting radius (as viewed at infinity) of $6^{+8}_{-3} d_{1.5}$ km.
" This is not especially constraining. being consistent 63d,both with cooling emission fromthe whole surface. and also withthe smaller radiseen for blackbody fits to neutronstar atmospheres (e.g..Pavlovetal.2001:Mc- 2004).."," This is not especially constraining, being consistent both with cooling emission fromthe whole surface, and also withthe smaller radiiseen for blackbody fits to neutronstar atmospheres \citep[e.g.,][]{pzs+01,mzc+04}. ."
 The inferred temperature and bolometric luminosity.as viewed at infinity. are AT“=1023:22 eV and L=5.035di.«10°? eres sl. respectively.," The inferred temperature and bolometric luminosity,as viewed at infinity, are $kT^{\infty}
= 102\pm22$ eV and $L^{\infty} = 5.0^{+5.0}_{-2.2}~d_{1.5}^2 \times10^{32}$ ergs $^{-1}$ , respectively."
 The age, The age
lis a neutron star Low Mass. X-ray Binary (LAINB) discovered. in 1996. hy /Wicle Fickle Camera (WEC) during a single tyvpe-bANY X-ray burs (intZandetal.1999).,is a neutron star Low Mass X-ray Binary (LMXB) discovered in 1996 by /Wide Field Camera (WFC) during a single type-I X-ray burst \cite{zand99}.
.. Llowever. no steady emission was detected from. the source. leading to an upper limit of about 5 mCrab (28 keV) for a total exposure of 300. ks (intZandetal.1999).," However, no steady emission was detected from the source leading to an upper limit of about 5 mCrab (2–8 keV) for a total exposure of 300 ks \cite{zand99}."
.. Cornelisse et al. (, Cornelisse et al. (
2004) propose bbeing member of a possible non-homogeneous class of LAINBs. the so-called *hurst-only” sources (see also Cocchi et al.,"2004) proposed being member of a possible non-homogeneous class of LMXBs, the so-called “burst-only” sources (see also Cocchi et al."
 2001)., 2001).
 These are a group of nine bursters discoverec by AN//WEC when exhibiting à. tvpe-L burst without any detectable persistent N-ray. emission., These are a group of nine bursters discovered by /WFC when exhibiting a type-I burst without any detectable persistent X-ray emission.
 Reconth. identified. two new members of this class.," Recently, identified two new members of this class."
 In. fact. photospherie radius expansion. (PRE) bursts have been caught in two previously unclassified sources. namely citebrandt06. and citechelo07..," In fact, photospheric radius expansion (PRE) bursts have been caught in two previously unclassified sources, namely \\cite{brandt06} and \\cite{chelo07}."
 Afterwards. both have been classified às 7quasi- Very Faint X-ray Transients (VENT). since they uncereo prolonged. accretion episodes of many vears at. low ((Del Santo et al.," Afterwards, both have been classified as ``quasi-persistent'' Very Faint X-ray Transients (VFXTs), since they undergo prolonged accretion episodes of many years at low (Del Santo et al."
 2007. Bassa et al.," 2007, Bassa et al."
 2008)., 2008).
 VENIS are transients showing outbursts with low peak luminosity (1075. 107 iin 210 keV). mainly discovered. with high. sensitivity instruments on-board aand cleuring survevs of the Galactic Center region," VFXTs are transients showing outbursts with low peak luminosity $10^{34}$ $^{36}$ in 2–10 keV), mainly discovered with high sensitivity instruments on-board and during surveys of the Galactic Center region"
is responsible for ~30% of the UV flux within (7111.,is responsible for $\sim$ of the UV flux within 14.
 The contribution from the uuderlving exteuded population is subsequently reduced., The contribution from the underlying extended population is subsequently reduced.
 The uucleus of NGC 205. at a comparable distance of T20kpc. is dn many ways simular to that of AD). but without the N-vayv source.," The nucleus of NGC 205, at a comparable distance of 720kpc, is in many ways similar to that of M33, but without the X-ray source."
 It is globular cluster like. has a comparable AZ and a low upper lunit of 9-10tAL.. on the mass within the contra pc (HeathJonesctal. 1996)).," It is globular cluster like, has a comparable $M_V$ and a low upper limit of $\cdot$ $^{4}$ on the mass within the central pc \cite{heath}) )."
" These authors fux that the uucleus is inore extended. with a. E555W FWIDI of Q(/22, Επ]. P and only Frigypyw=6.0-10 19 (within 072273)."," These authors find that the nucleus is more extended, with a F555W FWHM of 2, $F_{\rm F555W}$ $\cdot$ $^{-15}$ and only $F_{\rm F160BW}$ $\cdot$ $^{-16}$ (within 273)."
 Using the same aperture on the ND33 data we find for ΑΟ Frsssy=3 Ll 1 and Frioypw=2-- 10PoressFIGUD, Using the same aperture on the M33 data we find for M33 $F_{\rm F555W}$ $\cdot$ $^{-15}$ and $F_{\rm F160BW}$ $\cdot$ $^{-15}$.
"WExchiding the contribution from the poiut source. the flux of M33a0 is1 about 1.510.1?5 aud the ratios: of the fluxes between the bands become comparable,"," Excluding the contribution from the point source, the F160BW flux of M33 is about $\cdot$ $^{-15}$ and the ratios of the fluxes between the bands become comparable."
 If there is in a pot source at the center of he nucleus of M32.with factthe magnitude estimates in Table 2 are consistent a Ravleigh-Jeaus tail.," If there is in fact a point source at the center of the nucleus of M33, the magnitude estimates in Table 2 are consistent with a Rayleigh-Jeans tail."
 We estimate he total UV fux of the source wouldbe ~1.0-L0 i552 which would suggest Lor Lx ~0.05 whichix reasonable for the type of N-rav source we have xostulated Χ-δ tobe.," We estimate the total UV flux of the source would be $\sim$ $\cdot$ $^{-15}$, which would suggest $L_{\rm opt}$ $L_X$$\sim$ 0.05 which is reasonable for the type of X-ray source we have postulated X-8 to be."
 This source would be responsible or niost of the colour eradicut inCV-D audU-D aud lus the very compact appearance of the nucleus at these waveleneths., This source would be responsible for most of the colour gradient in and and thus the very compact appearance of the nucleus at these wavelengths.
 As it contributes ouly —1854 of the total FIGOBW fux and ~8% of the E300W fux. the spectral energv distribution of the nucleus is not changed much aud the consequences oin population svuthesis studies should be minor.," As it contributes only $\sim$ of the total F160BW flux and $\sim$ of the F300W flux, the spectral energy distribution of the nucleus is not changed much and the consequences on population synthesis studies should be minor."
 However. it does have a iiajor influence in that it relaxes the coustraiuts on e.g. nass scerceation to explain the very strong colour eradieuts inCTY-D andC," However, it does have a major influence in that it relaxes the constraints on e.g. mass segregation to explain the very strong colour gradients in and."
- The star located about 17 from the nucleus in M33D. has I—19.15. V —19.25. {πουν 15.05 aud Mgprioopw 1r.GU.," The star located about 1"" from the nucleus in M33 has $B$ =19.45, $V$ =19.25, $M_{\rm F300W}$ =18.05 and $M_{\rm F160BW}$ =17.60."
 The count rates ave compatible with an AO type spectrum which would male it similar. although fainter. to the two A superegiauts detected by Masseyetal.1996.," The count rates are compatible with an A0 type spectrum which would make it similar, although fainter, to the two A supergiants detected by \cite{massey}."
. The fluxes in the F300W and F160DW bands. at the distance of NES (SUO kpe). are ~ 1077 eres. 1.," The fluxes in the F300W and F160BW bands, at the distance of M33 $\sim$ 800 kpc), are $\sim$ $10^{37}$ $\cdot$ $^{-1}$."
 The positional accuracy of theIIRI does not rule out this star as a possible counterpart to the N-ray source N-8., The positional accuracy of the does not rule out this star as a possible counterpart to the X-ray source X-8.
 The ratio Ly Lo~ 100 aud the observed (Band BoV would agree with what is expected from a low-inass A-ray binary (vanParacdijs&MeCliutock19051)., The ratio $L_X$ $L_{\rm opt} \sim$ 100 and the observed $U-B$ and $B-V$ would agree with what is expected from a low-mass X-ray binary \cite{vanp}) ).
" But its absolute magnitude (AA,275.2) would make it more similar to a high mass X- binary.", But its absolute magnitude $M_V\approx$ -5.2) would make it more similar to a high mass X-ray binary.
 The main argument against this star beme the N-rav binary is N-8s unique character and special location at the nucleus., The main argument against this star being the X-ray binary is X-8's unique character and special location at the nucleus.
 Caven that etal.1996. found ~300 analogous UV sources in M32.theMassey it would be remarkable that the one near the nucleus is most Luminous X-ray source in the Local Group.," Given that \cite{massey} found $\sim$ 300 analogous UV sources in M33, it would be remarkable that the one near the nucleus is the most luminous X-ray source in the Local Group."
 The UV high resolution inages obtained with the IISTCamera show the nucleus of M33 is extremely compact., The UV high resolution images obtained with the HST show the nucleus of M33 is extremely compact.
 We have fitted convolved models to the radial profiles iu the different bands from which we find the FWIIM of the nucleus in UV to be —070035 and ο 0090 in V. About half of the UW flux comes from the inner Q'111 of the nucleus., We have fitted convolved models to the radial profiles in the different bands from which we find the FWHM of the nucleus in UV to be $\sim$ 035 and $\sim$ 090 in V. About half of the UV flux comes from the inner 14 of the nucleus.
 The UV and U profiles are also well fitted if one assumes a blue point source superposed on au extended population with the same EWITIM as in V. If this is the correct iiodoel for the nucleus. then this poiut source is likely to be the UV counterpart to the very Iuminous X-rav source X-8.," The UV and U profiles are also well fitted if one assumes a blue point source superposed on an extended population with the same FWHM as in V. If this is the correct model for the nucleus, then this point source is likely to be the UV counterpart to the very luminous X-ray source X-8."
 Such a counterpart would be responsible for most of the strong colour gradient secu in UV., Such a counterpart would be responsible for most of the strong colour gradient seen in UV.
 Its contribution to the total UV. dux of the uncleus would be aboutIs%., Its contribution to the total UV flux of the nucleus would be about.
. Models for the structure of the nucleus still need to account for the remainder of the UV fis but the constraints ou population segregation (minore compact blue star population) are reduced., Models for the structure of the nucleus still need to account for the remainder of the UV flux but the constraints on population segregation (more compact blue star population) are reduced.
 Wiel spatial resolution UV spectroscopy of the nucleus is the obvious next step. which we will undertake shortly.," High spatial resolution UV spectroscopy of the nucleus is the obvious next step, which we will undertake shortly."
 We thauk Sylvia Bagectt for help with the F160DW. filter reduction., We thank Sylvia Baggett for help with the F160BW filter reduction.
 G.D. and K.S.L wish to thauk the Oxford Astrophysics Departinent where part of this work was completed., G.D. and K.S.L wish to thank the Oxford Astrophysics Department where part of this work was completed.
 We acknowledge support by the British-Freuch joint research programAlMance aud by NASA eraut NAC 5-1539 to the STScI. Based ou observations with the NASA/ESA Hubble Space Telescope obtained at the STScI which is operated by the AURA under NASA contract No., We acknowledge support by the British-French joint research program and by NASA grant NAG 5-1539 to the STScI. Based on observations with the NASA/ESA Hubble Space Telescope obtained at the STScI which is operated by the AURA under NASA contract No.
 NÀS5-26555., NAS5-26555.
,.
2006).. Biggs&Ivison(2008) reported physical sizes of these SMGs ranging from 1—8 kpc with a median of 5 kpc., \citet{BiggsIvison08} reported physical sizes of these SMGs ranging from $1-8$ kpc with a median of 5 kpc.
" Lastly, Caseyetal.(2009) used MERLIN (0""3—0""5 resolution) to observe seven ULIRGs and a HyLIRG selected by optical and radio color criteria from GOODS and the Lockman Hole at redshifts 0.9<z2.4."," Lastly, \citet{Casey09} used MERLIN $0\farcs3-0\farcs5$ resolution) to observe seven ULIRGs and a HyLIRG selected by optical and radio color criteria from GOODS and the Lockman Hole at redshifts $0.9 < z <
2.4$."
" They resolved all eight and report an average physical diameter of 5.0+1.2 kpc, corresponding to the surface area within the 3c detection level."," They resolved all eight and report an average physical diameter of $5.0 \pm 1.2$ kpc, corresponding to the surface area within the $3\sigma$ detection level."
" Muxlowetal.(2005) used MERLIN+VLA to conduct a deep 1.4 GHz survey of the HDF and Hubble Flanking Field (HFF) and studied 92 radio sources at 0.18<z4 with 0/2—0/""5 resolution.", \citet{Muxlow05} used MERLIN+VLA to conduct a deep 1.4 GHz survey of the HDF and Hubble Flanking Field (HFF) and studied 92 radio sources at $0.18 < z < 4$ with $0\farcs2-0\farcs5$ resolution.
 All but one source was resolved., All but one source was resolved.
" For this study, we exclude sources that Muxlowetal.(2005) classified as AGN or AGN candidates."," For this study, we exclude sources that \citet{Muxlow05} classified as AGN or AGN candidates."
" We also exclude the Muxlow sources that are observed only by the VLA and hence have lower resolution that may affect the accuracy of size estimates, as well as those with complex morphology such that their size estimation requires visual inspection (Muxlowetal.2005;etal. 2007)."," We also exclude the Muxlow sources that are observed only by the VLA and hence have lower resolution that may affect the accuracy of size estimates, as well as those with complex morphology such that their size estimation requires visual inspection \citep{Muxlow05, Thrall07}."
". Out of 92 objects from the Muxlowet sample, 72 have spectroscopic redshifts and 27 pass (2005)the aforementioned criterion and have a 24 jum counterpart in GOODS."," Out of 92 objects from the \citet{Muxlow05} sample, 72 have spectroscopic redshifts and 27 pass the aforementioned criterion and have a 24 $\micron$ counterpart in GOODS."
 Muxlowetal.(2005) reported sizes in terms of the largest angular size determined by Gaussian fitting (Thralletal.2007)., \citet{Muxlow05} reported sizes in terms of the largest angular size determined by Gaussian fitting \citep{Thrall07}.
". Taking the largest angular size, which is the size of the major axis, as the diameter directly would overestimate the surface area; we assume that the sizes of the minor axes of these galaxies are 0.6 of the major-axis largest angular FWHM) and then calculate the surface(the areas and circularized diameters for the Muxlowetal.(2005) sample as reported in Table A2.."," Taking the largest angular size, which is the size of the major axis, as the diameter directly would overestimate the surface area; we assume that the sizes of the minor axes of these galaxies are 0.6 of the major-axis (the largest angular FWHM) and then calculate the surface areas and circularized diameters for the \citet{Muxlow05} sample as reported in Table \ref{table_highz}. ."
 This value is an average ratio of minor-to-major axes of SMGs from the Biggs&Ivison(2008) sample., This value is an average ratio of minor-to-major axes of SMGs from the \citet{BiggsIvison08} sample.
 An average physical diameter for the Muxlowetal.(2005) sample is 4.5 kpc., An average physical diameter for the \citet{Muxlow05} sample is $4.5$ kpc.
" ffor the Chapmanetal.(2004), the Muxlowetal.(2005), and the Caseyetal.(2009) samples in the GOODS field were estimated via the Riekeetal.(2009) formalism and theSpitzer 24 um imaging in GOODS."," for the \citet{Chapman04}, , the \citet{Muxlow05}, and the \citet{Casey09}
 samples in the GOODS field were estimated via the \citet{Rieke09}
 formalism and the 24 $\micron$ imaging in GOODS."
 ffor the Biggs&Ivison(2008) and the Caseyetal.(2009) samples in other fields are also estimated using the Riekeetal.(2009) formalism based on the 24 jum fluxes reported by Ivisonetal.(2007) and Caseyetal.(2009).," for the \citet{BiggsIvison08} and the \citet{Casey09}
 samples in other fields are also estimated using the \citet{Rieke09}
 formalism based on the 24 $\micron$ fluxes reported by \citet{Ivison07} and \citet{Casey09}."
. Apparently there are two approaches in reporting galaxy sizes from the radio measurements: the deconvolved FWHM etal.2005;Biggs&Ivi-son2008) and the circularized(Muxlow diameter enclosing the surface area of the >30 radio detections (Chapmanetal.2004;Caseyet 2009).," Apparently there are two approaches in reporting galaxy sizes from the radio measurements: the deconvolved FWHM \citep{Muxlow05,
 BiggsIvison08} and the circularized diameter enclosing the surface area of the $>3 \sigma$ radio detections \citep{Chapman04,Casey09}."
". Despite the methodological differences between these two approaches, they agree very well for the objects with overlapping observations, especially in the HDF where there are six such galaxies."," Despite the methodological differences between these two approaches, they agree very well for the objects with overlapping observations, especially in the HDF where there are six such galaxies."
" This agreement is expected, since the area detected above 3c is likely to be consistent with that of half-max power given the typical levels of signal to noise ratio in the observations."," This agreement is expected, since the area detected above $3 \sigma$ is likely to be consistent with that of half-max power given the typical levels of signal to noise ratio in the observations."
" We also note a very good agreement of sizes from 1.4 GHz and CO (3—2) observations of a HDF object, SMMJ123707--6214SW, where observations from Muxlowetal.(2005) and Tacconietal.(2006) overlap."," We also note a very good agreement of sizes from 1.4 GHz and CO $3-2$ ) observations of a HDF object, SMMJ123707+6214SW, where observations from \citet{Muxlow05} and \citet{Tacconi06} overlap."
" Although galaxies drawn from the aforementioned samples are likely dominated by star-forming activity, we independently confirm this by inspecting whether their ratio of 850 wm and 1.4 GHz fluxes follows the IR-radio relation for starburst galaxies given by Riekeetal.(2009)."," Although galaxies drawn from the aforementioned samples are likely dominated by star-forming activity, we independently confirm this by inspecting whether their ratio of 850 $\micron$ and 1.4 GHz fluxes follows the IR-radio relation for starburst galaxies given by \citet{Rieke09}."
 Galaxies with radio flux significantly stronger than predicted by the IR-radio relation likely harbor radio-loud AGN., Galaxies with radio flux significantly stronger than predicted by the IR-radio relation likely harbor radio-loud AGN.
" Indeed, we have found that the only two objects with Faso/Fi4>>2 are the most compact objects in our compilation (consistent with AGN-domination) and thus we exclude these objects (SMMJ163650.43+405734.5 and SMMJ105207.49+571904.0)."," Indeed, we have found that the only two objects with $_{850}/$ $_{1.4} \gg 2$ are the most compact objects in our compilation (consistent with AGN-domination) and thus we exclude these objects (SMMJ163650.43+405734.5 and SMMJ105207.49+571904.0)."
" It is worth noting that this test is unnecessary for the Muxlowetal. sample, where both radio spectral slope and (2005)radio morphology are already employed to identify AGN and AGN candidates."," It is worth noting that this test is unnecessary for the \citet{Muxlow05} sample, where both radio spectral slope and radio morphology are already employed to identify AGN and AGN candidates."
 In total our compilation has 48 intermediate and high-redshift starburst galaxies with a median redshift of 1.0 and a median size of 5.1 kpc., In total our compilation has 48 intermediate and high-redshift starburst galaxies with a median redshift of 1.0 and a median size of 5.1 kpc.
" Our primary sample comprises the 27 galaxies from Muxlowetal. (2005), which should provide an unbiased and complete sample, selected by radio flux alone."," Our primary sample comprises the 27 galaxies from \citet{Muxlow05}, which should provide an unbiased and complete sample, selected by radio flux alone."
 The results from the additional 21 galaxies are consistent with those from the Muxlowetal.(2005) observations.," The results from the additional 21 galaxies are consistent with those from the \citet{Muxlow05}
 observations."
" For local galaxies, there is no homogeneous set of radio or submillimeter images suitable forour needs."," For local galaxies, there is no homogeneous set of radio or submillimeter images suitable forour needs."
" Therefore, we use two additional measures of the star formation rate, Paschen-a images (Alonso-Herreroet and images at 24 jum from anumber of sources."," Therefore, we use two additional measures of the star formation rate, $\alpha$ images \citep{AH06} and images at 24 $\micron$ from anumber of sources."
" As with the radio, both measures are not strongly affected"," As with the radio, both measures are not strongly affected"
oobservatious of2853.,observations of.
.. deletected the source in outburst. which allowed us to deteriuue its location precisely (Munoctal.2003).," detected the source in outburst, which allowed us to determine its location precisely \citep{mun03}."
. Using this location. we study the flux from this source dn quiescence.," Using this location, we study the flux from this source in quiescence."
 We also report the detection of a thermonuclear X-rav burst while the source was in outburst., We also report the detection of a thermonuclear X-ray burst while the source was in outburst.
 To date. the location of hhas been observed ou five occasions withChandra: three times as part of a program to monitor the superauassive black hole at the ceuter of the Galaxy. citepbag0l.bag03. nmnu03.. and twice as part of a shallow survey of the iuner 300 pe of the Galaxy (Wane 2002).," To date, the location of has been observed on five occasions with: three times as part of a program to monitor the super-massive black hole at the center of the Galaxy, \\citep{bag01, bag03, mun03}, and twice as part of a shallow survey of the inner 300 pc of the Galaxy \citep{wang02}."
. The observations are listed in Table 2.., The observations are listed in Table \ref{tab:obs}.
 All five observations were taken using the Advanced CCD Duagiug Spectrometer imaging array CACIS-D., All five observations were taken using the Advanced CCD Imaging Spectrometer imaging array (ACIS-I).
 The ACTS-T is a set of four 102L-by-1021 pixel CCDs. covering a field of view of bbv17.," The ACIS-I is a set of four 1024-by-1024 pixel CCDs, covering a field of view of by."
" When placed on-axis at the focal plane of the erazine-Incidence X-rav τος, the imaging resolution is determined by the pixel size of the CCDs. 071192."," When placed on-axis at the focal plane of the grazing-incidence X-ray mirrors, the imaging resolution is determined by the pixel size of the CCDs, 492."
 The CCDs measure the cnerey of incident photons with a resolution of 50-300 eV (depending ou photon euergv and distance from the read-out node) within a calibrated ΟΠΟΙΟΥ baud of 0.58 keV. We reduced the data starting with the level 1 eveut files provided by the ANaav Counter (CNC)., The CCDs measure the energy of incident photons with a resolution of 50-300 eV (depending on photon energy and distance from the read-out node) within a calibrated energy band of 0.5–8 keV. We reduced the data starting with the level 1 event files provided by the X-ray Center (CXC).
 We first removed the pixel randoniüzation applied by the default processing software. and corrected the pulse heights of each eveut for the position-dependent cliarge-trausfer ineffcieneyv. caused. by radiation damage carly in the mission (Townsley 2000).," We first removed the pixel randomization applied by the default processing software, and corrected the pulse heights of each event for the position-dependent charge-transfer inefficiency caused by radiation damage early in the mission \citep{tow00}."
. We excluded most events flageed as possible backeround. but left iu possible cosmüc ray after- because they are difficult to distinguish from the strong diffuse cinission and nunerous poiut sources near the Calactie center.," We excluded most events flagged as possible non-X-ray background, but left in possible cosmic ray after-glows because they are difficult to distinguish from the strong diffuse emission and numerous point sources near the Galactic center."
 We applied the standard eeradoe filters to the events. as well as the eood-time filters supplied by the CNC.," We applied the standard grade filters to the events, as well as the good-time filters supplied by the CXC."
 We examined cach observation for backeround flares. aud removed 10 ks of strong flaring from ObsID 0212 when the count rate rose to 3 standard deviations above the mean.," We examined each observation for background flares, and removed 10 ks of strong flaring from ObsID 0242 when the count rate rose to 3 standard deviations above the mean."
" Finally, we corrected the pointing aspect for cach observation as specified iu the ddata caveats To search for point sources iu each observation. we eniploved the CIAO tool (Freciuanctal.2002) in the manner described in Mimoctal.(2003)."," Finally, we corrected the pointing aspect for each observation as specified in the data caveats To search for point sources in each observation, we employed the CIAO tool \citep{fre02} in the manner described in \citet{mun03}."
. We set the detection threshold to 10.7. which corresponds to the probability that a backerouncl fluctuation i ouv elven pixel will be identified as a source.," We set the detection threshold to $10^{-7}$, which corresponds to the probability that a background fluctuation in any given pixel will be identified as a source."
 The resulting source list for the nunonitoring campaign. which covers ObxIDs 0212. 1561a. and 1561b. is preseuted in Munoetal.(2003).," The resulting source list for the monitoring campaign, which covers ObsIDs 0242, 1561a, and 1561b, is presented in \citet{mun03}."
. The source lists for ObsIDs 2281 and 2287 will be reported elsewhere., The source lists for ObsIDs 2284 and 2287 will be reported elsewhere.
 These lists were used to identify candidate counterparts to and to exclude poiut sources from the estimate of the backeround flux., These lists were used to identify candidate counterparts to and to exclude point sources from the estimate of the background flux.
 We verified that the astrometry for cach observation was accurate to within (0755 by cross-correlating the positious of the half-dozen sources that were both detected im the soft N-rav baud (0.51.5 keV) and present in the USNO-D wwas evident by eve as a bright transient im the image from ObsID 1561a (2000 October 26: Figure 1))., We verified that the astrometry for each observation was accurate to within 5 by cross-correlating the positions of the half-dozen sources that were both detected in the soft X-ray band (0.5–1.5 keV) and present in the USNO-B was evident by eye as a bright transient in the image from ObsID 1561a (2000 October 26; Figure \ref{fig:obsid1561a}) ).
 Since the source is located ffrom the aim-poiut. of the observation. its location is uncertain by about citep|scealso][jimn03..," Since the source is located from the aim-point of the observation, its location is uncertain by about \\citep[see also][]{mun03}."
 Its position is within the ecrror circles of previous localizations of with aandASCA., Its position is within the error circles of previous localizations of with and.
. The trausicut source exhibited a thermonuclear burst caving this oobservation (see below). so we can confidently associate iowith 2853.," The transient source exhibited a thermonuclear burst during this observation (see below), so we can confidently associate it with."
.. We searched the source lists from the other observations. aud ideutified a faiut poiut source within our initial localization of aat ffrom the aim point of ObsID 2287.," We searched the source lists from the other observations, and identified a faint point source within our initial localization of at from the aim point of ObsID 2287."
 Its position was 17h in 2.338. ((J2000). with au uucertainty of citep{compare|||braQL..," Its position was 17h 45m 2.33s, (J2000), with an uncertainty of \\citep[compare][]{bra01}."
 No point source was identified by within oof in the other three observations., No point source was identified by within of in the other three observations.
 Uaving established oa precise position for28523.. we estimated its flux— iu each of— the five observations using the routine from the Tools for X-ray Analysis (Broosetal.2002).," Having established a precise position for, we estimated its flux in each of the five observations using the routine from the Tools for X-ray Analysis \citep{bro02}."
. We extracted event lists for the source from a polvgonal region chosen to match the coutour of encircled eucrev from tle poiut-spread function (PSF) at an energy of 1.5 keV. Dackerouud event lists were extracted from circular regions centered on2853.. excludiug from the event lists counts in circles circumscribing the contours of the. PSF around the point sources detected im cach field.," We extracted event lists for the source from a polygonal region chosen to match the contour of encircled energy from the point-spread function (PSF) at an energy of 4.5 keV. Background event lists were extracted from circular regions centered on, excluding from the event lists counts in circles circumscribing the contours of the PSF around the point sources detected in each field."
The munber of total. background. and net counts from cach observation are listed in Table 3..,"The number of total, background, and net counts from each observation are listed in Table \ref{tab:counts}. ."
 We computed, We computed
2007a).,.
". Besides the tentative velocity gradient, the H;80 line width is remarkably narrow: its width to zero intensity is about 1.7 wwith a full width half maximum (FWHM) from a Gaussian fit of 1.0s7!."," Besides the tentative velocity gradient, the $_2^{18}$ O line width is remarkably narrow: its width to zero intensity is about 1.7 with a full width half maximum (FWHM) from a Gaussian fit of 1.0."
". For comparison, a envelope toward a central star with even a very small mass of 0.1 Mo would have a characteristic infall velocity of 2.5 aat 25 AU -- ie., we would expect a line width significantly larger than that observed in this case."," For comparison, a free-falling envelope toward a central star with even a very small mass of 0.1 $M_\odot$ would have a characteristic infall velocity of 2.5 at 25 AU – i.e., we would expect a line width significantly larger than that observed in this case."
" A rotationally supported edge-on disk would show a similar line width, whereas a more face-on disk, would provide a lower width in agreement with what is observed, although a disk seen entirely face-on would not produce any observable velocity gradient."," A rotationally supported edge-on disk would show a similar line width, whereas a more face-on disk, would provide a lower width in agreement with what is observed, although a disk seen entirely face-on would not produce any observable velocity gradient."
" Alternatively, an inclined disk with sub-keplerian rotation (e.g., a ""pseudo-disk"" (Galli&Shu1993;HennebelleCiardi 2009))) or a density enhancement in the inner envelope due to a magnetic shock wall slowing down the infalling material 2008))) could result in a lower line width - although such phenomena usually occur on larger scales than 25 AU."," Alternatively, an inclined disk with sub-keplerian rotation (e.g., a “pseudo-disk” \citep{galli93a,hennebelle09}) ) or a density enhancement in the inner envelope due to a magnetic shock wall slowing down the infalling material \citep{tassis05,chiang08}) ) could result in a lower line width - although such phenomena usually occur on larger scales than 25 AU."
 Watsonetal.(2007) argued based on the detected mid-IR emission that the outflow cone must be close to pole-on providing a low opacity route for the IR emission to escape toward the observer., \cite{watson07} argued based on the detected mid-IR emission that the outflow cone must be close to pole-on providing a low opacity route for the IR emission to escape toward the observer.
 This would imply a nearly face-on disk., This would imply a nearly face-on disk.
" However, the proper motions and radial velocities of HyO masers suggest an outflow directed closer to the plane of the sky (Marveletal.2008;Desmurs2009).."," However, the proper motions and radial velocities of $_2$ O masers suggest an outflow directed closer to the plane of the sky \citep{marvel08,desmurs09}."
 Clearly observations of the dynamical structure of the inner regions of IRAS4B with even higher angular resolution arerequired to distinguish between these scenarios., Clearly observations of the dynamical structure of the inner regions of IRAS4B with even higher angular resolution arerequired to distinguish between these scenarios.
" The strength of the H;80 transition suggests a significant reservoir of water vapor toward IRAS4B. We estimate the column density of H2O by adopting an excitation temperature for p-H,80 of 170 K from the modeling results of Watsonetal.(2007) and assume that the emission is optically thin and uniform over its extent.", The strength of the $_2^{18}$ O transition suggests a significant reservoir of water vapor toward IRAS4B. We estimate the column density of $_2$ O by adopting an excitation temperature for $_2^{18}$ O of 170 K from the modeling results of \cite{watson07} and assume that the emission is optically thin and uniform over its extent.
" With these assumptions we estimate a column density for p-H,80 of 4x10 cm~?, which translates into a total HyO column density of 9x10/5 cm""? assuming a !°O/!8O ratio of 560 and an ortho-para ratio for H2O of 3 (i.e., an ortho-pararatio established at high temperatures)."," With these assumptions we estimate a column density for $_2^{18}$ O of $4\times
10^{15}$ $^{-2}$, which translates into a total $_2$ O column density of $9\times 10^{18}$ $^{-2}$ assuming a $^{16}$ $^{18}$ O ratio of 560 and an ortho-para ratio for $_2$ O of 3 (i.e., an ortho-pararatio established at high temperatures)."
" For comparison the column density estimated based on the Spitzer detections by Watsonetal.(2007) is two orders of magnitude lower, 9.2x10!© cm""?, over an emitting area of x0.24"" — i.e., comparable to the deconvolved size of the Hi*o emission here."," For comparison the column density estimated based on the Spitzer detections by \cite{watson07} is two orders of magnitude lower, $\times
10^{16}$ $^{-2}$, over an emitting area of $\times$ – i.e., comparable to the deconvolved size of the $_2^{18}$ O emission here."
 Our inferred column density is almost unchanged if the excitation temperature is lowered to 100 K and increases by up to a factor 5 if the temperature is increased to 1000 K. These temperatures cover the range of conditions expected for any of the scenarios for the origin of the H2O emission discussed in refintroduction.., Our inferred column density is almost unchanged if the excitation temperature is lowered to 100 K and increases by up to a factor 5 if the temperature is increased to 1000 K. These temperatures cover the range of conditions expected for any of the scenarios for the origin of the $_2$ O emission discussed in \\ref{introduction}.
" The total mass contained in the detected HO is 6.0x107? Mo, or 0.02 Mg."," The total mass contained in the detected$_2$ O is $\times
10^{-8}$ $M_\odot$ , or 0.02 $M_{\rm Earth}$ ."
" Assuming a typical H50 abundance relative to Hy of 1074, corresponding to sublimation of the H5O rich dust ice mantles 2004),, the totalH» mass of the H2O emitting material (dust+gas) is 7.5x107> Mo or"," Assuming a typical $_2$ O abundance relative to $_2$ of $^{-4}$ , corresponding to sublimation of the $_2$ O rich dust ice mantles \citep[e.g.,][]{pontoppidan04}, , the total$_2$ mass of the $_2$ O emitting material (dust+gas) is $\times 10^{-5}$ $M_\odot$ or"
ooffers the opportunity to tackle the above mentioned questions.,offers the opportunity to tackle the above mentioned questions.
band. aud 1.1 (0.9) Jv at QO baud.,"band, and 1.4 (0.9) Jy at Q band."
 Shown in Figure 1 are the same distributions of 5 for oulv those 97 sources that are conmuuon to all three bauds., Shown in Figure \ref{FIG:XKQ_97COMMON_FLUX} are the same distributions of $\bar{S}$ for only those 97 sources that are common to all three bands.
 Hore the mean (median) values are 1.5 (0.9). 1.5 (0.9) aud 1.5 (1.0) Jv at XN. Iv and Q bands. respectively.," Here the mean (median) values are 1.7 (0.9), 1.5 (0.9) and 1.5 (1.0) Jy at X, K and Q bands, respectively."
 It is not surprising that for the case of all sources and for the 97 sources in conunon that the mean flux densities at X band are the highest. especially since the strougest (>300 indy) most compact sources iu the S/N-bhaud ICRF program were selected for the RK/O-baud observations iu order to optimize the probabilitv of detection at the higher frequencies (Paper D.," It is not surprising that for the case of all sources and for the 97 sources in common that the mean flux densities at X band are the highest, especially since the strongest $>300$ mJy) most compact sources in the S/X-band ICRF program were selected for the K/Q-band observations in order to optimize the probability of detection at the higher frequencies (Paper I)."
 This selection effect is also. verv likely. the reason that. on average. the Q-hanud mean flux densities are larger thaw those at IX baud for all sources (Figure 3)).," This selection effect is also, very likely, the reason that, on average, the Q-band mean flux densities are larger than those at K band for all sources (Figure \ref{FIG:XKQ_ALL_FLUX}) )."
 The Q-baud sources were drawn strictly from this pool of stroug. compact. ICRF sources.," The Q-band sources were drawn strictly from this pool of strong, compact, ICRF sources."
 However. as Table 1. shows. a sienificaut fraction of the IX-baud sources were observed In a sinele-session survey that occurred ou 2003 Max: 22.," However, as Table \ref{TAB:OBS} shows, a significant fraction of the K-band sources were observed in a single-session survey that occurred on 2003 May 22."
 For this survey. the N-banud fiux deusitv criterion was relaxed from a nominal value of 300 αν to 200 dy in order to include more sources.," For this survey, the X-band flux density criterion was relaxed from a nominal value of 300 mJy to 200 mJy in order to include more sources."
 For the 97 sources iu conumion to all three bands. the selection effect is reduced since αμα of the RK-baud survey sources are excluded.," For the 97 sources in common to all three bands, the selection effect is reduced since many of the K-band survey sources are excluded."
 Figure d shows the I&- and Q-band mean flux deusities to be consistent with each other to within the —1054 sessilon-to-session flux density scale variatious mentioned previously., Figure \ref{FIG:XKQ_97COMMON_FLUX} shows the K- and Q-band mean flux densities to be consistent with each other to within the $\sim$ session-to-session flux density scale variations mentioned previously.
 The extent to which the total fiux density varies over time can be characterized bv the flix density variability index (o¢/S). which is the standard deviation of the measured flux deusities divided bv the mean.," The extent to which the total flux density varies over time can be characterized by the flux density variability index $\sigma_S/\bar{S}$ ), which is the standard deviation of the measured flux densities divided by the mean."
 A miuinniu value. σοο=0.0. indicates no variation in the flux density with time.," A minimum value, $\sigma_S/\bar{S}=0.0$, indicates no variation in the flux density with time."
 Since à ΠΙΟ of the sources at I& aud Q bands were observed im iuultiple sessions. it is possible to determine osο for these sources.," Since a number of the sources at K and Q bands were observed in multiple sessions, it is possible to determine $\sigma_S/\bar{S}$ for these sources."
 This was not the case for X baud where only 25 sources were observed iu both of the two N-baud sessions., This was not the case for X band where only 25 sources were observed in both of the two X-band sessions.
 Distributions for the flux. density variability index for the 235 sources at Iv band and 82 sources at Q baud observed in two or more sessions are shown iu Figure 5.., Distributions for the flux density variability index for the 235 sources at K band and 82 sources at Q band observed in two or more sessions are shown in Figure \ref{FIG:KQ_FLUX_VAR}.
 The mean aud median values of 24/5 at IN baud are 0.18 and 0.16. respectively.," The mean and median values of $\sigma_S/\bar{S}$ at K band are 0.18 and 0.16, respectively."
 The Q-band mean and median values of the flux density variabilitv for sources observed i two or nore sessionis are 0.19 and 0.15. respectively.," The Q-band mean and median values of the flux density variability for sources observed in two or more sessions are 0.19 and 0.15, respectively."
 Civeu of the relatively few nuuber of sessions thus far observed. the values above represent lower limits to the flux density variability.," Given of the relatively few number of sessions thus far observed, the values above represent lower limits to the flux density variability."
 Increasing the cutoff for the niniuunmi umuber of sessious per source to five. reduces the umber of available I&-baud sources to 8S7 aud vields a mean value of σοςο of 0.23.," Increasing the cutoff for the minimum number of sessions per source to five, reduces the number of available K-band sources to 87 and yields a mean value of $\sigma_S/\bar{S}$ of 0.23."
 Further iucreasing the miniumuiun cutoff to LO sessions. reduces the number of sources to 18 and results i a mean flux density variability of 0.2L.," Further increasing the minimum cutoff to 10 sessions, reduces the number of sources to 18 and results in a mean flux density variability of 0.24."
 At Q baud there were 22 sources with data available from all four observing sessious., At Q band there were 22 sources with data available from all four observing sessions.
 The mea flux density variability for these 22 sources was 0.28., The mean flux density variability for these 22 sources was 0.28.
 Our preliminary results therefore indicate that on average the flux deusity varied by at IS baud. and at Q baud.," Our preliminary results therefore indicate that on average the flux density varied by at K band, and at Q band."
 These results are limited to the ~5 vear period over which our data were recorded aud by the ~10% variations in the flux density scaling mentioned previously., These results are limited to the $\sim$ 5 year period over which our data were recorded and by the $\sim$ variations in the flux density scaling mentioned previously.
" Iu adcaition to the flux deusityv variability index. we also determined the ratio of the maxiumun to minii fiux deusitv (53,/794) for cach source observed iu more han one session."," In addition to the flux density variability index, we also determined the ratio of the maximum to minimum flux density $S_{\rm max}/S_{\rm min}$ ) for each source observed in more than one session."
" Distributious for S4,/iin at IX baud and OQ baud ave plotted in Figure 6.. with the mean aud uediau ratio for cach distribution imdicated."," Distributions for $S_{\rm max}/S_{\rm min}$ at K band and Q band are plotted in Figure \ref{FIG:KQ_SMAX_SMIN}, with the mean and median ratio for each distribution indicated."
 For the two distributions. the mean (median) values ave 1.7 (1.5) at Iv uad. and 1.2 (1.1) at Q baud.," For the two distributions, the mean (median) values are 1.7 (1.5) at K band, and 1.7 (1.4) at Q band."
 The figure shows that over he five-vear period in which the sources were nionitored. he 11axinimna to minima fux density ratio was ereater han a factor of 2 for roughly of the sources at Iv mand aud of the sources at Q band.," The figure shows that over the five-year period in which the sources were monitored, the maximum to minimum flux density ratio was greater than a factor of 2 for roughly of the sources at K band and of the sources at Q band."
 The ratio is eveater than a factor of 3 for ouly of the sources at Iv band and of the sources at C). Our flux density results cau be compared to other extragalactic source monitoring programs., The ratio is greater than a factor of 3 for only of the sources at K band and of the sources at Q. Our flux density results can be compared to other extragalactic source monitoring programs.
 The lm telescope at the AMetsahovi Radio Observatory has monitored the flux densities of active galactic nuclei at 22. 37. and BST (1 over the course of | vears CTerásrantaetal.2005.audreferencetherein)..," The 14-m telescope at the Metsahovi Radio Observatory has monitored the flux densities of active galactic nuclei at 22, 37, and 87 GHz over the course of $+$ years \citep[][and reference therein]{TERASRANTA:05}."
 Statistical analysis of this data usns several different methods (ovattactal.2008.2007) las shown that smaller flux. deusitv variations occur on shorter timescales of 12 wears. while larecr outbursts are rare occurring on average every 6 vears.," Statistical analysis of this data using several different methods \citep{HOVATTA:08, HOVATTA:07} has shown that smaller flux density variations occur on shorter timescales of 1–2 years, while larger outbursts are rare occurring on average every 6 years."
 Their sample consisted. of SO sources with a miununn flux density of l Jv aud a Παπ monitoring time of ten wears., Their sample consisted of 80 sources with a minimum flux density of 1 Jy and a minimum monitoring time of ten years.
 Although the time span of our data is shorter. the sunupliug much sparser. and the nüuimuun flux deusitv lower at 200300 1nJx. our Ἱν- aud Q-band results seen to be consistent with the Metsahovi data.," Although the time span of our data is shorter, the sampling much sparser, and the minimum flux density lower at 200–300 mJy, our K- and Q-band results seem to be consistent with the Metsahovi data."
 The variability index. results show that the variations are relatively sanall over the five-year monitoring period. while the Sax/min results indicate very few large-scale outbursts over the same time period.," The variability index results show that the variations are relatively small over the five-year monitoring period, while the $S_{\rm max}/S_{\rm min}$ results indicate very few large-scale outbursts over the same time period."
 As previously demonstrated by Charlot(1990). the contribution of intrinsic source structure fo a VLBI bandwidth svuthesis delay measurement can be significant.," As previously demonstrated by \cite{CHARLOT:90}, the contribution of intrinsic source structure to a VLBI bandwidth synthesis delay measurement can be significant."
 It depends on the exact form of the spatial brightuess distribution of the extended radio source relative to the ecometry of the VEDI baseline vector projected outo the plane of the sky., It depends on the exact form of the spatial brightness distribution of the extended radio source relative to the geometry of the VLBI baseline vector projected onto the plane of the sky.
 The overall magnitude of the source structure effect for a eiven source is nmiost casily estimated by calculating corrections to the banchwidth svuthesis delay based upon the observed source structure. for a range of (10) coordinates (a aud e are coordinates of the baseline vector projected ou) the plane of the sky aud are expressed im units of the observing wavelength).," The overall magnitude of the source structure effect for a given source is most easily estimated by calculating corrections to the bandwidth synthesis delay based upon the observed source structure, for a range of ) coordinates $u$ and $v$ are coordinates of the baseline vector projected onto the plane of the sky and are expressed in units of the observing wavelength)."
" Following such a scheme. Fev defined a source ""structure index according to the median value of the structure delay corrections. Tdi Calculated for all projected VLBI baselines that could be possibly observed with Earth-based VLBI (i.c.. for all baselines with Yu?|02 less thia the diameter of the Earth). separating the sources mto four classes as follows: Based on this definition. the structure iudex at each of the three bands CX. Ix. aud QJ) was obtained for cach source/session. thus providing au indication of the"," Following such a scheme, \cite{FC:97} defined a source “structure index” according to the median value of the structure delay corrections, $\tau_{\rm median}$, calculated for all projected VLBI baselines that could be possibly observed with Earth-based VLBI (i.e., for all baselines with $\sqrt{u^2+v^2}$ less than the diameter of the Earth), separating the sources into four classes as follows: Based on this definition, the structure index at each of the three bands (X, K, and Q) was obtained for each source/session, thus providing an indication of the"
by performing detailed GEANTA simulations of proton-photon interactions in the internal shocks (Murase&Nagataki2006)..,by performing detailed GEANT4 simulations of proton-photon interactions in the internal shocks \citep{2006PhRvD..73f3002M}.
 This body of work supports the hypothesis that neutrino detection by km? Cherenkov detectors will probably be in coincidence with a handful of bright GRBs., This body of work supports the hypothesis that neutrino detection by $^3$ Cherenkov detectors will probably be in coincidence with a handful of bright GRBs.
 Previous studies have correlated photon emission with neutrino emission by approximating the photon spectra with a broken power law., Previous studies have correlated photon emission with neutrino emission by approximating the photon spectra with a broken power law.
 We will show that this approximation overestimates the contribution of high energy photons (and conversely underestimates the contribution trom low energy photons)., We will show that this approximation overestimates the contribution of high energy photons (and conversely underestimates the contribution from low energy photons).
 This has consequences in the number of events expected by neutrino telescopes such as IceCube (Achterbergetal2007c) and KM3NET (Kappesetal2007) and in the characteristic energy of the neutrinos detected., This has consequences in the number of events expected by neutrino telescopes such as IceCube \citep{2007..IceCube..ICRC} and KM3NET \citep{2007..KM3NET..Faro} and in the characteristic energy of the neutrinos detected.
 In section $2. we describe the Band function., In section \ref{sec:band} we describe the Band function.
 In section 3 we follow the standard calculation of HE neutrinos emission from GRBs., In section \ref{sec:photopion} we follow the standard calculation of HE neutrinos emission from GRBs.
 In section $4. we compare the Band function with its broken power law approximation., In section \ref{sec:bandapprox} we compare the Band function with its broken power law approximation.
 In section 5 we calculate neutrino event rates., In section \ref{sec:eventrates} we calculate neutrino event rates.
 And in section 6 we discuss the consequences of the calculation shown here., And in section \ref{sec:discussion} we discuss the consequences of the calculation shown here.
" The time-averaged prompt GRB photon flux is often described by fitting a Band function (Bandetal 1993): The parameters of the Band function are the amplitude -ἰ.. the (asymptotic) low-energy spectral index α-. the high-energy spectral index 2, and the peak energy £y; of the Ff, distribution."," The time-averaged prompt GRB photon flux is often described by fitting a Band function \citep{1993ApJ...413..281B}: : The parameters of the Band function are the amplitude $A_\gamma$ , the (asymptotic) low-energy spectral index $\alpha_\gamma$, the high-energy spectral index $\beta_\gamma$ and the peak energy $E_{\mathrm{pk}}$ of the $\nu F_\nu$ distribution."
" Typical values for the parameters above are: a,zz —1. 3,zz—2.2, ezz 30 keV-1000 keV and A,22 0.001-1 photons ? ! +."," Typical values for the parameters above are: $\alpha_\gamma \approx -1$ , $\beta_\gamma \approx-2.2$, $\epsilon^b_\gamma
\approx$ 30 keV–1000 keV and $A_\gamma \approx$ 0.001–1 photons $^{-2}$ $^{-1}$ $^{-1}$."
 The break energy εὖ andthe peakenergy £j; are related by: There are many good descriptions ofthe processes that lead to HE neutrinos from prompt, The break energy $\epsilon^b_\gamma$ andthe peakenergy $E_{\mathrm{pk}}$ are related by: There are many good descriptions ofthe processes that lead to HE neutrinos from prompt
earlv-0tvpe companions.,early-type companions.
 These sources have spectra similar to black hole LAINB companions (vanParaclijs1999).. and probably are not well separated [rom black hole LAIXNB sources in the color-color diagram.," These sources have spectra similar to black hole LMXB companions \citep{vanP99}, and probably are not well separated from black hole LMXB sources in the color-color diagram."
 There are a small number of sources with II220 and II1—-1., There are a small number of sources with H2=0 and H1=-1.
 These sources essentially have no counts above keV. ancl some are undoubtedly classical supersolt sources llasinger.&Ixahabka 1991).," These sources essentially have no counts above 1keV, and some are undoubtedly classical supersoft sources \citep{greiner91}."
". These ""supersoft"" sources seem (o occur in both bulge ancl disk svstenms.", These “supersoft” sources seem to occur in both bulge and disk systems.
 They are frequently brighter (han tlie sources identilied as supernova remnants. ancl most are variable (see Section 6)).," They are frequently brighter than the sources identified as supernova remnants, and most are variable (see Section \ref{sec:LVSD}) )."
 Penceetal(2001) have identified 10 sources in MIOI as supersolt., \citet{pence01} have identified 10 sources in M101 as supersoft.
 We do not include 6 of these [unter sources because they have less (han 20 counts., We do not include 6 of these fainter sources because they have less than 20 counts.
 We detect the remaining four sources. three of which have 7/1<—0.9 and one of which has 71=—0.51.," We detect the remaining four sources, three of which have $H1\le -0.9$ and one of which has $H1=-0.81$."
 It is also possible Chat some (or even all) of the soft sources identified as thermal SNR. are absorbed supersoft sources., It is also possible that some (or even all) of the soft sources identified as thermal SNR are absorbed supersoft sources.
 Absorption will decrease the τις in the lowest energy bband. shifting the soft color vertically upwards on the color-color plot (see also 5)).," Absorption will decrease the flux in the lowest energy band, shifting the soft color vertically upwards on the color-color plot (see also \ref{sec:abs}) )."
 If this is (he case. we have identified à new population of faint. absorbed. accretion-powered soft sources in disk ealaxies. probably associated with the voung stellar population.," If this is the case, we have identified a new population of faint, absorbed, accretion-powered soft sources in disk galaxies, probably associated with the young stellar population."
 We favor the hypothesis (hat the soft source population is dominated bx SNR. but a new population of accretion-powered soft sources cannot be ruled out.," We favor the hypothesis that the soft source population is dominated by SNR, but a new population of accretion-powered soft sources cannot be ruled out."
 An in-depth study of the variability and spectral characteristics of the soft sources (see Section 6)) is required to distinguish between (hese (wo moc«els., An in-depth study of the variability and spectral characteristics of the soft sources (see Section \ref{sec:LVSD}) ) is required to distinguish between these two models.
 There will undoubtedly be some overlap in source (vpe in the color-color diagram. both because different sources have intrinsically similar spectra (e.g. black hole binaries with hieh and low mass companions) ancl because of absorption.," There will undoubtedly be some overlap in source type in the color-color diagram, both because different sources have intrinsically similar spectra (e.g. black hole binaries with high and low mass companions) and because of absorption."
 Photons in the lowest οποιον bband will be preferentially removed in an absorbed source. making Z1 “harder”.," Photons in the lowest energy band will be preferentially removed in an absorbed source, making $H1$ “harder”."
 This will move (he source vertically up the color-color diagram. as illustrated in Figure 4..," This will move the source vertically up the color-color diagram, as illustrated in Figure \ref{fig:xcol_models}."
 When the absorption becomes very severe (1547 1077) photons in the medium bband will also be removed. causing //2 to become harder.," When the absorption becomes very severe $n_H\sim 10^{22}$ ) photons in the medium band will also be removed, causing $H2$ to become harder."
 The source will then curve to the right in Figure 4.. as shown by the tracks of increasing absorption.," The source will then curve to the right in Figure \ref{fig:xcol_models}, as shown by the tracks of increasing absorption."
 Absorption will cause supersolt sources to be confused with SNR. move SNR into the LAINB region of the diagram. and blur the distinction between LAINBs and IIMXBDs.," Absorption will cause supersoft sources to be confused with SNR, move SNR into the LMXB region of the diagram, and blur the distinction between LMXBs and HMXBs."
 Color information in the bbands will be very limited for highly inclined. galaxies. where essentially all low energy photons are lost.," Color information in the bands will be very limited for highly inclined galaxies, where essentially all low energy photons are lost."
 In this case. sources distinguished primarily on (the basis of their soft color," In this case, sources distinguished primarily on the basis of their soft color"
The magnetic mountain is hvdromagnetically stable. even (hough the conlining magnetic field is heavily distorted.,"The magnetic mountain is hydromagnetically stable, even though the confining magnetic field is heavily distorted."
" Numerical simulations using ZEUS-3D. a mulüpurpose. tiime-dependent. ideal-MIID code for astrophysical {hid dynamics which uses staggered-mesh finite differencing and operator splitüng in three dimensions (Stone&Norman1992).. show that the equilibria from 82.1. are not disrupted by growing Parker or interchange modes over a wide range of accreted mass (LOM,<AL,105M. ) and intervals as long as 10 Alfvénn crossing times (Pavne&Melatos2005)."," Numerical simulations using ZEUS-3D, a multipurpose, time-dependent, ideal-MHD code for astrophysical fluid dynamics which uses staggered-mesh finite differencing and operator splitting in three dimensions \citep{sto92}, show that the equilibria from \ref{sec:gradshafranov}
 are not disrupted by growing Parker or interchange modes over a wide range of accreted mass $10^{-7}\Msun\lesssim\Ma\lesssim 10^{-3}\Msun$ ) and intervals as long as $10^{4}$ Alfvénn crossing times \citep{pay05}."
. The numerical experiments leading to this conclusion are perlormed bv loading the output (p and B) of the Grad-Shafranov. code described in Payne&Melatos(2004) into ZEUS-3D. with the time-step determined by the Courant. condition satisfied bv the [fast maegnetosonic mode.," The numerical experiments leading to this conclusion are performed by loading the output $\rho$ and $\vv{B}$ ) of the Grad-Shafranov code described in \citet{pay04}
 into ZEUS-3D, with the time-step determined by the Courant condition satisfied by the fast magnetosonic mode."
 The code was veriliel (Pavne&Melatos2005) by. reproducing the Classical Parker imstabilitv of a plane-parallel magnetic fiekl (Mouschovias1974) and the analvlic prolile of a static. spherical. isothermal atmosphere.," The code was verified \citep{pay05} by reproducing the classical Parker instability of a plane-parallel magnetic field \citep{mou74} and the analytic profile of a static, spherical, isothermal atmosphere."
" Coordinates are rescaled in ZEUS-3D to handle the disparate radial (22/6M, ) and latitudinal (42,) length scales.", Coordinates are rescaled in ZEUS-3D to handle the disparate radial $c_{\rm s}^2 R_*^2/GM_*$ ) and latitudinal $R_*$ ) length scales.
 The stability is confirmed by plotting the kinetic. gravitational potential. aud magnetic energies as Iunctions of Gime and observing that the total energy decreases back (ο its equilibrium value monotonically ie. the Gracd-Shalranoy equilibria are (local) energv minima.," The stability is confirmed by plotting the kinetic, gravitational potential, and magnetic energies as functions of time and observing that the total energy decreases back to its equilibrium value monotonically i.e. the Grad-Shafranov equilibria are (local) energy minima."
 Note that increasing p uniformly (e.g. five-fold) does lead to a transient Parker instability. (localized near the pole) in which <1% of the magnetic flux in the tutu escapes through the outer boundary. leaving the magnetic dipole and mass ellipticity essentially unaffected.," Note that increasing $\rho$ uniformly (e.g. five-fold) does lead to a transient Parker instability (localized near the pole) in which $\lesssim 1 \%$ of the magnetic flux in the `tutu' escapes through the outer boundary, leaving the magnetic dipole and mass ellipticity essentially unaffected."
 Although the mountain is stable. it does wobble when perturbed. as sound and Alfvénn waves propagate through it (Pavne&Melatos2005).," Although the mountain is stable, it does wobble when perturbed, as sound and Alfvénn waves propagate through it \citep{pay05}."
.. Consequently. the ellipticity € of the star oscillates about its mean value €.," Consequently, the ellipticity $\epsilon$ of the star oscillates about its mean value $\bar{\epsilon}$."
 The [frequency of the oscillation decreases with M. as described below.," The frequency of the oscillation decreases with $\Ma$, as described below."
" The mean value € increases will M, up to a critical mass M, and increases wilh (yw/esV. as described in 83.1.."," The mean value $\bar{\epsilon}$ increases with $\Ma$ up to a critical mass $\Mc$ and increases with $\psi_{\rm a}/\psi_{*}$, as described in \ref{sec:gwpolarization}."
 In ideal MIID. (here is no dissipation and the oscillations persist lor a long (ime even if undriven. decaving on the Alfvénn radiation üime-sceale (which is much longer than our longest simulation run).," In ideal MHD, there is no dissipation and the oscillations persist for a long time even if undriven, decaying on the Alfvénn radiation time-scale (which is much longer than our longest simulation run)."
 In realitw. the oscillations are also damped by ohmic dissipation. which is mimickecl (impreciselv)by. e@ric-related losses in our work.," In reality, the oscillations are also damped by ohmic dissipation, which is mimicked (imprecisely)by grid-related losses in our work."
 To investigate the oscillations quantitatively. we load slightly. perturbed versions of the Grad-Shalranov equilibria in 82.1. into ZEUS-3D and caleulate ε as a function of time /.," To investigate the oscillations quantitatively, we load slightly perturbed versions of the Grad-Shafranov equilibria in \ref{sec:gradshafranov}
 into ZEUS-3D and calculate $\epsilon$ as a function of time $t$."
 Figure 2. shows the result(s of these numerical experiments., Figure \ref{fig:hmean} shows the results of these numerical experiments.
 Grad-Shalranov equilibria are difficult to compute directly. from (2)) and (3)) for M42 1.6M.. because the magnetic topology changes and bubbles form. so instead we emplov a bootstrapping algorithm in," Grad-Shafranov equilibria are difficult to compute directly from \ref{eq:gradshafranov}) ) and \ref{eq:fpsi}) ) for $\Ma \gtrsim 1.6\Mc$ , because the magnetic topology changes and bubbles form, so instead we employ a bootstrapping algorithm in"
"heating rates for large and small values of @ could be expected in models involving the direct transfer of the dissipated energy to electrons: compare Aj. and Ac With Qy., shown in Figs Lt) and 2f).",heating rates for large and small values of $a$ could be expected in models involving the direct transfer of the dissipated energy to electrons; compare $\Lambda_{\rm ie}$ and $\Lambda_{\rm compr}$ with $Q_{\rm diss}$ shown in Figs \ref{fig:rates}( (f) and \ref{fig:rates_stel}( (f).
 Higher density implies also a larger optical depth for large e., Higher density implies also a larger optical depth for large $a$.
 Then. despite a stronger heating in these models. the energy balance in the innermost region is achieved at smaller electron temperatures than for small «.," Then, despite a stronger heating in these models, the energy balance in the innermost region is achieved at smaller electron temperatures than for small $a$."
 On the hand. the highly efficient Comptonization occurs at very small r. therefore. a much more centrally concentrated emissivity of Comptonized radiation is achieved in models with &=0.998 and the luminosity is subject to stronger reduction by GR ettects. discussed below.," On the hand, the highly efficient Comptonization occurs at very small $r$, therefore, a much more centrally concentrated emissivity of Comptonized radiation is achieved in models with $a=0.998$ and the luminosity is subject to stronger reduction by GR effects, discussed below."
 Different velocity fields result in a different strength of collimation toward the black hole horizon., Different velocity fields result in a different strength of collimation toward the black hole horizon.
 In principle. larger values of the radial velocity. [v- imply a stronger collimation for «=0. which could give rise to a larger fraction of photons being captured by the black hole.," In principle, larger values of the radial velocity, $|v^r|$, imply a stronger collimation for $a=0$, which could give rise to a larger fraction of photons being captured by the black hole."
 However. the radiation is produced. on average. at a smaller radial distance for &=0.998 and. therefore. the reduction of the observed luminosity due to GR effects. including the photon capture and the gravitational redshift. turns out to be larger in that case.," However, the radiation is produced, on average, at a smaller radial distance for $a=0.998$ and, therefore, the reduction of the observed luminosity due to GR effects, including the photon capture and the gravitational redshift, turns out to be larger in that case."
 Specifically. in models with «=0.998. {ομωι=0.2 for in=0.5 and 0.3 for rii= 0.1: the large reduction for fz=0.5 results from r>| (and trapping of photons) in the central region.," Specifically, in models with $a=0.998$, $L/Q_{\rm rad,tot} \approx 0.2$ for $\dot m=0.5$ and 0.3 for $\dot
m=0.1$ ; the large reduction for $\dot m=0.5$ results from $\tau > 1$ (and trapping of photons) in the central region."
" For a=0. L/Quau,90.7 (for ir= 0.1) and 0.9 (for rjr=0.5: in this case the emissivity profile is flatter)."," For $a=0$, $L/Q_{\rm rad,tot} \approx 0.7$ (for $\dot m=0.1$ ) and 0.9 (for $\dot
m=0.5$; in this case the emissivity profile is flatter)."
" For high values of a. etfects unique for the Kerr metric. bbending of photon trajectories to the equatorial plane combined with a dependence of the gravitational shift of energy on the direction of photon escape. result in an intrinsic. anisotropy of radiation produced within the innermost few A, (see Piran Shaham 1977. NiedZwwiecki 2005)."," For high values of $a$, effects unique for the Kerr metric, bending of photon trajectories to the equatorial plane combined with a dependence of the gravitational shift of energy on the direction of photon escape, result in an intrinsic anisotropy of radiation produced within the innermost few $R_{\rm g}$ (see Piran Shaham 1977, Niedźwwiecki 2005)."
" The Comptonized component from that region is softer hhas a larger spectral index and a smaller cut-off energy) and has a smaller normalization when observed at a smaller &,,,.", The Comptonized component from that region is softer has a larger spectral index and a smaller cut-off energy) and has a smaller normalization when observed at a smaller $\theta_{\rm obs}$.
" The gravitational ettects. underlying this property. are strong only within 4&,. then. the magnitude of the anisotropy in the total spectrum depends on the proportion between contributions from the innermost and the surrounding (beyond 4A.) regions."," The gravitational effects, underlying this property, are strong only within $4R_{\rm g}$ , then, the magnitude of the anisotropy in the total spectrum depends on the proportion between contributions from the innermost and the surrounding (beyond $4 R_{\rm g}$ ) regions."
" Figs 5tb) and 6(b) show the 8,,. dependence of the observed spectra in our models with &=0.998 Gin models with «=0 the dependence on &,,, is negligible).", Figs \ref{fig:spectra}( (b) and \ref{fig:spectra_stel}( (b) show the $\theta_{\rm obs}$ -dependence of the observed spectra in our models with $a=0.998$ (in models with $a=0$ the dependence on $\theta_{\rm obs}$ is negligible).
 As we can see. the spectra indeed show some anisotropy but the ettect is relatively moderate.," As we can see, the spectra indeed show some anisotropy but the effect is relatively moderate."
 For ri=0.1. the total flux emitted face-on (4.= 0-9-1) and edge- (fans 0-0.1) corresponds to the isotropic luminosity of |x and 3x10? erg |. respectively.," For $\dot m=0.1$, the total flux emitted face-on $\mu_{\rm obs}=0.9$ –1) and edge-on $\mu_{\rm obs}=0$ –0.1) corresponds to the isotropic luminosity of $1 \times$ and $3 \times 10^{43}$ erg $^{-1}$, respectively."
 In this model. the compressive heating results in a rather flat emissivity profile. then. mixing of relatively weak contribution from r«4 with radiation produced at more distant regions washes out the anisotropic properties.," In this model, the compressive heating results in a rather flat emissivity profile, then, mixing of relatively weak contribution from $r<4$ with radiation produced at more distant regions washes out the anisotropic properties."
 For fit=0.5. the face-on and edge-on fluxes correspond to 6.6%10° erg land L.5xI07 eres |. respectively.," For $\dot m=0.5$, the face-on and edge-on fluxes correspond to $6.6 \times 10^{36}$ erg $^{-1}$ and $1.5 \times 10^{37}$ erg $^{-1}$, respectively."
 In this model. the emissivity is strongly centrally concentrated: however. the optical depth of the inner region is large and a large fraction of photons produced there get reprocessed at larger distances or captured.," In this model, the emissivity is strongly centrally concentrated; however, the optical depth of the inner region is large and a large fraction of photons produced there get reprocessed at larger distances or captured."
 At both si. the spectral indices differ by AD“0.1 between a face-on and an on observer.," At both $\dot m$, the spectral indices differ by $\Delta \Gamma \la 0.1$ between a face-on and an edge-on observer."
 We expect that a much larger anisotropy would be produced in a model with moderate 5; (yielding 7«1) and strong direct heating of electrons (giving a very centrally concentrated emissivity)., We expect that a much larger anisotropy would be produced in a model with moderate $\dot m$ (yielding $\tau<1$ ) and strong direct heating of electrons (giving a very centrally concentrated emissivity).
" An optically thick dise may extend to relatively low radii in some objects observed at L/L, similar to these characterising our models.", An optically thick disc may extend to relatively low radii in some objects observed at $L/\ledd$ similar to these characterising our models.
 The best example concerns XTE 311182480. in which the derived truncation radius of the outer cold dise is >100A. (Esin et 22001).," The best example concerns XTE J1118+480, in which the derived truncation radius of the outer cold disc is $\simeq 100 R_{\rm g}$ (Esin et 2001)."
 We consider here the cooling of the inner hot flow due to Comptonization of thermal photons emitted by such a surrounding. cold disc.," We consider here the cooling of the inner hot flow due to Comptonization of thermal photons emitted by such a surrounding, cold disc."
 For this we take the innermost part. within r=100. of our self-consistent solutions described in Section 3.2..," For this we take the innermost part, within $r = 100$, of our self-consistent solutions described in Section \ref{final}."
" We use our MC method to compute the Comptonization of the seed photons emitted by a cold. Keplerian dise extending from rj,=100 to μι=1000."," We use our MC method to compute the Comptonization of the seed photons emitted by a cold, Keplerian disc extending from $r_{\rm in} = 100$ to $r_{\rm out} = 1000$."
 The local black-body temperature of its thermal emission is found using Page Thorne (1974): note that this approach gives the largest possible luminosity of the outer disc as it involves the assumption that the rate of the outward transport of mechanica energy. Which is the dominant source of the radiated energy a r>100. is the same as in an untruncated Keplerian disc.," The local black-body temperature of its thermal emission is found using Page Thorne (1974); note that this approach gives the largest possible luminosity of the outer disc as it involves the assumption that the rate of the outward transport of mechanical energy, which is the dominant source of the radiated energy at $r > 100$, is the same as in an untruncated Keplerian disc."
 We do not look for self-consistent solutions including this additiona source of seed photons., We do not look for self-consistent solutions including this additional source of seed photons.
 Instead. we compute the total rate of the Compton cooling by the outer dise photons. OUpat . and compare it with the total Compton cooling rate due to the synchrotron and bremsstrahlung emission. Qeonpro- of our self-consistent solutions.," Instead, we compute the total rate of the Compton cooling by the outer disc photons, $Q_{\rm Compt,tot}^{\rm th}$ , and compare it with the total Compton cooling rate due to the synchrotron and bremsstrahlung emission, $Q_{\rm Compt,tot}$, of our self-consistent solutions."
" We tind that in models with &=ϐ the effect is small. with QUnun!Ctompun=0.2. and for ¢=0.998 it is negligible. Qtinpa/Conpua,€107,"," We find that in models with $a=0$ the effect is small, with $Q_{\rm Compt,tot}^{\rm th}/Q_{\rm Compt,tot} \approx 0.2$, and for $a=0.998$ it is negligible, $Q_{\rm Compt,tot}^{\rm th}/Q_{\rm Compt,tot} < 10^{-2}$."
 We note that the weakness of this effects results from the following properties., We note that the weakness of this effects results from the following properties.
" First. the total luminosity of the dise truncated at rj,=100 is relatively small (specifically. it is approximately equal to the luminosity of the hot flow)."," First, the total luminosity of the disc truncated at $r_{\rm in}=100$ is relatively small (specifically, it is approximately equal to the luminosity of the hot flow)."
 Second. the solid angle subtended by the flow as seen by the outer dise is rather small.," Second, the solid angle subtended by the flow as seen by the outer disc is rather small."
 Third. the optical thickness of the outer parts of the flow is very small at the values of 77 considered here. so a small fraction of photons illuminating the flow get scattered.," Third, the optical thickness of the outer parts of the flow is very small at the values of $\dot m$ considered here, so a small fraction of photons illuminating the flow get scattered."
 Our results are consistent with the conclusion of Esin et (2001) that the outer dise is not an important source for Compton cooling in this range of parameters., Our results are consistent with the conclusion of Esin et (2001) that the outer disc is not an important source for Compton cooling in this range of parameters.
 Obviously. the etfect would be more important for a smaller truncation radius EEsin 1997). at which the luminosity of the dise would be larger. as well as at higher 1. at which a larger fraction of photons irradiating the flow would be scattered.," Obviously, the effect would be more important for a smaller truncation radius Esin 1997), at which the luminosity of the disc would be larger, as well as at higher $\dot m$, at which a larger fraction of photons irradiating the flow would be scattered."
 Most of the specific properties discussed in this paper have been studied in previous works., Most of the specific properties discussed in this paper have been studied in previous works.
 However. their quantitative importance could not be assessed properly as neither of previous studies considered self-consistently all relevant etfects.," However, their quantitative importance could not be assessed properly as neither of previous studies considered self-consistently all relevant effects."
 The solution of a hot-flow structure based on a self-similar model (eg. Narayan Yi 1995) is commonly used dueto its relative simplicity.," The solution of a hot-flow structure based on a self-similar model (e.g., Narayan Yi 1995) is commonly used dueto its relative simplicity."
 However. this simple solution introduces inaccuracies of up to an order of magnitude even for the basic flow parameters (e... the density given by 11 in Mahadevan Quataert (1997) is ~10 times larger than the densityin our solutions). which leads tofurther inaccuracies in papers based on," However, this simple solution introduces inaccuracies of up to an order of magnitude even for the basic flow parameters (e.g., the density given by 1 in Mahadevan Quataert (1997) is $\sim 10$ times larger than the densityin our solutions), which leads tofurther inaccuracies in papers based on"
Three ideal invariants exist: total enerey E—fdc[(Vc)?anm]: fux F—fdc? and cross-correlation 4.=[«Pene.,"Three ideal invariants exist: total energy $E = \int d^2x [(\nabla \psi)^2 +
\alpha n^2]$; flux $F = \int d^2x\; \psi^2$ and cross-correlation $H_c = \int
d^2 x\; n\psi$."
 Energy cascades to small scale (large ) while the flux and cross-correlation undergo an inverse eascade to large-scale (small 4) 1997)., Energy cascades to small scale (large $k$ ) while the flux and cross-correlation undergo an inverse cascade to large-scale (small $k$ ) \citep{fernandez97}.
. The inverse cascades require the initialized spectrum to peak al fyz0 to allow for buildup of magnetic flix αἱ large-scales Lor later times., The inverse cascades require the initialized spectrum to peak at $k_0 \neq 0$ to allow for buildup of magnetic flux at large-scales for later times.
" Linearizing the svstem vields a (dimensional) dispersion relation ic=Val.p,."," Linearizing the system yields a (dimensional) dispersion relation $\omega
= V_A k_z k_\perp \rho_s$."
 The mode combines perpendicular oscillation associated with a finite gvroradius with fIuctuations along a mean field. (z-direction)., The mode combines perpendicular oscillation associated with a finite gyroradius with fluctuations along a mean field $z$ -direction).
 The oscillating «quantities are magnetic field and density. out of phase by 7/2 radians.," The oscillating quantities are magnetic field and density, out of phase by $\pi/2$ radians."
 In the limit of strong mean field. quantities along the mean field (z-direction) equilibrate quickly. whieh allows 0/02—0. or |.—0.," In the limit of strong mean field, quantities along the mean field $z$ -direction) equilibrate quickly, which allows $\partial / \partial z
\rightarrow 0$, or $k_z \rightarrow 0$."
 Ixinetic Alfvénn waves still propagate. as long as there awe a broad range of scales that are excited. as in fully developed turbulence.," Kinetic Alfvénn waves still propagate, as long as there are a broad range of scales that are excited, as in fully developed turbulence."
 As lh.—0. all gradients are localized to the plane perpendicular to the mean field.," As $k_z \rightarrow 0$, all gradients are localized to the plane perpendicular to the mean field."
 Presuming a large-scale [Iuctuation al characteristic wavenumber ky. smaller-scale [Inctuations propagate linearly along this larger-scale. (lictuation so long as their characteristic scale f satisfies he>fy.," Presuming a large-scale fluctuation at characteristic wavenumber $\mathbf{k}_0$ , smaller-scale fluctuations propagate linearly along this larger-scale fluctuation so long as their characteristic scale $k$ satisfies $k \gg k_0$."
" Im this reduced. (wo dimensional svstem. (he above dispersion relation is moclilied to be w=οι:k/B)hp, which is sll Alfvénnie but with respect to aperturbed amplitude perpendicular to the mean field."," In this reduced, two dimensional system, the above dispersion relation is modified to be $\omega = V_A (\mathbf{b}_{k_0} \cdot \mathbf{k} / B) k \rho_s$ which is still Alfvénnic but with respect to aperturbed large-scale amplitude perpendicular to the mean field."
 Relaxing the scale separation criterion vields wxfh? for the general case., Relaxing the scale separation criterion yields $\omega \propto k^2$ for the general case.
 We evolve ((3)) and (4)) in a 2D periodic box. size [27]x22] on a mesh of resolution 512x512.," We evolve \ref{psi-eqn}) ) and \ref{den-eqn}) ) in a 2D periodic box, size $[2\pi] \times [2\pi]$ on a mesh of resolution $512 \times 512$."
 The & and » scalar fields are evolved in the Fourier domain. with the nonlinearities advanced pseudospectrally and with [ull 2/3 dealiasing in each dimension (Orszag1971).," The $\psi$ and $n$ scalar fields are evolved in the Fourier domain, with the nonlinearities advanced pseudospectrally and with full $2/3$ dealiasing in each dimension \citep{orszag71}."
 The diffusive and resistive terms normally introduce stilfiess into (he equations: using an integrating [actor removes any stability. constraints stemming from theseterms., The diffusive and resistive terms normally introduce stiffness into the equations; using an integrating factor removes any stability constraints stemming from theseterms.
 Following (he scheme outlined in Canutoetal. (1990)... we start with the semi-discrete formulation of," Following the scheme outlined in \citet{canuto90}, , we start with the semi-discrete formulation of"
200 Myr to several Gyr (see Israel. 1998 for a review).,"200 Myr to several Gyr (see Israel, 1998 for a review)."
 The gaseous component associated with the dust lane is in rapid rotation about the major axis of the galaxy. in (he same sense as GI. wilh a maximum rotation of 200 kim/s for the ionized gas (Blane et al.," The gaseous component associated with the dust lane is in rapid rotation about the major axis of the galaxy, in the same sense as G1, with a maximum rotation of 200 km/s for the ionized gas (Bland et al."
 1987) and 265 km/s for the HI gas (van Gorkonm οἱ al., 1987) and 265 km/s for the HI gas (van Gorkom et al.
 1990)., 1990).
 The gaseous disk is warped. ancl (he gas is still unstable in the outer regions. two indications that the merger is recent.," The gaseous disk is warped, and the gas is still unstable in the outer regions, two indications that the merger is recent."
 The stellar component of the galaxy. rotates around the minor axis. will a maximum velocity of about 40 km/s (Wilkinson et al.," The stellar component of the galaxy rotates around the minor axis, with a maximum velocity of about 40 km/s (Wilkinson et al."
 1936)., 1986).
 Considering first G2. it has no net rotation. is composed of old. massive. metal-rich ancl dvnamically evolved GCs. and. because of its high mean radial velocitv. might be associated with a high-velocity component of molecular gas (Israel 1998 and references therein).," Considering first G2, it has no net rotation, is composed of old, massive, metal-rich and dynamically evolved GCs, and, because of its high mean radial velocity, might be associated with a high-velocity component of molecular gas (Israel 1998 and references therein)."
 We propose that this eroup formed during the very. [ist merger that gave rise to the elliptical ealaxv., We propose that this group formed during the very first merger that gave rise to the elliptical galaxy.
 It is interesting to note that the metallicity distribution and age of G2 are very close to those of the outer halo stars of NGC! 5128. for which Rejkuba et al. (," It is interesting to note that the metallicity distribution and age of G2 are very close to those of the outer halo stars of NGC 5128, for which Rejkuba et al. ("
2005) derived a niean metallicity of -0.64 and a mean age of 8 Gyr.,2005) derived a mean metallicity of -0.64 and a mean age of 8 Gyr.
 Ht is thus tempting to assume that the GCs of G2 and the outer halo stars have the same origin., It is thus tempting to assume that the GCs of G2 and the outer halo stars have the same origin.
 The GCs of G2 are on average 9.4 Gvr old (see Table 4): however. considering that these are mostly massive GCs. (his age might be influenced by the unwanted presence in the spectra of horizontal-branch stars. and in fact much older.," The GCs of G2 are on average 9.4 Gyr old (see Table 4); however, considering that these are mostly massive GCs, this age might be influenced by the unwanted presence in the spectra of horizontal-branch stars, and in fact much older."
 Consequently. the merger may be even older than 9.4 Gyr. and at anv rate much older than (the one(s) that gave rise to the shells and dust lane.," Consequently, the merger may be even older than 9.4 Gyr, and at any rate much older than the one(s) that gave rise to the shells and dust lane."
 Mergers do produce a large number of super star clusters: for exanmple. a thousand such clusters have been detected in NGC 4038/4039 (Mengel et al.," Mergers do produce a large number of super star clusters; for example, a thousand such clusters have been detected in NGC 4038/4039 (Mengel et al."
 2005). a minority of which could later become massive GC's (Whitmore et al.," 2005), a minority of which could later become massive GCs (Whitmore et al."
 2007)., 2007).
 Super star clusters in (hat and other recent merger remnants. NGC 7252 (Schweizer seilzer 1993). NGC 1275 (Zepl et al.," Super star clusters in that and other recent merger remnants, NGC 7252 (Schweizer Seitzer 1998), NGC 1275 (Zepf et al."
 1995). are generally of solar metallicitv.," 1995), are generally of solar metallicity."
 The metallicity of G2 is lower than solar. because the merger occurred at least 9.4 Gvr ago.," The metallicity of G2 is lower than solar, because the merger occurred at least 9.4 Gyr ago."
Great care has been dedicated to manage the possible errors during mask bliaudliug (e.g. mask dropped or damaged diwiug manual displacements. wrong iusertiou into the ος or ICs. problems in bar code scanner or robot connection etc.).,"Great care has been dedicated to manage the possible errors during mask handling (e.g. mask dropped or damaged during manual displacements, wrong insertion into the SCs or ICs, problems in bar code scanner or robot connection etc.)."
 All the error messages are listed in the user's mauual and contain the suggestedMOD recovery procedure using the & provided recovery [unctious., All the error messages are listed in the user's manual and contain the suggested recovery procedure using the 8 provided recovery functions.
 The VIRMOS Mask Manuacturing Unit is now operational at Paranal Observatov. following verification that the specificaious set at the beginning of the project have been imei.," The VIRMOS Mask Manufacturing Unit is now operational at Paranal Observatory, following verification that the specifications set at the beginning of the project have been met."
 The laser cutting machine has been finely tuned with an optimal choice of the parameters. aid. although the MALU is largely assemblec [rom industrial products. every effort has been mace to desigu aud üinpleimenut software toos Which would make the mask flow nicely. fit in the VLT data flow system.," The laser cutting machine has been finely tuned with an optimal choice of the parameters, and, although the MMU is largely assembled from industrial products, every effort has been made to design and implement software tools which would make the mask flow nicely fit in the VLT data flow system."
 The slit quality that the VIROS MALU provides makes it the best operatioual tool for multi-object spectroscopy applicatio15., The slit quality that the VIRMOS MMU provides makes it the best operational tool for multi-object spectroscopy applications.
 The VIRMOS-NAU mauufacturiug is done τιider ESO contract 50079/INS/0T/7260/CWIL., The VIRMOS-MMU manufacturing is done under ESO contract 50979/INS/97/7569/GWI.
 The MMU developmen was aided by financial support from Italian CNR aud CNAA., The MMU development was aided by financial support from Italian CNR and CNAA.
 We warily thank LPIKF techuieal stall (T. Nagel. A. Mülller and S. Bonniek) for their appreciated collaboration during the tuuine of the laser cutting machine auk LN. D'Addea (CNR-ITIA) for lending us the laser interferometer system.," We warmly thank LPKF technical staff (T. Nagel, A. Mülller and S. Bönnigk) for their appreciated collaboration during the tuning of the laser cutting machine and N. D'Addea (CNR-ITIA) for lending us the laser interferometer system."
 We also thank C. Avia. aud the Paranal ESO Engineering stall lor their help during installation of the MMU at the Oservatlory.," We also thank G. Avila, and the Paranal ESO Engineering staff for their help during installation of the MMU at the Observatory."
the shape of the SED long-ward of the 2° band cannot be properly constrained (with the early-release UDS data). making it impossible to identify robust candidates (see Section 3.2).,"the shape of the SED long-ward of the $z^{\prime}-$ band cannot be properly constrained (with the early-release UDS data), making it impossible to identify robust candidates (see Section 3.2)."
 The requirements for non-detections in the D and V— bands. together with a red Roz'=3 colour. are designed to ensure our selection function has a sharp low-redshift cut-off at 2= 5.," The requirements for non-detections in the $B$ and $V-$ bands, together with a red $R-z^{\prime}\geq3$ colour, are designed to ensure our selection function has a sharp low-redshift cut-off at $z=5$ ."
 At >=5 both the 6 and \ bands sample rest-frame wavelengths short-ward of Lya. and any genuine high-redshift LBGs should not be detected at -2m significance in these bands.," At $z\geq5$ both the $B$ and $V-$ bands sample rest-frame wavelengths short-ward of $\alpha$, and any genuine high-redshift LBGs should not be detected at $>2\sigma$ significance in these bands."
 Moreover. in the redshift interval 5o«os<6 the Subaru /? and z/ filters bracket the onset of absorption by the Lya forest. and simulations demonstrate that a colour-cut of 2z>3 effectively excludes galaxies at z<5.," Moreover, in the redshift interval $5<z<6$ the Subaru $R$ and $z^{\prime}-$ filters bracket the onset of absorption by the $\alpha$ forest, and simulations demonstrate that a colour-cut of $R-z^{\prime}\geq3$ effectively excludes galaxies at $z<5$."
 Finally. it should be noted that during the optical selection process we did not actually impose the standard 7 drop criterion.," Finally, it should be noted that during the optical selection process we did not actually impose the standard $i-$ drop criterion."
 The reason for this is that the standard colour cut (18. £pizss0Ze1. Bunker et al., The reason for this is that the standard colour cut (i.e. $i_{814}-z_{850}>1.3$; Bunker et al.
 2004) can exclude galaxies in the redshift range 5.0«25.5., 2004) can exclude galaxies in the redshift range $5.0<z<5.5$.
 In contrast. for the purposes of this study. we are interested in including all possible candidates at >25. which is better achieved using the 22!>3 colour-cut.," In contrast, for the purposes of this study, we are interested in including all possible candidates at $z\geq5$, which is better achieved using the $R-z^{\prime}\geq3$ colour-cut."
 However. as can be seen from Table I. all of the final nine LBG candidates satisfy /ZOm]1.0. fully consistent with the expected colours of 2>5r LBGs in the Subaru filter set.," However, as can be seen from Table 1, all of the final nine LBG candidates satisfy $i^{\prime}-z^{\prime}\geq1.0$, fully consistent with the expected colours of $z\geq5$ LBGs in the Subaru filter set."
 Following the application of the initial optical selection criteria. the next step in the selection process was to analyse each potential candidate with our template-fitting photometric redshift code (Cirasuolo et al.," Following the application of the initial optical selection criteria, the next step in the selection process was to analyse each potential candidate with our template-fitting photometric redshift code (Cirasuolo et al."
 2006. in preparation). which is largely based on the public. package (Bolzonella. Miralles Pelló 2000).," 2006, in preparation), which is largely based on the public package (Bolzonella, Miralles Pelló 2000)."
 The photometrie redshift for each galaxy is computed by fitting the observed multi-wavelength (BVR7z AN) photometry with synthetic galaxy templates.," The photometric redshift for each galaxy is computed by fitting the observed multi-wavelength $
BVRi'z'JK$ ) photometry with synthetic galaxy templates."
 We adopt the stellar population synthesis models of Bruzual Charlot (2003) for the galaxy templates. assuming a Salpeter initial mass function (IMF) with a lower and upper mass cutoff and 100 A7. respectively.," We adopt the stellar population synthesis models of Bruzual Charlot (2003) for the galaxy templates, assuming a Salpeter initial mass function (IMF) with a lower and upper mass cutoff of 0.1 and 100 $M_{\odot}$ respectively."
 Both instantaneous. burst and exponentially.of0.1 declining star formation models — with e-folding times in the range O.Lr(Cvr)< — were included. all with fixed solar metallicity.," Both instantaneous burst and exponentially declining star formation models – with e-folding times in the range $\rm 0.1 \leq \tau (Gyr)\leq 15$ – were included, all with fixed solar metallicity."
 The code accounts for dust reddening by following the Calzetti et al. (, The code accounts for dust reddening by following the Calzetti et al. (
2000) obscuration law within the range 0—chy10. and accounts for Lyman series absorption due to the HI clouds in the inter galactic medium according to the Madau (1995) prescription.,"2000) obscuration law within the range $0 \leq A_V \le 10$, and accounts for Lyman series absorption due to the HI clouds in the inter galactic medium according to the Madau (1995) prescription."
 In addition to the best-titting photometric redshift. the code also returns best-fit values for: SED type. age. mass and reddening.," In addition to the best-fitting photometric redshift, the code also returns best-fit values for: SED type, age, mass and reddening."
 In addition to providing redshift and stellar mass esimates. the principle reason for performing the optical+near-infrared SED fits was to clean the sample of low-redshift interloVers.," In addition to providing redshift and stellar mass estimates, the principle reason for performing the optical+near-infrared SED fits was to clean the sample of low-redshift interlopers."
 A simple application of the optical selection criteria described above produced an initial sample of 74 candidates., A simple application of the optical selection criteria described above produced an initial sample of 74 candidates.
 By making full use of the extra information provided by the near-infrared UDS data. 65 members of this initial sample were eventually excluded as low-redshift contaminants.," By making full use of the extra information provided by the near-infrared UDS data, 65 members of this initial sample were eventually excluded as low-redshift contaminants."
 The names. positions and photometry for the final nine LBG candidates are listed in Table |. and the fits to their opticaltnear-infrared photometry are shown in Fig |.," The names, positions and photometry for the final nine LBG candidates are listed in Table 1, and the fits to their optical+near-infrared photometry are shown in Fig 1."
 In the next section We discuss in deail the potential sources of contamination. and the strategies adopted for excluding them from the final sample.," In the next section we discuss in detail the potential sources of contamination, and the strategies adopted for excluding them from the final sample."
 In the current study there are three types of contaminants which are of potential concern: ultra-cool galactic stars (M. Lo or T dwarfs). extremely red objects (EROS) at 2—1 and high-redshift quasars.," In the current study there are three types of contaminants which are of potential concern: ultra-cool galactic stars (M, L or T dwarfs), extremely red objects (EROs) at $z\simeq1$ and high-redshift quasars."
 It is important to develop strategies for excluding these contaminants based on their optical+near-infrared SEDs because. unlike previous HST studies. ground-based resolution makes it impossible to exclude stellar and QSO contaminants on the basis of being spatially unresolved.," It is important to develop strategies for excluding these contaminants based on their optical+near-infrared SEDs because, unlike previous HST studies, ground-based resolution makes it impossible to exclude stellar and QSO contaminants on the basis of being spatially unresolved."
 Due to the ~0.8’ FWHM resolution of the availabe ground-based imaging of the SXDS/UDS field. high-redshift LBGs are expected to be unresolved point sources.," Due to the $\simeq0.8$ FWHM resolution of the available ground-based imaging of the SXDS/UDS field, high-redshift LBGs are expected to be unresolved point sources."
 Consequently. the prospect of significant contamination of the final sample by ultra-cool galactic stars is potentially serious.," Consequently, the prospect of significant contamination of the final sample by ultra-cool galactic stars is potentially serious."
 Ultra-cool galactic stars (M. L or T. dwarfs) perfectly mimic the redshifted SEDs of ς=5 LBGs in the optical (Hewett et al.," Ultra-cool galactic stars (M, L or T dwarfs) perfectly mimic the redshifted SEDs of $z\geq5$ LBGs in the optical (Hewett et al."
 2006). and would certainly pass our initial optical selection criteria.," 2006), and would certainly pass our initial optical selection criteria."
 Consequently. in order to exclude stellar contaminants it is necessaryto rely on the fact that ultra-cool galactic stars exhibit redder 2’J colours than genuine z>5 LBGs.," Consequently, in order to exclude stellar contaminants it is necessaryto rely on the fact that ultra-cool galactic stars exhibit redder $z^{\prime}-J$ colours than genuine $z\geq5$ LBGs."
 To quantify this we used the library of spectra maintained, To quantify this we used the library of spectra maintained
The result of image segmentation is the detection list will the positions of candidate point source are determined.,The result of image segmentation is the detection list with the positions of candidate point source are determined.
 During the estimation stage a thorough fit of the data is performed to refine the positions. determine (he [Iuxes ancl deblend extracted point sources.," During the estimation stage a thorough fit of the data is performed to refine the positions, determine the fluxes and deblend extracted point sources."
 The background subtracted image is normally used [or the point source fitting., The background subtracted image is normally used for the point source fitting.
 A point source sublracted image can be produced (o assess the quality of point source extraction., A point source subtracted image can be produced to assess the quality of point source extraction.
" A point response function (PIE) is required to perform linear matched filtering aud data filling,", A point response function (PRF) is required to perform linear matched filtering and data fitting.
 MOPIEN included a separate task of PRE estimation., MOPEX included a separate task of PRF estimation.
 We illustrate (he intermediate products al various processing stages in Figure 2 using a small [ragment (90 by GO pixels) ol an MIPSTO mosaic image., We illustrate the intermediate products at various processing stages in Figure \ref{Illustration} using a small fragment (90 by 60 pixels) of an MIPS70 mosaic image.
 The mosaic is described in more details in Section ??.., The mosaic is described in more details in Section \ref{RealMIPS70}.
 In (his paper we present an overview of (he algorithm., In this paper we present an overview of the algorithm.
 One can create a processing chain to accomodate for any specilic features of the data. as descerived. for exaniple. in Section ??..," One can create a custom-tailored processing chain to accomodate for any specific features of the data, as descrived, for example, in Section \ref{RealMIPS24}."
 Careful tuning of various stages can be performed using a number of parameters., Careful tuning of various stages can be performed using a number of parameters.
 For details on running MOPEX. comprehensive listing of all the processing parameters. and sample namelists one can consult the online guide!.," For details on running MOPEX, comprehensive listing of all the processing parameters, and sample namelists one can consult the online guide."
. Dackground sublractecl images are used [or both point source detection ancl fitting., Background subtracted images are used for both point source detection and fitting.
 A common wav of estimating the background is to find lor each pixel the median in a window centered around (hat pixel., A common way of estimating the background is to find for each pixel the median in a window centered around that pixel.
 This method will inevitably overestimate the background in the vieinity of bright sources., This method will inevitably overestimate the background in the vicinity of bright sources.
 MOPEX provides a user controlled wav to alleviate the bias introduced bv the bright sources., MOPEX provides a user controlled way to alleviate the bias introduced by the bright sources.
 One can counter-bias the median by exeluding a lixed number of pixels with highest values in each window from median computation. (hus olfsetiing the bias introduced by the presence of bright sources.," One can counter-bias the median by excluding a fixed number of pixels with highest values in each window from median computation, thus offsetting the bias introduced by the presence of bright sources."
 This number is defined by (he user and is constant througout the image to which it is applied., This number is defined by the user and is constant througout the image to which it is applied.
 This step may require [further optimization bv automatically adjusting this number based on (he crowcliness in anv particular window., This step may require further optimization by automatically adjusting this number based on the crowdiness in any particular window.
 Ilowever. even in iis present state the quality of median subtracted images satisfies the needs of point source detection.," However, even in its present state the quality of median subtracted images satisfies the needs of point source detection."
 Anv bias in the background estimation introduced by the median fillering can be corrected. al the point source fitting stage as described below in Section ??.. since (he fitting includes among other fitting parameters a constant background for the fitting area.," Any bias in the background estimation introduced by the median filtering can be corrected at the point source fitting stage as described below in Section \ref{Extraction}, since the fitting includes among other fitting parameters a constant background for the fitting area."
"Using features derived at the time of discovery, we obtain a overall error rate (all error rates stated are found using 10-fold cross validation).","Using features derived at the time of discovery, we obtain a overall error rate (all error rates stated are found using 10-fold cross validation)."
" For the 1422 sources with SDSS coverage, the error rate is1.7%,, while for the other 531 sources with no SDSS coverage the error rate jumps to9."," For the 1422 sources with SDSS coverage, the error rate is, while for the other 531 sources with no SDSS coverage the error rate jumps to."
4%.. In Figure 11 we plot the ROC curves for both variable star and transient source classification., In Figure \ref{roc_ptfClass_discovery} we plot the ROC curves for both variable star and transient source classification.
" The ROC curves show that at purity, the random forest classifier attains efficiency of variable star classification and efficiency of transient classification."," The ROC curves show that at purity, the random forest classifier attains efficiency of variable star classification and efficiency of transient classification."
" Notably, for SDSS sources, we achieve a (100%)) efficiency of (TRANSIENT) classification at a purity level."," Notably, for SDSS sources, we achieve a ) efficiency of ) classification at a purity level."
" In robotclass classification, which divides the sources into five science classes, random forest obtains an error rate of6."," In robotclass classification, which divides the sources into five science classes, random forest obtains an error rate of."
"5%.. Figure 12 shows that for theAGN-cnSN-TDE,SN/Nova, and classes the classifier attains97%,,93%,, and recovery, respectively."," Figure \ref{confmat_robotClass_discovery} shows that for the, and classes the classifier attains, and recovery, respectively."
" Due to a large class imbalance, performance of the classifier suffers for the smaller classes andVarStar-misc."," Due to a large class imbalance, performance of the classifier suffers for the smaller classes and."
" Again, our classifier performs significantly better for sources in SDSS, attaining a error rate compared to error for sources with no SDSS coverage."," Again, our classifier performs significantly better for sources in SDSS, attaining a error rate compared to error for sources with no SDSS coverage."
" As more data are collected (post time of discovery), the robotclass random forest error rate decreases slightly: the error rate for objects without SDSS coverage drops to after 30 days and after 90 days, while the error rate for objects in SDSS does not change significantly with increased PTF observations."," As more data are collected (post time of discovery), the robotclass random forest error rate decreases slightly: the error rate for objects without SDSS coverage drops to after 30 days and after 90 days, while the error rate for objects in SDSS does not change significantly with increased PTF observations."
 This implies that additional PTF data only helps in classification when no SDSS features are available., This implies that additional PTF data only helps in classification when no SDSS features are available.
" Finally, the RF classification trees allow us to construct an estimate of the importance of each feature in the classifier."," Finally, the RF classification trees allow us to construct an estimate of the importance of each feature in the classifier."
" Using the prescription of Breiman(2001),, we compute the importance of each feature as the increased number of sources that are correctly classified when using that feature instead of a replacement feature of random noise."," Using the prescription of \citet{2001brei}, we compute the importance of each feature as the increased number of sources that are correctly classified when using that feature instead of a replacement feature of random noise."
" In Figure 13 we plot the importance of each feature for each of the robotclass classes was omitted due to a scarce amount of data), and the average importance across all classes."," In Figure \ref{featImp_robotClass_discovery} we plot the importance of each feature for each of the robotclass classes was omitted due to a scarce amount of data), and the average importance across all classes."
" Overall, the most important features are context based, while some light-curve-derived features (such as the ratio of the number of negative subtractions to positive subtractions) are important for distinguishing between certain classes."," Overall, the most important features are context based, while some light-curve-derived features (such as the ratio of the number of negative subtractions to positive subtractions) are important for distinguishing between certain classes."
" In the future, we may add more descriptive time-series features (such as those related to periodograms) which should also be useful in classification."," In the future, we may add more descriptive time-series features (such as those related to periodograms) which should also be useful in classification."
 There are some biases in the sample generation that require a careful interpretation of these ML results., There are some biases in the sample generation that require a careful interpretation of these ML results.
" For a source to be included in the training sample via existing catalogs, it must have a SIMBAD label (e.g., “RRLyr*” or *QSO"") that provides a “definitive” ground-truth statement about the nature of the variability."," For a source to be included in the training sample via existing catalogs, it must have a SIMBAD label (e.g., “RRLyr*” or “QSO”) that provides a “definitive” ground-truth statement about the nature of the variability."
" In some cases, that SIMBAD label comes from SDSS spectroscopy (particularly for quasars); since SDSS spectroscopy is used in the ML classification, the information in some of the training set is essentially known perfectly in the classifier (this is one explanation why classification is inferior in non-SDSS footprint fields)."," In some cases, that SIMBAD label comes from SDSS spectroscopy (particularly for quasars); since SDSS spectroscopy is used in the ML classification, the information in some of the training set is essentially known perfectly in the classifier (this is one explanation why classification is inferior in non-SDSS footprint fields)."
" Also, SIMBAD sources tend to be brighter than many PTF sources and so the above analysis can be thought of as applying to the brighter end of the distribution."," Also, SIMBAD sources tend to be brighter than many PTF sources and so the above analysis can be thought of as applying to the brighter end of the distribution."
 Spectroscopically confirmed SNe candidates found in PTF which are used in the training are obtained after humans in the PTF collaboration have vetted the PTF image- discoveries and decided to pursue spectroscopic followup., Spectroscopically confirmed SNe candidates found in PTF which are used in the training are obtained after humans in the PTF collaboration have vetted the PTF image-difference-based discoveries and decided to pursue spectroscopic followup.
 À bright supernova that, A bright supernova that
otherr wavelengths.,r wavelengths.
 Their true nature is. therefore. still unknown.," Their true nature is, therefore, still unknown."
Figure 7 compares also the contribution from the first generation of gamma rays with the full cascade radiation.,Figure \ref{mag} compares also the contribution from the first generation of gamma rays with the full cascade radiation.
 For low magnetic field (B.<5 G). all generations should be considered in the calculation.," For low magnetic field $B\la5~$ G), all generations should be considered in the calculation."
 For higher magnetic field (B>5 G). the first generation of gamma rays dominates the total cascade radiation.," For higher magnetic field $B>5~$ G), the first generation of gamma rays dominates the total cascade radiation."
 Only a few pairs ean radiate beyond the threshold energy for pair production and the cascade Is quenched., Only a few pairs can radiate beyond the threshold energy for pair production and the cascade is quenched.
 A non-uniform magnetic field was also investigated for a toroidal or dipolar magnetic structure generated by the massive star with a R7! or R? dependence)., A non-uniform magnetic field was also investigated for a toroidal or dipolar magnetic structure generated by the massive star with a $R^{-1}$ or $R^{-3}$ dependence).
 These configurations do not give different results compared with the uniform case., These configurations do not give different results compared with the uniform case.
 Most of the cascade radiation is produced close to the primary source (see Sect. 4.1)), Most of the cascade radiation is produced close to the primary source (see Sect. \ref{spatial}) )
 and depends mostly on the magnetic field strength at this location., and depends mostly on the magnetic field strength at this location.
 The full cascade radiation calculation is applied to LS 5039 and discussed below., The full cascade radiation calculation is applied to LS 5039 and discussed below.
 The black body spectrum and the spatial extension of the massive star are taken into account in this part., The black body spectrum and the spatial extension of the massive star are taken into account in this part.
 The primary source of gamma rays is computed here following the model described in ? where the pulsar is assumed to inject energetic electron-positron pairs with an isotropic power-law energy distribution at the shock front. expected to lie at the vicinity of the compact star.," The primary source of gamma rays is computed here following the model described in \citet{2008A&A...477..691D} where the pulsar is assumed to inject energetic electron-positron pairs with an isotropic power-law energy distribution at the shock front, expected to lie at the vicinity of the compact star."
" Taking v,=2400kms. M=1077M.yr! for the massive star wind (?).. and a pulsar spin-down luminosity L,=1076ergs7!. both wind momenta balance at a distance ος0.14. from the pulsar."," Taking $v_{\infty}=2400~\rm{km~s}^{-1}$, $\dot{M}=10^{-7}~M_{\odot}~\rm{yr}^{-1}$ for the massive star wind \citep{2004ApJ...600..927M}, and a pulsar spin-down luminosity $L_p=10^{36}~\rm{erg~s}^{-1}$, both wind momenta balance at a distance $r_{shock}\sim 0.1 d$ from the pulsar."
 Pairs generated by the pulsar emit inverse Compton scattering on stellar photons the primary gamma-ray photons., Pairs generated by the pulsar emit inverse Compton scattering on stellar photons the primary gamma-ray photons.
 Contrary to the previous section. the primary gamma-ray source is highly anisotropic.," Contrary to the previous section, the primary gamma-ray source is highly anisotropic."
 The orbital parameters of the system are taken from ?.., The orbital parameters of the system are taken from \citet{2005MNRAS.364..899C}.
 New optical observations of LS 5039 have been carried out recently by ? where slight corrections to theorbital parameters have been reported. but these do not change the results below.," New optical observations of LS 5039 have been carried out recently by \citet{2009ApJ...698..514A} where slight corrections to theorbital parameters have been reported, but these do not change the results below."
[rom observations.,from observations.
" The case with Zi=10""Ix and low density ny=0.0lema.*. which max be representative of the hot phase of the interstellar meclitim. does not lead to a shock as strongly modified as in the other cases because in such an environment an age older (han ~[OO0vr is needed to achieve a strong moclilication."," The case with $T_0=10^6\degK$ and low density $n_0=0.01\rmn{cm}^{-3}$, which may be representative of the hot phase of the interstellar medium, does not lead to a shock as strongly modified as in the other cases because in such an environment an age older than $\sim 1000\rmn{yr}$ is needed to achieve a strong modification."
" As discussed in Sec, ?2?.. "," As discussed in Sec. \ref{sec:model} ,"
the propagation of the accelerated particles in the shock region is described by Eq. 23.., the propagation of the accelerated particles in the shock region is described by Eq. \ref{conv-diff}.
 The terms u(r)—ο.) describe the velocity of (he scattering centers., The terms $u(x)-v_A(x)$ describe the velocity of the scattering centers.
 Ilere v4 represents the wave velocity. which. as discussed above. we assume to be the Allvénn speed ealeulated in the background field By.," Here $v_A$ represents the wave velocity, which, as discussed above, we assume to be the Alfvénn speed calculated in the background field $B_0$."
 Provided the waves remain Alfvénn waves even in the regime ουx1 (hiis result. holds in an exact wav., Provided the waves remain Alfvénn waves even in the regime $\delta B/B\gg 1$ this result holds in an exact way.
 It is however clear (hat (his can onlv be considered as an ideal case. in that turbulence may change such a picture in a considerable wav.," It is however clear that this can only be considered as an ideal case, in that turbulence may change such a picture in a considerable way."
 One wav in which the change may appear is in changing the wave speed., One way in which the change may appear is in changing the wave speed.
 In (his section we investigate the effects on the spectrum of accelerated particles and on the shock precursor which derive from calenlating the Alfvénn speed in the local amplified field. namely: We stress (hat in our opinion this assumption is totallv unjustified Irom the physical point of view. and we use (liis case only as à tov model to illustrate how sensitive (he results can be to unknown non-linear effects.," In this section we investigate the effects on the spectrum of accelerated particles and on the shock precursor which derive from calculating the Alfvénn speed in the local amplified field, namely: We stress that in our opinion this assumption is totally unjustified from the physical point of view, and we use this case only as a toy model to illustrate how sensitive the results can be to unknown non-linear effects."
 We notice however that. similar approaches have in fact been adopted. for instance by Zirakashvili&Ptuskin(2008b) ancl also. in some form. by Bell&Lucek(2001).," We notice however that similar approaches have in fact been adopted, for instance by \cite{poster} and also, in some form, by \cite{bl2001}."
. The net effect of this apparently harmless assumption is that the velocity of the scattering centers is ereatly enhanced and (his affects in an substantial wav Che effective compression factor al the subshock and therefore (he spectrum. as was already pointed out by in (he context of test particle theory.," The net effect of this apparently harmless assumption is that the velocity of the scattering centers is greatly enhanced and this affects in an substantial way the effective compression factor at the subshock and therefore the spectrum, as was already pointed out by \cite{bell78} in the context of test particle theory."
 In order to illustrate this effect in a quantitative wav we run our calculations Lor o. Aly=250. Ty=10? IK. ny=0.5em7. By=5uG and an age of the svstem ol ~1000vr.," In order to illustrate this effect in a quantitative way we run our calculations for $p_{inj}=3.7 p_{th,2}$ , $M_0=250$, $T_0=10^5$ K, $n_0=0.5\rmn{cm}^{-3}$, $B_0=5\umu\rmn{G}$ and an age of the system of $\sim1000\rmn{yr}$."
 The results are illustrated in Fig. 2.., The results are illustrated in Fig. \ref{fig:vAmod}.
 The solid lines are obtained using the Alfvénn speed calculated in the background magnetic field By. as in the previous section.," The solid lines are obtained using the Alfvénn speed calculated in the background magnetic field $B_0$, as in the previous section."
 In this case (he precursor is verv evident (top left panel in Fig. 2)), In this case the precursor is very evident (top left panel in Fig. \ref{fig:vAmod}) )
 and the spectrum of accelerated particles is concave (top right panel in Fig. 2))., and the spectrum of accelerated particles is concave (top right panel in Fig. \ref{fig:vAmod}) ).
SNRs.,SNRs.
 We choose proton and helium for our study and consider only those SNRs with distances <1 kpc from the Earth and ages yr., We choose proton and helium for our study and consider only those SNRs with distances $<1$ kpc from the Earth and ages $< 2\times 10^5$ yr.
" From the available literature, we found 10 SNRs listed as follows with their distances (kpc) and ages (yr) respectively given in parentheses: Cygnus Loop (0.54, 10*), HB21 (0.8,1.9x103), ΗΒΟ (0.8,6.6x 103), $147 (0.8,4.6x10°), Vela (0.3,1.1x10°), G299.2-2.9 (0.5,5x 103), SN185 (0.95,1.8x103), Monogem (0.3,1.1x10°), G114.3+0.3 (0.7,4.1x10*) and Vela Junior (0.75,3.5x103)."," From the available literature, we found 10 SNRs listed as follows with their distances (kpc) and ages (yr) respectively given in parentheses: Cygnus Loop $(0.54, 10^4)$ , HB21 $(0.8, 1.9\times 10^4)$, HB9 $(0.8, 6.6\times 10^3)$ , S147 $(0.8, 4.6\times 10^3)$, Vela $(0.3, 1.1\times 10^4)$ , G299.2-2.9 $(0.5, 5\times 10^3)$ , SN185 $(0.95, 1.8\times 10^3)$, Monogem $(0.3, 1.1\times 10^5)$, G114.3+0.3 $(0.7, 4.1\times 10^4)$ and Vela Junior $(0.75, 3.5\times 10^3)$."
" In addition to the contributions from the nearby SNRs, we assume that there exists a steady CR background in the Galaxy which dominates the overall CR spectrum."," In addition to the contributions from the nearby SNRs, we assume that there exists a steady CR background in the Galaxy which dominates the overall CR spectrum."
" For the CRs observed at the Earth, we assume that this background component consists of contributions from distant SNRs plus any other possible sources in the Galaxy."," For the CRs observed at the Earth, we assume that this background component consists of contributions from distant SNRs plus any other possible sources in the Galaxy."
" For our study, we obtain the background by fitting the observed CR spectrum between (20—200) GeV/n. This is the energy region where the contamination from the nearby sources is expected to be less and at the same time, not much affected by the Solar modulation."," For our study, we obtain the background by fitting the observed CR spectrum between $(20-200)$ GeV/n. This is the energy region where the contamination from the nearby sources is expected to be less and at the same time, not much affected by the Solar modulation."
" In fact, it has been shown in Thoudam 2008 that the presence of nearby sources can produce stronger density fluctuations at higher energies than at lower energies because of the energy dependent nature of CR diffusion."," In fact, it has been shown in Thoudam 2008 that the presence of nearby sources can produce stronger density fluctuations at higher energies than at lower energies because of the energy dependent nature of CR diffusion."
" Therefore, we believe that it is reasonable to assume that the low energy CRs that we observe at the Earth are not much affected by the presence of nearby SNRs and they largely represent the averaged background spectrum in the Galaxy."," Therefore, we believe that it is reasonable to assume that the low energy CRs that we observe at the Earth are not much affected by the presence of nearby SNRs and they largely represent the averaged background spectrum in the Galaxy."
 We will show in the following that this is indeed the most likely case., We will show in the following that this is indeed the most likely case.
" From the fit, the spectral indices of the background CRs are found to be 2.75+0.01 for the protons and 2.68+0.02 for the helium."," From the fit, the spectral indices of the background CRs are found to be $2.75\pm0.01$ for the protons and $2.68\pm0.02$ for the helium."
 The reason for the flatter helium spectrum is not properly understood., The reason for the flatter helium spectrum is not properly understood.
" Recently, Blasi & Amato 2011 showed that the flatter helium spectrum with respect to the protons above 1 TeV could be due to spallation effects."," Recently, Blasi $\&$ Amato 2011 showed that the flatter helium spectrum with respect to the protons above $1$ TeV could be due to spallation effects."
" Later, Vladimirov et al."," Later, Vladimirov et al."
 2011 showed that such effects can lead to boron-to-carbon ratios and the anti-proton fluxes which are inconsistent with the observed data., 2011 showed that such effects can lead to boron-to-carbon ratios and the anti-proton fluxes which are inconsistent with the observed data.
 Another possibility for the different spectral indices could be that the intrinsic source spectra itself are different., Another possibility for the different spectral indices could be that the intrinsic source spectra itself are different.
 It could be due to different acceleration sites of protons and helium (Biermann et al., It could be due to different acceleration sites of protons and helium (Biermann et al.
 2010) or inhomogeneous abundance of elements in superbubbles (Ohira & Ioka 2011)., 2010) or inhomogeneous abundance of elements in superbubbles (Ohira $\&$ Ioka 2011).
" For our present study, we assume that CRs are injected into the Galaxy with different source indices."," For our present study, we assume that CRs are injected into the Galaxy with different source indices."
 The index γ for an individual species is chosen such that (+a) is equal to the spectral index of the background obtain from the fit., The index $\gamma$ for an individual species is chosen such that $(\gamma+a)$ is equal to the spectral index of the background obtain from the fit.
" Before illustrating our results, we briefly discuss the choice of other model parameters involved in our calculations."," Before illustrating our results, we briefly discuss the choice of other model parameters involved in our calculations."
" Typically, tsea has values between ~(100—10?) yr depending on the gas density of the interstellar medium, mass of the ejecta and the energy output of the supernova explosion."," Typically, $t_{sed}$ has values between $\sim (100- 10^3)$ yr depending on the gas density of the interstellar medium, mass of the ejecta and the energy output of the supernova explosion."