ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 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."3 The uncertainties in relative positions are typically 10 mas., The uncertainties in relative positions are typically 10 mas.4" 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.5 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.6 comm.)., comm.).7 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.8 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 $^{-1}$.9 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.10 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 $-37^\circ$.11" 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.12 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.13 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.14 3C84 was used as the bandpass and amplitude. calibrator., 3C84 was used as the bandpass and amplitude calibrator.15 Its amplitude at the time of the, Its amplitude at the time of the16a 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).17 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."18 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."19 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 $\alpha_{\rm cond}=7$ .20 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.21 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."22 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."23 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."24 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.25 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}) ).26" 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."27" 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 $\eho=0.0$ ).28 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."29 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 $\beta_{\rm30SW}=1$, 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."31 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."32 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."33 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."34 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."35 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.36 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 $\sigma_8$."37 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)."38 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."39 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.40 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.41 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).42 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 $\epsilon_{\rm43SW}\approx 6$ ) and high conductivity $\alpha_{\rm cond}\approx 30$ ) for a model with $\sigma_8=0.93$."44 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 $\alpha_{\rm45cond}\approx 25$."46 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.47 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."48 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.49 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.50 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.51 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-standing52For 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 $^{-1}$."53 The simulations show a predominantly two-level bright frout., The simulations show a predominantly two-level bright front.54 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."55 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; )."56 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.57 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 )."58 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 )."59 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."60 The bright frout is approximately stationary af this time diff.mov)., The bright front is approximately stationary at this time ).61 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.62 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."63 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."64 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."65" 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."66 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."67 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."68 We note that a simular approach was also used in Patsourakosetal.(2009).., We note that a similar approach was also used in \cite{Patsourakos09a}.69 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."70 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}. .71 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."72 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.73 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}) )."74 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}) )."75 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."76 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."77 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.78 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.79 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. ("80The reader is referred to movie CO,The reader is referred to movie ).81R1joint1.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. ("82The 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 )."83 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.84 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.85 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}) ).86 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 1.1.87 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.88 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."89 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."90 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."91 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."92" 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)."93 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.94 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.95 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."96 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,"97system 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."98 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."99. 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."100 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."101 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."102spherical.,spherical.103 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."104 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."105 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."106 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."107 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.108 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."109 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.110 There are (wo main problems with this., There are two main problems with this.111 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."112 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."113 It is not clear what temperature to use., It is not clear what temperature to use.114 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)."115 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."116 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."117 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.118 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."119 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)."120 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"121The 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 $\mu \geq \mu_{\rm MM} = 1/4$ ).122 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.123" 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)."124 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 $z = 0.5$.125" 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 $r_{\rm p}^{\rm max}$."126 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}. .127" 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 $z = 0.5$."128" 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."129" 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}) )"130 does not bias our results with the following test., does not bias our results with the following test.131 We define a companion sample with Mg€—17.17—2.8z., We define a companion sample with $M_B \leq -17.17 - 2.8z$.132" 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 $z_{\rm r,2}$."133" 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."134 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.135 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.136" 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."137" 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."138" 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."139 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}.140 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 $q = 0.95\pm0.20$.141" 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."142. Westudy this evolution in more details in Sect. ??.., Westudy this evolution in more details in Sect. \ref{ffmmevol}. .143" 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."144" 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}) )."145 We assume wm=1/4 and Umm= 1/10., We assume $\mu_{\rm MM} = 1/4$ and $\mu_{\rm mm} = 1/10$ .146" 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$z = 0.5$ ."147" 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"148Grave&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.149 For fitting purposes a error was assumed on all fluxes., For fitting purposes a error was assumed on all fluxes.150" 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."151 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.152 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."153 The two straight lines in Fie., The two straight lines in Fig.154 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 7.155 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."156 Star 9 was not detected in the MIPS 24 yam band., Star 9 was not detected in the MIPS 24 $\mu$ m band.157 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."158 In Fie., In Fig.159 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.160 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."161 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.162 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."163 Also. the rich cluster of stars seen in Fig.," Also, the rich cluster of stars seen in Fig."164 1 is not visible in this composite image., 1 is not visible in this composite image.165 We will comment further on this issue in Sect., We will comment further on this issue in Sect.166 4., 4.167 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."168 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.169 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."170 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."171 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."172 307 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.173 The ΕΕΤΤ 77A. photometric data.," The UFTI $JHK$ photometric data,"174les 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."175 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)."176 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}.177. 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."178 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. ("1795.6.14a) in NT13.,5.6.14a) in NT73.180 Equation (7)) can be re-written as Lr PPA g," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"181 Equation (7)) can be re-written as Lr PPA ga," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"182 Equation (7)) can be re-written as Lr PPA gau," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"183 Equation (7)) can be re-written as Lr PPA gaus," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"184 Equation (7)) can be re-written as Lr PPA gauss," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"185 Equation (7)) can be re-written as Lr PPA gauss.," Equation\ref{mfdyn0}) ) can be re-written as 10^8 ,"186We 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}) ).187 For visual simplicity. we use the following denotations.," For visual simplicity, we use the following denotations."188 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"189"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."190 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.191. 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 $\sim0.4""$."192 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.193" 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""."194 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. (195"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."196 2003) and the Digitized Sky Survey with the improved X-ray position., 2003) and the Digitized Sky Survey with the improved X-ray position.197" 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."198 Monet et al., Monet et al.199 2003)., 2003).200 This confirms the result reported by Weisskopf et al. (, This confirms the result reported by Weisskopf et al. (2012006).,2006).202" 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."203 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.204" 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.205 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.206" 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."207 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.208". 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."209 Adopting the best-fit spectral parameters (cf., Adopting the best-fit spectral parameters (cf.210" Tab. 1)),"," Tab. \ref{spec_par}) ),"211 we found the count rates obtained from different detectors are consistent., we found the count rates obtained from different detectors are consistent.212" 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."213" 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 $^{-2}$ $^{-1}$."214" 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."215"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."216 Following Elsner et al. (, Following Elsner et al. (217"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)."218" 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."219 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.220 The results are plotted as the curves in Figure 5.., The results are plotted as the curves in Figure \ref{ccd}.221" 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 $S=M=H=0$."222" 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."223 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.224 Its hardness is found to be too hard for a field star., Its hardness is found to be too hard for a field star.225" 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."226 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.227 The same inference is obtained from the spectral analysis (see below)., The same inference is obtained from the spectral analysis (see below).228" 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."229 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.230" 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."231 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}. .232 All, All233accretion flow.,accretion flow.234" 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."235" 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)."236 Thus Py characterizes the gas temperature in our simulations., Thus $R'_S$ characterizes the gas temperature in our simulations.237 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.238" 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 $r= R'_C R_B$ )."239≺∢∐∎∐↲≼⇂≀↧↴⊳∖⊽↓⋟∪∐∪∖∖⇁⋟∖⊽⋅∶∖↥⊔∐↲↕↽≻∪↥≼↲⋟∖⇁⋅⊔⋅≼↲⋅⋅∣, The boundary conditions are specified as follows.240⇥∶∩≓≀↧↴∐≺⇂⊥⋖↽∖⋚∪≓ ∖∖⇁≼↲≀↧↴↕↽≻↕↽≻↥⋡∖↽≀↧↴∐≀↧↴⇀↸↕⋟∖⊽−∪↓⋟−⋟∖↕∖↽↕⊔∐∐↲↥↕⋅⋡∖↽∣↽≻∪∏∐≺⇂≀↧↴↕⋅∡∖↽≺∢∪∐≼∐∐∪∐⋅↼≚↥∣↽≻∪⊔↥⊔∐↲↕," At the poles, (i.e., $\theta=0^\circ$ and $180^\circ$ ), we apply an axis-of-symmetry boundary condition."241∐∐≼↲↕⋅≀↕↴↕∐⇂∪⋯≼↲↕⋅↕⋅≀↧↴≺∐≀↧↴↥ ∣↽≻∪∏∐≼⇂≀↧↴↕⋅↥≼↲⋝∖⊽⋅∖∖↽≼↲≀↧↴↕↽≻↕↽≻↥⋡∖↽≀↧↴∐∪∏⋯∪∖∖↽∣↽≻∪∏∐≼⇂≀↧↴↕⋅∡∖⇁≺∢∪∐≼∐∐∪∐↓⋟∪↕⋅≀↧↴∐≼⇂∡," At both the inner and outer radial boundaries, we apply an outflow boundary condition for all dynamical variables."242"∖↽∐≀↧↴∐↓↕≺∢≀↧↴↥∖↽≀↧↴↕⋅↕≀↧↴∣↽≻↥≼↲⊳∖⇁⋅↼≚⊳∖⊽↕↥ ↕↴∐∪⋮⋅↽⊰≀↧↪↥∪↕⋅≼↲↕↽≻↕⋅≼↲⋟∖⊽≼↲∐↥⋟∖⊽∩↲≀↧↴≼⇂⋡∖↽≺∢∪∐≼∐∐∪∐⋟∖⊽≀↧↴↥⊔∐↲∪∏∩↲↕⋅↕⋅⋯∐≀↧↴↥∣↽≻∪∏∐≺⇂≀↧↴↕⋅⋡∖⇁⋅≺⇂⋯⋅↕∐≸↽↔↴⊔∐↲≼↲∖↽∪↥∏∐∪∐∪↓⋟ 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."243" Note that we allow v, to float.", Note that we allow $v_r$ to float.244 To solve eqs.( 1))-( 3)), To solve eqs.( \ref{eq:con}) )-( \ref{eq:en}) )245 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."246 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 $10^{-3}$ )."247 The simulations for + —1.2 were performed only for RG=107., The simulations for $\gamma=$ 1.2 were performed only for $R'_S = 10^{-3}$.248 Summary. of all runs is presented in Tab 1.., Summary of all runs is presented in Tab \ref{tab:1}.249" 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"250The 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.251 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.252 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."253 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}."254 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}."255Given 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."256 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."257 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}."258 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.259 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.260 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."261 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."262 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."263 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.)."264 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.265 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."266 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.267 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.268 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.269 The paper is structured as follows., The paper is structured as follows.270 In Sect., In Sect.271 2 we describe HARPS anc the data-sets used in our investigation., $2$ we describe HARPS and the data-sets used in our investigation.272 Section 3 describes the principles of our method and data reduction., Section $3$ describes the principles of our method and data reduction.273 The results are presented 1n Sect., The results are presented in Sect.274 4 and discussed in Sect. 5., $4$ and discussed in Sect. $5$.275 We conclude in Sect., We conclude in Sect.276 6 with the lessons to learn from this campaign., $6$ with the lessons to learn from this campaign.277 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.278 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.279 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.280 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 $\AA$.281Thespectralresolutionwasmeasuredasbeingo 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).282 The sampling ts of 3.3 pixel per resolution element., The sampling is of 3.3 pixel per resolution element.283 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."284 A Th-Ar emission lamp is used for wavelength calibration., A Th-Ar emission lamp is used for wavelength calibration.285 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.286 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."287 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"288"and in turn. that 54=#03dxra) where for small enough5 ZZ. and Ü£,.","and in turn, that $\gamma_{\perp} = \pm (\beta \pm289 i\alpha )$ where for small enough $\Enu$ and $\Eeta$."290 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."291 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."292 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."293 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"294outflow. 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."295 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.296 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."297 1992)., 1992).298 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."299 JEG is also supported by Junta de Andaluctaa CE1C-126)., JFG is also supported by Junta de a (TIC-126).300 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.301with svuthetic colors of hydrogeu-deficicut stars. Evres et al. (,"with synthetic colors of hydrogen-deficient stars, Eyres et al. ("3021998) derive 1.15 from the observed versus expected II? flux.,1998) derive 1.15 from the observed versus expected $\beta$ flux.303 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."304 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."305 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."306 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)."307 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)."308 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)."309 Chemical abundances derived bv Asplund et al. (, Chemical abundances derived by Asplund et al. (3101997) are used as “normal input” for VI331 Ser.,1997) are used as “normal input” for V4334 Sgr.311 Asplund et. al. (, Asplund et al. (3121999) 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."313 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. ("3141997).,1997).315 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.316 Results are given in section 3.1., Results are given in section 3.4.317 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."318 Constauts for IDE computations were taken mainly from Tsuji (1973)., Constants for IDE computations were taken mainly from Tsuji (1973).319 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).320 The atomic line list was taken from the VALD database (Piskunov etal., The atomic line list was taken from the VALD database (Piskunov etal.321 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."322 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).323 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.324 By definition. JOLA overestimates molecular absorption produced by wea- nolecular bauds.," By definition, JOLA overestimates molecular absorption produced by weak molecular bands."325 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)."326 We used the DIGEL program of Yaremchuk (Nersisvan et al., We used the BIGF1 program of Yaremchuk (Nersisyan et al.327 1987) which realizes the method of Nameushikov et al. (, 1987) which realizes the method of Kamenshikov et al. (3281971).,1971).329 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."330 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}, 331v^{\prime\prime} \leq 9 $ (see Abia et al."332 1999 aud Pavleuko Yakovina 1999 for more details)., 1999 and Pavlenko Yakovina 1999 for more details).333 Couvection was computed using the ATLASS (Ixurucz 1993) scheme with 7/77 , Convection was computed using the ATLAS9 (Kurucz 1993) scheme with $l/H = 1.6$.334Convective overshooting was not considered., Convective overshooting was not considered.335 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.336 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 $P$."337 The table Tos=FOLP) was used for the temperature COYICCion iu SAMO.," The table $\tau_{\rm ross} = 338f (T,P)$ was used for the temperature correction in SAM941."339 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."340 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.341 the computatious for R. CrB. DPoavleunko 1999).," the computations for R CrB, Pavlenko 1999)."342iis 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,343 e.g.][]{prp02}."344 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.345 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}."346 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 $R_{\rm L} = 0.15 \pm 0.01{\rm347 \,R}_{\odot}$."348 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."349" 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."350" 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."351 The fractional changes of these measurements are shown in Figure 5.., The fractional changes of these measurements are shown in Figure \ref{fig:period}.352" 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."353" 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."354" 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 $\dot{P_{\rm b}} = 1.8(3) \times 10^{-11}$ ."355" 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."356 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 $m_{\rm357 p}=1.8$ $_{\odot}$ and a companion mass $m_{\rm c}=0.05$ $_{\odot}$ for an inclination angle of $i=40^{\circ}$."358" For wwe tind £4,GN~7.5.101."," For we find $\dot{P_{\rm359 b}}^{\rm GW} \simeq -7.5 \times 10^{-14}$."360 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 $\dot{P_{\rm b}}$.361 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."362 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}."363 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."364 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.365 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."366" 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 $\dot{P}_{\rm b}$ : $(\dot{P}_{\rm b}/P_{\rm b})^{\rm367 acc}=(\dot{P}/P)\simeq -3 \times 10^{-18}$ $^{-1}$ ."368 This is, This is369about50%.. which leaves some residual instrumental response of the order of[,"about, which leaves some residual instrumental response of the order of."370σοι 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).371 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."372 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.373" 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. “"374PLAIT™. 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."375 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.376 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.377 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."378 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."379 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.380 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."381 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}."382 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."383 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}. ."384 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."385 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."386 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 $\beta = -2.8$."387 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.388 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.389 We now use the five-year release of the WMAP observations (?).., We now use the five-year release of the WMAP observations \citep{Hinshaw09}.390" 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 $\degr$."391 We obtained a mean spectral index of p=-2.92+0.25., We obtained a mean spectral index of $\beta = -2.92\pm0.25$.392 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."393 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."394 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."395 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.396 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."397 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.398 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."399 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."400 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.401 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.402 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 $\Delta\beta = \pm$ 0.1.403 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."404 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}. .405 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"406correct 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.407 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}408 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."409 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.410 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."411" 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."412 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}."413. “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}."414. 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."415 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.416" 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)."417 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."418 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}."419. 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."420 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] $=-4.4$."421 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}."422. 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."423 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"424Tot 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 $\lambda \sim 2\,(1+z)\,{\rm \mu m}$."425&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."426 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."427 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}) )"428" 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, $t^{\rm429E}_{\rm ob} \equiv 10^6 t^{\rm E}_{\rm ob\,6}$ 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 $R_{\rm sub}$."430 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 $\alpha\sim1$."431 This is quite close to the spectral index of the observed afterelows., This is quite close to the spectral index of the observed afterglows.432" 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 $10^3\Delta t^{\rm433OT}_{\rm ob,3}$ 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."434" 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, $F^{\rm OT}_{\nu_{\rm ob}} = 435L_{\nu} f_{\rm ns} (4 \pi)^{-1} D_{\rm A}^{-2} (1+z)^{-3}$ , 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."436" For CRD 950326. an excess Ro fux FEO.LpJy was nmieasnred a tine fh,—20 d 19993)."," For GRB 980326, an excess R flux $F^{\rm E}_{\nu_{\rm ob}}[0.65{\rm \mu437m}]\sim0.4\,{\rm \mu Jy}$ was measured a time $t^{\rm E}_{\rm438ob}\sim20$ d )."439" 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 $R_{\rm440sub}\sim0.3(1+z)^{-1}{\rm pc}$ and $L\sim9\times10^{45}(1+z)^{-2}\;441{\rm erg\, s}^{-1}$."442 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."443 2)., 2).444" 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}) )."445" 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 $F^{\rm OT}_{\nu_{\rm446ob}}[0.65{\rm\mu m}] \sim200(1+z)^{-10}\;{\rm\mu Jy}$."447" 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 $F^{\rm OT}_{\nu_{\rm ob}}[0.65{\rm\mu m}] \sim10\;{\rm\mu448Jy}$ after 0.5 d), suffices to account for the observed excess after 20 d as a dust echo."449 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."450 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 $F^{\rm451OT}\propto t^{-2}$."452 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."453 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.454 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}) )."455 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}) )."456 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."457 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 $\sim0.1$ $_\odot$."458 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.459 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.460 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."461 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.462 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."463 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."464derived 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."465 sinibus:The problem stems frou au excess at high cucreics., The problem stems from an excess at high energies.466 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."467" 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."468 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."469 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}."470. 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}) ).471 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}."472(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.473 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)."474 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."475 Fieure Lo shows the correspouding color-color diagram., Figure \ref{ccd} shows the corresponding color-color diagram.476 Data were deadtimoe corrected. background subtracted. and N-ray bursts removed.," Data were deadtime corrected, background subtracted, and X-ray bursts removed."477 For each observation we calculated N-rav colors from the Standard? data., For each observation we calculated X-ray colors from the Standard2 data.478 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.479 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}) ).480 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.481 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.482 All previous bursts have been detected in harder states., All previous bursts have been detected in harder states.483 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 $^{-2}$ )."484 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 $^{-1}$ ."485 What causes burst oscillations is still not understood., What causes burst oscillations is still not understood.486 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.487 However. while this," However, while this"488GRBO030329/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."489 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}."490". 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}."491 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.492 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}."493. 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}."494. 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"495have incorporated this component.,have incorporated this component.496 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.497 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 $10^0498- 10^5$ keV background as a function of redshift."499" 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."500" 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 $\sim 20$."501". 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. ("502"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.)"503The paper is organized as follows., The paper is organized as follows.504" 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."505". 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."506" refsect:disc is a discussion of the implications of these results when taken collectively, and 855 is conclusion."," \\ref{sect:disc}507 is a discussion of the implications of these results when taken collectively, and 5 is conclusion."508" 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."509" 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."510" 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”."511" 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},"512 along with the data of mentioned above., along with the data of mentioned above.513" 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 $z>4$."514 We also implement an X-ray background., We also implement an X-ray background.515" 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 $EF_E$."516" 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."517" 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."518. 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 $z=0$.519" 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)."520 We call this the “New UV+X” model., We call this the “New UV+X” model.521 Notice in Fig., Notice in Fig.522 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}523 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."524. 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.525" 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."526" 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)."527" 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"528imply 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."529 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."530 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.531 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."532 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.533 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."534 Why is there a minimum mass in super-massive black holes?, Why is there a minimum mass in super-massive black holes?535 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 $10^{6} \, M_{\odot}$."536 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}."537 [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."538 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}."539 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."540 ? 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."541 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."542" 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."543 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."544" 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."545" 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 $N_{\star}$."546 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."547 ? 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."548 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.549 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."550 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."551 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."552 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?"553 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?"554 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."555 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."556 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}: $557\rho_{dm} \, \sim \, (x^{\gamma_{sp}} ( 1+ x)^{2})^{-1}558$ 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}."559" 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 $M_{R, dm} \, \sim \, R^{3 - \gamma_{sp}}$."560" 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 $\sigma_{\star} \, \sim \, x^{(2 - \gamma_{sp})/2}$ ."561The 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."562 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."563 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"564target galaxies in our program will be measured.,target galaxies in our program will be measured.565 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."566 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.567 This work deals with Cepheids in the Large Magellanic Cloud (LMC)., This work deals with Cepheids in the Large Magellanic Cloud (LMC).568 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.569 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}."570 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."571 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)."572 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."573 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."574 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."575 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."576 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."577 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.578 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}.579 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."580 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.581 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."582 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."583 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""."584 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}.585 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}.586 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}."587 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}. .588 Section 7 provides a summary., Section \ref{sec:summary} provides a summary.589"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 $dt$."590 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."591 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."592 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."593 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."594 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."595 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.596 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."597 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."598 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}) )"599 and the medium and low mass samples (Fig. 4))., and the medium and low mass samples (Fig. \ref{fig:resl}) ).600" 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."601 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 $\Delta t \approx 1 {\rm Gyr}$."602 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."603 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."604 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.605 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."606 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.607 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."608 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."609 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 of610he 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}.611 The predicted dux a energies X0.5 keV is 10 o LOO times lower than the observed flux.," The predicted flux at energies $\lesssim6120.5$ keV is $10$ to $100$ times lower than the observed flux."613 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.614 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."615 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."616 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.617 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.618 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}."619" 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."620 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 $I_{\rm X}\approx6215\times 10^{-10}\ {\rm ergs\ s^{-1} cm^{-2} sr^{-1}}$ (see his Table 1)."622" 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}) ):623107.2IkeVostemJyκο|: at chereies 25 keV (ie. zmOL of the total ARB at hese energies: sec ?)) the fux decreases steeplv at higher cucreics.," $10^{-2}-10^{-1}624{\rm keV\ s^{-1} cm^{-2} sr^{-1} keV^{-1}}$ 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."625 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."626 However. the cussion is pateliv auc is far roni beime uniforin.," However, the emission is patchy and is far from being uniform."627 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 $H_{\rm SC}\approx 0.25R_{\rm SC}$."628 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.629 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.630 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."631 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, $n_e\sim63210^{-5}-10^{-3}{\ \rm cm^{-3}}$ , regions within and around groups and clusters."633 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}. ."634The 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)."635 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.636 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 $\sim 1-10^{\circ}$."637 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."638 The kinetic SZ effect has simular inagnitude., The kinetic SZ effect has similar magnitude.639 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.640" 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."641 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.642 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"643The 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."644 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."645 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."646 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."647 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."648 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.649 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.650 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.,"651 in prep)., in prep).652 Iu Fig. L.," In Fig. \ref{f:fig4},"653 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."654 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}."655.. 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},"656 bottom pancl). as expected from their uiprocessed Chemical abiunudauces.," bottom panel), as expected from their unprocessed chemical abundances."657 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."658" 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},"659 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."660 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},"661 middle panel)., middle panel).662 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.663 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.664 Instead. the Ca abunudauices closely track those of Fe.," Instead, the Ca abundances closely track those of Fe."665 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.666 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."667 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?668 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."669 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."670 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."671 This linütatiou has finally been overcome by our study., This limitation has finally been overcome by our study.672 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."673 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 $\alpha-$."674", 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}) )."675 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."676 Finally. the dSpl moreed with the Milkv. Way.," Finally, the dSph merged with the Milky Way."677 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."678 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."679 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."680 We will rely on the observed chemistry., We will rely on the observed chemistry.681 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."682 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).683 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. ("684i) 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.685 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 normal686(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}.687" 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."688 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.689 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."690 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."691 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}.692 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}."693 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."694 It was determined that the light curves do not differ significantly., It was determined that the light curves do not differ significantly.695 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."696 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."697 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."698 The remaining would be primarily diffuse Galactic ganmia-ray enission., The remaining would be primarily diffuse Galactic gamma-ray emission.699 This is the estimated background for the EGRET lighteurve., This is the estimated background for the EGRET lightcurve.700 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.701 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}.702. 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"703Our 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.704 Ca is almost solar abundant., Ca is almost solar abundant.705 For C and Li we give only the upper limits., For C and Li we give only the upper limits.706 75(LIIUmi66385. {12 8161. 5840536 1102¢2632 is a spectroscopic binary star.," 6385, +12 3161, 84036, 102632, A1m) is a spectroscopic binary star."707 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 A1.708 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.709 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 $\vsini=33~$ $^{-1}$.710 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.711 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, $V_0=1.03~{\rm km\,s^{-1}}$, $\omega={\rm 114^{\circ}.86}$."712 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}) ).713 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."714 Ca is uncerabundant., Ca is underabundant.715 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.716 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.717 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 155375.718 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}) )."719 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."720 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."721 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.722 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.723 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."724 The angular separation between the components is 61.97., The angular separation between the components is ${\rm 61.9}\arcsec$.725 Both components of the binary system are Am stars., Both components of the binary system are Am stars.726 “Phere have been many determinationsof the stellar spectral class in the literature., There have been many determinationsof the stellar spectral class in the literature.727 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.728 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}) ).729 According to Abt&Carcona the star was ΕΣ. Cowl, According to \citet{ac84} the star was A4/F2V/F3.730eyctal.(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.731 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.732 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} - $\vsini=58~$ $^{-1}$ ."733opinally. Roveretal.(2002). determined the rotational velocity of the star as esing=68 kmss ," Finally, \citet{rgbgz02} determined the rotational velocity of the star as $\vsini=68~$ $^{-1}$ "734" 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}) )."735" 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.736" 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.737" 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}) )."738 '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."739" 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."740 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."741eached 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.742" 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 $t \sim 100\, \tau_A$."743"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."744 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."745" 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.746"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).747 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."748" the photosphere- mimicking boundary, and", The significance of their conclusions is then strongly diminished.749" 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.750 This verysimple, This is not a definition.751 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.752" recent simulations of ∙∙ In summary, the main results presented", Depending on the value of the perturbation the Parker angle so defined.753 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.754"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}."755 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}756 splits the domain into large and small scales."757sthe 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."758y ofthelarge X-pointsto linear magnetic takes withoutnon- , These laminar regimes correspond to our linear saturated diffusive regimes computed in \ref{par3}. .759"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.760 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."761ythe 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."762" 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}] ]."763rcingtoa 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."764ateisa simple and orderedapplied one-dimensional forcing. We have, Nevertheless itis useful to use \cite{hp92} model in order tounderstand why it is not applicable.765 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 simulated766 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 simulated767 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 simulated768 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 simulated769Thus 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."770 In Fig., In Fig.771 6 we show a trajectory with e=0.75., $6$ we show a trajectory with $e=0.75$.772 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.773 In Fig., In Fig.774 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 $q=2.6$."775 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.776 Clearly he length asvmumetry is completely disrupted: after. the second encounter., Clearly the length asymmetry is completely disrupted after the second encounter.777 In. Fig., In Fig.778 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."779 Thus as suspected only a single Lv-by can be allowed., Thus as suspected only a single fly-by can be allowed.780 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."781 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.782 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)."783 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."784 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. (7852000) report an age of 5 for the triple svstenm LID 141569.,2000) report an age of $5$ for the triple system HD 141569.786 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.787 Weinberger et al. (, Weinberger et al. (7882000) 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 $20$ )."789 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)."790 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.791 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.792 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."793 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.794 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.795 We summarise the coplanar results in Fig. 18.," We summarise the coplanar results in Fig. $18$,"796 and the inclined-orbit results in Fig. 19., and the inclined-orbit results in Fig. $19$.797 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).798 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."799 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.800" 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)."801" 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."802 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.803 LID 141569: Weinberecr et al., HD 141569; Weinberger et al.804 2000) that became, 2000) that became805 , 806dusty 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.807 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."808 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."809 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."810 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.811" 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 $\rho > 10^{-13}\,{\rm g/cm^2}$."812" 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 $i$."813 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."814" 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.815 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."816 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.817 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."818 A quantitative evaluation of the model images. however. is not so straightforward.," A quantitative evaluation of the model images, however, is not so straightforward."819 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.820 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."821 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.822" 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.823 —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."824" 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."825 In total. more than 0000 models were computed.," In total, more than 000 models were computed."826 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.827 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)."828 On the otherhand. we alsofoundseveral models that reproduced the near-infrared," On the otherhand, we alsofoundseveral models that reproduced the near-infrared"829be coustaut.,be constant.830 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.831 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."832 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."833 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.834 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: $\tau=(k_{\mathrm{IR,0}}/g)(P/P_{\mathrm{ref}})^\alpha$."835 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 $(\alpha+1)/4$."836" 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 $\alpha \geq 1/7$.837 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."838 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)."839 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: $(\alpha+1)/4 \geq R/c_p$ ."840 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.841For 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."842 1)., 1).843 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."844 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."845 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 $\mu$."846 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."847 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 $R = 10^7$ )."848 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 $\sim$ 3.849 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."850 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."851 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 $\mu$ ."852of 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}) )."853 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."854 Fig., Fig.855 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.856 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.857 “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.858 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."859 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)."860 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."861 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."862 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.863 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."864 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},"865 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."866 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}.867 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."868 The highest density of variables show periods xlow 12 hours., The highest density of variables show periods below 12 hours.869 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."870 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."871 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."872 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."873 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}."874 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)."875 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}) ),"876 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."877 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."878 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.879 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."880 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."881 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.882 Fig., Fig.883 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 variability884infrared.,infrared.885 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.886 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."887 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.888 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."889 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}.890. 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."891"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."892" 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."893 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}) )."894 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 $\beta$ )."895 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)."896 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.897 Our wavelength: bands are shown in Fig., Our wavelength bands are shown in Fig.898 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."899 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.900 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.901 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.902 These were identified and rejected through (heir FEWIIM.," These were identified and rejected through their FWHM,"903ancl the above arguments don't apply to them.,and the above arguments don't apply to them.904 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."905 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."906 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."907 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).908 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).909 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).910 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}."911.. A bin with less than five events is merged with its neighbor., A bin with less than five events is merged with its neighbor.912 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."913 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.914 Fie., Fig.915" 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."916" 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 ${dN}/{d\log917T}$ traditionally shown in such plots \citep[e.g.,][]{Kouveliotou93}."918. 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.919 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)."920 The extent of the platean varies slightly [rom one detector to another., The extent of the plateau varies slightly from one detector to another.921 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).922 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 $-4<\alpha<-3$."923 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.924 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."925 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)."926relativistic 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."927 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."928 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.929oxvgen-neon-WD.,oxygen-neon-WD.930 This Μα] example. and (he one presented by IIurley et al. (," This lurid example, and the one presented by Hurley et al. ("9312001). highlight an important point.,"2001), highlight an important point."932 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."933 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 $\sim 500$."934 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.935 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."936 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.937 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."938 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 104.939 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."940 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.941 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.942 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."943 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."944 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."945 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.946 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."947 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."948 Here we have to be careful with the terminology used., Here we have to be careful with the terminology used.949 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 a950(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}) )."951 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."952 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}) ).953 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."954 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}) )."955 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.956 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.957 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.958" 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.959 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}) )."960 In 82. we brielly describe the data and the analvsis method.," In 2, we briefly describe the data and the analysis method."961 We present and discuss the results in 33., We present and discuss the results in 3.962 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 = ) = _*"963"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."964 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."965 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."966 The procedure is described in more detail by Djuraseviéetal. (2004)., The procedure is described in more detail by \citet{djura04}.967. 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."968 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}.969 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."970 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."971 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.972 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}) ).973 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."974 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."975 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 $R_h\approx5 R_{\odot}$ )."976" 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 $d_e\approx1.6 R_{\odot}$ ."977 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},"978 with the temperature profile exponent ap=6.5., with the temperature profile exponent $a_T=6.5$.979 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.980 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)."981 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."982 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 $0^\circ-360^\circ$."983 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. $T_{hs}\approx 8700 K$."984" 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."985" 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."986 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.987 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."988" 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."989 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.990 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}."991.. 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.992 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 $\lambda_{bs1}\approx 170^\circ$ ."993This 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."994 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 $\lambda_{bs2}\approx 50^\circ$."995 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."996 It can be interpreted as the second spiral arm in the disk., It can be interpreted as the second spiral arm in the disk.997 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."998 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."999 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."1000 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."1001 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."1002 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.1003 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}. ."1004 Our previous investigationsof anomalously high values of gravity, Our previous investigationsof anomalously high values of gravity1005roughly 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}."1006. 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}."1007. 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.1008 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}."1009. 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.1010 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.1011 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.1012 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.1013 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}."1014.. 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}."1015. 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.1016 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:1017he 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."1018 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."1019" 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 $i$."1020 Model parameters for the different panels are oovided in Table 2.., Model parameters for the different panels are provided in Table \ref{tab2}.1021 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."1022 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.1023 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 $\vinf \approx10241000-3000$, 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."1025. 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.1026 There are four basic zones., There are four basic zones.1027 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."1028 In the cavity sector. there is no contribution.," In the cavity sector, there is no contribution."1029 Then there is the spherical zone and the compressed layer for rrog," Then there is the spherical zone and the compressed layer for $r\ge r_{\rm1030orb}$."1031 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."1032 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 $\sigma u_0/u_{\rm c} =1033\sigma r_{\rm c}/r_0$ becomes: If the binary orbit is exceedingly large, $\Lambda =1$ is again recovered because the CWIR is displaced to a location of irrelevance with $r_{\rm orb} \gg r_{\rm c}$ ."1034 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"1035Brightness variation of V1154CCyg was discovered by Strohmeieretal.(1963).,Brightness variation of Cyg was discovered by \citet{stretal63}.1036. 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.1037 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}.1038". 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}."1039. 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.1040 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.1041 None of these previous data can compete withKepler in photometric quality., None of these previous data can compete with in photometric quality.1042" 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}."1043". 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."1044" 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."1045" 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 $v \sin i = 12.3\pm 1.6$ $^{-1}$."1046" 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}. ."1047 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.1048 They published a value of [Fe/H]= —0.10., They published a value of ${\rm [Fe/H]} = - 0.10$ .1049FOC46192 bands.,$^{12}{\rmn C}^{12}{\rmn C}~\lambda6192$ bands.1050 Another is the complete. absorption oetween andA... whatever its precise cause.," Another is the complete absorption between and, whatever its precise cause."1051 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.1052 Fig., Fig.1053 I1. shows the relationship »etween these two measures., \ref{c_index} shows the relationship between these two measures.1054 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.1055 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."1056 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.1057 12 which is Fie., \ref{slopes_c} which is Fig.1058 4. redrawn this time with svmbols showing 10 c-index., \ref{slopes} redrawn this time with symbols showing the c-index.1059 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)."1060 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.1061 lig., Fig.1062 13 shows the c-index polted against (JJfy) emperature), \ref{c_t} shows the c-index plotted against $(J-K)$ (temperature).1063 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.1064 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.1065 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."1066 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.1067 Fig., Fig.1068 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.1069 Phe svmbols mark the (JA) colour., The symbols mark the $(J-K)$ colour.1070 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).1071 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."1072 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 spread1073(see AAVSO website).,(see AAVSO website).1074 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."1075 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."1076 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.1077 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.1078 Emission lines can thus help in determining the properties of this area., Emission lines can thus help in determining the properties of this area.1079 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.1080 Its temporal evolution follows the evolution of, Its temporal evolution follows the evolution of1081curve’ of the emission lines.,curve' of the emission lines.1082 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."1083 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."1084 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."1085 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."1086 1200)., \ref{cartoons}b b).1087 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."1088 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.1089 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."1090 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))."1091 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.1092 13ec)., \ref{cartoons}c c).1093 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."1094 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 PKS2250-41.1095 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.1096 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."1097 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."1098 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 $L_{IR}1099\sim 2\times 10^{11}L_{\odot}$ (Sanders Mirabel 1996)."1100 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."1101 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."1102 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.1103 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."1104 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.1105 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."1106 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."1107 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)."1108 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."1109 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."1110 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."1111 However. the similarities end there.," However, the similarities end there."1112 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."1113 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."1114 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."1115 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."1116 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."1117 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."1118 The key results are as follows:, The key results are as follows:1119therein).,.1120 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}."1121. 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."1122 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.1123 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."1124" 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."1125 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}.11261904)... 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.1127 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."1128 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)."1129 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."1130 “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}. ."11312006).., In Fig.1132 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.1133 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}."11342PO00).. 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."1135 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."1136 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}.1137. 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}. ."1138. 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."1139 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."1140 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."1141 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."1142 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.1143 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."1144 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."1145 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."1146 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."1147 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."1148 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."1149 The paper is organized as follows., The paper is organized as follows.1150 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.1151 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).1152 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 ."1153. 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 2.1154 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."1155 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 $^{-1}$ )."1156 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."1157 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."1158 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."1159 For wavelengths longer than 6500 A.. telluric," For wavelengths longer than 6500 , telluric"11600«z<3.,$0<z<3$.1161" 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."1162" 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."1163 Only masses above 2x10?A-!Mo were considered as clusters.," Only masses above $2 \times116410^{13}h^{-1}M_{\sun}$ were considered as clusters."1165 The observations include many systematic errors that are not in the simulation., The observations include many systematic errors that are not in the simulation.1166" 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."1167 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.1168" 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."1169" The line-of-sight pair selection criterion, |Azpnot|< 0.015, was chosen to balance competing considerations."," The line-of-sight pair selection criterion, $|\Delta1170 z_\mathrm{phot}| < 0.015$ , was chosen to balance competing considerations."1171" 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}1172 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."1173" 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."1174" 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."1175" 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)."1176 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.1177 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.1178" 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 $\sigma(z_\mathrm{phot}) =11790.015$ (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}) )."1180" 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 $\langle1181\cos^2\theta_{c,p}\rangle_{\Lambda\mathrm{CDM}}$, 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 $\xi\ll 1$."1182" When clustering is significant, the relative contamination becomes smaller."," When clustering is significant, the relative contamination becomes smaller."1183" 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."1184" 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."1185 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.1186" Fortunately, we also measure the projected correlation function, w(R) +1 = catalogue —"," Fortunately, we also measure the projected correlation function, w(R) +1 = = ."1187" 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 $\Gamma(R)=\beta=0.35$ ."1188" On the smallest scales,"," On the smallest scales,"1189",, and we select this for the ensemble flux.",", and we select this for the ensemble flux."1190" 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."1191" 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."1192" 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."1193" 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."1194 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.1195" 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."1196" 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."1197 This approach is since we are now using an incorrect spectral model., This approach is since we are now using an incorrect spectral model.1198" 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."1199" 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."1200 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}.

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