ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Therefore the total IR fluxes eiven by Flores et al., Therefore the total IR fluxes given by Flores et al.3 have been divided by a factor 2 in order to roughly represeut the flux at GO gan. The UV cnussion is taken at 0.28 juu as given bv Flores et al., have been divided by a factor 2 in order to roughly represent the flux at 60 $\mu$ The UV emission is taken at 0.28 $\mu$ m as given by Flores et al.4 It is difficult to estimate a correction factor to translate the UV data to 0.2;10 in the absence of observations of a large sample of galaxies at both waveleugths since the ratio depends on the star formation history aud tιο dust extinction., It is difficult to estimate a correction factor to translate the UV data to $\mu$ m in the absence of observations of a large sample of galaxies at both wavelengths since the ratio depends on the star formation history and the dust extinction.5 We can try to use svuthesis nodds for this estimate: assimuiuse a constant star formation rate over l Cr and using the models of Leitherer et al. (1999)), We can try to use synthesis models for this estimate: assuming a constant star formation rate over 1 Gyr and using the models of Leitherer et al. \cite{leitherer}) )6 for a solar iuetallicity we fiud Fu.2/Fues=1.7 the fux being defined as A+ fj.The difference of extinction between 0.28 aud 0.2 gam has been calculated using the extinction curves of the Milkv. Way and the LAIC (Pei 1992)) and that of Calzetti (1997)).," for a solar metallicity we find $\rm 7F_{0.2}/F_{0.28}=1.7$ the flux being defined as $\rm \lambda\cdot 8f_{\lambda}$ .The difference of extinction between 0.28 and 0.2 $\mu$ m has been calculated using the extinction curves of the Milky Way and the LMC (Pei \cite{pei}) ) and that of Calzetti \cite{calzetti}) )."9 The ratio Ayo/Ageoxs=1.2L1., The ratio $\rm A_{0.2}/A_{0.28}=1.2-1.4$.10 Therefore the two effects (star formation and extinction) roughly compensate cach other and we do not perform amy correction between 0.28 and 0.2 gan. The galaxies are plotted in figure 6 as crosscs or the true starbursts and empty circles for the Sevterts., Therefore the two effects (star formation and extinction) roughly compensate each other and we do not perform any correction between 0.28 and 0.2 $\mu$ m. The galaxies are plotted in figure 6 as crosses for the true starbursts and empty circles for the Seyferts.11 Thev all fill the eap between the IRAS/FOCA sample aud he ULICs., They all fill the gap between the IRAS/FOCA sample and the ULIGs.12 Therefore. they are not as extreme as Ες mut their extinction is larger than the nearby galaxies of he IRAS/FOCA sample.," Therefore, they are not as extreme as ULIGs but their extinction is larger than the nearby galaxies of the IRAS/FOCA sample."13 For the galaxies classified as stirbursts we have tentatively estimated this extinction roni their FIR to UV fiux ratio.," For the galaxies classified as starbursts, we have tentatively estimated this extinction from their FIR to UV flux ratio."14 The extinctions found span from 2 to 5.5 wili a 1iean at 2.3 mae (and a median at 3 mag)., The extinctions found span from 2 to 5.5 with a mean at 3.3 mag (and a median at 3 mag).15 This is much. lavecr than that estimated by Flores et al., This is much larger than that estimated by Flores et al.16 bv matching the global star formation rates deduced from the total FIR aud the UV Iuuinosities iu the observed field: they find extinctions around 2 mae at 0.28 ju. This discrepancy between the οκοΊο occurring in individual galaxies selected in infrared aud that deduced from the total FIR aud UV. Iuniünositv of a selected field (i.c. the stun of the Iuninositv of all galaxies detected im he waveleneth baud (FIR or UV)) is well illustrated iu he table 5 of Flores et al., by matching the global star formation rates deduced from the total FIR and the UV luminosities in the observed field: they find extinctions around 2 mag at 0.28 $\mu$ m. This discrepancy between the extinction occurring in individual galaxies selected in infrared and that deduced from the total FIR and UV luminosity of a selected field (i.e. the sum of the luminosity of all galaxies detected in the wavelength band (FIR or UV)) is well illustrated in the table 5 of Flores et al.17 where the ratio Li/L(O.28;421) calculated for individual objects observed at both 15 and 128 jun ds ~5 times larger than the ratio of the global tuuinositics IR and UV huninosities in the CERS field., where the ratio $\rm L_{IR}/L(0.28\mu m)$ calculated for individual objects observed at both 15 and 0.28 $\mu$ m is $\sim 5$ times larger than the ratio of the global luminosities IR and UV luminosities in the CFRS field.18 This is in ull agreement with our own results preseuted 1n section , This is in full agreement with our own results presented in section 4.19We lave constructed a sample of 102 nearby galaxies detected by IRAS at GO sau and for which UV observations at 0.2 gan are available down to muy~17.18.," We have constructed a sample of 102 nearby galaxies detected by IRAS at 60 $\mu$ m and for which UV observations at 0.2 $\mu$ m are available down to $\rm 20m_{UV}\sim 17-18$."21 Ouly five galaxies have no UV detection inplviug au extinction larger than 2-3 mae for these objects which are also very faint in FIR., Only five galaxies have no UV detection implying an extinction larger than 2-3 mag for these objects which are also very faint in FIR.22 The FIR aid UV properties of our sample have becu compared to the mean properties of the local Universe deduced: from. the hpuuimositv fictions aud densities at both wavelenghs., The FIR and UV properties of our sample have been compared to the mean properties of the local Universe deduced from the luminosity functions and densities at both wavelengths.23 As the galaxies become xiehter in FIR their FIR to UV flux ratio. Le. their extinction increases: d(log(Eoo/Lu.»D/d(ülosLey)&0.3 which translates to au increase of ~1.5 imag for the dust extiucion iu UV per decade of FIR Iuninosity.," As the galaxies become brighter in FIR their FIR to UV flux ratio, i.e. their extinction increases: $\rm d(\log( L_{60}/L_{0.2}))/d(\log 24L_{60})\simeq 0.3$ which translates to an increase of $\rm \sim 0.5$ mag for the dust extinction in UV per decade of FIR luminosity."25 The ratio of the FIR to UV local Iuuinositv densities is much lower than that found iu mdividual galaxies., The ratio of the FIR to UV local luminosity densities is much lower than that found in individual galaxies.26 It is also true for other samples of nearby galaxies usually considered as low redshift templates like the IUE sample, It is also true for other samples of nearby galaxies usually considered as low redshift templates like the IUE sample27Deuterium-bearimg molecules have become the target of many observations in recent vears and several models have been developed to account for them (e.g.Tielens1983:Millar 2000a.b)..,"Deuterium-bearing molecules have become the target of many observations in recent years and several models have been developed to account for them \citep[e.g.][]{tielens83,roberts00a,roberts00b}."28 Twenty-six such molecules have been detected to date in interstellar clouds., Twenty-six such molecules have been detected to date in interstellar clouds.29 The interest in this topic lies in (he unusual chemistry at work im the cold regions, The interest in this topic lies in the unusual chemistry at work in the cold regions30The soft gamma-ray repeaters (SCs) showcase [lux variability. on many cilferent timescales.,The soft gamma-ray repeaters (SGRs) showcase flux variability on many different timescales.31" Phe quiescent state. with persistent X-ray emission (Ly~10ergs Ly, punctuated by numerous sporadic short bursts of ganmiunia-ravs. with peak luminosities up to 107ergs and tvpical duration in the range ~0.011 s. mark the defining characteristics of SCGlts (see Mereghetti2008. and Woods&Thompson2006. [or à review)."," The quiescent state, with persistent X-ray emission $L_X\sim10^{35}\mbox{ erg s}^{-1}$ ), punctuated by numerous sporadic short bursts of gamma-rays, with peak luminosities up to $\sim 10^{42}\mbox{ erg s}^{-1}$ and typical duration in the range $\sim 0.01 - 1$ s, mark the defining characteristics of SGRs (see \citealt{Mereghetti2008} and \citealt{WoodsThompson2006} for a review)."32 Out of the seven conirmed SCAR sources. SCAR. 0525-66. σαν 1806-20. SCAR 190)|14. SGIU 1627-41. SCAR 1150-5418. SCR. O4IS|572). SG1 0501|4516 (with the last three acced recently to the SGto family: see Jxanekoetal.2010:vanderHorst.201:Ixumarctal. 2010)). the first three have been known to¢mt giant flares.," Out of the seven confirmed SGR sources, SGR 0525-66, SGR 1806-20, SGR 1900+14, SGR 1627-41, SGR 1150-5418, SGR 0418+5729, SGR 0501+4516 (with the last three added recently to the SGR family; see \citealt{Kanekoetal2010,Horstetal2010,Kumaretal2010}) ), the first three have been known to emit giant flares."33 A rare phenomenon compared to the commonlv occurring short bursts. the giant. Lares unleash a stupendous amount of⋅ energy (~10771l erg) in.. gamma-ravs over à timescale of ~0.20.5 s in a fast rising initial peak.," A rare phenomenon compared to the commonly occurring short bursts, the giant flares unleash a stupendous amount of energy $\sim 10^{44}$ erg) in gamma-rays over a timescale of $\sim 0.2-0.5$ s in a fast rising initial peak."34 The initial high energy. burst is followed. by a long (~200—400 s). exponcntially cecaving pulsating tail of hard. X-rav emission. the period. of which coincides with that of the rotation of the neutron star (INS).," The initial high energy burst is followed by a long $\sim 200-400$ s), exponentially decaying pulsating tail of hard X-ray emission, the period of which coincides with that of the rotation of the neutron star (NS)."35 Additionally. intermediate strength but rare outbursts lasting for few tens of seconds have been observed in the case of SCIL 1900|14.," Additionally, intermediate strength but rare outbursts lasting for few tens of seconds have been observed in the case of SGR 1900+14."36 ‘The first extremely energetic giant [are Tom a recurrent eamma-rayv source. SCAR 0525-66. was detecced on March 5. 1979 bv the gamma-ray burst. detector alxxwd the Venera 11 12 space probes and the nine interplanetary spacecraft of the burst sensor network (Mazetsetal.979:Lelfand&Long 1979).," The first extremely energetic giant flare from a recurrent gamma-ray source, SGR 0525-66, was detected on March 5, 1979 by the gamma-ray burst detector aboard the Venera 11 12 space probes and the nine interplanetary spacecraft of the burst sensor network \citep{Mazetsetal1979,HelfandLong1979}."37. Phe position of the source was found to be coincidentet with the supernova remnant N4ϱ ata distance of ~55 kpc in the Large Magellanic: Cloud., The position of the source was found to be coincident with the supernova remnant N49 ata distance of $\sim55$ kpc in the Large Magellanic Cloud.38 The [are consisted of a sharp rise (15 ms) to the peak eamma-pav luminosity. L-~IottCres subsequently followec bv an exponentially decaving tail with L«107ergs ," The flare consisted of a sharp rise $\sim 15$ ms) to the peak gamma-ray luminosity, $L_{\gamma}\sim10^{44}\mbox{ erg s}^{-1}$, subsequently followed by an exponentially decaying tail with $L_{\gamma}\sim10^{42}\mbox{ erg s}^{-1}$ ."39Remarkably. the initial burst only [lasted for ~0.1 s compared to the longer lasting (100 s) tail that pulsatec with a period of ~ Sos Clerrelletal.1980).," Remarkably, the initial burst only lasted for $\sim0.1$ s compared to the longer lasting $\sim100$ s) tail that pulsated with a period of $\sim8$ s \citep{Terrelletal1980}."40.. The tota emitted energy during the initial peak and the decaving tai amounted to an astonishing ~1014 erg., The total emitted energy during the initial peak and the decaying tail amounted to an astonishing $\sim10^{44}$ erg.41 An even more energetic [lare was detected. from SGR 1900]14 on August QT 1998 bv a multitude of space telescopes in the direction of a Galactic supernova remnant 42.8[0.6 (Llurlesetal. 1999a).. making it the second exceptionallyenergetic event. detected. in the pas century [rom a recurrent eamaa-ray source.," An even more energetic flare was detected from SGR 1900+14 on August 27, 1998 by a multitude of space telescopes in the direction of a Galactic supernova remnant G42.8+0.6 \citep{Hurleyetal1999a}, , making it the second exceptionallyenergetic event detected in the past century from a recurrent gamma-ray source."42 The burst hac, The burst had43ithium variations in the last 1 Cir.,lithium variations in the last 1 Gyr.44 Data for Galactic open clusters as well as old. elobular clusters are shown., Data for Galactic open clusters as well as old globular clusters are shown.45 The xe filled. cireles represent mean lithium. values. while the ig open circles represent the highest. lithium abundance observed in the cluster.," The big filled circles represent mean lithium values, while the big open circles represent the highest lithium abundance observed in the cluster."46 Different measurements of the ISM value are also shown., Different measurements of the ISM value are also shown.47 For 66397. the mean value for 12 urnoll stars corrected for non-LVLE and standard depletion is that by Bonifacio et al. (," For 6397, the mean value for 12 turnoff stars corrected for non-LTE and standard depletion is that by Bonifacio et al. ("482002).,2002).49 For T'Juc. the ‘Li abundance is the mean value for 2 turnoll stars determined » Pasquini Molaro (1997).," For Tuc, the $^7$ Li abundance is the mean value for 2 turnoff stars determined by Pasquini Molaro (1997)."50 For ALO92. the highest non-LYE Li abundance measured by Boesgaard et al. (," For 92, the highest non-LTE Li abundance measured by Boesgaard et al. ("511998) is shown.,1998) is shown.52 The age of all these clusters is taken from Carretta et al. (, The age of all these clusters is taken from Carretta et al. (532000).,2000).54 For 667. data ave from Pasquini. Rancdich Pallavicini (1997): for 7752. from Dalachandran (1995 average of the three stars with the largest Li abundances on the hot side of the Li dip).," For 67, data are from Pasquini, Randich Pallavicini (1997); for 752, from Balachandran (1995 – average of the three stars with the largest Li abundances on the hot side of the Li dip)."55" For 33680. we assume he highest. lithium. abundance from. Pasquini. Ranclich ""allavicini (2001): for the Livacles. the mean value is taken rom Dalachandran (1995 9 average of the three stars with he largest Li abundances on the cool side of the Li dip)."," For 3680, we assume the highest lithium abundance from Pasquini, Randich Pallavicini (2001); for the Hyades, the mean value is taken from Balachandran (1995 – average of the three stars with the largest Li abundances on the cool side of the Li dip)."56 The values for the Pleiades and a Per are from Socerblom et al. (, The values for the Pleiades and $\alpha$ Per are from Soderblom et al. (571993) and Itandich et al. (,1993) and Randich et al. (581998). respectively. while those or 2344 and 22602 are non-LVE values from Iandich et al. (,"1998), respectively, while those for 2391 and 2602 are non-LTE values from Randich et al. ("592001).,2001).60 For 44665 and 22264. non-L'TE values rom in Montes (1997) ancl Soclerblom et al. (," For 4665 and 2264, non-LTE values from n Montes (1997) and Soderblom et al. ("611999). respectively. are dcisplaved.,"1999), respectively, are displayed."62 “Phe value for σ Ori is the mean value from Zapatero Osorio ct al. (, The value for $\sigma$ Ori is the mean value from Zapatero Osorio et al. (632002).,2002).64 The ΝΕ values are from Lemoine ct al. (, The ISM values are from Lemoine et al. (651995: line of sight towards p Oph) and Ixnauth ct al. (,1993; line of sight towards $\rho$ Oph) and Knauth et al. (662000: lines of sight towards o Per and ὅ or).,2000; lines of sight towards $o$ Per and $\zeta$ Per).67 Data for field dwarls (Chen et al., Data for field dwarfs (Chen et al.68 2001) are also shown or comparison., 2001) are also shown for comparison.69 The thick continuous lines in Figs., The thick continuous lines in Figs.70" 5. and 6. represent he predictions of a model where Galactic. cosmic ray (CCR) and stellar lithium production are taken into account in order to rise the ‘Li content in the LSAL from. its oimordial value of logcU Li),~2.6 to the meteoritic log s(""Li) = 3.3] and local log σε) ~ 3.2] ones."," \ref{FigLiEv1} and \ref{FigLiEv2} represent the predictions of a model where Galactic cosmic ray (GCR) and stellar lithium production are taken into account in order to rise the $^7$ Li content in the ISM from its primordial value of $\log 71 \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.6$ to the meteoritic $\log 72 \varepsilon$ $^7$ Li) = 3.3] and local $\log \varepsilon$ $^7$ Li) $\sim$ 3.2] ones."73 In xwticular. lithium: is mostly. produced. by. low-mass. long-ived stellar sources (low-mass stars on the red giant branch and novae).," In particular, lithium is mostly produced by low-mass, long-lived stellar sources (low-mass stars on the red giant branch and novae)."74 Onlv a minor contribution comes from stars on he asymptotic giant branch and from Type LL supernovae (sce Romano et al., Only a minor contribution comes from stars on the asymptotic giant branch and from Type II supernovae (see Romano et al.75 1999. 2001 and Fig. 6))," 1999, 2001 and Fig. \ref{FigLiEv2}) )"76 contrary to what has been suggested by Fravaglio ct al. (, contrary to what has been suggested by Travaglio et al. (772001). who invoke a large contribution from AGB stars (see. however. Ventura. D'Xntona Mazzitelli 2002 for à critical analysis of their results).,"2001), who invoke a large contribution from AGB stars (see, however, Ventura, D'Antona Mazzitelli 2002 for a critical analysis of their results)."78" It is found that the amount of stellar »oduction required in order to match the meteoritic data does not change with changing the assumed. primordial ' Li abundance from logz( Li),~2.2 to logz( Li),~2.6. owing o the fact that the evolution of ‘Li during almost the whole Galaxy evolution is practically determined. solely by he amount of lithium produced through stellar ancl CT orocesses (οἱ."," It is found that the amount of stellar production required in order to match the meteoritic data does not change with changing the assumed primordial $^7$ Li abundance from $\log 79 \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.2$ to $\log 80 \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.6$, owing to the fact that the evolution of $^7$ Li during almost the whole Galaxy evolution is practically determined solely by the amount of lithium produced through stellar and GCR processes (cf."81 figure 6 of Romane ct al., figure 6 of Romano et al.82 2001)., 2001).83 The cillerent ines in Fig., The different lines in Fig.84 6 represent the predictions of models where ithium is contributed by only a single stellar Li factory: asymptotic elant branch (AGB) stars (short-clashecl line): Type HE supernovae (ΝΟ: clotted line): low-mass red giants (long-dashed line): novae (dot-dashed line)., \ref{FigLiEv2} represent the predictions of models where lithium is contributed by only a single stellar Li factory: asymptotic giant branch (AGB) stars (short-dashed line); Type II supernovae (SNeII; dotted line); low-mass red giants (long-dashed line); novae (dot-dashed line).85" We conclude that. even in the case in which the primordial abundance of lithium is as high as logz("" Li),~2.6 (as suggested. from recent data). our. previous conclusions on the ‘Li evolution in the solar neighbourhood are left unchanged: in particular. the rise olf the primordial plateau value is still explained as due to the same important. late ‘Li contribution [rom long-lived. stellar sources (Iow-mass red giants and novae: Romano et al."," We conclude that, even in the case in which the primordial abundance of lithium is as high as $\log \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.6$ (as suggested from recent data), our previous conclusions on the $^7$ Li evolution in the solar neighbourhood are left unchanged; in particular, the rise off the primordial plateau value is still explained as due to the same important, late $^7$ Li contribution from long-lived stellar sources (low-mass red giants and novae; Romano et al."86 1999. 2001).," 1999, 2001)."87 Llowever. if the logz(' Li); value is as high as ~2.6. we stronely need some non-stancare depletion mechanism able to explain the absence of scatter and the high level of Hatness observed in halo stars over a quite large range of metallicity.," However, if the $\log \varepsilon$ $^7$ $_{\mathrm{p}}$ value is as high as $\sim 88 2.6$, we strongly need some non-standard depletion mechanism able to explain the absence of scatter and the high level of flatness observed in halo stars over a quite large range of metallicity."89 The recent results from and their. direc consequences on the primordial abundances of the ligh elements have finally allowed us to check that GCE moclels able to reproduce all the major observed. properties of the Alilky Was are consistent also with SDBDN predictions., The recent results from and their direct consequences on the primordial abundances of the light elements have finally allowed us to check that GCE models able to reproduce all the major observed properties of the Milky Way are consistent also with SBBN predictions.90 1 has been shown several times in the last decade (e.g. Steigman Του 1992: Calli et al.," It has been shown several times in the last decade (e.g., Steigman Tosi 1992; Galli et al."91 1995: Prantzos 1996: ‘Tosi 1996: “Tosi ct al., 1995; Prantzos 1996; Tosi 1996; Tosi et al.92 1998: Chiappini et al., 1998; Chiappini et al.93 2002) tha only a verv moderate D epletion from. its. primordia xindance to the present one is allowed to let the models mfeproduce the observed: radial distributions of chemica »uindances. star and eas densities and star formation rates. as well as the age-metallicity relation and the C-dwarf metallicity. distribution.," 2002) that only a very moderate D depletion from its primordial abundance to the present one is allowed to let the models reproduce the observed radial distributions of chemical abundances, star and gas densities and star formation rates, as well as the age-metallicity relation and the G-dwarf metallicity distribution."94 Llowever. until now there was no definitive observational evidence on the primordial values. since one cannot completely exclude. that. even ueh-redshift. low-metallicity QSO absorbers might be already pollutecl by stellar nucleosvnthesis.," However, until now there was no definitive observational evidence on the primordial values, since one cannot completely exclude that even high-redshift, low-metallicity QSO absorbers might be already polluted by stellar nucleosynthesis."95 Now we have seen that the D. ο anc thle produced. during the big mane are in excellent. agreement with Galactic evolution requirements.," Now we have seen that the D, $^3$ He and $^4$ He produced during the big bang are in excellent agreement with Galactic evolution requirements."96 For ‘Li. some problems do arise. but they mainlv concern our understanding of the mechanisms of ithium cilution/destruction in stars rather than chemical evolution.," For $^7$ Li, some problems do arise, but they mainly concern our understanding of the mechanisms of lithium dilution/destruction in stars rather than chemical evolution."97" Indeed. by assuming the primordial abundance of Ίου e(* Li),2.6. we still need the same important. late contribution of Li from long-lived. stellar sources (low-mass red eiants and novae) required. in. order to explain the observations in meteorites when. starting from the lower value of logz( Li),2.2 suggested. [rom observations of halo stars under the hypothesis that they neither destroved their pristine. lithium nor sullerecl any pollution by early GCE (see also Romano ct al."," Indeed, by assuming the primordial abundance of $\log \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.6$, we still need the same important, late contribution of $^7$ Li from long-lived stellar sources (low-mass red giants and novae) required in order to explain the observations in meteorites when starting from the lower value of $\log 98 \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.2$ suggested from observations of halo stars under the hypothesis that they neither destroyed their pristine lithium nor suffered any pollution by early GCE (see also Romano et al."99 2001)., 2001).100 In [act. the amount of lithium production from stars ancl GChs required by our model in. order to explain the observations is such that it larecly overwelms the lithium primordial abundance.," In fact, the amount of lithium production from stars and GCRs required by our model in order to explain the observations is such that it largely overwelms the lithium primordial abundance."101" Llowever. ifthe adopted: primorclial value of logz( Li),2.6 is confirmed. a mechanism. able to deplete Li in halo stars while preserving the [latness of the and. producing almost. no scatter in"," However, ifthe adopted primordial value of $\log \varepsilon$ $^7$ $_{\mathrm{p}} \sim 2.6$ is confirmed, a mechanism able to deplete Li in halo stars while preserving the flatness of the and producing almost no scatter in"102prreviously applied have difficulties in detecting these objects.,reviously applied have difficulties in detecting these objects.103llow the polar field maximum of the next cxcle changes with the change in meridional flow is quite dillerent for the two profiles.,How the polar field maximum of the next cycle changes with the change in meridional flow is quite different for the two profiles.104 For the low latitude peak in flow. both Laster and slower meridional flow leads to a decrease in polar field strength. while for the high latitude flow peak. the faster(slower) the new flow. the larger (smaller) subsequent polar field peak.," For the low latitude peak in flow, both faster and slower meridional flow leads to a decrease in polar field strength, while for the high latitude flow peak, the faster(slower) the new flow, the larger (smaller) subsequent polar field peak."105 The changes in polar field (hat occur for an increase in meridional flow speed [or the two meridional [Iow profiles are completely consistent with the qualitative arguments made using Figure 1 of 322., The changes in polar field that occur for an increase in meridional flow speed for the two meridional flow profiles are completely consistent with the qualitative arguments made using Figure 1 of 2.106 So is the decline in polar field when the merilional circulation with hieh latitude peak is reduced., So is the decline in polar field when the meridional circulation with high latitude peak is reduced.107 Only the decline in polar fields when the flow pattern peaking al low latitudes is reduced requires a different explanation perhaps the dynamo cycle-period eels so long that there is more time for the surface flux moving to the poles to be diffused down aud not reach the pole., Only the decline in polar fields when the flow pattern peaking at low latitudes is reduced requires a different explanation – perhaps the dynamo cycle-period gets so long that there is more time for the surface flux moving to the poles to be diffused down and not reach the pole.108 These simulations of changes in polar fields clue to drastic changes in meridional circulation raise lead (o an additional question — how quickly does (he dvnanmo period adjust to the changed meridional flow., These simulations of changes in polar fields due to drastic changes in meridional circulation raise lead to an additional question – how quickly does the dynamo period adjust to the changed meridional flow.109 Figure 4 plots the period of the first live evcles computed following the abrupt change in meridional flow speed without altering the form of the streamlines., Figure 4 plots the period of the first five cycles computed following the abrupt change in meridional flow speed without altering the form of the streamlines.110 Not surprisnglv. (he new periods are about. what we would expect for an advection-dominated [Iux-transport dvnamo in which the dvuamo cvele-period is inversely proportional to the meridional flow-speed.," Not surprisingly, the new periods are about what we would expect for an advection-dominated flux-transport dynamo in which the dynamo cycle-period is inversely proportional to the meridional flow-speed."111 What is perhaps surprising about the results seen in Figure 4 is how quickly the model adjusts to the new period established by the changed meridional flow., What is perhaps surprising about the results seen in Figure 4 is how quickly the model adjusts to the new period established by the changed meridional flow.112 Except lor the extreme case when the [low peak is reduced [rom L8ms! to Gms+. we see that the adjustment occurs almost entirely within the first evele.," Except for the extreme case when the flow peak is reduced from $18\,{\rm m}{\rm s}^{-1}$ to $6\,{\rm m}{\rm s}^{-1}$, we see that the adjustment occurs almost entirely within the first cycle."113 A forthcoming paper on sequential data assimilation for solar dvuamo models is addressing this issue in more cetail: preliminary results indicate that the response Gime’ of a flux-transport dvnamo to a change in meridional flow is as short as about 8 months., A forthcoming paper on sequential data assimilation for solar dynamo models is addressing this issue in more detail; preliminary results indicate that the 'response time' of a flux-transport dynamo to a change in meridional flow is as short as about 8 months.114 The Figure 4 shows the settlement of the cdvnamo cvele-period for a calibrated dynamo as discussed in833.2., The Figure 4 shows the settlement of the dynamo cycle-period for a calibrated dynamo as discussed 3.2.115 For a pure Dabcock-Leighton flux-transport dvnamo as in DikpatiCharbonneau (1999).. without anv tachocline a-elfect. (he evele-period changes in a similar wav.," For a pure Babcock-Leighton flux-transport dynamo as in \citet{dc99}, , without any tachocline $\alpha$ -effect, the cycle-period changes in a similar way."116 However. for the same meridional flow-speed the eveles are little faster in (hat case.," However, for the same meridional flow-speed the cycles are little faster in that case."117 In diffusion-dominated dyvnamos the eveles are faster (han that in advection-dominated dvnamos. due to enhanced diffusive transport added (o (he advective (ransport of magnetic flix.," In diffusion-dominated dynamos the cycles are faster than that in advection-dominated dynamos, due to enhanced diffusive transport added to the advective transport of magnetic flux."118 An extensive analvsis by Hotta&Yokovama(2010) shows how the dvnamo period would change when (he magnetic diffusivity in the bulk of (he convection zone is increased., An extensive analysis by \citet{hotta2010} shows how the dynamo cycle-period would change when the magnetic diffusivity in the bulk of the convection zone is increased.119 From the above study we anticipate that cilfusion-cominated dyvnamos would respond (o a sudden change in meridional flow-speed in an analogous wav to that seen in Figure 3(b)., From the above study we anticipate that diffusion-dominated dynamos would respond to a sudden change in meridional flow-speed in an analogous way to that seen in Figure 3(b).120 However. if would be worthwhile in the future to do an investigation of the response of advection-dominateddvnamos of Munoz-Jaraamilloetal(2010). that used a," However, it would be worthwhile in the future to do an investigation of the response of advection-dominateddynamos of \citet{mnmy2010} that used a"121“short” and “lone” distance scales.,"""short"" and ""long"" distance scales."122 The former. based on the statistical parallax ancl the red. clump methocds. gives the distance modulus to the LMC as approximately 18.2-18.3 mae (Gould Popowski 1998. Uclalski 2000a. 10/0).," The former, based on the statistical parallax and the red clump methods, gives the distance modulus to the LMC as approximately 18.2-18.3 mag (Gould Popowski 1998, Udalski 2000a, 2000b)."123 This is 0.2-0.3 mag smaller than the value of about 18.50 mag resulting from the cepheid. period-Iuminosity relation and theoretical models of the horizontal branch or elobular clusters main sequence fitting., This is 0.2-0.3 mag smaller than the value of about 18.50 mag resulting from the cepheid period-luminosity relation and theoretical models of the horizontal branch or globular clusters main sequence fitting.124 Llowever. recent improvements of the red clump method sed. on the infrared photometry seem to indicate that this method also provides the distance modulus of the LMC at around. 18.50 mag.," However, recent improvements of the red clump method based on the infrared photometry seem to indicate that this method also provides the distance modulus of the LMC at around 18.50 mag."125" For example. three papers based. on he A-band photometry were recently published. giving the distance modulus to the LMC as equal to. 18.49+ 0.04. 18.54d:0.10 and 18.501+0.008,0.045. mag. respectively (Alves et."," For example, three papers based on the $K$ -band photometry were recently published giving the distance modulus to the LMC as equal to $18.49\pm0.04$ , $18.54\pm0.10$ and $18.501\pm0.008_r\pm0.045_s$ mag, respectively (Alves et."126 al 2002. Sarajedini ct al.," al 2002, Sarajedini et al."127 2002. Pietrzvsski (σος 2002).," 2002, Pietrzyńsski Gieren 2002)."128 Also. the newest determination of the cepheicl period-uminosity relation. based on 2600 stars. indicates i - distance to the LMC is around 18.50 mag (Sebo et al.," Also, the newest determination of the cepheid period-luminosity relation, based on $\sim$ 600 stars, indicates that the distance to the LMC is around 18.50 mag (Sebo et al."129"2""", 2002).130 Our estimate is in excellent agreement with recent results proving that our calibration of the absolute magnitudes of It Lyr stars produces valuable results., Our estimate is in excellent agreement with these recent results proving that our calibration of the absolute magnitudes of RR Lyr stars produces valuable results.131 We would like to thank Igor Soszvisski from the OGLE project for his valuable suggestions concerning the OGLE Catalogue of Variable Stars in the LAIC and. Cirzegorz Stachowski for reading the manuscript., We would like to thank Igor Soszyńsski from the OGLE project for his valuable suggestions concerning the OGLE Catalogue of Variable Stars in the LMC and Grzegorz Stachowski for reading the manuscript.132 This work was supported. by the WBN grant. number 2 PO3D 024 22 to A. Olech and 5 P03D 004 21to J. Waluzny., This work was supported by the KBN grant number 2 P03D 024 22 to A. Olech and 5 P03D 004 21to J. Kaluzny.133 zΕΤ zz150 sE zz30° (~1000 ~1.5 ~2 did not detect a notheru counterpart for the chain of aligned kuots extending ~307 southwards from the ΠΠ 3l source., $\approx 417$ $\approx 150$ $^{-1}$ $\approx 30^\circ$ $\sim 1000$ $\sim 1.5$ $\sim 2$ did not detect a nothern counterpart for the chain of aligned knots extending $\sim 30''$ southwards from the HH 34 source.134 This situation changed with the paper of ia Lóppez ct al. (, This situation changed with the paper of a Lóppez et al. (1352008). who obtained IR (1.6 aud 2.] jun) long-slit spectra im which the enmüssiou of the rorthern counter-jet was finally detected.,"2008), who obtained IR (1.6 and 2.1 $\mu$ m) long-slit spectra in which the emission of the northern counter-jet was finally detected."136 These authors rote that the cussion of the counterjet has intensity oeaks at positions (1. ¢.. distances frou the source) which approximately coincide with the knots along the southeru jet.," These authors note that the emission of the counterjet has intensity peaks at positions (i. e., distances from the source) which approximately coincide with the knots along the southern jet."137 Iu this paper. we present πο Spitzer IRAC images of ΠΠ 31.," In this paper, we present new Spitzer IRAC images of HH 34."138 These nuages show the southern jet aud jiortheru counterjets with comparable intensities. and with a surprising degree of svnuuetry.," These images show the southern jet and northern counterjets with comparable intensities, and with a surprising degree of symmetry."139 The observatious are deseribed iu section 2., The observations are described in section 2.140 Iu section 3 woe present an Ππασο of the central region of the ITII 31 outflow. quantitatively evaluate the degree of sviunietry. between the jet and the counterjet. and discuss the iuplications of the results for theviteejection mechauisi that has produced the outflow.," In section 3 we present an image of the central region of the HH 34 outflow, quantitatively evaluate the degree of symmetry between the jet and the counterjet, and discuss the implications of the results for the ejection mechanism that has produced the outflow."141 The re are stuunarized in section L., The results are summarized in section 4.142 The observations of TWD 31 are part of our original Spitzer Space Telescope (Werner et al., The observations of HH 34 are part of our original Spitzer Space Telescope (Werner et al.143 2001) Ceneral Observer (CO) program 3315 (PI Noviega-Crespo)} obtained with both the infrared camera TRAC (Fazio ct al., 2004) General Observer (GO) program 3315 (PI Noriega-Crespo) obtained with both the infrared camera IRAC (Fazio et al.144 2001) aud the iufrared plotometer MIPS (Ricke et al., 2004) and the infrared photometer MIPS (Rieke et al.145 2001) in March 28. 2005.," 2004) in March 28, 2005."146 The data have been recovered Bon the Spitzer Leeacy Archive aud the qualitv of the fnal nuages (Post Basic Calibrated Data or Post-BCD: S15.7 products) is outstaudiue. so that no further processing was required.," The data have been recovered from the Spitzer Legacy Archive and the quality of the final images (Post Basic Calibrated Data or Post-BCD; S18.7 products) is outstanding, so that no further processing was required."147 Iu this study we present the TRAC observations obtained in the four channels (1. 2. Bt) = (3.6. [5. 5.8 SU nu) covering a FOV of ~30ν30% (the result of a G\G6 array map with a stepsize) and with a total inteeration time per pixel of 30 sec.," In this study we present the IRAC observations obtained in the four channels (1, 2, 3, 4) = (3.6, 4.5, 5.8 8.0 ) covering a FOV of $\sim 30\times14830$ (the result of a $\times$ 6 array map with a stepsize) and with a total integration time per pixel of 30 sec."149 The final amages are sampled with per pixel. uearly one third of staudard ~2” IRAC aueular resolution.," The final images are sampled with per pixel, nearly one third of standard $\sim 2$ IRAC angular resolution."150 Figure 1 shows a three color nuage of TIT 31 using channels1. 2 aud3.," Figure 1 shows a three color image of HH 34 using channels 1, 2 and 3."151 Like withother protostellay outflowsobserved withIRAC (see c.g. Noricea-Crespo et al., Like withother protostellar outflowsobserved withIRAC (see e.g. Noriega-Crespo et al.152 200L:, 2004;153of detection.,of detection.154 Figure 6 shows contours around one such peak. showing the level at which the signal drops to .75. .5 and .25 its maximumi value.," Figure 6 shows contours around one such peak, showing the level at which the signal drops to $.75$ , $.5$ and $.25$ its maximum value."155 This figure demonstrates the point mentioned above regarding shapes of these contour., This figure demonstrates the point mentioned above regarding shapes of these contour.156" Even though we cannot suggest the most optimum. smoothing function for improving the signal to noise ratio. we have confirmed that the signal to noise ratio does improve for. smoothing. with. a square top hat window. of D.size zmgnij2 where nga ds the number of pixels enclosed. within the contour of 5,,,,/2."," Even though we cannot suggest the most optimum smoothing function for improving the signal to noise ratio, we have confirmed that the signal to noise ratio does improve for smoothing with a square top hat window of size $\approx n_{HM}^{1/2}$ where $n_{HM}$ is the number of pixels enclosed within the contour of $S_{max}/2$."157 We assumed that the noise scales as nL7 for small n., We assumed that the noise scales as $n^{-1/2}$ for small $n$.158 The signal to noise ratio for moclels 11 (MDM) and HE CLOCDAM) can be improved by a factor two in this manner., The signal to noise ratio for models II (MDM) and III (LCDM) can be improved by a factor two in this manner.159 The standard CDM model model I] has more small scale power and hence the peaks are much sharper and therefore the eain in signal to noise ratio by smoothing is somewhat limited in this case., The standard CDM model [model I] has more small scale power and hence the peaks are much sharper and therefore the gain in signal to noise ratio by smoothing is somewhat limited in this case.160 In order to estimate the integration time required. for detection and imaging of these peaks we need to know the root mean square amplitude of noise expected for the GMT receivers., In order to estimate the integration time required for detection and imaging of these peaks we need to know the root mean square amplitude of noise expected for the GMRT receivers.161 The expected noise for the central array is The svstem temperature for 327MIIz window is 110Ix and the corresponding number for 233MLI is 250K. ]t is clear from table 1 that the peak signal expected," The expected noise for the central array is The system temperature for $327$ MHz window is $110$ K and the corresponding number for $233$ MHz is $250$ K. \cite{swarup84}162It is clear from table 1 that the peak signal expected"163of the background contamination.,of the background contamination.164 To this optical catalogue. we added the positions of bright (V.« 11) stars found in this region from the SIAIBAD database.," To this optical catalogue, we added the positions of bright $V<11$ ) stars found in this region from the SIMBAD database."165 ‘To establish an appropriate eross-correlation radius to use between the X-ray and optical source Lists we mocelle the cumulative number of X-ray sources that were correlate with an optical source with V«19 (see Jeffries. Phurston Pye 1997 for details).," To establish an appropriate cross-correlation radius to use between the X-ray and optical source lists we modelled the cumulative number of X-ray sources that were correlated with an optical source with $V<19$ (see Jeffries, Thurston Pye 1997 for details)."166 Assuming a uniform spread of optica sources. we determined that there were 77 correlations (from S3 X-ray sources inside the CCL) survey) within 10 aresee of X-ray positions. that 66 of these would be true counterparts to X-ray sources. LL would be spurious correlations and tha the la X-ray error circle was 3.7 arcsecs.," Assuming a uniform spread of optical sources, we determined that there were 77 correlations (from 83 X-ray sources inside the CCD survey) within 10 arcsec of X-ray positions, that 66 of these would be true counterparts to X-ray sources, 11 would be spurious correlations and that the $1\sigma$ X-ray error circle was 3.7 arcsecs."167 Figure 2b shows 75 sources that have an optical counterpart within LO aresecs (another two are bright stars without Wo{ colours)., Figure 2b shows 75 sources that have an optical counterpart within 10 arcsecs (another two are bright stars without $V-I$ colours).168 Clearly the A-rayv emitting population coincides with the propose PAIS population in the CMD., Clearly the X-ray emitting population coincides with the proposed PMS population in the CMD.169 Indeed. if we were to consider just. a subset of the optical catalogue consisting of a broac strip containing all these PATS sources. we would only expec lofthese correlations to be spurious.," Indeed, if we were to consider just a subset of the optical catalogue consisting of a broad strip containing all these PMS sources, we would only expect 1 of these correlations to be spurious."170 We have calculated the X-rav properties. of this population and these along with the LIRL observations wil be reported in a subsequent publication., We have calculated the X-ray properties of this population and these along with the HRI observations will be reported in a subsequent publication.171 Brielly. the X-ray to bolometric Dux ratio of these objects lies in the range 10 to LO7. broadly what we would expect from a population of voung PALS stars.," Briefly, the X-ray to bolometric flux ratio of these objects lies in the range $10^{-5}$ to $10^{-2}$, broadly what we would expect from a population of young PMS stars."172 Phe cut-olf in the PATS X-ray correlations at VoxIN is almost certainly due to the X-ray sensitivity., The cut-off in the PMS X-ray correlations at $V\simeq18$ is almost certainly due to the X-ray sensitivity.173 ‘To be detected: in N-ravs. fainter objects would. have to have higher than feasible X-ray to bolometric lux ratios.," To be detected in X-rays, fainter objects would have to have higher than feasible X-ray to bolometric flux ratios."174 Llowever. it is clear that the PAIS sequence we have found extends down to the limits of our optical survey at V720.5.," However, it is clear that the PMS sequence we have found extends down to the limits of our optical survey at $V\sim20.5$."175 The appearance of FPig.2 should leave the reader in no doubt that we have found a voung and exceptionally rich population of low mass active stars in the of ? Vel., The appearance of Fig.2 should leave the reader in no doubt that we have found a young and exceptionally rich population of low mass active stars in the of $\gamma^{2}$ Vel.176 The central question to be answered is whether these sources are physically close to 57. Vel and/or whether they are background. members of the Vela OD2 association., The central question to be answered is whether these sources are physically close to $\gamma^{2}$ Vel and/or whether they are background members of the Vela OB2 association.177 We can tackle this problem in a number of wavs., We can tackle this problem in a number of ways.178 Figure 3a shows the spatial distribution. of PALS) stars selected from the €MD in a strip enclosing the bulk of the X-ray sources (marked with a dashed box in Fig.2)., Figure 3a shows the spatial distribution of PMS stars selected from the CMD in a strip enclosing the bulk of the X-ray sources (marked with a dashed box in Fig.2).179 This box was chosen to avoid background contamination., This box was chosen to avoid background contamination.180 We determined. the radial distribution of these stars. centred on +> Vel.," We determined the radial distribution of these stars, centred on $\gamma^{2}$ Vel."181 The distribution is normalized using the radial distribution of background stars with a similar V magnitude range. but à colour range of Q.S«V.fo «l.Y. under he assumption that the background: stars are. uniformly distributed.," The distribution is normalized using the radial distribution of background stars with a similar $V$ magnitude range, but a colour range of $<V-I_{\rm c}<$ 1.7, under the assumption that the background stars are uniformly distributed."182 Phe resulting racial distribution is shown in bie., The resulting radial distribution is shown in Fig.183 3b. which exhibits a small but significant (at the 30 level) oak within 5 aremin of the centre of the CCL) survey aux 5 Vel.," 3b, which exhibits a small but significant (at the $3\sigma$ level) peak within 5 arcmin of the centre of the CCD survey and $\gamma^{2}$ Vel."184 The point closest to 57. Vel is missing because the 30 arcsec τοσιο immeciatelv surrounding 25 Vel is SWelllpec w its light and no accurate photometry was obtained there., The point closest to $\gamma^{2}$ Vel is missing because the 30 arcsec region immediately surrounding $\gamma^{2}$ Vel is swamped by its light and no accurate photometry was obtained there.185 1n a similar fashion we can show that the N-rav sources are also mareinally concentrated toward the centre of the Lele even after correction for the PSPC vignetting function., In a similar fashion we can show that the X-ray sources are also marginally concentrated toward the centre of the field even after correction for the PSPC vignetting function.186 ‘Phere is thus some evidence that the PAIS stars and 572 Vel are spatially correlated. although this does not rule ou a chance alignment of 57. Vel with a background cluster o low mass stars in Vela OB2.," There is thus some evidence that the PMS stars and $\gamma^{2}$ Vel are spatially correlated, although this does not rule out a chance alignment of $\gamma^{2}$ Vel with a background cluster of low mass stars in Vela OB2."187 In particular. it is possible tha any concentration we see could be associated with «+ Vel. a 19111. sinele lined spectroscopic binary which is a common proper motion companion to 57. Vel.," In particular, it is possible that any concentration we see could be associated with $\gamma^{1}$ Vel, a B2III, single lined spectroscopic binary which is a common proper motion companion to $\gamma^{2}$ Vel."188 24 Vel is also a likely, $\gamma^{1}$ Vel is also a likely189each panel are for ο=6. 3. 1.5 and 1. respectively.,"each panel are for $\beta=6$ $3$ $1.5$ and $1$, respectively."190" It is seen that INS;—IN, can becomparable {ο :peak", It is seen that $N_{peak}-N^{ran}_{peak}$ can becomparable to $N^{ran}_{peak}$.191 For @¢;=]avemin. :peakc92. 0.3. 0.04. and0.004 at ιν=3.5. 4. 1.5 and 5. respectively.," For $\theta_G=1\hbox { arcmin}$ $N^{ran}_{peak}\sim 2$ $0.3$ , $0.04$ , and$0.004$ at $\nu_{ran}=3.5$ , $4$ , $4.5$ and $5$, respectively."192" The corresponding IN,—Nu, owe 0.25. 0.06. 0.013. and 0.0015 lor FENandA Ln.= 1.29."," The corresponding $N_{peak}-N^{ran}_{peak}$ are $0.25$ , $0.06$ , $0.013$, and $0.0015$ for $\beta=1.5$ and $A=1.29$ ."193"aAForD larger 7.9 thedJopoar odnumbers are larger⋅⋅↽⋅⋅ and IN,⇁ — .j—peakTNealVhen Mom.1.5μυ... and4.5Dlor 3 = diac6 and3.respectively."," For larger $\beta$, the numbers are larger and $N_{peak}-N^{ran}_{peak}>N^{ran}_{peak}$ when $\nu_{ran}>3.7$ and $4.5$ for $\beta=6$ and $3$, respectively."194" For 6(;- =.L58) 2aan avemin.:INS;—-Panos. peakiyDNMVeowhenοZ7>3.0 3. 3.8and5.3 for 9 = 6. 3and1.5. respectively,"," For $\theta_G=2\hbox{ arcmin}$, $N_{peak}-N^{ran}_{peak}>N^{ran}_{peak}$ when $\nu_{ran}>3$, $3.8$ and $5.3$ for $\beta=6$, $3$ and $1.5$, respectively."195 Therefore the existence of intrinsic alignments can result significant number of extra false peaks in lensing convergence maps., Therefore the existence of intrinsic alignments can result significant number of extra false peaks in lensing convergence maps.196" In Figure 4. we show the dependence ofthe ratio ria;=ρε)Nu, on the level of intrinsic alignments represented by the amplitude A for 6;=1arcmin."," In Figure 4, we show the dependence ofthe ratio $r_{peak}=N_{peak}/N^{ran}_{peak}$ on the level of intrinsic alignments represented by the amplitude $A$ for $\theta_G=1\hbox{ arcmin}$."197 The 5 value in each panel is written out explicitly., The $\beta$ value in each panel is written out explicitly.198" The solid. dotted. dashed and dash-dotted lines are respectivelv [orMpg,=5. 4.5. 4 and 3.5."," The solid, dotted, dashed and dash-dotted lines are respectively for$\nu_{ran}=5$, $4.5$, $4$ and $3.5$."199" For 9=1.5 and A=1.29. we have ryc1.36. 1.21. 1.2 and 1.14 lor v,,,,=5. 4.5. 4 and 3.5. respectively."," For $\beta=1.5$ and $A=1.29$ , we have $r_{peak}\sim 1.36$, $1.27$, $1.2$ and $1.14$ for $\nu_{ran}=5$, $4.5$, $4$ and $3.5$, respectively."200 With larger smoothing scales. the relative effect of intrinsic alignments is higher.," With larger smoothing scales, the relative effect of intrinsic alignments is higher."201 For 9;=2arcanin. (he corresponding ratios change to 1.6. 1.4. 1.3 and 1.2.," For $\theta_G=2\hbox{ arcmin}$, the corresponding ratios change to $1.6$, $1.4$, $1.3$ and $1.2$."202" For 2=6. the ratios for ιν=5 reach as high as 3.8 and 1.6 Dor θε,=1 and 2arcmin. respectively."," For $\beta=6$, the ratios for $\nu_{ran}=5$ reach as high as $3.8$ and $7.6$ for $\theta_G=1$ and $2\hbox{ arcmin}$, respectively."203 From eqs. (, From eqs. (204"25) and (26). it can be shown that IN. depends largely on the detection threshold with IN,Xrexp(—772/2) al p>3 (e.g.. van Waerbeke 2000).","25) and (26), it can be shown that $N_{peak}$ depends largely on the detection threshold with $N_{peak}\propto \nu\exp(-\nu^2/2)$ at $\nu>3$ (e.g., van Waerbeke 2000)."205" Given a detection threshold on ρω. the corresponding threshold for the (true significance is Which decreases with the increase of OF/05,."," Given a detection threshold on $\nu_{ran}$ , the corresponding threshold for the true significance is $\nu=\nu_{ran}/(1+\sigma^2_{0corr}/\sigma^2_{0ran})^{1/2}$ , which decreases with the increase of $\sigma^2_{0corr}/\sigma^2_{0ran}$ ."206" ο is largely determined by the ratio OF,oe./Thran", Thus $r_{peak}$ is largely determined by the ratio $\sigma^2_{0corr}/\sigma^2_{0ran}$.207" In Figure 5. we show ry4; with respect lO 05,05, lor 0c;=1 (upper panel) aud 2arcmin (lower panel)."," In Figure 5, we show $r_{peak}$ with respect to $\sigma^2_{0corr}/\sigma^2_{0ran}$ for $\theta_G=1$ (upper panel) and $2\hbox{ arcmin}$ (lower panel)."208" The four sets of lines from top to bottom in each panel correspond respectively (ο the threshold v,,,,=5. 4.5. 4 and 3.5."," The four sets of lines from top to bottom in each panel correspond respectively to the threshold $\nu_{ran}=5$, $4.5$, $4$ and $3.5$ ."209" Note that each set contains four lines with ο)—6. 3. 1.5. and 1. respectively,"," Note that each set contains four lines with $\beta=6$, $3$, $1.5$, and $1$, respectively."210" With σοιση~OM. the values o ry; are about 1.7. 1.55. L4 and 1.3 for pu,=5. 4.5. 4 and 3.5."," With $\sigma^2_{0corr}/\sigma^2_{0ran}\sim 5\%$, the values of $r_{peak}$ are about $1.7$, $1.55$, $1.4$ and $1.3$ for $\nu_{ran}=5$, $4.5$, $4$ and $3.5$."211" For 65,/05,,,»LOW. the corresponding ρω are 2.9. 2.3. L9 and 1.6."," For $\sigma^2_{0corr}/\sigma^2_{0ran}\sim 10\%$, the corresponding $r_{peak}$ are $2.9$, $2.3$, $1.9$ and $1.6$."212" A specilic value of 05,ση depends on the strength of the intrinsic alignment. the surface number densitv and the redshift distribution of source galaxies. and σι."," A specific value of $\sigma^2_{0corr}/\sigma^2_{0ran}$ depends on the strength of the intrinsic alignment, the surface number density and the redshift distribution of source galaxies, and $\sigma_{\epsilon}$."213" The dotted. vertical lines from left to right in each panel show ⊔∐↲≺∢∪↕⋅↕⋅≼↲⊳∖⇁↕↽≻∪∐≼⊔∐≸≟∖↽≀↧↴↥⋯↲⊳∖⊽∪↓⊔∣−∣↿⋅⋯≖∕≖∕∕∕⊔∎−∣∕≦⋯∣↓∪↕⋅↽≳↾∶⊥⋅⊥⋅↱≻⋅≡↽⊰≀↧↴∐≼⊔≻⋅∖∖↽↥∐↲↕⋅≼↲∖∖↽≼↲↥≀↧↴↳↽≼↲⇀∣↥∶⊥⋅−≻≤∍⋅0032 0, ⋅ ⋅ ⋅ ⊔⋠∶∩⋅∔⋅≀↕↴∐≼⇂∣∣∙↙∕∶⋮↽⊰∩≀↧↴↕⋅≺∢∐∐∐−⋅∐⋟∖⊽∐⋯∐≼⊓↽≻≼↲∐∪↥⋯⇂⊔⋯↴↥∣∣∙↙∕∏⋟∖⊽∏≀↧↴∐∡∖↽∖"," The dotted vertical lines from left to right in each panel show the corresponding values of $\sigma^2_{0corr}/\sigma^2_{0ran}$ for $\beta=1$ ,$1.5$ ,$3$ and $6$ , where we take $A=1.29$ ,$\sigma_{\epsilon}=0.4$ ,and$n_g=30\hbox{ arcmin}^{-2}$ ."214↽≀↧↴↕⋅↕≼↲⋟∖⊽∖∖⇁↕⊔↥⊔∐↲↕⋅≼↲≼⇂⋟∖⊽↥∐∐ ≼∐⊳∖⊽∏⋅↕∣↽≻∏∐∪∐∪↓⋟⊳∖⇁⋯∐⋅≺∢≼↲≸↽↔↴≀↧↴↥≀↧↴⇀↸↕≼↲⊳∖⊽⋅⊳↔⊲⋯⋅∖↽≼↲∡∖⇁⊳∖⊽, It shouldbe notedthat $n_g$ usually varieswith theredshift distribution of source galaxies.215⊔⋯↥≺∢≀↧↴∐↕⋅≼↲⋯∢∐∐↕≸↽↔↴∐↕⋅≼↲≺⇂⊳∖⇁∐∐≯↥⊳∖⇁↥⋡∖↽↕↽≻↕≺∢≀↧↴∐∡∖⇁∐≀↧↴∖↽≼↲↥≀, Surveys thatcan reach high redshifts typically have large $n_g$ .216"↧↴↕⋅≸↽↔↴≼↲∣∣∙↙∕⋅ ↴Thus our estimates− on 05,,,,/05,,,>2 with. fixed, n,=30.arcmin2~ may overestimate. the ratio. for ο= 6.", Thus our estimates on $\sigma^2_{0corr}/\sigma^2_{0ran}$ with fixed $n_g=30\hbox{ arcmin}^{-2}$ may overestimate the ratio for $\beta=6$ .217" On the other hand. fordeep surveys with large n, (e.g. 7,~LOO for SNAP. and n,~300 for SNAP Deep). we can divide thesource galaxies into different bins with ny~30.arcmin: LEin each bin."," On the other hand, fordeep surveys with large $n_g$ (e.g., $n_g\sim 100$ for SNAP, and $n_g\sim 300$ for SNAP Deep), we can divide thesource galaxies into different bins with $n_g\sim 30 \hbox{ arcmin}^{-2}$ in each bin."218": In (his:case. the narrow distribution. with. ο=6. can be one of theseebins.bins and ourP above estimateexti on 05,,,/05,,,2442 WilhE ny—=30arcminarem72 can be a representative value for galaxies within the bin."," In thiscase, the narrow distribution with $\beta=6$ can be one of thesebins, and our above estimate on $\sigma^2_{0corr}/\sigma^2_{0ran}$ with $n_g=30\hbox{ arcmin}^{-2}$ can be a representative value for galaxies within the bin."219"In this section we describe the procedure we follow to assess the temperature. rotation rate and chemical abundances of the accreting WD in VW νι,","In this section we describe the procedure we follow to assess the temperature, rotation rate and chemical abundances of the accreting WD in VW Hyi."220 This procedure consists of comparing the observed FUSE spectrum of VW Lyi with a grid of theoretical spectra obtained. assuming different assumptions and for different. values of the parameters of the svslel as explained below., This procedure consists of comparing the observed FUSE spectrum of VW Hyi with a grid of theoretical spectra obtained assuming different assumptions and for different values of the parameters of the system as explained below.221 For this purpose. we use a combination of svntlietic stellar. disk and accretion belt spectra.," For this purpose, we use a combination of synthetic stellar, disk and accretion belt spectra."222 The best fit models are then obtained using a 47 minimization filling procedure., The best fit models are then obtained using a $\chi^{2}$ minimization fitting procedure.223 The stellar atinosphere models aud the accretion belt models are both generated using ihe TLUSTY code (IIubeny.1983) for different. values of the parameters of the accreting WD. such as temperature. composition ancl surface gravitv.," The stellar atmosphere models and the accretion belt models are both generated using the TLUSTY code \citep{hub88} for different values of the parameters of the accreting WD, such as temperature, composition and surface gravity."224 A spectrum svnthesis code (SYNSPEC) is then used (o generate the spectra of the stellar atmospheres obtained in TLUSTY (IIubenyetal.1994:IIubeny.&Lanz1995).," A spectrum synthesis code (SYNSPEC) is then used to generate the spectra of the stellar atmospheres obtained in TLUSTY \citep{hub94,hub95}."225. For the accretion disk spectrum. we used the grid of accretion disk spectra computed by Wade&IInbeny(1993).. who use a slightlv different. version of the TLUSTY code (TLUSDISIX) to generate the theoretical spectrum of an accretion disk.," For the accretion disk spectrum, we used the grid of accretion disk spectra computed by \citet{wad98}, who use a slightly different version of the TLUSTY code (TLUSDISK) to generate the theoretical spectrum of an accretion disk."226 The accretion disk model is really made of a collection of rings., The accretion disk model is really made of a collection of rings.227 The disk models are computed assuming LTE and vertical hydrostatic ecquilibrium., The disk models are computed assuming LTE and vertical hydrostatic equilibrium.228 Dradiation from external sources is neglected., Irradiation from external sources is neglected.229 Local spectra of disk annuli are computed taking into account line (ranusiGion from elements 1-28 (I1 through. Ni)., Local spectra of disk annuli are computed taking into account line transition from elements 1-28 (H through Ni).230 Limb darkening as well as Doppler broadening aud blending of lines are taken into account., Limb darkening as well as Doppler broadening and blending of lines are taken into account.231 Then. to carry out the model fits. we masked the following wavelength regions where several narrow enission-like features occur in theFUSE spectrum: [048-952|]. [972-914]]. (089-993]]. ]]. [1078-1082|]. [1092-1096]]. [LLGS-1170A]].," Then, to carry out the model fits, we masked the following wavelength regions where several narrow emission-like features occur in the spectrum: ], ], ], ], ], ], ]."232 Most of these emission lines are due (o air glow. when part of the observations are carried out during dav time.," Most of these emission lines are due to air glow, when part of the observations are carried out during day time."233 In particular we masked the OVI emission lines around 1033 aand 1037, In particular we masked the OVI emission lines around 1033 and 1037.234 Details ofthe codes and the 47 (4? per degree of freedom) minimization fitting procedures are discussed in detail in Sionetal.(1995) and παςetal. (1996a).. and will not be repeated here.," Details ofthe codes and the $\chi^{2}_{\nu}$ $\chi^2$ per degree of freedom) minimization fitting procedures are discussed in detail in \citet{sio95} and \citet{hua96a}, , and will not be repeated here."235is on a leneth scale siualler than about one fifth of the beau.,is on a length scale smaller than about one fifth of the beam.236 Iun Fig. 10..," In Fig. \ref{f3:can_shape},"237 the steepest canals found iu the data are shown., the steepest canals found in the data are shown.238 In this figure the top plots give a onc- cross-cut of P across three canals agaiust position. for all frequencies.," In this figure the top plots give a one-dimensional cross-cut of $P$ across three canals against position, for all frequencies."239" The frequencies in which the canals were defined were 311 MIIz. 355 MIIZ and 319 ΑΠΣ respectively, and the canals were selected. for their steepness."," The frequencies in which the canals were defined were 341 MHz, 355 MHz and 349 MHz respectively, and the canals were selected for their steepness."240 The bottom plots give ouly the P distribution across the canal at the frequency at which it was defined (solid. line)., The bottom plots give only the $P$ distribution across the canal at the frequency at which it was defined (solid line).241 Superimposed in dashed lines is the P distribution of the moce of Fig., Superimposed in dashed lines is the $P$ distribution of the model of Fig.242 9 for the steepest angle chauge convolved with the svuthesized eam., \ref{f3:can_the} for the steepest angle change convolved with the synthesized beam.243 Less steep anele changes give less steep 2 profiles aud worse fits to the data., Less steep angle changes give less steep $P$ profiles and worse fits to the data.244 An interpretation of these (specifically selected) steep canals in terms of differcutial Faraday dispersion is difficult. because the canals would have to be much more closely spaced than observed.," An interpretation of these (specifically selected) steep canals in terms of differential Faraday dispersion is difficult, because the canals would have to be much more closely spaced than observed."245 Bea depolarization predicts a change iu depth of the canals across the frequency bands of about20%.. in aereciment with the observations.," Beam depolarization predicts a change in depth of the canals across the frequency bands of about, in agreement with the observations."246 We conclude that the dominant process creating o1e-bean wide canals of almost complete depolarization is most likely beau depolarization., We conclude that the dominant process creating one-beam wide canals of almost complete depolarization is most likely beam depolarization.247 In this case. abrupt RAL changes have to be present iu the medii.," In this case, abrupt $RM$ changes have to be present in the medium."248 It may secur ortuitous that only RAL eracieuts of the right magnitude o 1unake canals would exist., It may seem fortuitous that only RM gradients of the right magnitude to make canals would exist.249 However. this is not the case: RAL eradieuts of auy magnitude are likely to occur in the neci. bi oulv the RM eradieuts that cause Aco907 vield a visible signature in P.," However, this is not the case: RM gradients of any magnitude are likely to occur in the medium, but only the RM gradients that cause $\Delta\phi\approx 90\dg$ yield a visible signature in $P$."250 Because RAL is an iutegra along the line of sight. it is difficult to see what physica xocess would be responsible for this.," Because $RM$ is an integral along the line of sight, it is difficult to see what physical process would be responsible for this."251 However. umnerica nodels of a magueto-ionized ISM show that RM eracicuts steep enough to produce canals at 350 MITz are couuinion (IIwverkorn Ποιοι 2001).," However, numerical models of a magneto-ionized ISM show that RM gradients steep enough to produce canals at 350 MHz are common (Haverkorn Heitsch 2004)."252 The relatively low RA eradicut needed to make a canal at 350 MITz. às conmiparec to Ll GIIz observations. could explain why canals are abundaut in these WSRT observations. but are 11uch less conumuon at 1.I GITz (Uvauiker et 11998).," The relatively low RM gradient needed to make a canal at 350 MHz, as compared to 1.4 GHz observations, could explain why canals are abundant in these WSRT observations, but are much less common at 1.4 GHz ker et 1998)."253 Nevertheless. Figs.," Nevertheless, Figs."254 7 and & show that beam depolarization certainly is uot the whole explanation., \ref{f3:can_l2} and \ref{f3:can_rm} show that beam depolarization certainly is not the whole explanation.255 Tf differcutial Faraday rotation were the main cause of the canals. it would be hard to understiud why all canals are exactly oue beam wide. aux why we do rot observe any siguiücaunt change iu he position of re canals with frequency.," If differential Faraday rotation were the main cause of the canals, it would be hard to understand why all canals are exactly one beam wide, and why we do not observe any significant change in the position of the canals with frequency."256 Furthermore. the existeuce f canals in which P egocs down to aliuos zero woul ren indicate a very 1aiform medium in both magnetic field ancl electron deusitv.," Furthermore, the existence of canals in which $P$ goes down to almost zero would then indicate a very uniform medium in both magnetic field and electron density."257 Sokoloff e ((1998) «lose ji an cexponcutial asvunuetre slab. causes llOll-ZOTO uiiminia for the canals. which even disappear conipletelv in a turbulent iiediuu.," Sokoloff et (1998) showed that an exponential asymmetric slab causes non-zero minima for the canals, which even disappear completely in a turbulent medium."258 Small-scale structure m observe RAL indicates tha siuall-scale structure in magnetic fik and/or electron density is abundant. so that a uniform meditm necded for deep canals iu the differcutial Faraday rotation interpretation is unlikely.," Small-scale structure in observed $RM$ indicates that small-scale structure in magnetic field and/or electron density is abundant, so that a uniform medium needed for deep canals in the differential Faraday rotation interpretation is unlikely."259 However. Sliukurov ane Berkhuijsen (2003) argue that the canals they observe at Ll GIIz in M2I are best explaimed as due to depth depoluwization.," However, Shukurov and Berkhuijsen (2003) argue that the canals they observed at 1.4 GHz in M31 are best explained as due to depth depolarization."260 Suiall-scale structure in the linearly polarized componcut of the diffuse Calactic svuchrotroun enmüsson Is seen iu almost everv direction., Small-scale structure in the linearly polarized component of the diffuse Galactic synchrotron emission is seen in almost every direction.261 Mosthv. this structure is not correlated with otal e1uission. and therefore cannot be due to sinall-scale structure dn enmission.," Mostly, this structure is not correlated with total emission, and therefore cannot be due to small-scale structure in emission."262 lustead. the polarization angle o is Faradax-rotated in the maenueto-ionic iuediun through which the linearly polarized radiation propagates.," Instead, the polarization angle $\phi$ is Faraday-rotated in the magneto-ionic medium through which the linearly polarized radiation propagates."263 However. the structure in polarized intensityi 2 cannot be produce by Faraday rotation (which only rotates ©). but there aro several Xocesses respousible for this.," However, the structure in polarized intensity $P$ cannot be produced by Faraday rotation alone (which only rotates $\phi$ ), but there are several other processes responsible for this."264 First. instruiment-related effects produce structure in P. such as large-scale conponents in the radiation that are undetectable with an interferometer. depolarization duc to variation in anele within the telescope beam. or over the frequency baud width.," First, instrument-related effects produce structure in $P$, such as large-scale components in the radiation that are undetectable with an interferometer, depolarization due to variation in angle within the telescope beam, or over the frequency band width."265 Furthermore. physical depolarization processes iu the ISM can cause depolarization if Faraday rotation and svuchrotron emission occur in the same medium.," Furthermore, physical depolarization processes in the ISM can cause depolarization if Faraday rotation and synchrotron emission occur in the same medium."266TC and ΑΛ are thankful to IUCAA for assistance through. Associateship Programme.,TC and AKC are thankful to IUCAA for assistance through Associateship Programme.267 MES acknowledges of a partial support of a Russian Foundation Basic Research grant 08-02-00627., MES acknowledges of a partial support of a Russian Foundation Basic Research grant 08-02-00627.268 The authors are extremely grateful to Ethan T. Vishniac. Editor in Chief aud Richarel de Grijs (Scientific Editor) for their aclive cooperation ancl support.," The authors are extremely grateful to Ethan T. Vishniac, Editor in Chief and Richard de Grijs (Scientific Editor) for their active cooperation and support."269 The authors are also verv thankful to the referee for valuable suggestions and technical guidance in improving (he quality of the work., The authors are also very thankful to the referee for valuable suggestions and technical guidance in improving the quality of the work.270 This method was developed by Salibianu-Barrera et al.(2006)., This method was developed by Salibiánn-Barrera et al.(2006).271 This Principal Component Analvsis (PCA) is new in the sense that it is based on multivariate of shape instead. of sample covariances., This Principal Component Analysis (PCA) is new in the sense that it is based on multivariate MM-estimator of shape instead of sample covariances.272 MM-estimator gives a robust estimate having the properties of maximum likelihood estimator., MM-estimator gives a robust estimate having the properties of maximum likelihood estimator.273 A robust estimate is not affected by outliers or small deviations [rom the model assumptions., A robust estimate is not affected by outliers or small deviations from the model assumptions.274 The definitions of MM-estimators of Mulüvariate location and shape are given in the above mentioned paper., The definitions of MM-estimators of Multivariate location and shape are given in the above mentioned paper.275 The-MM estimators were computed by applying the fast bootstrap procedure of Salibiiun-Darrera Zamar (2002)., The-MM estimators were computed by applying the fast bootstrap procedure of Salibiánn-Barrera Zamar (2002).276" IE there are p parameters in (he data set and we denote the estimated variances ofp principal components by λι>""mucXp. then to find the optimum number of principal components the following proportion has been used (Salibiáun-Darrera el al."," If there are $p$ parameters in the data set and we denote the estimated variances of$p$ principal components by $\widehat{\lambda_{1}} > \widehat{\lambda_{2}} > ... >277\widehat{\lambda_{p}}$, then to find the optimum number of principal components the following proportion has been used (Salibiánn-Barrera et al."278 2006): One should consider that value of k as optimum for which the value of 100p; exceeds some cul off value., 2006): One should consider that value of k as optimum for which the value of $\widehat{p}_{k}$ exceeds some cut off value.279 In our case. this cut off value has been chosen as 90%..," In our case, this cut off value has been chosen as ."280 One can also, One can also281pixel which we believe to be a fair estimate of the true limits of our data.,pixel which we believe to be a fair estimate of the true limits of our data.282 The last column of Table 3. gives the value of systematic uncertainty that we have added in quadrature in cach case., The last column of Table \ref{res:prob} gives the value of systematic uncertainty that we have added in quadrature in each case.283 We then calculate the probability of the true orbita periods being further than 1 and LO per cent from the values we obtained - see Morales-Itueda et al., We then calculate the probability of the true orbital periods being further than 1 and 10 per cent from the values we obtained - see Morales-Rueda et al.284 (20032). and. Marsh. Dhillon Duck (1995) for an explanation of the methoc used to caleulate these probabilities - and. present them in ‘Table 3..," \shortcite{lmr03} and Marsh, Dhillon Duck \shortcite{mdd95} for an explanation of the method used to calculate these probabilities - and present them in Table \ref{res:prob}."285 In all cases. the probabilities of the quoted periods being wrong are very low and. we are certain that the values given in Table 2. correspond to the true orbital solution.," In all cases, the probabilities of the quoted periods being wrong are very low and we are certain that the values given in Table \ref{res:rv:tab1} correspond to the true orbital solution."286 In the cases where the probability of the orbital period beinge further than 1 and 10 per cent from our favoured value is the same. the significant probability lies within a very small range around the best period. with all the significant competition (i.c. next best alias) placed outside the 10 per cent region around the best alias.," In the cases where the probability of the orbital period being further than 1 and 10 per cent from our favoured value is the same, the significant probability lies within a very small range around the best period, with all the significant competition (i.e. next best alias) placed outside the 10 per cent region around the best alias."287 By knowing the radial velocity semiamplitude of one of the components of the binary (the observable component). A. and the orbital period of the system. we can then calculate the mass function of the unseen component by using: where the subscripts 71 anc 727 refer to the brighter and the dimmer components respectively.," By knowing the radial velocity semiamplitude of one of the components of the binary (the observable component), $K$, and the orbital period of the system, we can then calculate the mass function of the unseen component by using: where the subscripts $1$ ” and $2$ ” refer to the brighter and the dimmer components respectively."288 Phe mass function is the lower limit of the mass of the unseen component., The mass function is the lower limit of the mass of the unseen component.289 Table 4 gives the mass functions of the unseen components for the four DDs studied., Table \ref{res:mass} gives the mass functions of the unseen components for the four DDs studied.290 In two cases aand 188)) the companion's mass function is ereater than wwhich corresponds to the mass of à late M dwarf if it is a main sequence star., In two cases and ) the companion's mass function is greater than which corresponds to the mass of a late M dwarf if it is a main sequence star.291 Phe masses of the brighter components of the svstems have been measured by fitting their hydrogen line profiles to stellar atmosphere models using the tracks by Althaus Benvenuto (1997) and can be substituted. together with the assumption of the orbital inclination of the system being 907. in the mass function equation to give a larger lower limit for the masses of the unseen coniponents.," The masses of the brighter components of the systems have been measured by fitting their hydrogen line profiles to stellar atmosphere models using the tracks by Althaus Benvenuto \shortcite{ab97} and can be substituted, together with the assumption of the orbital inclination of the system being $^\circ$, in the mass function equation to give a larger lower limit for the masses of the unseen components."292 These revised lower limits (also presented in Table 4)) are all greater than wwhich indicates that the unseen companions cannot be main sequence stars because if they were we should be able to deteet them (Alarshetal.1995)., These revised lower limits (also presented in Table \ref{res:mass}) ) are all greater than which indicates that the unseen companions cannot be main sequence stars because if they were we should be able to detect them \cite{mdd95}.293. The unseen companions must therefore. be also compact. objects. probably white cbwarfs.," The unseen companions must therefore be also compact objects, probably white dwarfs."294 We searched for the signature of the faint companions bv shifting out the fitted: radial. velocity for. cach binary and then looking for cdillerences in the line profiles. at the quacrature phases (Marshetal.1995)... Le. 0.25. and 0.75.," We searched for the signature of the faint companions by shifting out the fitted radial velocity for each binary and then looking for differences in the line profiles at the quadrature phases \cite{mdd95}, i.e. 0.25 and 0.75."295 “Phe spectra at quacdrature phases were obtained. by averaging the spectra contained in two separate phase ranges. Le. the spectra in the range from 0.1 to OA were averaged to obtain the phase 0.25 spectrum. ancl the spectra in the range from 0.6 to 0.9 to obtain the phase 0.75 spectrum.," The spectra at quadrature phases were obtained by averaging the spectra contained in two separate phase ranges, i.e. the spectra in the range from 0.1 to 0.4 were averaged to obtain the phase 0.25 spectrum, and the spectra in the range from 0.6 to 0.9 to obtain the phase 0.75 spectrum."296 The results are. plotted. in Fig. 4.., The results are plotted in Fig. \ref{res:comp}.297 Any contribution from the companion white dwarl should be seen as an asvmmetry in the line profile at phase 0.25 that is mürrored at phase 0.75 with respect to the rest wavelength, Any contribution from the companion white dwarf should be seen as an asymmetry in the line profile at phase 0.25 that is mirrored at phase 0.75 with respect to the rest wavelength298with the aromatic features for SINGS ITn galaxies. but that Sevterts and LINERs often exhibit excess IT» cluission. which they attribute to shocks.,"with the aromatic features for SINGS $\hii$ galaxies, but that Seyferts and LINERs often exhibit excess $_2$ emission, which they attribute to shocks."299" We explore the hypothesis that shocks cause both the excess IT, enission and the anomalous aromatic ratios for ACiNs in Section ??..", We explore the hypothesis that shocks cause both the excess $_2$ emission and the anomalous aromatic ratios for AGNs in Section \ref{discussion:h2}.300 Among the sources excluded. from the above analysis due to a lack of 6.2. 7.7. or 8.6 pau aromatic feature detections. there are a sigmificant nuuber with clearly detected 11.3 san features aud IT» $63) lines.," Among the sources excluded from the above analysis due to a lack of 6.2, 7.7, or 8.6 $\mu$ m aromatic feature detections, there are a significant number with clearly detected 11.3 $\mu$ m features and $_2$ S(3) lines."301 In Figure 7.. we show the nuclear spectra for a dozen of these sources. sorted by the equivalent width of the 11.5 pau feature.," In Figure \ref{fig:cut}, we show the nuclear spectra for a dozen of these sources, sorted by the equivalent width of the 11.3 $\mu$ m feature."302 These spectra exhibit the szinall L(7.7 ;nu)/L(11.3. jan) ratios and stroug IT $(3) lines characteristic of sources in the top-left of Figure 6.., These spectra exhibit the small L(7.7 $\mu$ m)/L(11.3 $\mu$ m) ratios and strong $_2$ S(3) lines characteristic of sources in the top-left of Figure \ref{fig:h2}.303" Due to uncertainties associated with estimating the streneth of weal. broad features aud determine robust upper limits (οσοι, proper continu placement). we do not imcelude auv of these sources in our subsequent analvsis."," Due to uncertainties associated with estimating the strength of weak, broad features and determining robust upper limits (e.g., proper continuum placement), we do not include any of these sources in our subsequent analysis."304 However. them behavior is consistent with that in Figure 6 aud supports the reality of the trend between aromatic feature characteristics and IL. line streneth.," However, their behavior is consistent with that in Figure \ref{fig:h2} and supports the reality of the trend between aromatic feature characteristics and $_2$ line strength."305 The sources with the largest τι ;n)/L(11.3. pan) ratios. NGC15 and NGC3079. also have the strougest silicate absorption19 features.," The sources with the largest L(7.7 $\mu$ m)/L(11.3 $\mu$ m) ratios, NGC4945 and NGC3079, also have the strongest silicate absorption features."306 This suggests that the 11.9 youn feature is being siguificauth attenuated. consistent with previous results for starburst and huninous infrared galaxies (e.e..Brandletal.2006:Pereira-Sanutacllaetal. 2010).. and implies that a sjenificaut fraction of the silicate-absorbing material is extended relative. to the regions that produce the aromatic features.," This suggests that the 11.3 $\mu$ m feature is being significantly attenuated, consistent with previous results for starburst and luminous infrared galaxies \citep[e.g.,][]{bra06,per10}, and implies that a significant fraction of the silicate-absorbing material is extended relative to the regions that produce the aromatic features."307 Although the aromatic feature measurements im PATIFIT are corrected for extinction. ln cases as extreme as these two galaxies the resulting feature strengths are highlv uncertain.," Although the aromatic feature measurements in PAHFIT are corrected for extinction, in cases as extreme as these two galaxies the resulting feature strengths are highly uncertain."308 For all other ealaxies in our sample. the inferred extinctious are «50% for all features.," For all other galaxies in our sample, the inferred extinctions are $<50\%$ for all features."309 The result that Sevfert ealaxies exhibit weak 6.2. 7.7. and 8.6 421. aromatic features relative to the 11.3 juu feature could be explained bv radiative or mechanical processing of the molecular carriers by the active nucleus.," The result that Seyfert galaxies exhibit weak 6.2, 7.7, and 8.6 $\mu$ m aromatic features relative to the 11.3 $\mu$ m feature could be explained by radiative or mechanical processing of the molecular carriers by the active nucleus."310 Tere we explore the relevant plivsical aud chemical effects that could modify the observed feature streugtlis., Here we explore the relevant physical and chemical effects that could modify the observed feature strengths.311" Previous experinenutal (οἱ,Szezepauski&Vala1993:Thidgins&Allamandola1995) and theoretical (e.¢..De-Freesetal.1993:Laughotf1996) work ou PAs has shown that the ο stretching modes that produce the 6.2 and 7.7 gan features. as well as the CII in-plane bending modes that produce the &.6 jan feature. are more ficiently excited i ionized molecules."," Previous experimental \citep[e.g.,][]{szc93,hud95} and theoretical \citep[e.g.,][]{def93,lan96} work on PAHs has shown that the C–C stretching modes that produce the 6.2 and 7.7 $\mu$ m features, as well as the C–H in-plane bending modes that produce the 8.6 $\mu$ m feature, are more efficiently excited in ionized molecules."312 The ratios of these features to the 11.3 pau feature. which is produced by €J out-of-plane bending modes. are lower for neutral," The ratios of these features to the 11.3 $\mu$ m feature, which is produced by C–H out-of-plane bending modes, are lower for neutral"313Methods using the spatial separation of the source of RFI and the radio source can effectively address this issue. see (Ellingson&HHampson2003.. Jeff et al.2005.. Kesteven2007.. Comwell et al.2004)).,"Methods using the spatial separation of the source of RFI and the radio source can effectively address this issue, see \cite{ellingson}, Jeff et \cite{jeff}, \cite{kesteven}, , Cornwell et \cite{corn}) )."314 Several types of correlators that are able to mitigate the impulse-like. strong RFI in both temporal and frequency domains have been studied in this paper.," Several types of correlators that are able to mitigate the impulse-like, strong RFI in both temporal and frequency domains have been studied in this paper."315 They are used in applications where input data are contaminated with outliers. and these correlators are statistically more stable thai correlator (1).," They are used in applications where input data are contaminated with outliers, and these correlators are statistically more stable than correlator (1)."316 Some of them could be used in radio astronomy. especially in radio interferometric systems withο correlators where the caleulation of visibilities i8. carriec out on general-purpose computers. as in (LOFAR2009— anc JIVE. Kruithof2009)).," Some of them could be used in radio astronomy, especially in radio interferometric systems with correlators where the calculation of visibilities is carried out on general-purpose computers, as in \cite{lofar} and JIVE, \cite{kruithof}) )."317 Software correlators are. by definition. much more flexible than traditional hardware correlators: any algorithm adapted or modified for a particular observation cai be optionally downloaded as a There are two operations in the numerator of (1): multiplication of the input samples of X and Y and summation (averaging).," Software correlators are, by definition, much more flexible than traditional hardware correlators: any algorithm adapted or modified for a particular observation can be optionally downloaded as a There are two operations in the numerator of (1): multiplication of the input samples of $X$ and $Y$ and summation (averaging)."318 Here it is proposed that they be modified to make the correlator more The new features can be introduced in the first operation to analyze the statistics of X and Y and to introduce a type of editing in order to eliminate The second operation. summation. which is considered às post-correlation averaging. can be divided into three steps: primary averaging over a time interval which is not too long to smooth RFI bursts appearing at this stage above the noise RFI mitigation and secondary averaging to the required time interval depending on the observational specifications (wavelength. baseline. radio source properties). see end of section 3 and Figs.," Here it is proposed that they be modified to make the correlator more The new features can be introduced in the first operation to analyze the statistics of $X$ and $Y$ and to introduce a type of editing in order to eliminate The second operation, summation, which is considered as post-correlation averaging, can be divided into three steps: primary averaging over a time interval which is not too long to smooth RFI bursts appearing at this stage above the noise RFI mitigation and secondary averaging to the required time interval depending on the observational specifications (wavelength, baseline, radio source properties), see end of section 3 and Figs."319 13 - 15., 13 - 15.320 Different types of correlators described 1n the following section are compared with Pearson’s correlator (1) using two |., Different types of correlators described in the following section are compared with Pearson's correlator (1) using two 1.321 The of the estimate p produced by RFI compared to the input correlation coefficient po: 2., The of the estimate $\widehat{\rho}$ produced by RFI compared to the input correlation coefficient $\rho_{0}$ 2.322 The effectiveness of an estimator is judged by the at the output of the correlator in both the presence and in the absence of outliers., The effectiveness of an estimator is judged by the at the output of the correlator in both the presence and in the absence of outliers.323 Computer simulations were performed to estimate these values. of the bias andrs.," Computer simulations were performed to estimate these values, of the bias and."324 Also some results of the processing of CSI data will be shown., Also some results of the processing of CS1 data will be shown.325" There are two classic estimators of correlation. coefficients using of samples instead of the samples themselves. 1970)). Let two input sample sequences xj....v, and γιη be sorted in ascending order: wy)<XsX..€x, and XpSMaSsXovg, The ith v"," There are two classic estimators of correlation coefficients using of samples instead of the samples themselves, \cite{kendall}) Let two input sample sequences $x_{1},...x_{n}$ and $y_{1},...y_{n}$ be sorted in ascending order: $x_{(1)}\leq x_{(2)}\leq ...\leq x_{(n)}$ and $y_{(1)}\leq y_{(2)}\leq ...\leq y_{(n)}$ ."326"alue xj, 1s called ith-order statistic.", The $i$ th value $x_{(i)}$ is called $i$ .327 Each sample x; has its kth position in the sorted series NMojeeVun.," Each sample $x_{j}$ has its $k$ th position in the sorted series $x_{(1)},...x_{(n)}$ ."328 This number& is the of the sample v; and is denoted by pj(2A), This number$k$ is the of the sample $x_{j}$ and is denoted by $p_{j}(=k)$.329 Similarly. the rank of y; is denoted by qj.," Similarly, the rank of $y_{j}$ is denoted by $q_{j}$."330 Let Gr.vp and Gv with i=ἐν and j=id... be two data pairs from the original data sequences.," Let $(x_{i},y_{i})$ and $(x_{j},y_{j})$ with $i=1,...n$ and $j=i+1,...n$ be two data pairs from the original data sequences."331 If p;— and qj;—q; have the same sign. these two data pairs are said to beconcordant. otherwise. they arediscordant.," If $p_{j}-p_{i}$ and $q_{j}-q_{i}$ have the same sign, these two data pairs are said to be, otherwise, they are."332 Let n. be the number of concordant pairs and ny the number of discordant pars., Let $n_{c}$ be the number of concordant pairs and $n_{d}$ the number of discordant pairs.333" It is clear that 2,+ay=n(i—1)/2.", It is clear that $n_{c}+n_{d}=n(n-1)/2$.334 The Spearman’s correlation coefficient is calculated by The bivariate correlation coefficient corresponding to Pearson’s rp can be restored using the relationship: Kendall's correlation coeficient is calculated by The bivariate correlation coefficient corresponding to Pearson’s rp can be restored using the relationship: One of the first constructions of a robust correlator Is based on the quarter square identity 1972.. Huber198 1)):where cov denotes covariance and var denotes variance.," The Spearman's correlation coefficient is calculated by The bivariate correlation coefficient corresponding to Pearson's $r_{P}$ can be restored using the relationship: Kendall's correlation coefficient is calculated by The bivariate correlation coefficient corresponding to Pearson's $r_{P}$ can be restored using the relationship: One of the first constructions of a robust correlator is based on the quarter square identity \cite{gnad}, \cite{huber}) ):where $cov$ denotes covariance and $var$ denotes variance."335 The correlation coefficient can be obtained with Therefore. any robust estimators of variance can be used for this correlator (Fridman2008)).," The correlation coefficient can be obtained with Therefore, any robust estimators of variance can be used for this correlator \cite{fridman}) )."336 Several of them are applied here to (7)., Several of them are applied here to (7).337" Samples Z,=X+Yand Z=X—Yare sorted in ascendingorder: zii)... zu."," Samples $Z_{1}=X+Y$and $Z_{2}=X-Y$are sorted in ascendingorder: $z_{(1)},...z_{(n)}$ ."338 Let y denote the chosen amount of trimming. οςy€OS and k= [yn].," Let $\gamma$ denote the chosen amount of trimming, $0 \le \gamma\le 0.5$ and $ k=[ \gamma n]$ ."339 Sample-trimmed variance is computed by removing κ of thelargest and & of the smallest, Sample-trimmed variance is computed by removing $k$ of thelargest and $k$ of the smallest340collision. strengths.,collision strengths.341" Using the line emissivities given by?. we deduce log(C.7/O7)2.0.6502. for 7,=10000-Ε1000 K. The ionization correction factor to deduce C/O from 7/O ? is expected to be small(2).and therefore we find (C/Oloswasc 0.3-0.7(CΟΥ. wherelog(C/0). =0.26(2)."," Using the line emissivities given by, we deduce $\log ({\rm C}^{+2}/{\rm O}^{+2}) = -0.6342\pm 0.2$, for $T_{\rm e} = 10\,000 \pm 1000$ K. The ionization correction factor to deduce C/O from $^{+2}$ $^{+2}$ is expected to be small,and therefore we find $({\rm C}/{\rm343 O})_{\rm CSWA\,5} \simeq 0.3$ $0.7 ({\rm C}/{\rm O})_{\odot}$, where $\log ({\rm C}/{\rm O})_\odot = -0.26$."344 As for nitrogen. the sub-solar C/O ratio in CSWA5 conforms to the established behaviour of C/O vs. O/H in metal-poor regions(22). Galactic stars(22). and Lyman break galaxies(22).," As for nitrogen, the sub-solar C/O ratio in CSWA 5 conforms to the established behaviour of C/O vs. O/H in metal-poor regions, Galactic stars, and Lyman break galaxies."345 Finally. we can obtain estimates of the age and stellar mass of the lensed galaxy by comparing its observed spectral energy distribution (SED). from the rest frame UV to the near-IR (from  2000A to 1 μπι where our X-shooter spectrum has the highest S/N ratio). to those of synthetic spectra computed with simple stellar population models23.," Finally, we can obtain estimates of the age and stellar mass of the lensed galaxy by comparing its observed spectral energy distribution (SED), from the rest frame UV to the near-IR (from $\sim3462000$ Å to $\sim 1.2 \mu$ m where our X-shooter spectrum has the highest S/N ratio), to those of synthetic spectra computed with simple stellar population models."347 For this purpose. we used the population synthesis code of with IMF. metallicity Z=Ü.4Z. (Section 5.4). and no internal reddening (Section 5.1).," For this purpose, we used the population synthesis code of with IMF, metallicity $Z = 0.4 Z_\odot$ (Section ), and no internal reddening (Section )."348 We generated two families of models with these parameters. for the two limiting cases of continuous star formation and an instantaneous burst. in each case varying the age from | to 70Myr.," We generated two families of models with these parameters, for the two limiting cases of continuous star formation and an instantaneous burst, in each case varying the age from 1 to 70 Myr."349" The best fitting models and the corresponding values of age and assembled stellar mass were determined by minimizing the value of 47. given by: where fi;y and a...) are the observed flux at wavelength À and its error respectively. and. fi,41, is the flux of the? model spectrum at the same wavelength."," The best fitting models and the corresponding values of age and assembled stellar mass were determined by minimizing the value of $\chi^2$, given by: where $f_{\mathrm{obs}, \lambda}$ and $\sigma_{\mathrm{obs}, \lambda}$ are the observed flux at wavelength $\lambda$ and its error respectively, and $f_{\mathrm{model}, \lambda}$ is the flux of the model spectrum at the same wavelength."350 The normalization factor b gives the stellar mass of the galaxy., The normalization factor $b$ gives the stellar mass of the galaxy.351 In fitting the model spectra to the X-shooter spectrum of image iml. we excluded regions affected by strong telluric absorption or by prominent residuals in the subtraction ofsky emission lines.," In fitting the model spectra to the X-shooter spectrum of image im1, we excluded regions affected by strong telluric absorption or by prominent residuals in the subtraction ofsky emission lines."352" We found that the best fitting models (see Figure 6) are those with ages of 40+10 Myr and stellar masses AJ;=(2.8+1.0)10""Lffieus M... where fii; is the unknown magnification factor of iml. and the error includes the 20% uncertainty in the absolute flux calibration (Section 2)."," We found that the best fitting models (see Figure ) are those with ages of $40 \pm 10$ Myr and stellar masses $M_\ast = (2.8 \pm3531.0) \times 10^9 \times 1/f_{\rm lens} \, {\rm M}_\odot$ , where $f_{\rm lens}$ is the unknown magnification factor of im1, and the error includes the 20% uncertainty in the absolute flux calibration (Section )."354 These values appear to be robust to the choice of star formation mode. with the single burst and continuous star formation models converging to. similar. solutions.," These values appear to be robust to the choice of star formation mode, with the single burst and continuous star formation models converging to similar solutions."355 We also investigated the possibility that the stellar continuum suffers a greater extinction than the emission lines. since the former could in principle sample a different (presumably older) stellar population than the latter.," We also investigated the possibility that the stellar continuum suffers a greater extinction than the emission lines, since the former could in principle sample a different (presumably older) stellar population than the latter."356 However. without recourse to the rest frame far-UV spectral range (which is inaccessible from the ground at >= 1.0686). we come up against the well-known age-extinction degeneracy.," However, without recourse to the rest frame far-UV spectral range (which is inaccessible from the ground at $z = 1.0686$ ), we come up against the well-known age-extinction degeneracy."357 For example. models in which the stellar continuum is reddened with a colour excess (713=0.3 also provide satisfactory fits to the X-shooter spectrum of CSWA5. albeit with younger preferred ages of 25+7 Myr and lower stellar masses of (2340.7).10°.14fiM.," For example, models in which the stellar continuum is reddened with a colour excess $E(B-V) = 0.3$ also provide satisfactory fits to the X-shooter spectrum of CSWA 5, albeit with younger preferred ages of $25 \pm 7$ Myr and lower stellar masses of $(2.3 \pm 0.7) \times 10^9 \times 1/f_{\rm lens} \, {\rm358 M}_\odot$."359 We can obtain an independent estimate of the age of the starburst from consideration of the equivalent. width of the HJ emission line. Wo(ll3) 100À. using the values of Wo(13) as a function of time caleulated with the spectral synthesis code(2).," We can obtain an independent estimate of the age of the starburst from consideration of the equivalent width of the $\beta$ emission line, $W_0({\rm H}\beta) \simeq 100$ Å, using the values of $W_0({\rm360 H}\beta)$ as a function of time calculated with the spectral synthesis code."361 The equivalent widths of the Balmer recombination lines fall rapidly following a burst of star formation and Wo(ll?)« 100À for all ages greater than 5Myr.," The equivalent widths of the Balmer recombination lines fall rapidly following a burst of star formation and $W_0({\rm H}\beta) <362100$ Å for all ages greater than 5 Myr."363 On the other hand. the time dependence is less steep for continuous star formation and ages of ~25 Myr are indicated. in better agreement the age inferred from the analysis of the UV to near-IR SED.," On the other hand, the time dependence is less steep for continuous star formation and ages of $\sim 25$ Myr are indicated, in better agreement the age inferred from the analysis of the UV to near-IR SED."364 Summarizing our findings. we have established that CSWA5 is an actively star-forming galaxy at 2=—1.0686 lensed by an apparent foreground group of massive red galaxies. at least one of which is at 2=0.3877.," Summarizing our findings, we have established that CSWA 5 is an actively star-forming galaxy at $z = 1.0686$ lensed by an apparent foreground group of massive red galaxies, at least one of which is at $z = 0.3877$."365 The lensed source is forming stars at arate SERc20M. yr+. suffers negligible reddening. has an oxygen abundance of approximately half-solar. and sub-solar N/O and C/O ratios (as expected for low metallicity galaxies).," The lensed source is forming stars at a rate ${\rm SFR}366\simeq 20\, {\rm M}_\odot$ $^{-1}$, suffers negligible reddening, has an oxygen abundance of approximately half-solar, and sub-solar N/O and C/O ratios (as expected for low metallicity galaxies)."367 Its blue spectral energy distribution. from the rest-frame UV to the IR. and high H.? equivalent width are indicative of a young age. only 25-50 Myr. during which time the galaxy has assembled a stellar mass of ~310M...," Its blue spectral energy distribution, from the rest-frame UV to the near-IR, and high $\beta$ equivalent width are indicative of a young age, only $\sim 25$ –50 Myr, during which time the galaxy has assembled a stellar mass of $\sim 3 \times 10^9 {\rm M}_\odot$."368 The star-formation rate and stellar mass areuncertain by- an unknown magnification factor. which may be of order unity for each of the four gravitationally lensed images of CWSA5.," The star-formation rate and stellar mass areuncertain by an unknown magnification factor, which may be of order unity for each of the four gravitationally lensed images of CWSA 5."369 These physical characteristics are broadly in line with those of the population of UV-selected galaxies at redshifts ;=1—3 and referencestherein). as we now discuss.," These physical characteristics are broadly in line with those of the population of UV-selected galaxies at redshifts $z = 1$ –3 \citep[e.g.][and references, as we now discuss."370 The most relevant comparison here is with the metallicity relation determined by for galaxies at a mean i25——(0.7 selected from the Gemini Deep Deep Survey (GDDS) and shown in Figure7., The most relevant comparison here is with the mass-metallicity relation determined by for galaxies at a mean $\langle z \rangle \simeq 0.7$ selected from the Gemini Deep Deep Survey (GDDS) and shown in Figure.371 It must be borne in mind here that such comparisons are fraught with pitfalls for the unwary. because of systematic differences between the methods used to determine both the stellar mass and the metallicity.," It must be borne in mind here that such comparisons are fraught with pitfalls for the unwary, because of systematic differences between the methods used to determine both the stellar mass and the metallicity."372 In this particular case. the study by should be compatible with ours. as both used a IMF in arriving at the stellar masses and the #28 index for the oxygen abundance (although the reliability of the latter at apparently super-solar metallicities is?).," In this particular case, the study by should be compatible with ours, as both used a IMF in arriving at the stellar masses and the $R23$ index for the oxygen abundance <cit.>[although the reliability of the latter at apparently super-solar metallicities is."373 An additional complication is the unknown magnification factor which applies to our derivation of the stellar mass., An additional complication is the unknown magnification factor which applies to our derivation of the stellar mass.374 If. fi... in Section is of order unity. then Figure shows that CSWA5 is somewhat metal-poor for its stellar mass. compared to GDDS galaxies at 2~0.7.," If $f_{\rm lens}$ in Section is of order unity, then Figure shows that CSWA 5 is somewhat metal-poor for its stellar mass, compared to GDDS galaxies at $z \sim 0.7$."375 Detinite conclusions are difficult at this stage. given the seatter of the GDDS galaxies about the mean relation in Figure7. and the lack of objects as metal-poor as CSWA5 in the sample considered by?.," Definite conclusions are difficult at this stage, given the scatter of the GDDS galaxies about the mean relation in Figure, and the lack of objects as metal-poor as CSWA 5 in the sample considered by."376 Taken at face value. the offset of CSWAS in Figure is in the same sense as the more general redshift evolution of the mass-metallicity relation proposed by?.," Taken at face value, the offset of CSWA 5 in Figure is in the same sense as the more general redshift evolution of the mass-metallicity relation proposed by."377 Since the lens magnification factor enters into the calculation of both SFR and Ad... the ratio of these two quantities. commonly referred to as the specific star formation rate (SSFR). should not be affected.," Since the lens magnification factor enters into the calculation of both SFR and $M_\ast$, the ratio of these two quantities, commonly referred to as the specific star formation rate (SSFR), should not be affected."378" For CSWA5 we deduce a specitie SER.=(8.2+ ""vr +."," For CSWA 5 we deduce a specific ${\rm SFR} = (8.2 \pm 2.2)379\times 10^{-9}$ $^{-1}$ ."380 The reciprocal of this value gives a timescale for the star formation activity of ~120+45 Mvr. which is 3 times higher than the age deduced in Section from SED fitting.," The reciprocal of this value gives a timescale for the star formation activity of $\sim 120 \pm 45$ Myr, which is $\sim 3$ times higher than the age deduced in Section from SED fitting."381 In Figure we compare the value of SSFR in CSWAS with, In Figure we compare the value of SSFR in CSWA 5 with382iustabiliVv erows towards the very latest stages of stellar evolution. however he ZOUCS OVCE which it is acting are extremely narrow and there is not enough time left yofore the supernova explosion for a siguificant musing QO OCCUTY.,"instability grows towards the very latest stages of stellar evolution, however the zones over which it is acting are extremely narrow and there is not enough time left before the supernova explosion for a significant mixing to occur."383 Thus. even when the inhibiting effect of the Hosradieut is reduced by horizontal turbulence. the GSF iustabili vds unable to smooth the steep €) eradicuts aud o significantly transport matter.," Thus, even when the inhibiting effect of the $\mu$ –gradient is reduced by horizontal turbulence, the GSF instability is unable to smooth the steep $\Omega$ –gradients and to significantly transport matter."384 We conclude that the wuplitude aud spatial extension of the CSF instability makes it tunable ο reduce he aueular miolüentuni οἳ the stellar cores di the pre-supernova stages by two orders of lagutucde., We conclude that the amplitude and spatial extension of the GSF instability makes it unable to reduce the angular momentum of the stellar cores in the pre-supernova stages by two orders of magnitude.385" Therefore. other mechauisus such as naguctic fields (Sprit 2002.. Alacder&Alevnet2001. Matlis&Zahn 2005... Zahuetal. 2007)) and eyavity Waves (Talon&Charbonnel 2005... Mathisetal. 2008)) 11115 be further The coefficient δες requires. because οἱ the horizontal turbulence. the knowledge of the compoucuts U, aud Y5 of the meridional circulation."," Therefore, other mechanisms such as magnetic fields \cite{Sp02}, , \cite{ROTMII}, \cite{MZ05}, , \cite{ZBM07}) ) and gravity waves \cite{TC05}, , \cite{MTPZ08}) ) must be further The coefficient $D_{\mathrm{GSF}}$ requires, because of the horizontal turbulence, the knowledge of the components $U_2$ and $V_2$ of the meridional circulation."386 If the solutious of the order system of equations governing meridional circulation are not available. some approximations may be considered.," If the solutions of the $^{th}$ order system of equations governing meridional circulation are not available, some approximations may be considered."387 We note that the same problem would occur for Eq. €1)), We note that the same problem would occur for Eq. \ref{vgsf}) )388 bx Endal aud Sofia (1978))., by Endal and Sofia \cite{EndalS78}) ).389 As shown hv stellar models. the orders of mmaenitude of (> aud V5 are the same.," As shown by stellar models, the orders of magnitude of $U_2$ and $V_2$ are the same."390 The muuerical models give in general Y5—05/3 and |2Y5al|~Y»., The numerical models give in general $V_2 \sim U_2/3$ and $\left|2V_2-\alpha U_2 \right| \sim V_2$.391 Using these orders of magnitude iu Eq. (8)).," Using these orders of magnitude in Eq. \ref{o2}) ),"392 we ect For £5. various expressions can be usec taking iuto account the amount of differential rotation Maedoer 20093).," we get For $U_2$, various expressions can be used taking into account the amount of differential rotation \cite{Maeder09}) )."393 We can also ect an order of magnitude using the approxinatiou for a mixture of perfect eas and radiation with a local angular velocity O(r). ignoring the effects of differcutial rotation ou the circulation velocity andthe Cratton-Oppik terii which is laree only in the outer lavers. where the various quautitieshavetheir usual meaniue.," We can also get an order of magnitude using the approximation for a mixture of perfect gas and radiation with a local angular velocity $\Omega(r)$, ignoring the effects of differential rotation on the circulation velocity andthe Gratton-Öppik term which is large only in the outer layers, where the various quantitieshavetheir usual meaning."394model of jet/uozzle aud the comparison between model predictions and observations: in Sect.,model of jet/nozzle and the comparison between model predictions and observations; in Sect.395 [ we discuss the results and draw our conchisious., 4 we discuss the results and draw our conclusions.396" We preseut the new Chandra/ACIS-S data of HII 151 (PI Schucider: ObsID 11016: f£,,,,=65.2 ks) ceutered at (0101121005. | 18:07:51.95). CRISS). performed on 29 Dec. 2009. that provides δ vears of time base between the first (2001. Favatactal. 2002)) aud the last Chandra observations of TIT 15L. with the itermeciate observation of 2005 (Favataetal. 2006))."," We present the new Chandra/ACIS-S data of HH 154 (PI Schneider; ObsID 11016; $t_{exp} = 65.2$ ks) centered at (04:31:34.998, $+$ 18:07:51.95) (FK5), performed on 29 Dec. 2009, that provides $8$ years of time base between the first (2001, \citealt{ffm02}) ) and the last Chandra observations of HH 154, with the intermediate observation of 2005 \citealt{fbm06}) )."397 To derive a nuiform comparison between the three Chandra data-sets; we reprocessed all the data with the same method. using the latest CIAO L3 package. aud filtered the data to restrict the energy to the 0.3.LO keV. band. as in Favatactal.(2006).," To derive a uniform comparison between the three Chandra data-sets, we reprocessed all the data with the same method, using the latest CIAO 4.3 package, and filtered the data to restrict the energy to the $0.3 - 4.0$ keV band, as in \citet{fbm06}."398.. Events were extracted. for all observations from regions around the source and the background. near the position of the ΠΠ 151 jet (L:31:5010. | 8:08:01.9). following Ballyetal. (2003)..," Events were extracted for all observations from regions around the source and the background, near the position of the HH 154 jet (4:31:34.10, $+$ 18:08:04.9), following \citet{bfr03}. ."399" We defiued a box 3.5%«5"" in the 2001 data-set. a box 3.5""45.5"" in the 2005 data-set. aud a box 3”<6"" in the 2009data-sct?.. to extract tle source eveuts. and four 10 pixel radius circles for the background. frour source-free parts of the image around the source."," We defined a box $3.5''\times5''$ in the 2001 data-set, a box $3.5''\times5.5''$ in the 2005 data-set, and a box $3''\times6''$ in the 2009, to extract the source events, and four $10$ pixel radius circles for the background, from source-free parts of the image around the source."400 We extracted the spectra of the three data-sets aud produced the ancillary response files. arf and zxiuf. bv usine SPECENTRACT.," We extracted the spectra of the three data-sets and produced the ancillary response files, arf and rmf, by using SPECEXTRACT."401 We grouped the spectra to have a iinimaun of 5 counts per bin., We grouped the spectra to have a minimum of $5$ counts per bin.402 We also applied the sub-pixel repositioning aleorithin available in CIAO. L3 (EDSER) to the Chandra images to refine the event positions (Lictal. 2001))., We also applied the sub-pixel repositioning algorithm available in CIAO 4.3 (EDSER) to the Chandra images to refine the event positions \citealt{lkp04}) ).403 We resample the improved ππασος at 0.25% pixel size. obtaining images with one-half of the native ACTS pixel scale.," We resample the improved images at $0.25''$ pixel size, obtaining images with one-half of the native ACIS pixel scale."404 Fig., Fig.405 dl. shows the XN-rav source associated with III 151 in 2001 (first panels). 2005 (second panels). aud 2009 (third panels). aud the X-ray maps svuthesized from our hvdrodyvnuanuie model (diseussed in Sect. 3))," \ref{mappa-X-bin} shows the X-ray source associated with HH 154 in 2001 (first panels), 2005 (second panels), and 2009 (third panels), and the X-ray maps synthesized from our hydrodynamic model (discussed in Sect. \ref{The model}) )"406 with the same spatial resolution as ACIS (last panels)., with the same spatial resolution as ACIS (last panels).407" The upper panels show the images with the native ACTS pixel size (0.57): the central paucls show the nuages with sub-pixel repositioning algorithi applied aud resampled with a pixel size 0.25"": the lower panels show the improved resainpled images smioothed using a Gaussian keruel of width 0.5"".", The upper panels show the images with the native ACIS pixel size $0.5''$ ); the central panels show the images with sub-pixel repositioning algorithm applied and resampled with a pixel size $0.25''$; the lower panels show the improved resampled images smoothed using a Gaussian kernel of width $0.5''$.408 The three Chandra data-sets show that the X-rav Cluission is αλα located at the base of the IIIT 151 jet im all epochs. uear the driving source.," The three Chandra data-sets show that the X-ray emission is mainly located at the base of the HH 154 jet in all epochs, near the driving source."409" Ballyetal.(2003). found that the N-vayv source is displaced by 0.5""1” with respect to the L1551 IRS5 drivius source.", \citet{bfr03} found that the X-ray source is displaced by $0.5''-1''$ with respect to the L1551 IRS5 driving source.410 The X-ray source cousists of a bright knot which appears to be stationary in the time period covered by the observations Gu the following the “stationary” component) and an clongatecd structure directed: away from the driving source Gu the following the elongated conrponeut] showing variability iu the three data-scts (sce Fig 1)).," The X-ray source consists of a bright knot which appears to be stationary in the time period covered by the observations (in the following the ""stationary"" component) and an elongated structure directed away from the driving source (in the following the ""elongated"" component) showing variability in the three data-sets (see Fig \ref{mappa-X-bin}) )."411 In particular. the latter component appears as a knot in unaeges with EDSER applied (ceutral and lower panels in Fie. 1))," In particular, the latter component appears as a knot in images with EDSER applied (central and lower panels in Fig. \ref{mappa-X-bin}) )"412 which is much. fainter than the snot of the stationary component., which is much fainter than the knot of the stationary component.413 The position of the aut kuot appears to be different iu the three data-scts., The position of the faint knot appears to be different in the three data-sets.414" Tn particular. for the 2001 and 2005 data. we coufim he results of Favataetal.(2006).. 1.0. an clongation of he XN-rav source corresponding to a proper motion of V1"" vy approximately westward away from the driving source of the ΠΠ 151 jet."," In particular, for the 2001 and 2005 data, we confirm the results of \citet{fbm06}, i.e. an elongation of the X-ray source corresponding to a proper motion of $0.7''$ /yr approximately westward away from the driving source of the HH 154 jet."415 As for the 2009 observations. here is an hint of a faint knot closer to the brighter stationary source than in 2005.," As for the 2009 observations, there is an hint of a faint knot closer to the brighter stationary source than in 2005."416" Note that recentlv an asviunietrv in the Chandra PSF has beendiscovered”?.. ocated at PA.=1ου.rollangle(425)"". corresponding o PLA.=269.1"" in 2009 observations. i.c. roughly the direction of the extension of the X-ray source detected in ΠΠ 151."," Note that recently an asymmetry in the Chandra PSF has been, located at $P.A. = 195 -rollangle (\pm25)^{\circ}$, corresponding to $P.A. = 269.4^{\circ}$ in 2009 observations, i.e. roughly the direction of the extension of the X-ray source detected in HH 154."417 We checked if this iustruueutal effect may influence the observed morphology of tlhe N-rav source vl found that the asvuunetrv of the PSF does uot affect our images on scales lavecr than 1 arcsec., We checked if this instrumental effect may influence the observed morphology of the X-ray source and found that the asymmetry of the PSF does not affect our images on scales larger than 1 arcsec.418 The ongated structure visible in the tages. therefore. is uot an artifact of the iustrmucut.," The elongated structure visible in the images, therefore, is not an artifact of the instrument."419" The maxima lenetl the whole N-rav source is 2:5"" (=TOO AU at D=110 pe).", The maximum length of the whole X-ray source is $\approx5''$ $\approx700$ AU at $D=140$ pc).420 We performed a spectral analysis ofthe individual datas, We performed a spectral analysis of the individual data.421ets! All spectra are well fitted by an absorbed thermal plasina (APEC in NSPEC)., All spectra are well fitted by an absorbed thermal plasma (APEC in XSPEC).422 The count rates of the three observations (0.76+0.10cuts/ks in 2001. 0.65+0.08 cuts/ks in 2005. aud 0.8940.12 cuts/ks in 2009) are compatible within the Poisson error. mdicatiug that 1ο N-rav huuiuositv is constant.," The count rates of the three observations $0.76\pm0.10$cnts/ks in 2001, $0.65\pm0.08$ cnts/ks in 2005, and $0.89\pm0.12$ cnts/ks in 2009) are compatible within the Poisson error, indicating that the X-ray luminosity is constant."423 Moreover the source shows no spectral variability in terms of temperature (T7) and enmisson measure (LAL) in the timescale analyzed rere., Moreover the source shows no spectral variability in terms of temperature $T$ ) and emission measure $EM$ ) in the timescale analyzed here.424 We have thus fit the three data-sets sinmltauecouslv &uding the values reported in Table 1.., We have thus fit the three data-sets simultaneously finding the values reported in Table \ref{fit}.425 The Ay is well eoustrained (more than in the analysis of NADENewtou data: Favatactal. 2002)) by fitting he three data-sets simultaneously but it cannot he constrained by the iudividual data-scts., The $N_{\rm H}$ is well constrained (more than in the analysis of XMM-Newton data; \citealt{ffm02}) ) by fitting the three data-sets simultaneously but it cannot be constrained by the individual data-sets.426 In the latter case we fixed Nyy to the value derived in the simultaneous fittine of the three data-sets., In the latter case we fixed $N_{\rm H}$ to the value derived in the simultaneous fitting of the three data-sets.427 The three data-sets aud he best-fit model are shown in Fie. 2.., The three data-sets and the best-fit model are shown in Fig. \ref{spettri}. .428 To investigate possible spatial variations of the spectral properties. we selected two regions. confining the stationary component (blue box in Fie. 1))," To investigate possible spatial variations of the spectral properties, we selected two regions, confining the stationary component (blue box in Fig. \ref{mappa-X-bin}) )"429 aud the clongated component (ereen box in Fie. 1)).," and the elongated component (green box in Fig. \ref{mappa-X-bin}) ),"430 aud coniputed the median photon energy. MPE. within cach reeion in 2001. 2005. and 2009.," and computed the median photon energy, $MPE$, within each region in 2001, 2005, and 2009."431 ALPE has been proved to be a robust indicator of the spectral properties of a source iu the case of low statistics (ITlongotal.200 £))., $MPE$ has been proved to be a robust indicator of the spectral properties of a source in the case of low statistics \citealt{hsg04}) ).432 Tn each data-set the total number of counts considered is zDO. in the stationary component. and ~10. in the clongated component.," In each data-set the total number of counts considered is $\approx50$, in the stationary component, and $\sim10$, in the elongated component."433" The stationary component shows no temporal variability in the three epochs. its AIPE varvingin the rauge MPE,=1.35.1.12 keV. The median energy of the elongated component. MPE. is always lower than M PE..."," The stationary component shows no temporal variability in the three epochs, its MPE varyingin the range $MPE_{\rm s} = 1.35434- 1.42$ keV. The median energy of the elongated component $MPE_{\rm e}$ is always lower than $MPE_{\rm s}$ ."435" By considering the threeobservations altogether ALPE,z1.1 keV. while the elongated component appears softer with A/PE,z1.0 keV. This result is supported by the I&oliiogorov-Siiirnov test we performed to check if the photon cucrev"," By considering the threeobservations altogether $MPE_{\rm s} \approx 1.4$ keV, while the elongated component appears softer with $MPE_{\rm e} \approx 1.0$ keV. This result is supported by the Kolmogorov-Smirnov test we performed to check if the photon energy"436"signal showed no correlation with RV, and a planetary origin was attributed to the signal.","signal showed no correlation with RV, and a planetary origin was attributed to the signal."437" We obtained a set of 28 CORALIE RV measurements, spanning 55 days."," We obtained a set of 28 CORALIE RV measurements, spanning 55 days."438 CORALIE is the fiber-fed echelle spectrograph mounted at the Swiss telescope at La Silla observatory., CORALIE is the fiber-fed echelle spectrograph mounted at the Swiss telescope at La Silla observatory.439 High-precision RV measurements are obtained by cross-correlating the spectra with a template mask (??).. ," High-precision RV measurements are obtained by cross-correlating the spectra with a template mask \citep{1996A&AS..119..373B, 2002A&A...388..632P}."440"The observations are reduced online, allowing a real-time calculation of the RV and photon noise estimation."," The observations are reduced online, allowing a real-time calculation of the RV and photon noise estimation."441 Previous campaigns showed that CORALIE can reach a long-term precision of mm/s (e.g.?).., Previous campaigns showed that CORALIE can reach a long-term precision of m/s \citep[e.g.][]{2009arXiv0908.1479S}.442" We also calculated the bisector velocity span of the cross-correlation function, following the procedure described in ?.."," We also calculated the bisector velocity span of the cross-correlation function, following the procedure described in \cite{2001A&A...379..279Q}."443" To do so we calculate the line that bisects the cross-correlation function; the and the are defined as the average bisector values for the ranges of 10 to and 60 to of the line depth, respectively."," To do so we calculate the line that bisects the cross-correlation function; the and the are defined as the average bisector values for the ranges of 10 to and 60 to of the line depth, respectively."444 The bisector span (henceforth BIS) is then the inverse of the slope of the line that connects the to thebottom., The bisector span (henceforth BIS) is then the inverse of the slope of the line that connects the to the.445 Our data points were obtained during two observing campaigns., Our data points were obtained during two observing campaigns.446" The first set contains 20 measurements spanning 21 days, from the 21 December 2009 to the 10 January 2010; the second has 8 measurements obtained in a mission of 10 days, from the 4 to the 14 February 2010."," The first set contains 20 measurements spanning 21 days, from the 21 December 2009 to the 10 January 2010; the second has 8 measurements obtained in a mission of 10 days, from the 4 to the 14 February 2010."447" The precision on the RVs is dominated by the photon noise contribution, with an average precision of mm/s. The weighted r.m.s."," The precision on the RVs is dominated by the photon noise contribution, with an average precision of m/s. The weighted r.m.s."448 of the data is of mm/s and the peak-to-peak amplitude is of mm/s. The data are presented in Fig., of the data is of m/s and the peak-to-peak amplitude is of m/s. The data are presented in Fig.449 1 (top panel)., \ref{RV_time} (top panel).450 The first point to note is the different amplitude of the RV variation on the two campaigns., The first point to note is the different amplitude of the RV variation on the two campaigns.451 On the first campaign the RVs are characterized by a weighted r.m.s., On the first campaign the RVs are characterized by a weighted r.m.s.452" of mm/s and a peak-to-peak amplitude is of mm/s, while for the second these values are of mm/s and mm/s. The average uncertainties of the measurements are very similar, of 12.02 and mm/s respectively, and cannot account for the discrepancy."," of m/s and a peak-to-peak amplitude is of m/s, while for the second these values are of m/s and m/s. The average uncertainties of the measurements are very similar, of 12.02 and m/s respectively, and cannot account for the discrepancy."453" In both cases the RV variations are smaller than the ones reported by ? by a factor of 10 and 5, respectively, and cover a time span of three times the published orbital period for the first campaign, and once for the second."," In both cases the RV variations are smaller than the ones reported by \cite{2009arXiv0912.2773H} by a factor of 10 and 5, respectively, and cover a time span of three times the published orbital period for the first campaign, and once for the second."454" A phase-folded plot on the announced orbit depicts very well this discrepancy (Fig. 1,,"," A phase-folded plot on the announced orbit depicts very well this discrepancy (Fig. \ref{RV_time},"455 bottom panel)., bottom panel).456" We only considered the first of the two proposed orbits, but they only differ slightly and the same conclusions hold."," We only considered the first of the two proposed orbits, but they only differ slightly and the same conclusions hold."457" Still, the weighted r.m.s."," Still, the weighted r.m.s."458 is well in excess of the average measurement precision., is well in excess of the average measurement precision.459 In order to evaluate the presence of a periodic signal in the data we used two different approaches., In order to evaluate the presence of a periodic signal in the data we used two different approaches.460" The first was the ""string-length"" method described in ?.."," The first was the “string-length"" method described in \cite{1983MNRAS.203..917D}."461" This method delivers the orbital period that minimizes the sum of the lengths of line segments in a (RV;, ¢;) diagram."," This method delivers the orbital period that minimizes the sum of the lengths of line segments in a $_i$, $\phi_i$ ) diagram."462" It is suitable for randomly spaced observations in small data sets, such as ours, and is very efficient in detecting single planets in eccentric orbits."," It is suitable for randomly spaced observations in small data sets, such as ours, and is very efficient in detecting single planets in eccentric orbits."463 Using it we find a period of ddays (and a To of ==22455238.6)., Using it we find a period of days (and a $_0$ of 2455238.6).464" If only the data from the first campaign are used the period changes slightly, to 2.765 days."," If only the data from the first campaign are used the period changes slightly, to 2.765 days."465 The phase-folded RV measurements for both data sets and periods are shown in Fig. 2.., The phase-folded RV measurements for both data sets and periods are shown in Fig. \ref{RV_folded}.466" Note that while the two data sets show different RV amplitude variations, the periodicity of the signal is not significantly affected by the inclusion of the data from the second campaign."," Note that while the two data sets show different RV amplitude variations, the periodicity of the signal is not significantly affected by the inclusion of the data from the second campaign."467" In particular, the data from the first campaign show a well-defined variation when phase-folded (bottom panel)."," In particular, the data from the first campaign show a well-defined variation when phase-folded (bottom panel)."468" We also computed the generalized Lomb-Scargle periodogram (asimplementedby?),, which revealed the presence of a 1.55 and ddays signal in both the RVs and in the BIS (Fig. 3))."," We also computed the generalized Lomb-Scargle periodogram \citep[as implemented by ][]{2009A&A...496..577Z}, which revealed the presence of a 1.55 and days signal in both the RVs and in the BIS (Fig. \ref{GLS}) )."469 The ddays signal is the alias of the 2.8ddays signal created by the ἆάαγ sampling., The days signal is the alias of the days signal created by the day sampling.470" The BIS shows a significant variation, of the same order of the RV, and correlated with it."," The BIS shows a significant variation, of the same order of the RV, and correlated with it."471 The BIS-RV plot is present in Fig.4;; note that the data from the second campaign show extreme values of both RV and BIS., The BIS-RV plot is present in \ref{BIS_RV}; note that the data from the second campaign show extreme values of both RV and BIS.472 A linear least squares fit delivers a slope of -0.826., A linear least squares fit delivers a slope of -0.826.473" The Pearson's correlation factor is of -0.747, and Monte Carlo simulations show that the probability of obtaining this value or lower from two random non-correlated distributions with 28 points is « 1107* and thus negligible."," The Pearson's correlation factor is of -0.747, and Monte Carlo simulations show that the probability of obtaining this value or lower from two random non-correlated distributions with 28 points is $<$ $^{-4}$ and thus negligible."474source JL206+3941. approximately 2° away from the galaxy.,"source J1206+3941, approximately $2^\circ$ away from the galaxy."475 This calibrator now has an absolute position determined to 1 mas accuracy (Beasleyοἱal.2002) in the International Celestial Relerence Frame (Maetal.1998)., This calibrator now has an absolute position determined to 1 mas accuracy \citep{bea02} in the International Celestial Reference Frame \citep{ma98}.476. That position uncertainty dominates (he absolute position errors for the NGC 4151 VLBI components. since ionospheric errors al 5 GlIz and hieher should be less than 1 mas for a 2° phase-relerence distance (Chatterjeeetal.2004).," That position uncertainty dominates the absolute position errors for the NGC 4151 VLBI components, since ionospheric errors at 5 GHz and higher should be less than 1 mas for a $2^\circ$ phase-reference distance \citep{cha04}."477. Therefore. the absolute position error for NGC 4151 is taken to be | mas.," Therefore, the absolute position error for NGC 4151 is taken to be 1 mas."478 The evele (ime was about 12 minutes in 1998. and 2.54.5 minutes in 2002.," The phase-referencing cycle time was about 12 minutes in 1998, and 2.5–4.5 minutes in 2002."479 A journal ol the observations. including the total integration time on the galaxy. is provided in Table 1..," A journal of the observations, including the total integration time on the galaxy, is provided in Table \ref{tab:obs}."480 Note that the need for frequent calibration meant that only about of the scheduled 10 hours was spent actually integrating on the galaxy in 2002., Note that the need for frequent calibration meant that only about of the scheduled 10 hours was spent actually integrating on the galaxy in 2002.481 All data calibration was carried out in NRAQO’s Astronomical Image Processing Svstem. AIPS (Greisen2003).," All data calibration was carried out in NRAO's Astronomical Image Processing System, AIPS \citep{gre03}."482. Amplitude calibration for the VLBA antennas in 1998 and 2002 was done by using the values for (he antenna gains maintained by VLBA personnel as well as nmeasurenienis of svsteni temperature mace once per minute during (he observations: (his calibration is believed to have an accuracy of5%., Amplitude calibration for the VLBA antennas in 1998 and 2002 was done by using the values for the antenna gains maintained by VLBA personnel as well as measurements of system temperature made once per minute during the observations; this calibration is believed to have an accuracy of.483.. A similar method was used for calibration ol Eb during 2002: approximately once per hour. that telescope was removed [rom the VLBI array in order (ο spend 15 minutes on pointing observations.," A similar method was used for calibration of Eb during 2002; approximately once per hour, that telescope was removed from the VLBI array in order to spend 15 minutes on pointing observations."484 The 27 VLA antennas were made mutually coherent by adjusting their individual phases periodically during short observations of the compact sources J12064-3941 (in 1993) and J11464-3958 (in 2002)., The 27 VLA antennas were made mutually coherent by adjusting their individual phases periodically during short observations of the compact sources J1206+3941 (in 1998) and J1146+3958 (in 2002).485 The amplitude response of the individual VLA antennas was calibrated with respect to the standard amplitude calibrator J13312-3030 (3C 286). on the scale of Baarsetal.(1977).. as modified slightly bv the most recent measurements at the VLA.," The amplitude response of the individual VLA antennas was calibrated with respect to the standard amplitude calibrator J1331+3030 (3C 286), on the scale of \citet{baa77}, as modified slightly by the most recent measurements at the VLA."486 This amplitude calibration then was transferred to the reference source J12062-3941. enabling ealeulation of the complex gain. which was interpolated in time to derive (he VLA calibration for NGC 4151.," This amplitude calibration then was transferred to the reference source J1206+3941, enabling calculation of the complex gain, which was interpolated in time to derive the VLA calibration for NGC 4151."487 For the GBT. we initially assumed a flat gain curve and aperture efficiency based on commissioning measurements.," For the GBT, we initially assumed a flat gain curve and aperture efficiency based on commissioning measurements."488 This initial calibration was assessed using the amplitude check calibrators J09274-3902. and. J1310+3220. and we found it necessary {ο reduce the amplitude (voltage) gain by to bring the 15 GlIz calibration in line with the VLBA scale.," This initial calibration was assessed using the amplitude check calibrators J0927+3902 and J1310+3220, and we found it necessary to reduce the amplitude (voltage) gain by to bring the 15 GHz calibration in line with the VLBA scale."489 Similar analvsis resulted in a eain reduction for Eb. and small adjustments for several VLBA antennas. but no modifications to (he a priori calibration were necessary al 5 and 8 Gllz.," Similar analysis resulted in a gain reduction for Eb, and small adjustments for several VLBA antennas, but no modifications to the a priori calibration were necessary at 5 and 8 GHz."490 For both 1998, For both 1998491(Dm 20pt 20pt 20pt Molecular regions m the Galaxy exist within a wide range of cuviromuental conditions.,.75in 20pt 20pt 20pt Molecular regions in the Galaxy exist within a wide range of environmental conditions.492 There are massive eiut molecular clouds near the Galactic Center with lavee mean densities (Bally 1988). highly excited molecular gas associated with ionization frouts and supernova remuauts (Elmeercen Lada 1977). quiesceut clouds aud globules (Clemens Darvaiuis 1988). aud diffuse. hieh latitude clouds with low coluun densities in the solar neighborhood (Alagnani. Blitz. Muudy 1985).," There are massive giant molecular clouds near the Galactic Center with large mean densities (Bally 1988), highly excited molecular gas associated with ionization fronts and supernova remnants (Elmegreen Lada 1977), quiescent clouds and globules (Clemens Barvainis 1988), and diffuse, high latitude clouds with low column densities in the solar neighborhood (Magnani, Blitz, Mundy 1985)."493 In. addition to local sources of perturbation. the molecular gas responds to large scale variations in the Galaxy such as spiral potentials aud the surface density of stars aud eas (Elucercen 19893.," In addition to local sources of perturbation, the molecular gas responds to large scale variations in the Galaxy such as spiral potentials and the surface density of stars and gas (Elmegreen 1989)."494 These ditfereut euvironmnenuts aud conditions regulate the stability of the eas aud therefore. modulate the formation of stars.," These different environments and conditions regulate the stability of the gas and therefore, modulate the formation of stars."495 Therefore. it is iniportaut to evaluate the molecular eas properties over a wide range of environmoeuts.," Therefore, it is important to evaluate the molecular gas properties over a wide range of environments."496 A general description of the molecular interstellar τος requires surveys of molecular line ciission over laree voluues of the Calaxy with high augular aud spectral resolution aud sampling., A general description of the molecular interstellar medium requires surveys of molecular line emission over large volumes of the Galaxy with high angular and spectral resolution and sampling.497 Such surveys provide a census of the molecular eas without an undue bias toward bright enission or association with active sites of star formation., Such surveys provide a census of the molecular gas without an undue bias toward bright emission or association with active sites of star formation.498 The subsequent large uuuboer of molecular regious identified in wide field surveys euable a statistical evaluation of gas properties aud classification with respect to the local environment., The subsequent large number of molecular regions identified in wide field surveys enable a statistical evaluation of gas properties and classification with respect to the local environment.499" There have rycen, several important wide field surveys of CO cinission from the Calaxy.", There have been several important wide field surveys of CO emission from the Galaxy.500 The laree scale distribution of uolecular gas iu the Milkv Way has been determined from the combiued North-South surveys sumunarizec w Dine (1987)., The large scale distribution of molecular gas in the Milky Way has been determined from the combined North-South surveys summarized by Dame (1987).501 However. the large effective beam size limits the description of gas properties to the argest elant molecular cloud complexcs.," However, the large effective beam size limits the description of gas properties to the largest giant molecular cloud complexes."502 The Massachusetts-Stouv Brook Survey imaged the iuner Galaxy with an effective resolution of ((Sanders 1985)., The Massachusetts-Stony Brook Survey imaged the inner Galaxy with an effective resolution of (Sanders 1985).503 Aualvsis of the data by Solomon (19857) aud Scoville (1987) ideutified a iuauberof eiat molecular clouds and cloud complexes., Analysis of the data by Solomon (1987) and Scoville (1987) identified a number of giant molecular clouds and cloud complexes.504 Maux ofthe accepted characteristics of the molecular interstellar medium are derived from these studies., Many of the accepted characteristics of the molecular interstellar medium are derived from these studies.505 These include the self gravitational equilibrium state of he eiut iiolecular clouds aud the relationship between the velocity dispersion aud size of the cloud., These include the self gravitational equilibrium state of the giant molecular clouds and the relationship between the velocity dispersion and size of the cloud.506 Iowevor. cloud properties determined from iuuer Galaxy surveys are compronuscd due to the ligh deeree of confusion along the line of sight due to velocity crowding which precludes a complete accouutiug of the euission (Lizst Durton 1981).," However, cloud properties determined from inner Galaxy surveys are compromised due to the high degree of confusion along the line of sight due to velocity crowding which precludes a complete accounting of the emission (Lizst Burton 1981)."507 To reduce the blending of emission from unrelated clouds. molecular regions are identified," To reduce the blending of emission from unrelated clouds, molecular regions are identified"508significantly from the value originally derived by Alves (2000). which is almost 0.1 mag brighter. when his value of —1.61 on the Bessell Brett (1988) system 1s converted to the 2MASS system (Carpenter 2001).,"significantly from the value originally derived by Alves (2000), which is almost 0.1 mag brighter, when his value of $-1.61$ on the Bessell Brett (1988) system is converted to the 2MASS system (Carpenter 2001)."509" There is no dependence of Afi, on metallicity. and the dependence on colour is weak: (OLS EMOTE W)y 232)"," There is no dependence of $M_{\rm K}$ on metallicity, and the dependence on colour is weak: $(-0.15 \pm 0.07)$ $(V-K)_0 - 2.32$ )."510 The RC stars m clusters have also been used to derive Af. originally by Grocholski Sarajedini (2002).," The RC stars in clusters have also been used to derive $M_{\rm K}$, originally by Grocholski Sarajedini (2002)."511 Absolute magnitudes are derived by combining the observed mean magnitude in a box in colour and magnitude with a reddening estimate and main-sequence fitting distances., Absolute magnitudes are derived by combining the observed mean magnitude in a box in colour and magnitude with a reddening estimate and main-sequence fitting distances.512 Grocholski Sarajedini (2002) used the second incremental data release of 2MASS and 14 clusters to find an absolute A'-magnitude that is in agreement with the value in Alves (2000)., Grocholski Sarajedini (2002) used the second incremental data release of 2MASS and 14 clusters to find an absolute $K$ -magnitude that is in agreement with the value in Alves (2000).513 Using the AII-Sky data release and increasing the sample to 24 clusters. Van Helshoecht Groenewegen (2007; hereafter VHG) find a value fainter by 0.05 mag. or Mg;=1.57£0.05 on the Bessell Brett system.," Using the All-Sky data release and increasing the sample to 24 clusters, Van Helshoecht Groenewegen (2007; hereafter vHG) find a value fainter by 0.05 mag, or $M_{\rm K} = -1.57 \pm 0.05$ on the Bessell Brett system."514 Both Grocholski Sarajedini (2002) and vHG used averages over the cluster sample to arrive at the quoted means., Both Grocholski Sarajedini (2002) and vHG used averages over the cluster sample to arrive at the quoted means.515" To compare these values to the absolute magnitude of the local Hipparcos sample. a “population correction"" has to be made. to account for the difference in metallicity and age of the RC population."," To compare these values to the absolute magnitude of the local Hipparcos sample, a “population correction” has to be made, to account for the difference in metallicity and age of the RC population."516 The calculation of this correction ts outlined and tabulated in Girardi Salaris (2001) and Salaris Girardi (2002)., The calculation of this correction is outlined and tabulated in Girardi Salaris (2001) and Salaris Girardi (2002).517 In the A-band. Salaris Girardi (2002) show that the correction is a strong function of age for ages below 3 Gyr.," In the $K$ -band, Salaris Girardi (2002) show that the correction is a strong function of age for ages below 3 Gyr."518 The top panel in their Fig., The top panel in their Fig.519 3 shows that the correction is well behaved for ages above about 4 Gyr. and for each metallicity a linear function was fitted.," 3 shows that the correction is well behaved for ages above about 4 Gyr, and for each metallicity a linear function was fitted."520 From vHG the five clusters older that 4 Gyr were taken and the population correction was determined by. linear interpolation in age and [Fe/H] in the results in Salaris Girardi., From vHG the five clusters older that 4 Gyr were taken and the population correction was determined by linear interpolation in age and [Fe/H] in the results in Salaris Girardi.521 The error in AAg is based on a 1 Gyr error in age and 0.1 dex in metallicity.," The error in $\Delta522M_{\rm K}$ is based on a 1 Gyr error in age and 0.1 dex in metallicity."523 The results are listed in Table 4.. which lists the age and metallicity (see vHG for the references for age and metallicity determination). the 1j; value and its error that takes into account the error in reddening and the assumed distance to the cluster (from vHG). the population correction with error. and the corrected value.," The results are listed in Table \ref{Tab-Clusters}, which lists the age and metallicity (see vHG for the references for age and metallicity determination), the $M_{\rm K}$ value and its error that takes into account the error in reddening and the assumed distance to the cluster (from vHG), the population correction with error, and the corrected value."524 The weighted mean of the corrected A/j; values is 1.39-+ 0.06 with a dispersion of 0.2 mag., The weighted mean of the corrected $M_{\rm K}$ values is $-1.39 \pm$ 0.06 with a dispersion of 0.2 mag.525 Correcting to the 2MASS system (Carpenter 2001) this becomes 1.2 0.06., Correcting to the 2MASS system (Carpenter 2001) this becomes $-1.43 \pm$ 0.06.526 Although the dispersion is large. the mean value is also fainter than the straight mean for the cluster sample and agrees within 2-sigma with the determination from the Hipparcos sample.," Although the dispersion is large, the mean value is also fainter than the straight mean for the cluster sample and agrees within 2-sigma with the determination from the Hipparcos sample."527 To settle the issue on the importance of the bias and the absolute A-magnitude of RC stars would require accurate NIR magnitudes of a 100 to a few hundred (cf., To settle the issue on the importance of the bias and the absolute $K$ -magnitude of RC stars would require accurate NIR magnitudes of a 100 to a few hundred (cf.528 Table 2) bright (A 2:5) RC stars., Table 2) bright $K \less 5$ ) RC stars.529 Given the brightness. this represents a challenge to modern instrumentation because of saturation.," Given the brightness, this represents a challenge to modern instrumentation because of saturation."530 The absolute magnitudes derived here are in. better agreement with theory than previously., The absolute magnitudes derived here are in better agreement with theory than previously.531" Using a plausible star formation history for the solar neighbourhood. Salaris Girardi (2002) derived absolute magnitudes of Mj—/—(0.17 and Af,=—1.581 (on the 2MASS system). which they compared to the Hipparcos-based results quoted in Alves et al. ("," Using a plausible star formation history for the solar neighbourhood, Salaris Girardi (2002) derived absolute magnitudes of $M_{\rm I}= -0.17$ and $M_{\rm K}= -1.584$ (on the 2MASS system), which they compared to the Hipparcos-based results quoted in Alves et al. ("5322002) of Aj.—0.2640.03 and Ay;-L.GLI+0.03 (on the 2MASS system). which indicates differences at the 2-3 sigma level.,"2002) of $M_{\rm I}= -0.26 \pm 0.03$ and $M_{\rm K}= -1.644 \pm 0.03$ (on the 2MASS system), which indicates differences at the 2-3 sigma level."533 The new results of Aj=—0.22+0.03 and My=15L4£0.01 agree with theory at the 1-2 sigma level., The new results of $M_{\rm I}= -0.22 \pm 0.03$ and $M_{\rm K}= -1.54 \pm 0.04$ agree with theory at the 1-2 sigma level.534 Finally. some implications for existing distance determinations using the RC are discussed.," Finally, some implications for existing distance determinations using the RC are discussed."535 The derived value of the absolute magnitude in the /-band of 1M—0.220.03 is not very different from previously adopted values in the literature of 0.20 or 0.26. so the impact on derived distances is not so great.," The derived value of the absolute magnitude in the $I$ -band of $M_{\rm I} = -0.22 \pm 0.03$ is not very different from previously adopted values in the literature of $-0.23$ or $-0.26$, so the impact on derived distances is not so great."536 The distance to the LMC quoted by Pietrzynsski Giere (2002) of 18.501 + 0.049 (random-systematic error anc assuming no population correction) based on JA observations of two fields in the bar would become 18.35 + 0.05 for a assumed population correction of —0.03. (Salaris Girard 2002)., The distance to the LMC quoted by Pietrzyńsski Gieren (2002) of 18.501 $\pm$ 0.049 (random+systematic error and assuming no population correction) based on $JK$ observations of two fields in the bar would become 18.35 $\pm$ 0.05 for an assumed population correction of $-0.03$ (Salaris Girardi 2002).537 Babusiaux Gilmore (2005) obtained infrared data o some fields in the direction of the Galactic bulge., Babusiaux Gilmore (2005) obtained infrared data on some fields in the direction of the Galactic bulge.538 Their distance scale is tied to one field for which they obtain Ag;=—1.72 for an assumed distance of 8 kpe., Their distance scale is tied to one field for which they obtain $M_{\rm K} = -1.72$ for an assumed distance of 8 kpc.539 Taking My=1.5L would result in a distance of 7.4 kpe (0.2 kpe judging from the, Taking $M_{\rm K} = -1.54$ would result in a distance of 7.4 kpc $\pm$ 0.2 kpc judging from the540we can see that the pair annihilation process dominates the electron/positron capture process.,we can see that the pair annihilation process dominates the electron/positron capture process.541" Therefore Q, is approximately proportional to T?HcTD}.", Therefore $Q_{\nu}$ is approximately proportional to $T^9 H \propto T^{13}\Sigma^{-1}$.542" By equating the viscous heating rate and the neutrino cooling rate, Qt and Qz;, we can easily see that the temperature T' is proportional to a!/®M!/15p-3/19 and does not depend on X, which is also found in Fig."," By equating the viscous heating rate and the neutrino cooling rate, $Q^+$ and $Q_{\nu}^-$, we can easily see that the temperature $T$ is proportional to $\alpha^{1/5} M^{1/15} r^{-3/10}$ and does not depend on $\Sigma$, which is also found in Fig."543" 2, because both Q* and Q; are proportional to X! when described in terms of (X, T)."," 2, because both $Q^+$ and $Q_{\nu}^-$ are proportional to $\Sigma^{-1}$ when described in terms of $\Sigma$, $T$ )."544" Considering the same part of the thermal equilibrium curves on the (©, M) plane, as McQ-?Qtος T5X:-!, we can see that Mοςa19/8M which is seen in Fig."," Considering the same part of the thermal equilibrium curves on the $\Sigma$, $\dot{M}$ ) plane, as $\dot{M} \propto \Omega^{-2} Q^+ \propto T^8 \Sigma^{-1}$ , we can see that $\dot{M} \propto \alpha^{13/8} M^{-7/10} r^{21/10} \Sigma^{-1}$, which is seen in Fig."545" 1 as a negative gradient part in each -77/10,21/10y-1.thermal equilibrium curve.", 1 as a negative gradient part in each thermal equilibrium curve.546" We can interpret the physical reason of the viscously unstable branch in the convective NDAFs as follows; due to the vertical convection in the accretion disk, the advective energy transport would be significantly suppressed, and thus the disk temperature would become higher than in the case without the convection."," We can interpret the physical reason of the viscously unstable branch in the convective NDAFs as follows; due to the vertical convection in the accretion disk, the advective energy transport would be significantly suppressed, and thus the disk temperature would become higher than in the case without the convection."547" In such a hot disk, electrons would not be degenerate and hence the electron-positron annihilation process would become an efficient cooling process."," In such a hot disk, electrons would not be degenerate and hence the electron-positron annihilation process would become an efficient cooling process."548" It can dominate the emissivity due to the electron/positron capture by nucleons (ος T9), which was the dominant neutrino emission process in the NDAF models without the convection (e.g. see KM02)."," It can dominate the emissivity due to the electron/positron capture by nucleons $\propto T^6 \rho$ ), which was the dominant neutrino emission process in the NDAF models without the convection (e.g. see KM02)."549" Therefore the temperature-dependence of the cooling rate would be changed and make the unstable branch appear in the (X, M) plane."," Therefore the temperature-dependence of the cooling rate would be changed and make the unstable branch appear in the $\Sigma$, $\dot{M}$ ) plane."550 We can see in Fig., We can see in Fig.551 1 that with smaller viscous parameter o the unstable branch appears with smaller mass accretion rate M., 1 that with smaller viscous parameter $\alpha$ the unstable branch appears with smaller mass accretion rate $\dot{M}$.552" This is because the advective energy transport is more inefficient with smaller o, which makes the disk temperature high enough for the electron-positron annihilation process to become dominant even with smaller mass accretion rate."," This is because the advective energy transport is more inefficient with smaller $\alpha$, which makes the disk temperature high enough for the electron-positron annihilation process to become dominant even with smaller mass accretion rate."553" In particular, the typical timescale for this instability is the order of the viscous timescale, and therefore when r/H<O(1), which is often realized in the innermost region of NDAF, this timescale is short enough to account for the time variability observed in the prompt emissions of GRBs."," In particular, the typical timescale for this instability is the order of the viscous timescale, and therefore when $r/H \lesssim {\cal O}(1)$, which is often realized in the innermost region of NDAF, this timescale is short enough to account for the time variability observed in the prompt emissions of GRBs."554" In this study, we show that if the convective energy transfer along the vertical direction, which is naturally expected in the radiation-pressure dominated accretion disk, the thermal equilibrium solution for a hyperaccretion disk would have a viscously unstable branch."," In this study, we show that if the convective energy transfer along the vertical direction, which is naturally expected in the radiation-pressure dominated accretion disk, the thermal equilibrium solution for a hyperaccretion disk would have a viscously unstable branch."555" When this branch is realized, highly time-dependent mass accretion onto a black hole would occur."," When this branch is realized, highly time-dependent mass accretion onto a black hole would occur."556 Such instabilities have been studied in the contexts of dwarf novae (Meyer Meyer-Hofmeister 1981; Hoshi 1979; Smak 1982) and microquasars which have radiation pressure-dominated accretion disks in themselves (Lightman Eardley 1974; Shibazaki Hoshi 1975; Pringle 1976; Shakura Sunyaev 1976; Abramowicz et al., Such instabilities have been studied in the contexts of dwarf novae (Meyer Meyer-Hofmeister 1981; Hoshi 1979; Smak 1982) and microquasars which have radiation pressure-dominated accretion disks in themselves (Lightman Eardley 1974; Shibazaki Hoshi 1975; Pringle 1976; Shakura Sunyaev 1976; Abramowicz et al.557" 1988), and the time variable features observed in those sources are theoretically interpreted with the model of those instabilities."," 1988), and the time variable features observed in those sources are theoretically interpreted with the model of those instabilities."558" In the context of GRBs, the process discussed in this study may lead to the intermittent energy release from a hyperaccretion disk, and then the highly inhomogeneous jet would be launched, which can be the origin of the short-term variability in the prompt emission of GRBs."," In the context of GRBs, the process discussed in this study may lead to the intermittent energy release from a hyperaccretion disk, and then the highly inhomogeneous jet would be launched, which can be the origin of the short-term variability in the prompt emission of GRBs."559" Judging from isotropic luminosities of observed long GRBs (Liso9?erg s~'), the accretion rate adopted here M~107?109?!—107Mas!10997?!ergs71/c? is relatively small."," Judging from isotropic luminosities of observed long GRBs $L_{\rm iso} \sim 10^{51-52}560{\rm erg}~{\rm s}^{-1}$ ), the accretion rate adopted here $\dot{M}\sim 10^{-3}-10^{-4}M_{\odot}{\rm s}^{-1}\sim 10^{50-51}{\rm erg}~{\rm s}^{-1}/c^2$ is relatively small."561" However, recent observations show that there is a population of low luminosity GRBs (Liso~10*°~*%ergs; e.g. Galama et al."," However, recent observations show that there is a population of low luminosity GRBs $L_{\rm iso}\sim 10^{46-49}{\rm erg}~{\rm s}^{-1}$; e.g. Galama et al."562 1998; Campana et al., 1998; Campana et al.563" 2006), and the variabilities of their prompt emissions may be explained with our model."," 2006), and the variabilities of their prompt emissions may be explained with our model."564 Our model may be also appropriate to describing the late time activity of GRBs such as X-ray flares in afterglows (e.g. Burrows et al., Our model may be also appropriate to describing the late time activity of GRBs such as X-ray flares in afterglows (e.g. Burrows et al.565 2005)., 2005).566" We should note that the absorption optical depth for neutrinos is smaller than unity, according to Eqs. ("," We should note that the absorption optical depth for neutrinos is smaller than unity, according to Eqs. ("56755)-(57) in Kohri et al. (,55)-(57) in Kohri et al. (5682005).,2005).569" This means that we do not have to consider the neutrino trapping effect in the current situation (T$10"" XS10?gcm ?)."," This means that we do not have to consider the neutrino trapping effect in the current situation $T\lesssim 10^{11}{\rm K}$, $\Sigma \lesssim 10^{15}{\rm g}~{\rm cm}^{-2}$ )."570" We shouldK, mention the effect of magnetohydrodyanmic turbulence on the convection in our accretion disks.", We should mention the effect of magnetohydrodyanmic turbulence on the convection in our accretion disks.571" A hyperaccretion disk should have a turbulent structure due to the magnetorotational instability (MRI), and it would help the mass accretion as the turbulent viscosity (cf."," A hyperaccretion disk should have a turbulent structure due to the magnetorotational instability (MRI), and it would help the mass accretion as the turbulent viscosity (cf."572 Balbus and Hawley 1991)., Balbus and Hawley 1991).573" Here, using a parameter,we can describe the turbulent viscosity νι as vi~acsH where c, and H are the speed of sound and the scale height of an accretion flow, respectively."," Here, using $\alpha$ parameter,we can describe the turbulent viscosity $\nu_t$ as $\nu_t\sim \alpha c_s H$ where $c_s$ and $H$ are the speed of sound and the scale height of an accretion flow, respectively."574" In our study o parameter is assumed to be relatively small value, < 10?, which may be due to the"," In our study $\alpha$ parameter is assumed to be relatively small value, $\lesssim 10^{-2}$ which may be due to the"575If the resulting apparent axis ratio. as given by equation (12)). was q<0.866. I retained il in mv sample.,"If the resulting apparent axis ratio, as given by equation \ref{eq:qform}) ), was $q \leq 0.866$, I retained it in my sample."576 Hit was flatter than this limit. I discarded it.," If it was flatter than this limit, I discarded it."577 By repeating this procedure until I had νι=28 axis ratios. I created one possible realization.," By repeating this procedure until I had $N_{\rm gal} = 28$ axis ratios, I created one possible realization."578 After creating 16.000 realizations. | computed the mean ancl standard deviation in each of the bins in Figure 3..," After creating 000 realizations, I computed the mean and standard deviation in each of the bins in Figure \ref{fig:loge_ab}."579 A search through parameter space revealed that the best fit. as measured by a X7 test. was eiven bv yo=—2.29. στ1.04.," A search through parameter space revealed that the best fit, as measured by a $\chi^2$ test, was given by $\mu = -2.29$, $\sigma = 1.04$."580 The probability of the fit. illustrated bythe points and error bars in Figure 3.. was 2?=0.98.," The probability of the fit, illustrated bythe points and error bars in Figure \ref{fig:loge_ab}, was $P = 0.98$."581 The 28 galaxies of Andersen&Dershady(2003) do not. by themselves. provide a strong constraint on the distribution of intrinsic ellipticilies.," The 28 galaxies of \citet{an03} do not, by themselves, provide a strong constraint on the distribution of intrinsic ellipticities."582 Figure 4. shows the goodness of fit. as measured by the 4? probability Py(ec).," Figure \ref{fig:prob_ab} shows the goodness of fit, as measured by the $\chi^2$ probability $P_{\rm AB} (\mu,\sigma)$."583 The cross indicates the best fit. and the dotted aad solid lines show the isoprobalbility contours.," The cross indicates the best fit, and the dotted and solid lines show the isoprobability contours."584 Note that the 2=0.1 contour the innermost dotted line encloses a large area stretching off to the upper right of the plot., Note that the $P = 0.1$ contour – the innermost dotted line – encloses a large area stretching off to the upper right of the plot.585 That is. a distribution with a large value of ji. corresponding to a very. flattened average shape. is acceptable as long as il is paired with a laree value of σ. signifving a wide spread in shapes.," That is, a distribution with a large value of $\mu$, corresponding to a very flattened average shape, is acceptable as long as it is paired with a large value of $\sigma$, signifying a wide spread in shapes."586 Because the Anclersen-Bershady contains only galaxies which are nearly circular in projection. it is strongly weighted toward galaxies which are nearly circular in (heir intrinsic shape. aud (hus cannot effectively constrain the high-ellipticitv end of (2).," Because the Andersen-Bershady contains only galaxies which are nearly circular in projection, it is strongly weighted toward galaxies which are nearly circular in their intrinsic shape, and thus cannot effectively constrain the high-ellipticity end of $f(\varepsilon)$."587 A Gaussian peaking al 7=0 (see equation (1))) doesn't provide a good lit to the Andersen-Dershady sample., A Gaussian peaking at $\varepsilon = 0$ (see equation \ref{eq:apm}) )) doesn't provide a good fit to the Andersen-Bershady sample.588 The best-fitting Gaussian. with o-=0.143. had a 4? probability of only P=0.009.," The best-fitting Gaussian, with $\sigma_\varepsilon = 0.143$, had a $\chi^2$ probability of only $P = 0.009$."589 Thus. although the data of Andersen&Dershady(2003) doesn't constrain the high-ellipticitv end of (ο) its discriminatory power at low values of 2 weighs strongly against a distribution peaking al =0.," Thus, although the data of \citet{an03} doesn't constrain the high-ellipticity end of $f(\varepsilon)$, its discriminatory power at low values of $\varepsilon$ weighs strongly against a distribution peaking at $\varepsilon = 0$."590 The kinematic and photometric information exploited by Andersen&Bershady(2003) is in somewavs complementary (to the purely photometric information included in (he SDSS DRI axis ratios., The kinematic and photometric information exploited by \citet{an03} is in someways complementary to the purely photometric information included in the SDSS DR1 axis ratios.591 The nearly face-on galaxies of the Andersen-Bershacly sample constrain the low-ellipticity end of (2): thescarcity. of nearly cireular galaxies in the ο) DRI exponential sample (see Figure 1)) constrains the high-ellipticity endl of /(2)., The nearly face-on galaxies of the Andersen-Bershady sample constrain the low-ellipticity end of $f(\varepsilon)$; thescarcity of nearly circular galaxies in the SDSS DR1 exponential sample (see Figure \ref{fig:q_sdss}) ) constrains the high-ellipticity end of $f(\varepsilon)$.592" The kinematic measurements of Andersene£al(2001) typically οὐ ont to 243 scale lengths (1.241.8r, ).", The kinematic measurements of \citet{an01} typically go out to $2 \to 3$ scale lengths $1.2 \to 1.8 r_e$ ).593 The ellipticities determined by Andersenefa£(2001) and Andersen&Bershacly(2003). can be thought of as average ellipticities over the inner region of the galaxy., The ellipticities determined by \citet{an01} and \citet{an03} can be thought of as average ellipticities over the inner region of the galaxy.594" Thus. the Andersen- ellipticities ave more directly comparable to the ellipticities found from qi, than [rom the outer axis ratios qo."," Thus, the Andersen-Bershady ellipticities are more directly comparable to the ellipticities found from $q_{\rm am}$ than from the outer axis ratios $q_{25}$."595 As shown in Figure 4.. the best fit using qu in the ? band. indicated bv the filled circle. is marginally consistent with the ÀAndersen-Dershady. results.," As shown in Figure \ref{fig:prob_ab}, the best fit using $q_{\rm am}$ in the $i$ band, indicated by the filled circle, is marginally consistent with the Andersen-Bershady results."596 Alultiplxing together the probability fields in Figure 2 ancl Figure 4. vields a best joint fit of fe=— 1.89. 8= 0.96.," Multiplying together the probability fields in Figure \ref{fig:prob_sdss} and Figure \ref{fig:prob_ab}597 yields a best joint fit of $\mu = -1.89$ , $\sigma = 0.96$ ."598 For this set of parameters. the 4? fit to the Andersen-Dershady. data has P= 0.44. and the fit to the qu; data in the / band has P—9.5x10 .," For this set of parameters, the $\chi^2$ fit to the Andersen-Bershady data has $P = 0.44$ , and the fit to the $q_{am}$ data in the $i$ band has $P = 9.5 \times 10^{-5}$ ."599 <500 (~1.5 δὲ f=LOPts E=1077Ex»org»=προ? ~LOW E=10%ere D=0σα.) #=105s 9«10.foreeui2 R=(fet Rzπο/ng)teu. RPPee=Εν. Fy ," $\lesssim 500$ $\sim 1.5 $ $\delta t$ $t=10^3t_3 \,{\rm s}$ $E=10^{52}E_{52}{\rm erg}$$n=1{\rm n_0 \,cm^{-3}}$ $8\times 10^{-7} {\rm erg cm^{-2} }$ $\sim10\%$ $E=10^{52}{\rm erg}$ $D=10^{28}\rm cm$ $t=10^3~ {\rm s}$ $9\times 10^{-7} {\rm600erg~cm^{-2} }$ $R=4601\Gamma^2ct$ $\Gamma\simeq30(E_{52}/n_0)^{1/8}t_3^{-3/8}$ $ R\simeq10^{17}(602E_{52} t_3/n_0)^{1/4}\cm$ $R^2\Gamma^2U'_Xc=D^2F_X$ $U'_X=D^2F_X/(\Gamma^2R^2c)$ $F_X$ "603"instrumental broadening and a Gaussian noise in order to reproduce the observed average SNR (per pixel): $NR—5 for the Triplet, the Sextet and Pair A; SNR=9 for Pair U and SNR=15 for Pair Q. In this section, we discuss the flux correlations in the absorption spectra of our sample of QSO pairs.","instrumental broadening and a Gaussian noise in order to reproduce the observed average SNR (per pixel): $SNR=5$ for the Triplet, the Sextet and Pair A; $SNR=9$ for Pair U and $SNR=15$ for Pair Q. In this section, we discuss the flux correlations in the absorption spectra of our sample of QSO pairs."604" The statistical quantities are the same as those already computed in Paper I, here we verify the effect of having an increased SNR."," The statistical quantities are the same as those already computed in Paper I, here we verify the effect of having an increased SNR."605" On the basis of the interpretation of the fforest as due to a continuous density field with a one-to-one correspondence between density and transmitted flux, we computed the correlation properties of the transmitted flux in QSO aand regarded them as indicators of the correlation properties of matter in the IGM."," On the basis of the interpretation of the forest as due to a continuous density field with a one-to-one correspondence between density and transmitted flux, we computed the correlation properties of the transmitted flux in QSO and regarded them as indicators of the correlation properties of matter in the IGM."606" We selected in each normalized spectrum the region between the eemission (or the shortest observed wavelength, when the wwas not included in the spectrum) and 5000 ffrom the eemission (to avoid proximity effect due to the QSO)."," We selected in each normalized spectrum the region between the emission (or the shortest observed wavelength, when the was not included in the spectrum) and 5000 from the emission (to avoid proximity effect due to the QSO)."607" Absorption lines due to ions of elements heavier than hydrogen contaminate the fforest and can give spurious contributions to the clustering signal (seeKimetal.2004,foradiscussioninthecaseofsinglelinesof sight).."," Absorption lines due to ions of elements heavier than hydrogen contaminate the forest and can give spurious contributions to the clustering signal \citep[see][for a discussion in the case of single 608lines of sight]{kim04}."609 We flagged and removed the spectral regions where metal lines and aabsorptions of damped and sub-damped systems occurred inside the fforest., We flagged and removed the spectral regions where metal lines and absorptions of damped and sub-damped systems occurred inside the forest.610" Given the normalized transmitted flux, f, as a function of the velocity along the aand the angularvj position 0 on the sky, we define (f- f), where the average flux, f, is computed for every spectrum as the mean of the transmitted flux over all the considered pixels in that spectrum."," Given the normalized transmitted flux, $f$, as a function of the velocity $v_{\pa}$ along the and the angular position $\theta$ on the sky, we define $\delta_f = (f - \bar{f})$ , where the average flux, $\bar{f}$, is computed for every spectrum as the mean of the transmitted flux over all the considered pixels in that spectrum."611" We neglected the redshift evolution of the average transmitted flux in the fforest of the individual spectra, which translates into the redshift evolution of the mean oopacity of the Universe (Kimetal.2002;Schaye2003; 2008),, because we verified that its effect on the correlation function is negligible."," We neglected the redshift evolution of the average transmitted flux in the forest of the individual spectra, which translates into the redshift evolution of the mean opacity of the Universe \citep{kim02,schaye03,viel2004,kirkman05,fg2008}, because we verified that its effect on the correlation function is negligible."612" By means of this new field, 9;, we could then define and compute three useful tools for the investigation of the ccorrelation properties: the Auto and Cross correlation function and the set of the Cross correlation coefficients."," By means of this new field, $\delta_f$, we could then define and compute three useful tools for the investigation of the correlation properties: the Auto and Cross correlation function and the set of the Cross correlation coefficients."613 The unnormalized Auto correlation function (Auto CF) of the flux along the lis defined as: following previous studies on the same subject (e.g.Mc-Donaldetal.2000;Rollinde2003;Becker 2004).," The unnormalized Auto correlation function (Auto CF) of the flux along the is defined as: following previous studies on the same subject \citep[e.g.][]{mcdonald00,rollinde03,becker04}."614. The Auto CF for our sample of QSO spectra was obtained by averaging over all the pixels of all the QSOs., The Auto CF for our sample of QSO spectra was obtained by averaging over all the pixels of all the QSOs.615 The results were binned in 50 vvelocity bins., The results were binned in 50 velocity bins.616 The Auto CF for the simulated spectra was computed as the arithmetic mean of the correlation functions obtained for 50 realizations of the observed sample and the error is the corresponding standard deviation., The Auto CF for the simulated spectra was computed as the arithmetic mean of the correlation functions obtained for 50 realizations of the observed sample and the error is the corresponding standard deviation.617 It is important to recognize that the computed error bars for the simulated £ are strongly correlated., It is important to recognize that the computed error bars for the simulated $\xi^f_{\pa}$ are strongly correlated.618 This is due to the fact that every pixeli contributes to the correlation function in several velocity bins., This is due to the fact that every pixel contributes to the correlation function in several velocity bins.619 In Fig., In Fig.620 2 we show the comparison of the observed and simulated Auto CF in velocity space., \ref{fig:auto_CF} we show the comparison of the observed and simulated Auto CF in velocity space.621" With the improved SNR with respect to Paper I, the agreement between the two correlation functions has increased, weakening any evidence for extra clustering in the regions occupied by QSO pairs."," With the improved SNR with respect to Paper I, the agreement between the two correlation functions has increased, weakening any evidence for extra clustering in the regions occupied by QSO pairs."622 In this section we exploit the capabilities of our sample of QSO pairs by determining the clustering properties of the IGM across thesight., In this section we exploit the capabilities of our sample of QSO pairs by determining the clustering properties of the IGM across the.623". The great advantage with respect to the correlation function along thesight,, in particular for a sample like ours showing a large variety of pair separations, is that we have the guarantee of sampling true spatial separations between the pixels, the effect of peculiar velocities being negligible or absent."," The great advantage with respect to the correlation function along the, in particular for a sample like ours showing a large variety of pair separations, is that we have the guarantee of sampling true spatial separations between the pixels, the effect of peculiar velocities being negligible or absent."624" As a first approach, we computed the Cross-correlation function (Cross CF) extending in a natural way the procedure adopted for the Auto CF."," As a first approach, we computed the Cross-correlation function (Cross CF) extending in a natural way the procedure adopted for the Auto CF."625 Every pixel along the iis considered as an element of the density field at the QSO angular position in the sky and at a distance from the observer (comoving along the sight)) corresponding to the redshift of the pixel:, Every pixel along the is considered as an element of the density field at the QSO angular position in the sky and at a distance from the observer (comoving along the ) corresponding to the redshift of the pixel:626Here it is assumed that the electrons are accelerated into a power-law distribution. N(5)x57 with p=2.0.,"Here it is assumed that the electrons are accelerated into a power-law distribution, $N(\gamma) \propto \gamma^{-p}$ with $p=2.0$."627 These equations. reveal that the strength of the non-relativistic emission is strongly dependent on the density of the circumstellar medium (especially in the case of a wind) and is best probed at low frequencies., These equations reveal that the strength of the non-relativistic emission is strongly dependent on the density of the circumstellar medium (especially in the case of a wind) and is best probed at low frequencies.628 While this analytic model provides robust predictions for the afterglow emission at ¢>tN. 1t does not describe the early evolution or the transition from relativistic to sub-relativistic expansion.," While this analytic model provides robust predictions for the afterglow emission at $t>t_{\rm NR}$, it does not describe the early evolution or the transition from relativistic to sub-relativistic expansion."629 At early time. the observed emission from an off-axis GRB is strongly dependent on the viewing angle and dynamties of the jet.," At early time, the observed emission from an off-axis GRB is strongly dependent on the viewing angle and dynamics of the jet."630 To investigate this early afterglow evolution and the transition to sub-relativistic expansion. we developed a detailed semi-analytic model. described below.," To investigate this early afterglow evolution and the transition to sub-relativistic expansion, we developed a detailed semi-analytic model, described below."631 In modeling the afterglow emission from an off-axis GRB jet. we adopt the standard framework for a adiabatic blastwave expanding into either à uniform or wind stratified medium (Sart1997:Granot&Sart 2002).," In modeling the afterglow emission from an off-axis GRB jet, we adopt the standard framework for a adiabatic blastwave expanding into either a uniform or wind stratified medium \citep{sar97,gs02}."632. We assume a uniform. sharp- jet such that Lorentz factor and energy are constant over the jet surface.," We assume a uniform, sharp-edged jet such that Lorentz factor and energy are constant over the jet surface."633 The hydrodynamic evolution of the jet 1s fully deseribed in Oren.Nakar&Piran (2004).., The hydrodynamic evolution of the jet is fully described in \citet{onp04}. .634 As the bulk Lorentz factor of the ejecta approaches [—I. the jets begin to spread laterally at the sound speedcurves.," As the bulk Lorentz factor of the ejecta approaches $\Gamma\sim 1$, the jets begin to spread laterally at the sound speed."635.. Our off-axis light-curves are obtained by integrating the afterglow emission over equal arrival time surface., Our off-axis light-curves are obtained by integrating the afterglow emission over equal arrival time surface.636 We note that these resulting light-curves are in broad agreement with Model 2 of Granotetal...(2002) and are consistent with Waxman's analytic model refsec:waxman)) on timescales. f>νι.," We note that these resulting light-curves are in broad agreement with Model 2 of \citet{gpk+02} and are consistent with Waxman's analytic model \\ref{sec:waxman}) ) on timescales, $t\gtrsim t_{\rm NR}$."637" Over-plotted in Figure 1. are our off-axis models caleulated for both wind-stratified and homogeneous media at an observing frequency of i44,=8.46 GHz.", Over-plotted in Figure \ref{fig:lum_limits_oa} are our off-axis models calculated for both wind-stratified and homogeneous media at an observing frequency of $\nu_{\rm obs}=8.46$ GHz.638" We assume standard GRB parameters of Es;2A nzlegz2e,20..pz222 and 0;=5"". consistent with the typical values inferred from broadband modeling of GRBs (Panaitescu&Kumar2002;Yostetal,2003:Chevalier.Li&Fransson 2004)."," We assume standard GRB parameters of $E_{51}=A_*=n=1$, $\epsilon_B=\epsilon_e=0.1$, $p=2.2$ and $\theta_j=5^{\rm o}$, consistent with the typical values inferred from broadband modeling of GRBs \citep{pk02,yhs+03,clf04}."639. We compute model light-curves for off-axis viewing angles between 30 and 90 degrees., We compute model light-curves for off-axis viewing angles between 30 and 90 degrees.640" As clearly shown in the figure. the majority of our late-time SNe Ibe limits are significantly fainter than«// of the model light-curves. constraining even the extreme case where Pay,=907."," As clearly shown in the figure, the majority of our late-time SNe Ibc limits are significantly fainter than of the model light-curves, constraining even the extreme case where $\theta_{\rm obs}=90^{\rm o}$."641 Based on the double-peaked profiles observed for the nebular lines of neutral oxygen and magnesium. Μαζζαetal.(2005) argue that 2200514 was an aspherical. axisymmetric explosion viewed near the equatorial plane.," Based on the double-peaked profiles observed for the nebular lines of neutral oxygen and magnesium, \citet{mkm+05} argue that 2003jd was an aspherical, axisymmetric explosion viewed near the equatorial plane."642 They suggest that this asymmetry may be explained if the SN explosion was accompanied by a tightly collimated and relativistic GRB jet. initially directed ~70 degrees from our line-of-sight.," They suggest that this asymmetry may be explained if the SN explosion was accompanied by a tightly collimated and relativistic GRB jet, initially directed $\sim 70$ degrees from our line-of-sight."643 This hypothesis is consistent with the observed lack of prompt gamma-ray emission (Hurleyetαἱ.2003) as well as the absence of strong radio and X-ray emission at early time (Soderberg.Kulkarni&Frail2003;Watsonetal...2003).," This hypothesis is consistent with the observed lack of prompt gamma-ray emission \citep{hcm+03} as well as the absence of strong radio and X-ray emission at early time \citep{skf03,wpr+03}."644. Our radio observation of 2200314 at r—1.6 years imposes strong constraints on the putative off-axis GRB jet., Our radio observation of 2003jd at $t\sim 1.6$ years imposes strong constraints on the putative off-axis GRB jet.645" While the early data constrain only mildly off-axis jets (ay,<30°). our late-time epoch constrains even those jets initially directed perpendicular to our line-of-sight."," While the early data constrain only mildly off-axis jets $\theta_{\rm obs}\lesssim64630^{\rm o}$ ), our late-time epoch constrains even those jets initially directed perpendicular to our line-of-sight."647" As shown in Figure ].. our radio limit is a factor of 200 (>20) fainter than that predicted for a typical GRB expanding into a homogeneous (wind-stratified) medium. even in the extreme case where (44,~90°."," As shown in Figure \ref{fig:lum_limits_oa}, our radio limit is a factor of $\gtrsim 200$ $\gtrsim 20$ ) fainter than that predicted for a typical GRB expanding into a homogeneous (wind-stratified) medium, even in the extreme case where $\theta_{\rm obs}\sim 90^{\rm o}$."648 Given the assumption of typical GRB parameters. we conclude that our late-time radio limit is inconsistent with the presence of an off-axis GRB Jet.," Given the assumption of typical GRB parameters, we conclude that our late-time radio limit is inconsistent with the presence of an off-axis GRB jet."649" We note that the model assumptions and physical parameters of our off-axis afterglow light-curves are identical to those adopted by Mazzalietal,(2005).", We note that the model assumptions and physical parameters of our off-axis afterglow light-curves are identical to those adopted by \citet{mkm+05}.650. We next explore the range of parameters ruled out by our deep radio limits., We next explore the range of parameters ruled out by our deep radio limits.651 As shown in Equations | and 2. the luminosity of the late-time emission is a function. of the ejecta energy. the density of the circumstellar medium and the equipartition fractions.," As shown in Equations 1 and 2, the luminosity of the late-time emission is a function of the ejecta energy, the density of the circumstellar medium and the equipartition fractions."652" To investigate the effect of energy and density on the late-time radio luminosity. we fix the equipartition fractions to e,2ey=0.1. chosen to be consistent with the values typically inferred from afterglow modeling of cosmological GRBs (Panaitescu&Kumar2002;2003)."," To investigate the effect of energy and density on the late-time radio luminosity, we fix the equipartition fractions to $\epsilon_e=\epsilon_B=0.1$, chosen to be consistent with the values typically inferred from afterglow modeling of cosmological GRBs \citep{pk02,yhs+03}."653 In Figure 2.. we illustrate how each radio epoch for 22003jd maps to a curve within the two-dimensional parameter space of kinetic energy and circumstellar density for an off-axis GRB.," In Figure \ref{fig:SN2003jd_ed}, we illustrate how each radio epoch for 2003jd maps to a curve within the two-dimensional parameter space of kinetic energy and circumstellar density for an off-axis GRB."654" Here we adopt our semi-analytic model refsec:our,,ode/)) forawind —strati Fiedmedium.alongwithatypicalelecti 90°: the faintest model for a given set of equipartition fractions."," Here we adopt our semi-analytic model \\ref{sec:our_model}) ) for a wind-stratified medium, along with a typical electron index of $p=2.2$ and a viewing angle of $\theta_{\rm obs}=90^{\rm o}$ ; the faintest model for a given set of equipartition fractions."655 By comparing the luminosity limit for 2200514 at a particular epoch with the off-axis model prediction for that time. we exclude the region of parameter space rightward of the curve since this region produces a jet which tsbrighter than the observed limit.," By comparing the luminosity limit for 2003jd at a particular epoch with the off-axis model prediction for that time, we exclude the region of parameter space of the curve since this region produces a jet which is than the observed limit."656 The union of these regions represents the total parameter space ruled out for an associated GRB., The union of these regions represents the total parameter space ruled out for an associated GRB.657" As shown in this figure. the total excluded parameter space extends from A.>0.03 and E—1077 to 10°,"," As shown in this figure, the total excluded parameter space extends from $A_*\gtrsim 0.03$ and $E \sim 10^{47}$ to $10^{52}$."658 We compare these constraints with the beaming-corrected kinetic energies and CSM densities for 18 cosmological GRBs (Table 3))., We compare these constraints with the beaming-corrected kinetic energies and CSM densities for 18 cosmological GRBs (Table \ref{tab:grb}) ).659 Here we make the rough approximation that A.z70: a reasonable assumption for circumstellar radii near ~10'5 em., Here we make the rough approximation that $A_*\approx n_0$; a reasonable assumption for circumstellar radii near $\sim 10^{18}$ cm.660 As shown in Figure 2.. these GRBs span the region of parameter space roughly bracketed by A..~0.002 to 100 andEος10” to 4«I0?!.," As shown in Figure \ref{fig:SN2003jd_ed}, these GRBs span the region of parameter space roughly bracketed by $A_*\sim6610.002$ to 100 and $E\sim 2\times 10^{49}$ to $4\times 10^{51}$."662 The majority of the bursts (13 out of 18) fall within the excluded region of parameter space for 22003514., The majority of the bursts (13 out of 18) fall within the excluded region of parameter space for 2003jd.663 We conclude that 22003jd was not likely associated with a typical GRB at a confidence level of ~72%., We conclude that 2003jd was not likely associated with a typical GRB at a confidence level of $\sim 72\%$.664 While physical parameters atypical of the cosmological GRB population can be invoked to hide an off-axis GRB for 22003jd. it is exceedingly unlikely foratypical parameters to dominate a large statistical sample of SNe Ibe.," While physical parameters atypical of the cosmological GRB population can be invoked to hide an off-axis GRB for 2003jd, it is exceedingly unlikely foratypical parameters to dominate a large statistical sample of SNe Ibc."665 Motivated thus. we searched for off-axis GRBsin the 67 local Type Ibe SNe for which we have compiled late-time (¢~0.5—30 yr) radio observations.," Motivated thus, we searched for off-axis GRBsin the 67 local Type Ibc SNe for which we have compiled late-time $t\sim 0.5-30$ yr) radio observations."666 Applying the method described in we produce exclusion regions in the Es;—A parameter space for each SN., Applying the method described in \\ref{sec:SN2003jd} we produce exclusion regions in the $E_{51}-A_*$ parameter space for each SN.667 Figure 3. shows the resulting," Figure \ref{fig:all_ed}668 shows the resulting"669The initial eas clistribution can be deduced. from. the present day light distribution. by assuming that the stars map the density distribution of the eas from which they formed.,"The initial gas distribution can be deduced from the present day light distribution, by assuming that the stars map the density distribution of the gas from which they formed."670" From the observed exponential surface density light profile seen in dSph's we infer for the gas distribution: where A, is the observed core radius of the luminous galaxies and £2. is the exponential scale length of. the surface brightness distribution.", From the observed exponential surface density light profile seen in dSph's we infer for the gas distribution: where $R_g$ is the observed core radius of the luminous galaxies and $R_e$ is the exponential scale length of the surface brightness distribution.671 In all that follows we will be assuming that the form of both gas ancl stellar profiles is the same. with the respective normalizations given by the relative amounts of both components.," In all that follows we will be assuming that the form of both gas and stellar profiles is the same, with the respective normalizations given by the relative amounts of both components."672 Notice that as the potential energy is determined overwhelmingly by the dark halo. the onset of the galactic wind is sensitive to the total gas present. but depends very little on the details of the gas distribution.," Notice that as the potential energy is determined overwhelmingly by the dark halo, the onset of the galactic wind is sensitive to the total gas present, but depends very little on the details of the gas distribution."673 All structural parameters (2). 2). 2.) are taken from Mateo (1998) and are assumed fixed during the evolution of the systems we study.," All structural parameters $R_{t}$, $R_g$, $R_e$ ) are taken from Mateo (1998) and are assumed fixed during the evolution of the systems we study."674 Since current studies show diSph galaxies contain only stars. the present-clay dynamical mass is formed by. DM ane stellar mass.," Since current studies show dSph galaxies contain only stars, the present-day dynamical mass is formed by DM and stellar mass."675 The stellar mass is computed. from the tota Iuminosities of the galaxies ancl the respective SFIgea assuming a constant IME.," The stellar mass is computed from the total luminosities of the galaxies and the respective $SFR_{HGV}$, assuming a constant IMF."676" Mateo (1998) gives the dvnamica mass value at Ze. so when A,7B; the dvnamical mass is sealed by (Reon/;)* since we have assumed. constan density for the core regions."," Mateo (1998) gives the dynamical mass value at $R_t$, so when $R_{core} > R_t$ the dynamical mass is scaled by $(R_{core}/ R_t)^3$ since we have assumed constant density for the core regions."677 The initial barvonic (gas|stars) to non-barvonic mass ratio is indicated. by λος. which is an. inpu parameter.," The initial baryonic (gas+stars) to non-baryonic mass ratio is indicated by $M_{lum}/M_{DM}$, which is an input parameter."678" By default. Ap,,;/Mpa=0.09. that is the niaximum barvonic content of late (ype galaxies obtained by llernandez Cilmore (1998a)."," By default, $M_{lum}/M_{DM}=0.09$, that is the maximum baryonic content of late type galaxies obtained by Hernandez Gilmore (1998a)."679 Notice that this value refers only to material within the core region. large amounts of dark matter would have been contained in the initial halo. much of which has been shaven olf by the Galactic tical fickle. not so the barvons. which cooled. anc contracted at very early times.," Notice that this value refers only to material within the core region, large amounts of dark matter would have been contained in the initial halo, much of which has been shaven off by the Galactic tidal field, not so the baryons, which cooled and contracted at very early times."680" We define £v as the potential energy. of the gas. in this case. the energv necessary to carry all the gas in the system out to a radius =e, (Martinelli. Colafrancesco 1998). at which point it will be stripped by the tidal field of our Galaxy as: where P(£) is the force between πρ) and. the dynamical mass for rr«€. and From the above. we see that the exponential cut olf in the dark matter density is not considered in the dvnamical determination of the problem. and has only a small relevance in determining the total amount of barvonic matter. through the Adon,λέω ralio."," We define $E_{GRAV}$ as the potential energy of the gas, in this case, the energy necessary to carry all the gas in the system out to a radius $= R_{core}$ (Martinelli, Colafrancesco 1998), at which point it will be stripped by the tidal field of our Galaxy as: where $F(\xi)$ is the force between $dm_{g}(r)$ and the dynamical mass for $r < \xi$, and From the above, we see that the exponential cut off in the dark matter density is not considered in the dynamical determination of the problem, and has only a small relevance in determining the total amount of baryonic matter, through the $M_{lum}/M_{DM}$ ratio."681 We assumed that Mpa 15 not allected by internal or external galactic processes. as ib is dynamically dominant. and is subject only to gravitational interactions.," We assumed that $M_{DM}$ is not affected by internal or external galactic processes, as it is dynamically dominant, and is subject only to gravitational interactions."682" We further assume that the shape of AZ,(r) and νι). is conserved during the whole evolution.", We further assume that the shape of $M_{g}(r)$ and $M_{s}(r)$ is conserved during the whole evolution.683 The gravitational potential of the gas evolves according to the evolution of AZj(r) and Αν, The gravitational potential of the gas evolves according to the evolution of $M_{g}(r)$ and $M_{s}(r)$.684 Ht will be the total gravitational potential energy of the gas that will establish its stability. when compared to the thermal energy. of this component.," It will be the total gravitational potential energy of the gas that will establish its stability, when compared to the thermal energy of this component."685 In this way. the determinant factor is the escape velocity. which is an integral property of the dark halo Eq.," In this way, the determinant factor is the escape velocity, which is an integral property of the dark halo Eq."686 3)., 3).687 This makes our results rather insensitive to the details of the clark matter profile taken., This makes our results rather insensitive to the details of the dark matter profile taken.688 We note the recent results of IxIevna et al. (, We note the recent results of Kleyna et al. (6892001) who use a new maximum likelihood analysis together. with recent measurements of stellar kinematics in Draco to model the dark matter halo of this galaxy.,2001) who use a new maximum likelihood analysis together with recent measurements of stellar kinematics in Draco to model the dark matter halo of this galaxy.690 Phese authors find again large amounts of dark matter. and a density profile well fitted by an isothermal pU)xRH almost into the very centre.," These authors find again large amounts of dark matter, and a density profile well fitted by an isothermal $\rho(R) \propto R^{-2}$ almost into the very centre."691 Phe resulting escape velocities being much the same as obtained [rom our Eq., The resulting escape velocities being much the same as obtained from our Eq.692 3., 3.693 ‘To compute the thermal energy of the eas (eragg). M is necessary to know the number of SN events occurring in our galaxies.," To compute the thermal energy of the gas $E_{THER}$ ), it is necessary to know the number of SN events occurring in our galaxies."694 We will use a chemical evolution code. which computes the SN rates. abundances. gas and stellar masses. as a function of gas accretion and star formation rates. as a funetion of time.," We will use a chemical evolution code, which computes the SN rates, abundances, gas and stellar masses, as a function of gas accretion and star formation rates, as a function of time."695" We have considered a one-zone chemical evolution moclel under the following assumptions: 1) The barvonic component of dSph galaxies is formed by an infall. X. of primordial material ((NX,=0.76. 3,— 0.24)."," We have considered a one-zone chemical evolution model under the following assumptions: 1) The baryonic component of dSph galaxies is formed by an infall, A, of primordial material $X_o=0.76$, $Y_o=0.24$ )."696 a) The initial infall rate decreases exponentially with tine., a) The initial infall rate decreases exponentially with time.697 ely is determined by requiring that the present clay total luminosity of the models matches observed values., $A_0$ is determined by requiring that the present day total luminosity of the models matches observed values.698 T is another input parameter ancl it is determined. by the beeinine of the first star formation episode., $\tau$ is another input parameter and it is determined by the begining of the first star formation episode.699 In cases with more than one episode. a secondary. infall begining | Car before the second. burst is assumed.," In cases with more than one episode, a secondary infall begining 1 Gyr before the second burst is assumed."700 b) Since Carina ane Leo L have two main star formation episodes. à constan secondary infall is considered.," b) Since Carina and Leo I have two main star formation episodes, a constant secondary infall is considered."701 The curation and. intensity of this infall are free parameters ancl are chosen such tha the second. galactie wind occurs after the maximum of the second SE burst., The duration and intensity of this infall are free parameters and are chosen such that the second galactic wind occurs after the maximum of the second SF burst.702 2) When μοι=Leaeay a wind develops anc, 2) When $E_{THER}=E_{GRAV}$ a wind develops and703where ~ and sy stand for the characteristic values of these quantities for a given particle clistribution.,where $\gamma$ and $\gamma_0$ stand for the characteristic values of these quantities for a given particle distribution.704 Lt should. be noted that the svstem (18) along with cquation (14) are literally the same as those for the scattering olf the rectilincarly moving particles. except for the factor 5 entering equation (14) (οἱ.eqs.(8)-(9)inPetrova2008a).," It should be noted that the system (13) along with equation (14) are literally the same as those for the scattering off the rectilinearly moving particles, except for the factor $\gamma_0^2$ entering equation (14) \citep[cf. eqs. (8)-(9) in][]{p07a}."705". Generally speaking. one can write that 55/5""E=1” and conclude that the scattering by spiraling particles with[o the longitudinal Lorentz-[actor | is τὔ times less ellicient than that by the particles streaming along the magnetic field with the same Lorentz-[actor 5j."," Generally speaking, one can write that $\gamma_0^2/\gamma^5=1/\gamma_\Vert^5\gamma_0^3$ and conclude that the scattering by spiraling particles with the longitudinal Lorentz-factor $\gamma_\Vert$ is $\gamma_0^3$ times less efficient than that by the particles streaming along the magnetic field with the same Lorentz-factor $\gamma_\Vert$."706 However. in pulsar case the situation is somewhat cdillerent.," However, in pulsar case the situation is somewhat different."707 As has been shown in Lyubarskii&Petrova(1998).. the svnchrotron absorption. which determines the evolution of the particle distribution function. acts to slow down the longitudinal motion and increase the transverse momenta in such a way that =το keeps constant as long as the particle pitch-angle is small enough. 56/7«0.," As has been shown in \citet{lp98}, the synchrotron absorption, which determines the evolution of the particle distribution function, acts to slow down the longitudinal motion and increase the transverse momenta in such a way that $\gamma=\gamma_0\gamma_\Vert$ keeps constant as long as the particle pitch-angle is small enough, $\gamma_0/\gamma <\theta$."708 Thus. it is reasonable to compare the scatterings by the original and evolved distributions of particles. in which case * is the same.," Thus, it is reasonable to compare the scatterings by the original and evolved distributions of particles, in which case $\gamma$ is the same."709 Then the spiraling particles scatter +6 times more ellicientlv than the streaming particles., Then the spiraling particles scatter $\gamma_0^2$ times more efficiently than the streaming particles.710 Note also that in the two processes the directions of predominant. scattering of the radio beam photons clilfer. oy~l/|=cu and 0177~L/*5. respectively. implving cilferent relations between the interacting frequencies. and vycO7257.," Note also that in the two processes the directions of predominant scattering of the radio beam photons differ, $\theta_1^{\rm max}\sim7111/\gamma_\Vert=\gamma_0/\gamma$ and $\theta_1^{\rm max}\sim7121/\gamma$, respectively, implying different relations between the interacting frequencies, $\nu_1\sim\nu\theta^2\gamma^2/\gamma_0^2$ and $\nu_1\sim\nu\theta^2\gamma^2$."713 ‘The system (13) has the first integral. and the solution is written as where.r—{μιiU)p££.gil]exp(Eo). and .Vo=Zar.," The system (13) has the first integral, and the solution is written as where $x\equiv [I_{\nu_1}^{(0)}/I_\nu^{(0)}]\exp (\Gamma_0)$ and $\Gamma_0\equiv Ia_0r$."714 ENThe latter quantity characterizes the scattering elliciency. whereas the former one the extent of intensity transfer from the racio beam to the background.," The latter quantity characterizes the scattering efficiency, whereas the former one the extent of intensity transfer from the radio beam to the background."715" As long as or<1l. £4,~{μιiUdexp(b'u)EB andl,©ἐνi1."," As long as $x\ll 1$, $I_{\nu_1}\sim I_{\nu_1}^{(0)}\exp (\Gamma_0)$ and $I_\nu\approx I_\nu^{(0)}$."716" M c~ {νι becomes comparable with+ 45""iU. and enters the stage of saturation.", At $x\sim 1$ $I_{\nu_1}$ becomes comparable with $I_\nu^{(0)}$ and enters the stage of saturation.717". Given(t that wousSed. {μιe""i1 and d,~ἐνiUdfre."," Given that $x\gg 1$, $I_{\nu_1}\approx I_\nu^{(0)}$ and $I_\nu\sim718I_\nu^{(0)}/x$."719 Sincem initiallyEM the intensity ratio of the background and the beam is extremely small. 42/41271cc dun10 is necessary to provide aoc Loin which case a substantial part of the beam intensity is transferred. to the background.," Since initially the intensity ratio of the background and the beam is extremely small, $I_{\nu_1}^{(0)}/I_\nu^{(0)}<<<1 $, $\Gamma_0\sim n\times 10$ is necessary to provide $x\sim 1$, in which case a substantial part of the beam intensity is transferred to the background."720 As has been shown in Petrova(2008a).. in case of the scattering by the streaming particles this condition may well be satisfied. in. pulsars.," As has been shown in \citet{p07a}, in case of the scattering by the streaming particles this condition may well be satisfied in pulsars."721 For the scattering by the spiraling particles Ly is 55 larger and. corresponcinely. the growth of the scattered component should be even more significant.," For the scattering by the spiraling particles $\Gamma_0$ is $\gamma_0^2$ larger and, correspondingly, the growth of the scattered component should be even more significant."722 Now we turn to induced scattering between such two states that one of the frequencies is close to the resonance. sey. and another one is well below the resonance.wg.," Now we turn to induced scattering between such two states that one of the frequencies is close to the resonance, $\omega\gamma\eta\sim s\omega_H$ , and another one is well below the resonance,."723" The intensity evolution in the two states is given by where the subscripts 7.j denote the polarization states of the initial ancl final photons. (der.,,,1/dO0,) is the section of the scattering from £ to £, in the guiding centre [rame (it is given by eq. ("," The intensity evolution in the two states is given by where the subscripts $i,j$ denote the polarization states of the initial and final photons, $({\rm d}\sigma_{\nu\to\nu_1}^{ij}/{\rm724d}O_1)^{(\rm c )}$ is the cross-section of the scattering from $\nu$ to $\nu_1$ in the guiding centre frame (it is given by eq. ("7257) and. should. be expressed. via the quantities of the laboratory system). the components of the momentum increment are given by equations (8) and (9) and it is taken into account that 457xwyoi.,"7) and should be expressed via the quantities of the laboratory system), the components of the momentum increment are given by equations (8) and (9) and it is taken into account that $\omega_1\gamma\eta_1\ll\omega_H,\omega\gamma\eta$."726" The symmetry of the scattering. probability (4) with respect to direct ancl inverse scatterings implies that ενfdOY-—(vtvi)(det4/0041)7. and one can(dez! see that [1αναι(ιάep,πι where à=euge 1."," The symmetry of the scattering probability (4) with respect to direct and inverse scatterings implies that $({\rm727d}\sigma_{\nu_1\to\nu}^{ji}/{\rm d}O)^{(\rm c )}=(\nu^4/\nu_1^4)728({\rm d}\sigma_{\nu\to\nu_1}^{ij}/{\rm d}O_1)^{(\rm c )}$, and one can see that $\vert{\rm d}I_\nu^i/{\rm d}r\vert=\alpha\vert{\rm729d}I_{\nu_1}^j/{\rm d}r\vert$, where $\alpha\equiv\nu\eta/\nu_1\eta_1\gg 1$ ."730 Comparing the increments of the transverse and longitudinal momenta we find that apiOf/üpy~ Op).," Comparing the increments of the transverse and longitudinal momenta we find that $\Delta731p_\perp\partial f/\partial732p_\perp\sim(\theta^2p_\Vert^2/p_\perp^2)\Delta p_\Vert\partial733f/\partial p_\Vert$ ."734 ln our consideration «f. Le. 1ο particleaos pitch-anele is much less thanpi/pj the photon wopagation angle. and hence the contribution of the ongitudinal increment can be neglected.," In our consideration $p_\perp/p_\Vert\ll\theta$, i.e. the particle pitch-angle is much less than the photon propagation angle, and hence the contribution of the longitudinal increment can be neglected."735" Note that pj enters equation (17) via f as well as via the argument. s£ of the Bessel function and its derivative in the cross-sections (7). €&—Josin8=pysind/pyyo2p,/pje1."," Note that $p_\perp$ enters equation (17) via $f$ as well as via the argument $s\xi$ of the Bessel function and its derivative in the cross-sections (7), $\xi\equiv\beta_0\sin\theta^{(c)}=p_\perp\sin\theta/p_\Vert\eta\sim7362p_\perp/p_\Vert\theta\ll 1$."737 One can see that integration of the kinetic equations (17) over pi w parts results in a change of the sign of the right-hand sides., One can see that integration of the kinetic equations (17) over $p_\perp$ by parts results in a change of the sign of the right-hand sides.738 Εις. in the course of the scattering the photons are mainly transferred from high harmonics. 7sey/2371. bo he zeroth one. £i«wau/2zin.," Thus, in the course of the scattering the photons are mainly transferred from high harmonics, $\nu\sim s\omega_H/2\pi\gamma\eta$, to the zeroth one, $\nu_1\ll\omega_H/2\pi\gamma\eta_1$."739" Phe harmonic number s enters cf,fdr via τς(Εξ) or sod.2By;(8E). which. peak at high. harmonics.. 5BITESag."," The harmonic number $s$ enters ${\rm d}I_\nu/{\rm d}r$ via $s^2J_s^2(s\xi)$ or $s^2J_S^{\prime^2}(s\xi)$, which peak at high harmonics, $s^{\rm max }\sim\gamma_0^3$."740 Note. rowever. that inside the light. evlinder ον corresponds to he frequencies ing in the optical or soft X-ray range (see.e.g..Petrova 2008c).. and the intensity of pulsar radiation in these bands is much weaker than the radio intensity.," Note, however, that inside the light cylinder $s^{\rm max}$ corresponds to the frequencies lying in the optical or soft X-ray range \citep[see,741e.g.,][]{p07c}, and the intensity of pulsar radiation in these bands is much weaker than the radio intensity."742 Therefore the induced: seattering from. high harmonics. So8Inax is expected. to. be inellicient.," Therefore the induced scattering from high harmonics, $s\sim s^{\rm max}$, is expected to be inefficient."743 In. the present consideration. we leave aside this problem and concentrate on the induced. scattering of the radio emission.," In the present consideration, we leave aside this problem and concentrate on the induced scattering of the radio emission."744 As the magnetic field. strength. decreases with distance from the neutron star. a given radio frequeney. passes through the resonances of increasingly higher order.," As the magnetic field strength decreases with distance from the neutron star, a given radio frequency passes through the resonances of increasingly higher order."745 At the same tine. the number density of the scattering particles anc the radio intensity strongly. decrease with distance. so that the scatteringbecomes much weaker.," At the same time, the number density of the scattering particles and the radio intensity strongly decrease with distance, so that the scatteringbecomes much weaker."746 Thus. we are interested in the radio beam scattering at several first harmonics of the σονοοσον. which takes place within the light cvlinder.," Thus, we are interested in the radio beam scattering at several first harmonics of the gyrofrequency, which takes place within the light cylinder."747 Then s€ is still a small quantity and one can take approximately that, Then $s\xi$ is still a small quantity and one can take approximately that748ΑΦ—XA.,"$A\,^{3}\Phi -749X\,^{3}\Delta$."750 We believe that the (1-1) band of the same system at ((not marked in Fig. 1) , We believe that the (1-1) band of the same system at (not marked in Fig. \ref{spectr_f}) )751"is also visible in our spectrum,", is also visible in our spectrum.752 Another prominent band of TiO is that of (0-0) of the y’ system X?A) nearAA.," Another prominent band of TiO is that of (0-0) of the $\gamma'$ system $B\,^{3}\Pi - X\,^{3}\Delta$ ) near."753. Some other emission features can also be identified with excited bands of the TIO y system., Some other emission features can also be identified with excited bands of the TiO $\gamma$ system.754 Following Herbig (1974)) we have identified the prominent emission features at 6040 and wwith ScO. AIO. which dominates the near-IR spectrum of BVAO3.. is also present in emission m our spectrum.," Following Herbig \cite{herbig}) ) we have identified the prominent emission features at 6040 and with ScO. AlO, which dominates the near-IR spectrum of \cite{baner}, is also present in emission in our spectrum."755 Four bands of the BX-—XX- electronic transitions are clearly seen near aand ΑΑ.," Four bands of the $B\,^{2}\Sigma^{+} - X\,^{2}\Sigma^{+}$ electronic transitions are clearly seen near and ."756". The strong feature at wwas initially attributed to H,. but we subsequently identified it as the IL intercombination transition."," The strong feature at was initially attributed to $\mbox{H}_\alpha$, but we subsequently identified it as the I intercombination transition."757" H, was seen in the spectrum of V4332 Ser during its 1994 eruption (MWT99)). but at that time its mean position wasAA."," $\mbox{H}_\alpha$ was seen in the spectrum of V4332 Sgr during its 1994 eruption \cite{martini}) ), but at that time its mean position was."758". Since our position of the Nall D lines differs only by ffrom that of MWT99.. it 1s difficult to reconcile a difference of im the H, position."," Since our position of the I D lines differs only by from that of \cite{martini}, it is difficult to reconcile a difference of in the $\mbox{H}_\alpha$ position."759 BAO4+ suggest that this feature might be due to TiO γ΄ (0.1) as the wavelength matches well.," \cite{banash} suggest that this feature might be due to TiO $\gamma'$ (0,1) as the wavelength matches well."760 In that case. other emissions originating from the same transition should be seen at aand but they are not present.," In that case, other emissions originating from the same transition should be seen at and but they are not present."761 Furthermore. Table 2 shows that the ffeature is significantly narrower than other molecular lines. suggesting that it is an atomic line.," Furthermore, Table \ref{spectr_t} shows that the feature is significantly narrower than other molecular lines, suggesting that it is an atomic line."762 Given the very low excitation of the spectrum of V4332 Ser we consider I (SP?) as the most probable identification for the ffeature.," Given the very low excitation of the spectrum of V4332 Sgr, we consider I $^1$ $-^3$ $^0$ ) as the most probable identification for the feature."763 As can be seen from Fig. |..," As can be seen from Fig. \ref{spectr_f},"764 the spectrum shows a clear continuum rising toward longer wavelengths., the spectrum shows a clear continuum rising toward longer wavelengths.765 The rather low signal-to-noise ratio does not allow unambiguous identification of absorption. features., The rather low signal-to-noise ratio does not allow unambiguous identification of absorption features.766 Yet à comparison. of the observed spectrum with synthetic ones suggest an early M type., Yet a comparison of the observed spectrum with synthetic ones suggest an early M type.767 An analysis of the strengths and widths of molecular bands allows us to estimate excitation conditions in the medium where the bands are produced., An analysis of the strengths and widths of molecular bands allows us to estimate excitation conditions in the medium where the bands are produced.768 For the TIO bands we have used the data base described by Schwenke (1998)) and distributed by Kurucz (1999))., For the TiO bands we have used the data base described by Schwenke \cite{schw}) ) and distributed by Kurucz \cite{kur}) ).769 The observed widths of the emission peaks of the (0-0) band of the y system of TiO AA)) has allowed us to estimate the rotational temperature. Ty.= 200KK. On the otherhand. the relative equivalent widths of these emission peaks point to the electronic temperature. Το2500 KK. Following the analysis of Herbig," The observed widths of the emission peaks of the (0-0) band of the $\gamma$ system of TiO ) has allowed us to estimate the rotational temperature, $T_\mathrm{rot} \simeq770200$ K. On the otherhand, the relative equivalent widths of these emission peaks point to the electronic temperature, $T_\mathrm{el} \ga771500$ K. Following the analysis of Herbig"772optical counterparts of several Galactic X-ray novae which almost certainly harbor BHs (Remillard&MeClintock2006).,optical counterparts of several Galactic X-ray novae which almost certainly harbor BHs \citep{rem06}.773 The optical modulation is induced mostly by the ellipsoidal effect from the nondegenerate component. with an amplitude depending mainly on the Roche lobe filling factor and the inclination of the orbit.," The optical modulation is induced mostly by the ellipsoidal effect from the nondegenerate component, with an amplitude depending mainly on the Roche lobe filling factor and the inclination of the orbit."774 Close degenerate binaries can be unambiguously identified as single-line systems with large orbital velocities (K>150 km s7!)., Close degenerate binaries can be unambiguously identified as single-line systems with large orbital velocities $K >150$ km $^{-1}$ ).775 Among them. those harboring a black hole are distinguished by a mass function fn=GNSINE[Up>2M. where mp; and n stand for the total mass of the binary and the mass of one of its components (Remillard&MeClintock2006).," Among them, those harboring a black hole are distinguished by a mass function $f_m = (m\sin i)^3/(m_{bin})^2 >2M_{\odot}$, where $m_{bin}$ and $m$ stand for the total mass of the binary and the mass of one of its components \citep{rem06}."776. Our sample and observational data are introduced in Sect., Our sample and observational data are introduced in Sect.777 2., 2.778 A detailed analysis of the data is reported in Sect., A detailed analysis of the data is reported in Sect.779 3. and the results are discussed in Sect.," 3, and the results are discussed in Sect."780 4., 4.781 Optical counterparts of active X-ray sources in GCs have been found essentially everywhere in color-magnitude diagrams: to the left of the main sequence. on the main sequence. to the right of it. on subgiant. giant and horizontal branches. and also above the main-sequence turnoff in the regions occupied by blue stragglers and EHB stars (e.g. Heinke et al.," Optical counterparts of active X-ray sources in GCs have been found essentially everywhere in color-magnitude diagrams: to the left of the main sequence, on the main sequence, to the right of it, on subgiant, giant and horizontal branches, and also above the main-sequence turnoff in the regions occupied by blue stragglers and EHB stars (e.g. Heinke et al."782 2005; Servillat et al., 2005; Servillat et al.783 2008)., 2008).784 Thus. while preparing this study we felt it justified to pick photometrically suitable targets without paying attention to their location with respect to the main sequence.," Thus, while preparing this study we felt it justified to pick photometrically suitable targets without paying attention to their location with respect to the main sequence."785 Since large-amplitude variables are more likely to be ordinary contact binaries. we avoided objects with full V-band amplitudes in excess of ~0.35 mag.," Since large-amplitude variables are more likely to be ordinary contact binaries, we avoided objects with full $V$ -band amplitudes in excess of $\sim$ 0.35 mag."786 Suspected pulsating variables were excluded: we also took care to exclude systems with appreciable X-ray emission., Suspected pulsating variables were excluded; we also took care to exclude systems with appreciable X-ray emission.787 Where it was possible. the membership status of the target was established based on the data of Bellinietal.(2009) for ω Cen. and Anderson (priv.," Where it was possible, the membership status of the target was established based on the data of \cite{bel09} for $\omega$ Cen, and Anderson (priv."788 comm.), comm.)789 and Strickleretal.(2009) for NGC 6397.," and \cite{str09}790 for NGC 6397."791 In the remaining cases membership was decided based on the systemic velocity with respect to the cluster., In the remaining cases membership was decided based on the systemic velocity with respect to the cluster.792 Alltogether. we selected seven objects in w Cen and four in NGC 6397.," Alltogether, we selected seven objects in $\omega$ Cen and four in NGC 6397."793 In color-magnitude diagrams two of them are located to the left of the main sequence. seven above the main-sequence turnoff. and two at the turnoff itself (see Figs.," In color-magnitude diagrams two of them are located to the left of the main sequence, seven above the main-sequence turnoff, and two at the turnoff itself (see Figs."794 | and 2)), \ref{fig: omegacmd} and \ref{fig: n6397cmd}) ).795 A summary of their basic data is given in Table | (the labels are the same as those introducedby Kaluznyetal. for w Cen and Kaluznyetal.(2006) for NGC 6397).," A summary of their basic data is given in Table \ref{tab: objects}796 (the labels are the same as those introducedby \cite{kal04} for $\omega$ Cen and \cite{kal06} for NGC 6397)."797 Our paper is based oi photometric measurements. of Kaluznyetal.(2004.2006). supplemented by those of Weldrakeetal.(2007) and by our unpublished data for NGC 6397 obtained on the du Pont telescope at Las Campanas Observatory. Chile (a total of 771 V-band frames was acquired between 2009 June 20 and June 30. and reduced in the same way as in Kaluznyetal. (2006))).," Our paper is based on photometric measurements of \cite{kal04,kal06}, supplemented by those of \cite{wel07} and by our unpublished data for NGC 6397 obtained on the du Pont telescope at Las Campanas Observatory, Chile (a total of 771 $V$ -band frames was acquired between 2009 June 20 and June 30, and reduced in the same way as in \cite{kal06}) )."798 The newer data demonstrated that the observed light variations were coherent. although in some cases the light curve was changing from symmetric to asymmetric or vice versa.," The newer data demonstrated that the observed light variations were coherent, although in some cases the light curve was changing from symmetric to asymmetric or vice versa."799 Spectroscopic data were collected during the nights 20/21 and 21/22 of May 2009 with the MagE (Magellan Echellette) spectrograph attached to the 6.5-m Magellan-Clay telescope at Las Campanas Observatory., Spectroscopic data were collected during the nights 20/21 and 21/22 of May 2009 with the MagE (Magellan Echellette) spectrograph attached to the 6.5-m Magellan-Clay telescope at Las Campanas Observatory.800 The seeing varied betweeen 045 and 08 on the first night. and betweeen 074 and 079 on the second one.," The seeing varied betweeen $0\farcs5$ and $0\farcs8$ on the first night, and betweeen $0\farcs4$ and $0\farcs9$ on the second one."801 A 0785 slit was used. providing a resolution R= 4820.," A $0\farcs85$ slit was used, providing a resolution $R=4820$ ."802 During the observations pairs of scientific spectra taken for the same target were separated by an exposure of a thorium-argon hollow-cathode lamp., During the observations pairs of scientific spectra taken for the same target were separated by an exposure of a thorium-argon hollow-cathode lamp.803 The exposure times per spectrum ranged from 120 s to 1080 s. depending on object brightness and observing conditions.," The exposure times per spectrum ranged from 120 s to 1080 s, depending on object brightness and observing conditions."804 After bias and flat-field correction each pair of frames was combined into a single frame. allowing for the rejection of cosmic ray hits.," After bias and flat-field correction each pair of frames was combined into a single frame, allowing for the rejection of cosmic ray hits."805 The observations were reduced with the ECHELLE package., The observations were reduced with the ECHELLE package.806 In order to determine the mass function of a given target it was sufficient to take Just a few spectra at different phases. and fit a simple sinusoid to its velocity curve phased with the photometric ephemeris.," In order to determine the mass function of a given target it was sufficient to take just a few spectra at different phases, and fit a simple sinusoid to its velocity curve phased with the photometric ephemeris."807 Altogether. 96 spectra were obtained (from 5 to 10 per object).," Altogether, 96 spectra were obtained (from 5 to 10 per object)."808 The useful range of the reduced spectra extended from 4050 tto 6800A.. where the spectra had 13<S/N23 (with a few cases of lower quality).," The useful range of the reduced spectra extended from 4050 to 6800, where the spectra had $13<\mathrm{S/N}<23$ (with a few cases of lower quality)."809 Radial velocities were measured in that range with the help of IRAF routines FXCOR and XCSAO. using synthetic templates from the library compiled by Munart (2005)...," Radial velocities were measured in that range with the help of IRAF routines FXCOR and XCSAO, using synthetic templates from the library compiled by \cite{mun05}. ."810 The results were verified by applying the broadening function formalism described by Rucinski (2002).., The results were verified by applying the broadening function formalism described by \cite{ruc02}. .811enhancements at high latitudes (b=—30°: Snowden et al.,enhancements at high latitudes $b\gtrsim-30^\circ$; Snowden et al.812 1997)., 1997).813 Within |b)<10. the interstellar absorption is severe. little can be inferred reliably about the properties of the hot eas.," Within $|b| \lesssim 10^\circ$, the interstellar absorption is severe, little can be inferred reliably about the properties of the hot gas."814 It is in (his corresponding region in (he bbulge (|:<6’) that the diffuse soft X-ray intensity shows the steepest increase (by about one order of magnilite) toward the ealactie center., It is in this corresponding region in the bulge $|z| \lesssim 6^\prime$ ) that the diffuse soft X-ray intensity shows the steepest increase (by about one order of magnitute) toward the galactic center.815 Such a mid-plane concentration of diffuse soft X-ray emission may also be present instrinsically in our Galactic bulge., Such a mid-plane concentration of diffuse soft X-ray emission may also be present instrinsically in our Galactic bulge.816 Clealv. a more careful comparison ancl modeling of the X-ray data sets are needed in order to understand the similaritv and dillerence in the hot eas characteristics and their relationship to other ealactic properties (e.g.. the effect of recent active star lormation in (he Galactic center).," Clealy, a more careful comparison and modeling of the X-ray data sets are needed in order to understand the similarity and difference in the hot gas characteristics and their relationship to other galactic properties (e.g., the effect of recent active star formation in the Galactic center)."817 We thank D. Calzett. M. Fardal. and S. Tang for helpful comments and discussions.," We thank D. Calzetti, M. Fardal, and S. Tang for helpful comments and discussions."818 This work is supported bv the SAO grant. ART-3006X., This work is supported by the SAO grant AR7-8006X.819"also use only one stage, 10.7% of the training examples is misclassified (see Table 4)).","also use only one stage, $10.7\%$ of the training examples is misclassified (see Table \ref{cm_sg}) )."820" We can thus conclude that our multi-stage classification tree with Gaussian mixtures at its nodes, is a significant improvement."," We can thus conclude that our multi-stage classification tree with Gaussian mixtures at its nodes, is a significant improvement."821 We analyzed 25947 light curves in the TrES Lyr1 field.," We analyzed $25\,947$ light curves in the TrES Lyr1 field."822" TrES, the Trans-atlantic Exoplanet Survey, is a network of three ten-centimeter optical telescopes searching the sky for transiting planets (Alonsoetal.2007;O’Donovan2008)."," TrES, the Trans-atlantic Exoplanet Survey, is a network of three ten-centimeter optical telescopes searching the sky for transiting planets \citep{Alonso:2007, Donovan:2008}."823". This network consisted of Sleuth (Palomar Observatory, Southern California), the PSST (Lowell Observatory, Northern Arizona) and STARE (Observatorio del Teide, Canary Islands, Spain), as TrES now excludes Sleuth and STARE, but includes WATTS."," This network consisted of Sleuth (Palomar Observatory, Southern California), the PSST (Lowell Observatory, Northern Arizona) and STARE (Observatorio del Teide, Canary Islands, Spain), as TrES now excludes Sleuth and STARE, but includes WATTS."824" The TrES Lyrl field is a 5.7?x field, centered on the star 16 Lyr and is part of the field (Alonsoetal. 2007).."," The TrES Lyr1 field is a $5.7^{\circ} \times8255.7^{\circ}$ field, centered on the star 16 Lyr and is part of the field \citep{Alonso:2007}. ."826 Most light curves have about 15000 observations spread with a total time span of aproximately 75 days.," Most light curves have about $15\,000$ observations spread with a total time span of aproximately 75 days."827 A small fraction has less than 5000 observations with a total time span of around 62 days.," A small fraction has less than $5\,000$ observations with a total time span of around 62 days."828 Observations are given in either the Sloan r (Sleuth) or the Kron-Cousins R magnitude (PSST) and the mean R magnitude ranges from 9.2 to 16.3., Observations are given in either the Sloan r (Sleuth) or the Kron-Cousins R magnitude (PSST) and the mean R magnitude ranges from $9.2$ to $16.3$.829" With the use of the variability detection algorithm, described in section ??,, we searched for frequencies in the range 3/Tiot to 50 c/d, with Tio, the total timespan of the observations in days."," With the use of the variability detection algorithm, described in section \ref{sec:vd}, we searched for frequencies in the range $3/T_{tot}$ to 50 c/d, with $T_{tot}$ the total timespan of the observations in days."830" In order to avoid the problem of daily aliasing in an automated way, small frequency intervals around multiples of 1 c/d were flagged as “unreliable”."," In order to avoid the problem of daily aliasing in an automated way, small frequency intervals around multiples of 1 c/d were flagged as “unreliable”."831" Using a false alarm probability of a=0.005 (the null-hypothesis of only having noise in the light curves is rejected when P«a, with P the probability of finding such a peak in the power spectrum of a time series that only contains noise.),"," Using a false alarm probability of $\alpha = 0.005$ (the null-hypothesis of only having noise in the light curves is rejected when $P<\alpha$, with $P$ the probability of finding such a peak in the power spectrum of a time series that only contains noise.),"832 about 180000 objects were found non-constant., about 000 objects were found non-constant.833" The stars for which we could not find significant frequencies were used to determine the RMS level of the time series as a function of the mean magnitude, which is plotted in Fig. 3,,"," The stars for which we could not find significant frequencies were used to determine the RMS level of the time series as a function of the mean magnitude, which is plotted in Fig. \ref{noise},"834 indicating to what level we can detect variability., indicating to what level we can detect variability.835 The upward trend can be explained in terms of photon noise., The upward trend can be explained in terms of photon noise.836" We used the multi-stage tree presented in section ??,, where we excluded the stars with activity and variables with rotational modulation."," We used the multi-stage tree presented in section \ref{sec:tc}, where we excluded the stars with activity and variables with rotational modulation."837" As already mentioned earlier, these classes were included inthe multi-stage tree in view of the mission."," As already mentioned earlier, these classes were included inthe multi-stage tree in view of the mission."838" However, we do not expect to"," However, we do not expect to"839the origins of rocky planets in our solar system.,the origins of rocky planets in our solar system.840 10 may be dillicult. for rocky exoplanets to form as close to their stars as ColtoT-7 b and ο because high temperatures. in the protoplanetary gas disk may inhibit condensation ancl aeeretion of solid. materials(e.g... 2)).," It may be difficult for rocky exoplanets to form as close to their stars as CoRoT-7 b and c because high temperatures in the protoplanetary gas disk may inhibit condensation and accretion of solid materials, \citealt{1994ApJ...421..615P}) )."841 At the very least. solid planets that form close-in will probably be depleted in volatile materials(e.g... water 2].," At the very least, solid planets that form close-in will probably be depleted in volatile materials, water \citealt{2008MNRAS.384..663R}] ])."842 Instead. of forming in heir current orbits. close-in rocky exoplanets like ColtoT-7 » may have coalesced. in orbits farther away from their lost stars. where disk temperatures were lower. ancl were ought to their current orbits.," Instead of forming in their current orbits, close-in rocky exoplanets like CoRoT-7 b may have coalesced in orbits farther away from their host stars, where disk temperatures were lower, and were brought to their current orbits."843 A variety of processes: Lor ringing planets to such close-in orbits have been proposed (??77).. and the orbital architecture of svstems with close-in ralanets and the planets! physical properties. provide. clues o their histories (?)..," A variety of processes for bringing planets to such close-in orbits have been proposed \citep{2004ApJ...614..955M, 2007Sci...318..210G, 2008MNRAS.384..663R}, and the orbital architecture of systems with close-in planets and the planets' physical properties provide clues to their histories \citep{2008MNRAS.384..663R}."844 One process that may have been important during ColtoT-7 bs history is evaporative mass loss., One process that may have been important during CoRoT-7 b's history is evaporative mass loss.845 Evaporation is thought to inlluence gaseous exoplanets made primarily of I1 and Le., Evaporation is thought to influence gaseous exoplanets made primarily of H and He.846" For example. observations of LLL 2094558 b suggest the planet's atmosphere is evaporating at a rate of order 0.1 Adega, (ve (72).. in line with theoretical expectationsfrom studies of evaporative mass loss (22277). "," For example, observations of HD 209458 b suggest the planet's atmosphere is evaporating at a rate of order 0.1 $M_{Earth}$ /Gyr \citep{2003Natur.422..143V,2004ApJ...604L..69V}, in line with theoretical expectationsfrom studies of evaporative mass loss \citep{2003ApJ...598L.121L, 2004A&A...418L...1L, 2004Icar..170..167Y, 2005ApJ...621.1049T, 2008SSRv..139..437Y}."847Aloreover. gaps in the distributions of orbital and. physical properties of close-in planets have also been interpreted to indicate that many gaseous. planets have had. their atmospheres completely. evaporated. (2)...," Moreover, gaps in the distributions of orbital and physical properties of close-in planets have also been interpreted to indicate that many gaseous planets have had their atmospheres completely evaporated \citep{2009MNRAS.396.1012D}."848 Phe complete removal of a eas giants atmosphere would. likely leave behind its rocky core. with a mass of perhaps several Miss CUT).," The complete removal of a gas giant's atmosphere would likely leave behind its rocky core, with a mass of perhaps several $M_{Earth}$ \citep{2004A&A...419L..13B, 2008MNRAS.384..663R}."849 These considerations suggest. there may be a large population of close-in) planetary bodies. that are remnant cores of evaporated gas giants (22)...," These considerations suggest there may be a large population of close-in planetary bodies that are remnant cores of evaporated gas giants \citep{2005A&A...436L..47B, 2008PhDT........15H}."850 On the other hand. other studies argue that the observations. of Η 209458 b do not imply significant loss of mass (??).. and a theoretical study by 2? suggested. that complete evaporation of a gas giant’s atmosphere is unlikely.," On the other hand, other studies argue that the observations of HD 209458 b do not imply significant loss of mass \citep{2007ApJ...671L..61B, 2008Natur.451..970H}, and a theoretical study by \citet{2009ApJ...693...23M} suggested that complete evaporation of a gas giant's atmosphere is unlikely."851 These competing hypotheses may now be testable., These competing hypotheses may now be testable.852 With the detection capabilities of the anc missions. rocky planets arising from a variety of histories may be detected.," With the detection capabilities of the and missions, rocky planets arising from a variety of histories may be detected."853 Evaporation of mass may continue even after the planet loses its original Hl/lle envelope., Evaporation of mass may continue even after the planet loses its original H/He envelope.854 ? suggest the strong insolation received by ColioT-7 b may be sullicient to have evaporated: ancl removed several Earth masses of material. if the planet has always been solid. (made of ice and/or rock).," \citet{2009arXiv0907.3067V} suggest the strong insolation received by CoRoT-7 b may be sufficient to have evaporated and removed several Earth masses of material, if the planet has always been solid (made of ice and/or rock)."855 ? showed that thermal evaporation of rocky material from ColtoT-7 b's surface may produce a tenuous exosphere. similar to Mercurys.," \citet{2009ApJ...703L.113S} showed that thermal evaporation of rocky material from CoRoT-7 b's surface may produce a tenuous exosphere, similar to Mercury's."856 In this study. we consider evaporation both of a putative gaseous. envelope. and. of solid material during ColtoT-7 b's history. including effects ofits evolving In addition to mass loss. tides have plaved an important and interrelated role in ColtoT-7 bs history.," In this study, we consider evaporation both of a putative gaseous envelope and of solid material during CoRoT-7 b's history, including effects of its evolving In addition to mass loss, tides have played an important and interrelated role in CoRoT-7 b's history."857 Tides have shaped the distribution. of orbital elements. for. close-in exoplanets. reducing eccentricities (2277). ancl semi-major axes (οοτι. ," Tides have shaped the distribution of orbital elements for close-in exoplanets, reducing eccentricities \citep{1996ApJ...470.1187R, 2006ApJ...638L..45F, 2008ApJ...678.1396J} and semi-major axes \citep{2009MNRAS.395.2268B, 2009ApJ...698.1357J, 2009ApJ...692L...9L}."858"""ides may also help to bring rocky planets into close-in orbits (2)..", Tides may also help to bring rocky planets into close-in orbits \citep{2008MNRAS.384..663R}.859 In fact. going forward in time. ColtoT-7 b may migrate into its host star in less than a few billion wears. and it may have undergone past orbital migration (?)..," In fact, going forward in time, CoRoT-7 b may migrate into its host star in less than a few billion years, and it may have undergone past orbital migration \citep{2009ApJ...698.1357J}."860 Also. ColtoT-7 b likely has undergone or currently undergoes strong tidal heating. either because its original orbit was eccentric or interactions with nearby ColtoT-7 c keep the eccentricity non-zero (?)..," Also, CoRoT-7 b likely has undergone or currently undergoes strong tidal heating, either because its original orbit was eccentric or interactions with nearby CoRoT-7 c keep the eccentricity non-zero \citep{2010ApJ...709L..95B}."861 Vidal evolution of its orbit may have been even faster in the past if Coho1-7 b once had a massive gaseous envelope. because the rate of evolution of a circular orbit. scales with its mass., Tidal evolution of its orbit may have been even faster in the past if CoRoT-7 b once had a massive gaseous envelope because the rate of evolution of a circular orbit scales with its mass.862 As the planet lost mass. though. the tical evolution may have slowed as the planet reached its current In this paper. we study the coupling of evaporative mass loss and tidal evolution. on ColtoT-7 b. Using parameterized models for mass loss and. orbital decay. due to tides. we evolved ColtoT-7 b's mass ancl orbit backwarc (and. forward) in time to determine the range of origina masses and orbits consistent with the planets current mass and orbit.," As the planet lost mass, though, the tidal evolution may have slowed as the planet reached its current In this paper, we study the coupling of evaporative mass loss and tidal evolution on CoRoT-7 b. Using parameterized models for mass loss and orbital decay due to tides, we evolved CoRoT-7 b's mass and orbit backward (and forward) in time to determine the range of original masses and orbits consistent with the planet's current mass and orbit."863 Although the exact initial conditions. depen: sensitively on the chosen model parameters. we considerec a wide range of parameters to map out all the possibilities.," Although the exact initial conditions depend sensitively on the chosen model parameters, we considered a wide range of parameters to map out all the possibilities."864 We find that. if ColtoT-7 b has always been a solid planet. it mav have lost as much as half its original mass ancl hac its semi-major axis reduced by as much as by tida migration alone.," We find that, if CoRoT-7 b has always been a solid planet, it may have lost as much as half its original mass and had its semi-major axis reduced by as much as by tidal migration alone."865" 1 C'oltoT-7 b was originally a gas gian planet ancl only recently lost its gaseous envelope. its origina mass οἱ not exceed 200 375,45. and its original semi-major axis could have been almost twice as large as the curren value."," If CoRoT-7 b was originally a gas giant planet and only recently lost its gaseous envelope, its original mass did not exceed 200 $M_{Earth}$, and its original semi-major axis could have been almost twice as large as the current value."866 In both scenarios. ColtoT-7 b may have encounterec and migrated through a 3:1 mean motion resonance with nearby ColtoT-7 c. as tides caused planet bs orbit to decay.," In both scenarios, CoRoT-7 b may have encountered and migrated through a 3:1 mean motion resonance with nearby CoRoT-7 c, as tides caused planet b's orbit to decay."867 Consideration of the resulting dvnamical interactions may place constraints on the orbital history of the C'oltoT-7 system and the original mass of ColtoT-7 b. We also discuss how the coupling between orbital decay and. mass loss may play an important in the evolution and survival of gaseous close-in planets., Consideration of the resulting dynamical interactions may place constraints on the orbital history of the CoRoT-7 system and the original mass of CoRoT-7 b. We also discuss how the coupling between orbital decay and mass loss may play an important in the evolution and survival of gaseous close-in planets.868 Several recent studies of evaporation. of exoplanetary atmospheres provide. simple. parameterized models for mass loss rates(e.g... 2??)).," Several recent studies of evaporation of exoplanetary atmospheres provide simple, parameterized models for mass loss rates, \citealt{2004A&A...419L..13B,2005A&A...439..771J, 2007A&A...472..329E}) )."869 Other authors have applied conventional mocels for tidal damping to study the orbital evolution of close-in cxoplancts (77777) A," Other authors have applied conventional models for tidal damping to study the orbital evolution of close-in exoplanets \citep{1963MNRAS.126..257G, 1966Icar....5..375G, 1996ApJ...470.1187R, 2008CeMDA.101..171F, 2008ApJ...678.1396J}."870lthough the wo processes may be coupled. few studies have considered o)Qh processes together.," Although the two processes may be coupled, few studies have considered both processes together."871 2?/— investigated: the coupling tween. tical heating and Roche lobe overllow. but they only considered the orbital ellects of the tide raised on the λαοί by the host star.," \citet{2003ApJ...588..509G} investigated the coupling between tidal heating and Roche lobe overflow, but they only considered the orbital effects of the tide raised on the planet by the host star."872 For close-in exoplanets on nearly circular orbits. like ColtoT-7 b. the tide raised on the host star by the planet can dominate the orbital evolution.," For close-in exoplanets on nearly circular orbits, like CoRoT-7 b, the tide raised on the host star by the planet can dominate the orbital evolution."873 In his section. we describe our mocel for coupled mass-orbital evolution and how we can use it to constrain the original orbit and mass of ColtoT-7 For energv-limited evaporative mass loss. the rate at which escaping gas molecules carry away energy from the planet is roughly proportional to the rate of input of energv [from stellar insolation (?7)..," In this section, we describe our model for coupled mass-orbital evolution and how we can use it to constrain the original orbit and mass of CoRoT-7 For energy-limited evaporative mass loss, the rate at which escaping gas molecules carry away energy from the planet is roughly proportional to the rate of input of energy from stellar insolation \citep{2008SSRv..139..437Y}. ."874 Relating the rate of energy input from insolation to the change in gravitational energy required for, Relating the rate of energy input from insolation to the change in gravitational energy required for875 e-0.7 Dor the low-J transitions to 21 for the CO(7-G) transition relative to the brighest line.,$\sim$ 0.7 for the low-J transitions to $\sim$ 1 for the CO(7-6) transition relative to the brighest line.876 For Model HI which has some of the physical conditions necessary for producing a top-heavy IMF. the ratios go [from ~0.9 for the low-J (transitions to ~0.7 for the CO(7-6) line relative to the brightest line.," For Model III which has some of the physical conditions necessary for producing a top-heavy IMF, the ratios go from $\sim$ 0.9 for the low-J transitions to $\sim$ 0.7 for the CO(7-6) line relative to the brightest line."877 Thus. the trencls predicted by Model HI seem (o match (he observations for this source better than those predicted by. Model HH.," Thus, the trends predicted by Model III seem to match the observations for this source better than those predicted by Model II."878 Once again. we emphasise (hat we have not tried. in our models. to mateh the physical conditions for this source in any wav whatsoever or to (rv ancl reproduce (he observed SED of this galaxy.," Once again, we emphasise that we have not tried, in our models, to match the physical conditions for this source in any way whatsoever or to try and reproduce the observed SED of this galaxy."879 However. our theoretical work does present some evidence. albeit speculative. that the Cloverleal source at redshilt 02.5 has roughly solar metallicity ancl potentially a top-heavy IME.," However, our theoretical work does present some evidence, albeit speculative, that the Cloverleaf source at redshift $\sim$ 2.5 has roughly solar metallicity and potentially a top-heavy IMF."880 It is expected Chat future observational studies will be able to test the validity of this prediction., It is expected that future observational studies will be able to test the validity of this prediction.881 Although the focus of the paper so [ar has been on the CO SED associated with different physical conditions. the PPDR code includes a network of 83 chemical species whose abundances are also output by the code.," Although the focus of the paper so far has been on the $^{12}$ CO SED associated with different physical conditions, the PDR code includes a network of 83 chemical species whose abundances are also output by the code."882 In this section we consider how the fractional abundances of some of (hese different molecular species that could potentially be observed with future survevs. change lor dillerent ΙΔΙΕΣ at. high redshilt.," In this section we consider how the fractional abundances of some of these different molecular species that could potentially be observed with future surveys, change for different IMFs at high redshift."883 The fractional abundance is defined (io be the abundance of the molecule relative to the total hydrogen abundance. ny=(ll)+n(IHs).," The fractional abundance is defined to be the abundance of the molecule relative to the total hydrogen abundance, $n_H=n(H)+n(H_2)$."884 Once again we consider Models I. HE and HII of BOS.," Once again we consider Models I, II and III of B09."885" In Table 4 we list the fractional abundances of some kev species [ον the three models both at A,3 and A,~8 typical of the (translucent gas component and dense PDR. eas component of galaxies.", In Table \ref{tab:mol} we list the fractional abundances of some key species for the three models both at $A_v \sim 3$ and $A_v \sim 8$ typical of the translucent gas component and dense PDR gas component of galaxies.886 In Figure 16. we also plot the variation of the molecular abundances of some kev species with depth inside the cloud for the different high-redshift moclels., In Figure \ref{fig:abundance} we also plot the variation of the molecular abundances of some key species with depth inside the cloud for the different high-redshift models.887 From Table - and Figure 16.. we observe the following trends.," From Table \ref{tab:mol} and Figure \ref{fig:abundance}, we observe the following trends."888 Assuming a limit οἱ detectability of LO! in fractional abundance which is arbitrary but roughly satisfied in our own galaxy (?).. all the species apart from CO are below (his limit for Model I where the extremely high cosmic rav [hixes and FUV fields serve to destroy molecules effectively even deep within the cloud.," Assuming a limit of detectability of $^{-12}$ in fractional abundance which is arbitrary but roughly satisfied in our own galaxy \citep{Bayet:09}, all the species apart from CO are below this limit for Model I where the extremely high cosmic ray fluxes and FUV fields serve to destroy molecules effectively even deep within the cloud."889 Note (hat our qualitative results do not change if we assume a slightly higher or lower detectability. limit., Note that our qualitative results do not change if we assume a slightly higher or lower detectability limit.890" For Model HI. the fractional abundances of all molecules considered here are above the detectability limit even at low A, where we normally do not expect complex molecules to be present."," For Model III, the fractional abundances of all molecules considered here are above the detectability limit even at low $A_v$ where we normally do not expect complex molecules to be present."891" ILowever. (he high density in this model means (hat the svstenis considered here are alwavs likely to be compact rather than diffuse and the size of the cloud is smaller (han for any of the other high redshift models at a given sl, as is illustrated by Figure 15.."," However, the high density in this model means that the systems considered here are always likely to be compact rather than diffuse and the size of the cloud is smaller than for any of the other high redshift models at a given $A_v$ as is illustrated by Figure \ref{fig:size}. ."892 The high fractional abundance of CO in this model also traces a high abundance of [fo even at, The high fractional abundance of CO in this model also traces a high abundance of $H_2$ even at893from 10A4. which 5give a tegood mateh to the line.,"from 5–10, which give a good match to the line."894 In the simulation the upperDPI 1pole is at the 73 o'clock” position at phase zero and rotates counterclockwise., In the simulation the upper pole is at the “3 o'clock” position at phase zero and rotates counterclockwise.895 At phases 0.20.3 we are seeing the upper pole furthest from us and the lower pole towards us., At phases 0.2–0.3 we are seeing the upper pole furthest from us and the lower pole towards us.896 Phe material at 510 iis beginning to fall back to the plane., The material at 5–10 is beginning to fall back to the plane.897 Ehe inclination σεtilts this motion towards the plane of the sky. reducing the projected: velocities. and the profile is thus narrowest.," The inclination tilts this motion towards the plane of the sky, reducing the projected velocities, and the profile is thus narrowest."898 Half a evele later (upper pole towards us). the inclination tilts the infall motion towards the linc-ol-sight. and the observed. velocities are higher.," Half a cycle later (upper pole towards us), the inclination tilts the infall motion towards the line-of-sight, and the observed velocities are higher."899 The profiles [rom aare noisier than those fromA4686:: alter subtraction of the phasc-invariant profile only 3 per cent of the original Iline remained., The profiles from are noisier than those from: after subtraction of the phase-invariant profile only 3 per cent of the original line remained.900 However. the profile can be interpreted. as being similar to that ofA4686.. with the exception that the blue wing seen at phases 0.40.5 is missing.," However, the profile can be interpreted as being similar to that of, with the exception that the blue wing seen at phases 0.4–0.8 is missing."901 In the previous sections we have compared: observed. line profiles with models caleulated for emission 5.10 {from the white dwarf., In the previous sections we have compared observed line profiles with models calculated for emission 5–10 from the white dwarf.902 We should emphasise that we do not regard the models as fits. since the model contains too many assumptions and adjustable parameters to lead to à unique solution.," We should emphasise that we do not regard the models as fits, since the model contains too many assumptions and adjustable parameters to lead to a unique solution."903 Instead. we describe here the result of varving the parameters.," Instead, we describe here the result of varying the parameters."904 Since the simulation calculates infall velocities. changing oor the white-cdwarf mass simply scales the velocities without altering the shape of the model profiles.," Since the simulation calculates infall velocities, changing or the white-dwarf mass simply scales the velocities without altering the shape of the model profiles."905 Similarly. adding an injection velocity at wwould increase the velocities at all points.," Similarly, adding an injection velocity at would increase the velocities at all points."906 Increasing the rrange of azimuth over which accretion occurs would smear out the profiles. but retain the overall shape.," Increasing the range of azimuth over which accretion occurs would smear out the profiles, but retain the overall shape."907 Llowever. the shape of the spin-resolvecl profiles. is allected by the dipole offset. combined with the inclination.," However, the shape of the spin-resolved profiles is affected by the dipole offset, combined with the inclination."908 At zero inclination. of course. our view would not change with spin phase and there would be no spin-evcle mocdulation.," At zero inclination, of course, our view would not change with spin phase and there would be no spin-cycle modulation."909 This rules out the extreme inclinations of « clcliscussedl in Section 5 only by adopting an inclination ofat least dedo we reproduce the observed. spin-resolved: profiles. in which the narrow profile at phase 0.2 changes to a broader. split. profile at. phase 70.7.," This rules out the extreme inclinations of $<$ discussed in Section 5 — only by adopting an inclination of at least do we reproduce the observed spin-resolved profiles, in which the narrow profile at phase 0.2 changes to a broader, split profile at phase 0.7."910 Even with an inclination of10.. the profile is only reproduced with 6 753307: any lower and the change in viewing angle over the spin evcle is insullicient to reproduce the observed profiles.," Even with an inclination of, the profile is only reproduced with $\delta$ $>$: any lower and the change in viewing angle over the spin cycle is insufficient to reproduce the observed profiles."911 Of course. we can also reproduce the spin-resolved: profiles using a higher inclination. but that only increases the problem of the lack of orbital modulation.," Of course, we can also reproduce the spin-resolved profiles using a higher inclination, but that only increases the problem of the lack of orbital modulation."912 We have shown that the emission lines in, We have shown that the emission lines in913"ol magnetic structures overlving (he filament (the large-scale magnetic structures denoted by character ""a and the smaller magnetic structures beneath the large-scale ones denoted bv ""b in top panels of Figure 1).","of magnetic structures overlying the filament (the large-scale magnetic structures denoted by character “a"" and the smaller magnetic structures beneath the large-scale ones denoted by “b"" in top panels of Figure 1)."914 lere. we concentrate on the eruption process of (he large polar crown filament.," Here, we concentrate on the eruption process of the large polar crown filament."915 The filament eruption was recorded in all ALA channels. and the 301 images with a time cadence of 12 s were used.," The filament eruption was recorded in all AIA channels, and the 304 images with a time cadence of 12 s were used."916 The eruption was also observed as a limb event at 304 by the two EUVIs aboard STEREO., The eruption was also observed as a limb event at 304 by the two EUVIs aboard STEREO.917" The pixel size of the EUVI image on the solar disk was 1.6 aresec, and the time cadence for 304 images was 10 minutes."," The pixel size of the EUVI image on the solar disk was 1.6 arcsec, and the time cadence for 304 images was 10 minutes."918 For the first time. we combine data from SDO and STEREO to reconstruct the erupting filament.," For the first time, we combine data from SDO and STEREO to reconstruct the erupting filament."919 The angle of separation between STEREO A and D during our observations is around 1497.6., The angle of separation between STEREO A and B during our observations is around $\degr$ .6.920 It is diffieult to deline the same feature in EUVI A and D images because of the large separation angle., It is difficult to define the same feature in EUVI A and B images because of the large separation angle.921 The separation angle between STEREO A and SDO is 737.8. and that between STEREO D aud SDO 107.8.," The separation angle between STEREO A and SDO is $\degr$ .8, and that between STEREO B and SDO $\degr$ .8."922 So observations from the three different viewpoints are used here to derive the 3D erupting process of the polar crown filament., So observations from the three different viewpoints are used here to derive the 3D erupting process of the polar crown filament.923 The western part of the filament was observed clearly bv STEREO A and SDO. and thus we reconstruct ihe western part by using data from STEREO A and SDO (Figures 2 and 3).," The western part of the filament was observed clearly by STEREO A and SDO, and thus we reconstruct the western part by using data from STEREO A and SDO (Figures 2 and 3)."924 Similarly. the eastern part is reconstructed by using observations of STEREO D and SDO (Figures 2 and 1).," Similarly, the eastern part is reconstructed by using observations of STEREO B and SDO (Figures 2 and 4)."925 In order to reconstruct the 3D shape of the eruptive filament. we use a routine called MEASURE (developed by W. Thompson). which is available in the STEREO package ol the Solar Software library.," In order to reconstruct the 3D shape of the eruptive filament, we use a routine called $_{-}$ MEASURE (developed by W. Thompson), which is available in the STEREO package of the Solar Software library."926 The routine uses triangulation to determine the 3D coordinate of the Gepointed feature., The routine uses triangulation to determine the 3D coordinate of the tiepointed feature.927 It is a widget based application that allows the user to locate (and select with a cursor) the same feature in (wo images from different views., It is a widget based application that allows the user to locate (and select with a cursor) the same feature in two images from different views.928" This routine makes use of an approximate ""epipolar constraint in locating the same [feature in both images (Inhester 2006)."," This routine makes use of an approximate “epipolar constraint"" in locating the same feature in both images (Inhester 2006)."929 The two observer positions aud any object point to be reconstructed exaclly define a plane. which is known as an epipolar plane.," The two observer positions and any object point to be reconstructed exactly define a plane, which is known as an epipolar plane."930 By definiion. epipolar planes are projected on both observer;s images as lines and (hese lines are called the epipolar lines.," By definition, epipolar planes are projected on both observer¡¯s images as lines and these lines are called the epipolar lines."931 Once we identily a feature in one image. it is possible to determine the projection of the epipolar plane (ie.. epipolar line) passing the same feature in the second image.," Once we identify a feature in one image, it is possible to determine the projection of the epipolar plane (i.e., epipolar line) passing the same feature in the second image."932 Alter selecting the same feature in both images. the 3D coordinates are then determined as longitude. latitude. ancl radial distance from the center of the Sun.," After selecting the same feature in both images, the 3D coordinates are then determined as longitude, latitude, and radial distance from the center of the Sun."933 It must be mentioned that there exists a limiting [actor in reconstructing filaments.," It must be mentioned that there exists a limiting factor in reconstructing filaments,"934and oh.©(0.5.3).,"and $A_* \in (0.5,3)$."935" The dependence of the mass-loss rate on stellar mass and metallicity inferred by Nugis Lamers (2000) lead to AD~10.PCAZZAL.:0Y22AL.vr1 for WN stars and A10PCMΛΙ.HTY?ZAL.vr1 for WES, which may suggest that the tenuous winds required by the reverse-orward shock scenario for the afterglows 900123 anc 021211 arise Wt stars which are less massive and less meta rich than Galactic Ws (Wijers 2001. Chevalier 2004)."," The dependence of the mass-loss rate on stellar mass and metallicity inferred by Nugis Lamers (2000) lead to $\dot{M} \sim 10^{-6} (M/M_\odot)^{1.1} Y^{2.2}\, \Msunyear$ for WN stars and $\dot{M} \sim 10^{-5} (M/M_\odot)^{1.1} Y^2 Z\, \Msunyear$ for WCs, which may suggest that the tenuous winds required by the reverse-forward shock scenario for the afterglows 990123 and 021211 arise WR stars which are less massive and less metal rich than Galactic WRs (Wijers 2001, Chevalier 2004)."936 If such stars do not exist. then either the microphysica parameters must be cilferent behind the RS crossing the GRB ejecta and the FS sweeping-up the CDM or the fas declining carly optical emission. of the afterglows 990123 and 021211 is not arising in the RS.," If such stars do not exist, then either the microphysical parameters must be different behind the RS crossing the GRB ejecta and the FS sweeping-up the CBM or the fast declining early optical emission of the afterglows 990123 and 021211 is not arising in the RS."937 One possibility is tha the carly optical afterglow emission is produced in interna shocks occurring in an unsteady wind Rees 1999). a scenario which was not investigated in this work.," One possibility is that the early optical afterglow emission is produced in internal shocks occurring in an unsteady wind Rees 1999), a scenario which was not investigated in this work."938 , 939ALDultiple independent lines of evidence. Πιοπιαπας observations of the large-scale matter distribution (ce. 7). cosmic microwave backeround fluctuations (og.?).. and type Ia superuovae (27).. Sugeest that our uuiverse is dominated by two components: dark matter. which is probably a form of nounbarvonic matter. and dark energy. which is a nane for the presently unknown cause of the observed acceleration of expansion.,"Multiple independent lines of evidence, including observations of the large-scale matter distribution \citep[eg.][]{Percival}, cosmic microwave background fluctuations \citep[eg.][]{WMAP5}, and type Ia supernovae \citep{Perlmutter,Riess}, suggest that our universe is dominated by two components: dark matter, which is probably a form of nonbaryonic matter, and dark energy, which is a name for the presently unknown cause of the observed acceleration of expansion."940 However. we still lack an understanding of auv possible interactions between these two principal constituents (7)..," However, we still lack an understanding of any possible interactions between these two principal constituents \citep{Bean}. ."941 Tn a previous paper (7:hereafterSROS).. we examined the role that interacting dark matter and dark energy would play iu the development of one-dimensional Zeldovich pancakes. an important idealized case useful for uuderstaudiug structure formation.," In a previous paper \citep[][hereafter SR08]{Sutter}, we examined the role that interacting dark matter and dark energy would play in the development of one-dimensional Zel'dovich pancakes, an important idealized case useful for understanding structure formation."942 Tere we extend that preliminary work to a more realistic threc-dimenusional simulation of the erowth of dark matter halos iu an attempt to fiud wavs to distinguish these models from standard cosimolosv., Here we extend that preliminary work to a more realistic three-dimensional simulation of the growth of dark matter halos in an attempt to find ways to distinguish these models from standard cosmology.943 Iu this paper. we study the effects of a Yukawa interaction between a single family of nonrcelativistic dark matter (DM) particles aud a scalar field that is responsible for the dark cucrey (DE).," In this paper, we study the effects of a Yukawa interaction between a single family of nonrelativistic dark matter (DM) particles and a scalar field that is responsible for the dark energy (DE)."944 We follow closely 1ο forialigi developed bv ?.., We follow closely the formalism developed by \cite*{Farrar}.945 Such an iuteractiou is initially attractive because it is motivated by particle owesies (2) and unight provide a way to explain ιο apparent eniptiuess of the voids. as demonstrated miuerically by ?..," Such an interaction is initially attractive because it is motivated by particle physics \citep{Amendola2} and might provide a way to explain the apparent emptiness of the voids, as demonstrated numerically by \cite*{Nusser}."946 There has been considerable imterest recently iu studving the effects of these interactions on structure. th. using an analytic approach (7) and using direct siuulatious (eg.?7:audothers).," There has been considerable interest recently in studying the effects of these interactions on structure, both using an analytic approach \citep{Mainini} and using direct simulations \citep[eg.][and others]{Maccio, Manera}."947" Towever. the current iunuerical studies suffer frou, poor resolution. aud these results cannot reliably be compared to simulations of standard cosinological structure formation."," However, the current numerical studies suffer from poor resolution, and these results cannot reliably be compared to simulations of standard cosmological structure formation."948 In this paper. we use hieh spatial resolution aud careful analysis to accurately capture many dark matter halos for use in conarison.," In this paper, we use high spatial resolution and careful analysis to accurately capture many dark matter halos for use in comparison."949 The following is a brief sununiurw of the equations we solve and our nuuerical techuiques., The following is a brief summary of the equations we solve and our numerical techniques.950 Tn Section 2 we discuss modifications to the halo mass function., In Section \ref{sec:massFunction} we discuss modifications to the halo mass function.951 We use these mass functious to distinguish interacting DALDE from standard concordance cosmology using a 4? likehhood test., We use these mass functions to distinguish interacting DM-DE from standard concordance cosmology using a $\chi^2$ likelihood test.952" Additionally, we discuss the feasibility of using the Dark Encrey Survey (7) to detect this coupling within its mass aud redshift coutraiuts."," Additionally, we discuss the feasibility of using the Dark Energy Survey \citep{Annis} to detect this coupling within its mass and redshift contraints."953 Finally. we determine the exteut to which two specific models of interacting DALDE can be differentiated from cach other.," Finally, we determine the extent to which two specific models of interacting DM-DE can be differentiated from each other."954 Compared to simulations of the full non-linear theory. we found in SROs that the perturbation theory presentec by Farrar and Peebles is very accurate in determining the evolution of structure. and hence we will maintain the perturbative approach aud assume fuctuatious im the scalar field are small.," Compared to simulations of the full non-linear theory, we found in SR08 that the perturbation theory presented by Farrar and Peebles is very accurate in determining the evolution of structure, and hence we will maintain the perturbative approach and assume fluctuations in the scalar field are small."955" Uuder perturbation theory. the homogenous part of the dark energy scalar field. 05. evolves as where Q,,9 is the dark matter particle fraction of the critical deusitv aud Z/y is the IIubble coustaut."," Under perturbation theory, the homogenous part of the dark energy scalar field, $\phi_b$, evolves as where $\Omega_{m0}$ is the dark matter particle fraction of the critical density and $H_0$ is the Hubble constant."956 A subscript of 0 denotes the presceut-day value., A subscript of $0$ denotes the present-day value.957 The dark matter particle equation of notion is Tere ® is the normal comoving eravitational potential. ais the scale factor. v dis the comoving particle peculiar," The dark matter particle equation of motion is Here $\Phi$ is the normal comoving gravitational potential, $a$is the scale factor, $\bf{v}$ is the comoving particle peculiar"958quite modest value of 4 will produce a very small value of Lowμις.,quite modest value of $\eta$ will produce a very small value of $T_{\mathrm{CW}}/T_{\mathrm{max}}$.959 We have studied the propagation of the vector potential in the plasma of the early Universe. assuming a EI. space of negative curvature. and. conclude that £A75 will slowly acquire a non-thermal part. Απ. due to distant. matter.," We have studied the propagation of the vector potential in the plasma of the early Universe, assuming a FRW space of negative curvature, and conclude that $\langle {\mathbf A}^2 \rangle$ will slowly acquire a non-thermal part, $\langle {\mathbf A}_{\mathrm{nt}}^2 \rangle$, due to distant matter."960 Even though this is a direct. effect. of the vector potential. and cannot be attributed to fields. the theory is nonetheless eauge invariant.," Even though this is a direct effect of the vector potential, and cannot be attributed to fields, the theory is nonetheless gauge invariant."961 This elfect is therefore in the same category as the Aharonov-Bohm ellect (?2)..," This effect is therefore in the same category as the Aharonov-Bohm effect \citep{ahar,tono}."962 Further research will be needed to determine whether this increase of {Αι is sullicient to cause a CW transition and generate a mass scale., Further research will be needed to determine whether this increase of $\langle {\mathbf A}_{\mathrm{nt}}^2 \rangle$ is sufficient to cause a CW transition and generate a mass scale.963 several additional questions are raised by the present paper. among them the following: ‘These considerations are beyond the scope of this paper. which is solely concerned. with the interplay. in a simple model. of field theory (density matrix. the Lagrangian for a scalar field. the Coleman-Weinbere transition) and the classical equations of propagation of the ordinary vector potential in a FRAY space of negative curvature.," Several additional questions are raised by the present paper, among them the following: These considerations are beyond the scope of this paper, which is solely concerned with the interplay, in a simple model, of field theory (density matrix, the Lagrangian for a scalar field, the Coleman-Weinberg transition) and the classical equations of propagation of the ordinary vector potential in a FRW space of negative curvature."964 A well-known text (7). suggests a connection between large and small scales similar to the one explored here., A well-known text \citep{pesk1} suggests a connection between large and small scales similar to the one explored here.965 After surveving the cdillieulties faced by current theories of the mass scale. the authors write: 7...1t may be that the overall scale of cnerey-momentum is genuinely ambiguous and is set by a cosmological boundary condition.”," After surveying the difficulties faced by current theories of the mass scale, the authors write: it may be that the overall scale of energy-momentum is genuinely ambiguous and is set by a cosmological boundary condition.”"966 We have presented a mechanism for such a connection., We have presented a mechanism for such a connection.967 It is based on the fanuliar electromagnetic interaction. and nothing radically new seems to be required.," It is based on the familiar electromagnetic interaction, and nothing radically new seems to be required."968" We acknowledge helpful correspondence with Bryce DeWitt, Leonard. Parker. Stephen Fulling. Don Melrose... David Alontgomery and Philip Mannheim."," We acknowledge helpful correspondence with Bryce DeWitt, Leonard Parker, Stephen Fulling, Don Melrose, David Montgomery and Philip Mannheim."969 We also wish to thank the chairman and. faculty of the department of physics at Washington University for providing an ollice and computer support for a retired colleague., We also wish to thank the chairman and faculty of the department of physics at Washington University for providing an office and computer support for a retired colleague.970 Cosmic space may be infinite. but ollice space is at a premium.," Cosmic space may be infinite, but office space is at a premium."971 In this appendix we will work in the usual spherical polar coordinates. and. make connection with Riemannian coordinates when necessarv.," In this appendix we will work in the usual spherical polar coordinates, and make connection with Riemannian coordinates when necessary."972 Suppose we have a dipole at the origin. oscillating with time dependence exp( ief) in the z direction.," Suppose we have a dipole at the origin, oscillating with time dependence $\exp (-{\mathrm i} \omega t)$ in the $z$ direction."973 The surrounding medium is of uniform conductivity. 0. so that the current density. J. is given by J=6E.," The surrounding medium is of uniform conductivity, $\overline{\sigma}$ so that the current density, ${\mathbf j}$, is given by ${\mathbf j} = \overline{\sigma} {\mathbf E}$."974 We will analyze this system by imagining a small sphere of radius rj cut out of the medium surrounding the dipole., We will analyze this system by imagining a small sphere of radius $r_1$ cut out of the medium surrounding the dipole.975 Induced currents Ilowing in the medium will cause surface charges to appear on the sphere. and the total dipole moment will be the sum of the original dipole moment and that due to theinduced charges.," Induced currents flowing in the medium will cause surface charges to appear on the sphere, and the total dipole moment will be the sum of the original dipole moment and that due to theinduced charges."976 We assume the permittivity and magnetic susceptibility are essentially unity. so D=E and B=H.," We assume the permittivity and magnetic susceptibility are essentially unity, so ${\mathbf D} = {\mathbf E}$ and ${\mathbf B} = {\mathbf H}$."977 In such a svstem Vj=0 follows from Alaxwell’s equations. so there are no volume charges in the medium.," In such a system $\nabla \cdot {\mathbf j} = 0$ follows from Maxwell's equations, so there are no volume charges in the medium."978 Denote by Dune=D(io)exp(il) the dipole moment ab the center of the sphere. where (o) is the true dipole strength at angular frequeney c.," Denote by $D_{\mathrm{true}} = D (\omega)\exp(-{\mathrm i} \omega t)$ the dipole moment at the center of the sphere, where $D (\omega)$ is the true dipole strength at angular frequency $\omega$."979 The induced dipole moment due to the surface charges is Ding. so the total dipole moment is Dy=Di|Di.," The induced dipole moment due to the surface charges is $D_{\mathrm{ind}}$, so the total dipole moment is $D_{\mathrm{tot}} = D_{\mathrm{true}} + D_{\mathrm{ind}}$."980 Just outside the sphere the electrostatic potential is given by where Pi(cos9)=cos@., Just outside the sphere the electrostatic potential is given by where $P_1 (\cos \theta ) = \cos \theta$.981" The radial component of E is given by The surface charge density. s. obeys the relation where j, is evaluatedjust outside the sphere."," The radial component of ${\mathbf E}$ is given by The surface charge density,$s$ , obeys the relation where $j_r$ is evaluated just outside the sphere."982 This gives, This gives983"previous treatments of weak MHD turbulence (Galtieretal.2000, 2002), it was assumed that q-(kj,ki) are smooth functions of kj, that is, their dynamically important components with kj=0 have the same k,-scaling as the components with kj#0.","previous treatments of weak MHD turbulence \citep{galtier_nnp00,galtier_nnp02}, it was assumed that $q^\pm(k_\|, k_\perp)$ are smooth functions of $k_\|$, that is, their dynamically important components with $k_\| = 0$ have the same $k_\perp$ -scaling as the components with $k_\|\neq 0$."984 Such an assumption would be self-consistent if the residual energy were absent., Such an assumption would be self-consistent if the residual energy were absent.985" However, as we have demonstrated, the residual energy is spontaneously generated by interacting Alfvénn waves."," However, as we have demonstrated, the residual energy is spontaneously generated by interacting Alfvénn waves."986" One can argue that the presence of the residual energy modifies the spectra of the Elsasser with their smooth parts, energiesj(ky)ki? |, that,now togetheracquire their own singular partsA*(k)k;?, where A+(ky) are concentrated in region (11))."," One can argue that the presence of the residual energy modifies the spectra of the Elsasser energies that, together with their smooth parts, $q^{\pm}=f^{\pm}(k_\|)k_\perp^{-3}$ , now acquire their own singular parts, where $\Delta^{\pm}(k_\|)$ are concentrated in region \ref{wedge}) )."987" When multiplied by ó(kj) and integrated over kj, both the smooth and the singular parts provide comparable contributions to the integrals."," When multiplied by ${\tilde \delta}(k_\|)$ and integrated over $k_\|$, both the smooth and the singular parts provide comparable contributions to the integrals."988 We established that the spontaneously generated residual energy is always negative., We established that the spontaneously generated residual energy is always negative.989" Although this result is obtained in the framework of weak MHD turbulence, it provides the first analytic explanation for the observational and numerical findings that magnetic energy exceeds kinetic energy in the inertial interval of MHD turbulence, Müller&Grappin2005;Podestaet,al.2007;Boldyrev(e.g.,&Perez2009)."," Although this result is obtained in the framework of weak MHD turbulence, it provides the first analytic explanation for the observational and numerical findings that magnetic energy exceeds kinetic energy in the inertial interval of MHD turbulence, \citep[e.g.,][]{muller_g05,podesta_rg07,boldyrev_p09}."990. We derived that the residual energy has the field-perpendicular spectrum ekj! (12)), We derived that the residual energy has the field-perpendicular spectrum $E_r(k_\perp) \propto k_\perp^{-1}$ \ref{erperp}) ).991" This relatively shallow spectrumE,(ki) holds in the inertial interval and breaks down at sufficiently large k,; when the nonlinear broadening of the residual energy spectrum in the field-parallel direction (11)) becomes comparable to the width of the field-parallel energy spectra of the Alfvénn waves, q(kj)."," This relatively shallow spectrum holds in the inertial interval and breaks down at sufficiently large $k_\perp$ when the nonlinear broadening of the residual energy spectrum in the field-parallel direction \ref{wedge}) ) becomes comparable to the width of the field-parallel energy spectra of the Alfvénn waves, $q^{\pm}(k_\|)$."992" It is easy to see, however, that this is precisely the scale beyond which the weak interaction approximation breaks down, and the turbulence becomes strong."," It is easy to see, however, that this is precisely the scale beyond which the weak interaction approximation breaks down, and the turbulence becomes strong."993 In our future work we will extend our analysis of residual energy to the case of strong MHD turbulence., In our future work we will extend our analysis of residual energy to the case of strong MHD turbulence.994" This work was supported by the US DoE Awards DE-FG02-07ER54932, DE-SC0003888, DE-SC0001794, the NSF Grant PHY-0903872, the NSF/DOE Grant AGS-1003451, and in part by the NSF Grant No."," This work was supported by the US DoE Awards DE-FG02-07ER54932, DE-SC0003888, DE-SC0001794, the NSF Grant PHY-0903872, the NSF/DOE Grant AGS-1003451, and in part by the NSF Grant No."995" NSF PHY05-51164, and the NSF Center for Magnetic Self-organization in Laboratory and Astrophysical Plasmas at U. Wisconsin-Madison."," NSF PHY05-51164, and the NSF Center for Magnetic Self-organization in Laboratory and Astrophysical Plasmas at U. Wisconsin-Madison."996 High Performance Computing resources were providedby the Texas Advanced Computing Center (TACC) at the University of Texas at Austin under the NSF-Teragrid Project TG-, High Performance Computing resources were providedby the Texas Advanced Computing Center (TACC) at the University of Texas at Austin under the NSF-Teragrid Project TG-PHY080013N.997"Hence, the specific angular momentum of accreted material at the surface of the star, of radius R, is This specific angular momentum of accreted material would be equal to that at the equator of the gainer if it were critically rotating at brake up.","Hence, the specific angular momentum of accreted material at the surface of the star, of radius $R$, is This specific angular momentum of accreted material would be equal to that at the equator of the gainer if it were critically rotating at brake up."998 It is much larger than found in normal stars., It is much larger than found in normal stars.999" When accreting at a rate Mace the rate of angular momentum transferred from the disc to the star is Assuming a negligible change in stellar radius, we can determine the amount of mass AM that must be transferred through the disc to spin the star up to its critical angular velocity Q, from an initial Oy when it had mass Mo."," When accreting at a rate $\dot{M}_{\rm acc}$ the rate of angular momentum transferred from the disc to the star is Assuming a negligible change in stellar radius, we can determine the amount of mass $\Delta M$ that must be transferred through the disc to spin the star up to its critical angular velocity $\Omega_{\rm k}$ from an initial $\Omega_0$ when it had mass $M_0$."1000 Let the radius of gyration of the star be kR so that its total angular momentum is k?MR?Q when spinning rigidly at Q then A more precise formula was derived by (1981) who took account of the change in the mass of the star but this is unnecessary for our purposes because AM is always small., Let the radius of gyration of the star be $kR$ so that its total angular momentum is $k^2MR^2\Omega$ when spinning rigidly at $\Omega$ then A more precise formula was derived by \citet{packet1981} who took account of the change in the mass of the star but this is unnecessary for our purposes because $\Delta M$ is always small.1001 For main-sequence stars k?zz0.1 and varies little., For main-sequence stars $k^{2}\approx 0.1$ and varies little.1002 Thus when 0.1«Qo/Qy<0.4 we find 0.1>AM/My0.06.," Thus when $0.1 < \Omega_0/\Omega_{\rm k} < 0.4$ we find $0.1 > \Delta M/M_0 >10030.06$."1004 This is very small when we consider that all classical Algols have a mass ratio of q<0.7 (mostlyqz0.2accordingto which indicates that the losers in the classical Algols have transferred more on the order of 1Mo.," This is very small when we consider that all classical Algols have a mass ratio of $q < 0.7$ \citep[mostly $q \approx 0.2$ according1005 to][]{ibanoglu2006} which indicates that the losers in the classical Algols have transferred more on the order of $1\,\rm{M}_\odot$."1006 The shaded area in Fig., The shaded area in Fig.1007 |4| shows the amount of material that must be accreted from a disc to spin the star up to its critical rotational velocity., \ref{figacc} shows the amount of material that must be accreted from a disc to spin the star up to its critical rotational velocity.1008" Despite having high spin velocities, 0.1«2/Q,0.4, observations show that the detached components in most of the Algols do not actually attain their critical rotational velocity (Table m Fig."," Despite having high spin velocities, $0.1 <1009\Omega/\Omega_{\rm k} < 0.4$, observations show that the detached components in most of the Algols do not actually attain their critical rotational velocity (Table \ref{table}, Fig."1010 ha)., \ref{figprot}a a).1011" The only exception, with the high ratio of Q/Q,=0.72, is RZ Sct."," The only exception, with the high ratio of $\Omega/\Omega_{\rm k} = 0.72$, is RZ Sct."1012 This star’s radial velocity curve is distorted (e.g∙∙, This star's radial velocity curve is distorted \citep[e.g][]{mcnamara1957}.1013 It shows emission in Πα outside eclipse and its light curve displays distortions due to an accretion stream ∙∙," It shows emission in $H_\alpha$ outside eclipse \citep{mcnamara1957,hansen1959} and its light curve displays distortions due to an accretion stream \citep{olson1989}."1014 These observed phenomena may be taken as the signature of a rapid mass transfer phase., These observed phenomena may be taken as the signature of a rapid mass transfer phase.1015" In all cases a mechanism is needed to dissipate this excess angular momentum, along with associated energy."," In all cases a mechanism is needed to dissipate this excess angular momentum, along with associated energy."1016 Here we examine various mechanisms for angular momentum loss and compare with the observations., Here we examine various mechanisms for angular momentum loss and compare with the observations.1017Chakrabarty 2001).,Chakrabarty 2001).1018 The binary spends most of the time undergoing mass transfer to the outer disk. but little to no aceretion onto the NS.," The binary spends most of the time undergoing mass transfer to the outer disk, but little to no accretion onto the NS."1019 Bildsten Chakrabarty (2001) showed that the X-ray emission always detected from NS binaries in quiescence (see Bildsten Rutledge 2001 for an overview) heat the companion on the side facing the NS., Bildsten Chakrabarty (2001) showed that the X-ray emission always detected from NS binaries in quiescence (see Bildsten Rutledge 2001 for an overview) heat the companion on the side facing the NS.1020" Presuming the NS thermal emission is at the level predicted by Brown. Bildsten Rutledge (1998) (giving Ly4,7:5«107erg s!) then the WDs in aand wwill have Τηz:5000—6500K.! The scale height is H/R.~3<107+. yielding τι~10°yr."," Presuming the NS thermal emission is at the level predicted by Brown, Bildsten Rutledge (1998) (giving $L_{X,q}\approx 5\times102110^{32} \ {\rm erg \ s^{-1}}$ ) then the WDs in and will have $T_{\rm Eff}\approx 5000-6500 \ {\rm K}$ The scale height is $H/R_c\approx 3\times 10^{-4}$, yielding $\tau_g\approx 10^6 \ {\rm yr}$."1022 This is still much longer than the time for viscosity to finally play a role in moving material outwards in the outer accretion disk. most likely allowing for the angular momentum to get back to the donor and move it out. so that gj2—5/3.," This is still much longer than the time for viscosity to finally play a role in moving material outwards in the outer accretion disk, most likely allowing for the angular momentum to get back to the donor and move it out, so that $n_{\rm1023R}=-5/3$."1024 The adiabatic index is never less than —5/3., The adiabatic index is never less than $-5/3$.1025" However. the isothermal WD response is <—5/3 and mass transfer would be unstable for M,<4(9.5)ΙΟΛΜ. ina He WD of 7.24(10)ς10?K as long as internal heating can keep the WD tsothermal as mass is lost."," However, the isothermal WD response is $<-5/3$ and mass transfer would be unstable for $M_c<4(9.5)\times 10^{-3}M_\odot$ in a He WD of $T_c=4(10)\times102610^5 \ {\rm K}$ as long as internal heating can keep the WD isothermal as mass is lost."1027" These masses are near the values I noted earlier for the ""evaporation"". namely where the mass-radius relation turns over."," These masses are near the values I noted earlier for the “evaporation”, namely where the mass-radius relation turns over."1028 The timing of two accreting millisecond pulsars (Markwardt et al., The timing of two accreting millisecond pulsars (Markwardt et al.1029 2002: Galloway et al., 2002; Galloway et al.1030 2002) in ultracompact binaries has probed the WD donor properties to new levels and shown that they are of finite entropy., 2002) in ultracompact binaries has probed the WD donor properties to new levels and shown that they are of finite entropy.1031 This motivated my calculations of low-mass WDs of finite 7). that allow for 7. to be constrained., This motivated my calculations of low-mass WDs of finite $T_c$ that allow for $T_c$ to be constrained.1032 For He WDs. the implied Τι. are nearly that expected just from adiabatic expansion of the initially hot WD that filled the RL.," For He WDs, the implied $T_c$ are nearly that expected just from adiabatic expansion of the initially hot WD that filled the RL."1033 Only a small amount of tidal heating is needed., Only a small amount of tidal heating is needed.1034 More tidal heating is needed to make a C/O WD fill the RL., More tidal heating is needed to make a C/O WD fill the RL.1035 These finite 7 solutions allowed for a re-evaluation of Ruderman Shaham's (1983) scenario for making isolated millisecond radio pulsars via a mass transfer instability., These finite $T$ solutions allowed for a re-evaluation of Ruderman Shaham's (1983) scenario for making isolated millisecond radio pulsars via a mass transfer instability.1036 [ find that the adiabatic mass transfer instability can occur for a hot He WD as long as the angular momentum leaving the RL filling star is not returned to the orbit., I find that the adiabatic mass transfer instability can occur for a hot He WD as long as the angular momentum leaving the RL filling star is not returned to the orbit.1037 I have thus eliminated one criticism of their model. though the question of angular momentum elimination remains a serious one.," I have thus eliminated one criticism of their model, though the question of angular momentum elimination remains a serious one."1038 The physics of hot. low-mass WDs also yields a minimum mass WD solution for a fixed 7. so that a mass transfer instability can occur if the donor remains isothermal under mass loss.," The physics of hot, low-mass WDs also yields a minimum mass WD solution for a fixed $T_c$ so that a mass transfer instability can occur if the donor remains isothermal under mass loss."1039 I thus speculate that an evaporative or mass transfer instability endpoint might occur as long as tidal (or other) heating persists at 40-80 minute orbital periods., I thus speculate that an evaporative or mass transfer instability endpoint might occur as long as tidal (or other) heating persists at 40-80 minute orbital periods.1040 This work also impacts AM CVn binaries. where a low mass He star donates material to a more massive WD (see Solheim 1995 for a review).," This work also impacts AM CVn binaries, where a low mass He star donates material to a more massive WD (see Solheim 1995 for a review)."1041 The larger WD radii lead to more GW emission and a higher M. than expected for a given Psp., The larger WD radii lead to more GW emission and a higher $\dot M_c$ than expected for a given $P_{\rm orb}$.1042 Hence. models which track the WD entropy will fall between the degenerate and non-degenerate models in Nelemans et al. (," Hence, models which track the WD entropy will fall between the degenerate and non-degenerate models in Nelemans et al. ("10432001) and depend on both the age of the system when RL filling occurs (as this fixes the initial WD entropy) and any tidal heating that occurs during the mass transfer.,2001) and depend on both the age of the system when RL filling occurs (as this fixes the initial WD entropy) and any tidal heating that occurs during the mass transfer.1044 If either of the instabilities discussed above occur. then the endpoint of AM CVn's could well be a DB WD (e.g. Tutukov Yungelson 1996).," If either of the instabilities discussed above occur, then the endpoint of AM CVn's could well be a DB WD (e.g. Tutukov Yungelson 1996)."1045 I thank Deepto Chakrabarty for alerting me to the discovery of these transients and for many conversations., I thank Deepto Chakrabarty for alerting me to the discovery of these transients and for many conversations.1046 [ra Wasserman provided great physics insights during the progress of this work. which was supported by NASA via grant NAG 5-8658 and by the NSF under grants PHY99-07949 and ASTO1-9642.," Ira Wasserman provided great physics insights during the progress of this work, which was supported by NASA via grant NAG 5-8658 and by the NSF under grants PHY99-07949 and AST01-9642."1047 L. B. is a Cottrell Scholar of the Research Corporation., L. B. is a Cottrell Scholar of the Research Corporation.1048was either fixed to 90° (to consider the case of coplanar orbits) or randomly chosen using a uniform probability distribution.,was either fixed to $90^{\circ}$ (to consider the case of coplanar orbits) or randomly chosen using a uniform probability distribution.1049 For each orbit (characterized by P. e. q. co. i) we numerically solved the disequation (7) in function of f (the epoch of the planetary period determination).," For each orbit (characterized by $P$ , $e$ , $q$, $\omega$, $i$ ) we numerically solved the disequation (7) in function of $t_0$ (the epoch of the planetary period determination)."1050 We subdivided the period P in 10000 equal intervals of time and evaluated disequation (7) at the extremes these intervals., We subdivided the period $P$ in 10000 equal intervals of time and evaluated disequation (7) at the extremes these intervals.1051 Then we isolated the intervals in which disequation (7) changed sign. and using the secant method imposing a threshold for the convergence equal to O.1 sec we obtained the roots of the correspondent equation.," Then we isolated the intervals in which disequation (7) changed sign, and using the secant method imposing a threshold for the convergence equal to 0.1 sec we obtained the roots of the correspondent equation."1052 Then we determined the intervals Afy where disequation (7) was satisfied.," Then we determined the intervals $\Delta\,t_0$ where disequation (7) was satisfied."1053 Summing up together these intervals of time anc dividing by the period of the binary gave the probability to observe the requested transit timing variation for that fixec orbit over the given timescale (Af=f— either 5 yr or 10 yr in our simulations) assuming to determine the perioc of the transiting planet in correspondence of a random orbital phase of the binary.," Summing up together these intervals of time and dividing by the period of the binary gave the probability to observe the requested transit timing variation for that fixed orbit over the given timescale $\Delta\,t=t-t_0$ either 5 yr or 10 yr in our simulations) assuming to determine the period of the transiting planet in correspondence of a random orbital phase of the binary."1054" In such a way we assigned to each simulated orbit a transit timing detection probability (P,,,).", In such a way we assigned to each simulated orbit a transit timing detection probability $P_{det}$ ).1055 I Fig., In Fig.1056 | (lower figure). we show the graphical representatio of disequation (7) in function of fo. the epoch of the planetary period determination. for the case of the orbit considered 1 Fig.," \ref{fig:typical_binary_curve}1057 (lower figure), we show the graphical representation of disequation (7) in function of $t_0$, the epoch of the planetary period determination, for the case of the orbit considered in Fig."1058 1. (upper figure). assuming a transit timing thresholc Tr=50 sec. and a timescale of 10 years.," \ref{fig:typical_binary_curve} (upper figure), assuming a transit timing threshold $\tau=50$ sec, and a timescale of 10 years."1059 The horizontal black solid lines denote the time intervals where disequatio (7) is satisfied., The horizontal black solid lines denote the time intervals where disequation (7) is satisfied.1060 We performed 100 runs of 10000 simulations each. calculating the mean detection probabilities and their c uncertainties. as reported in the next Section.," We performed 100 runs of 10000 simulations each, calculating the mean detection probabilities and their $\sigma$ uncertainties, as reported in the next Section."1061 The final detection probability histograms are shown in Fig. 2..," The final detection probability histograms are shown in Fig. \ref{fig:fig_hist},"1062 obtained from the entire sample of 10° simulated orbits., obtained from the entire sample of $10^6$ simulated orbits.1063 While most orbits imply null detection probabilities over the assumed timescales. in each one of the different situations we considered. the histograms present a probability tail extended toward large detection probabilities.," While most orbits imply null detection probabilities over the assumed timescales, in each one of the different situations we considered, the histograms present a probability tail extended toward large detection probabilities."1064" The number of orbits having Py,>99% 1s equal to 3.9%+0.2% after 5 years since the period determination (10.4%+0.3% after 10 years) for the case of coplanar orbits. and is equal to 2.2%+0.1% after 5 years (6.2%+0.2% after 10 years) for the case of random inclinations."," The number of orbits having $P_{det}>99\%$ is equal to $3.9\%\pm0.2\%$ after 5 years since the period determination $10.4\%\pm0.3\%$ after 10 years) for the case of coplanar orbits, and is equal to $2.2\%\pm0.1\%$ after 5 years $6.2\%\pm0.2\%$ after 10 years) for the case of random inclinations."1065 The fact that the histograms are extended toward large probabilities is a consequence of the period distribution and the adopted timescales., The fact that the histograms are extended toward large probabilities is a consequence of the period distribution and the adopted timescales.1066" Orbits having Py,>99% have also periods smaller than P<1.67-10° (considering a timescale of 10 years). which means that transit timing can allow discovery of stellar companions up to separations equal to &~75 AU after 10 years since the discovery of the planet (a.~36 AU after 5 years)."," Orbits having $P_{det}>99\%$ have also periods smaller than $P<1.67\,\cdot\,10^5$ (considering a timescale of 10 years), which means that transit timing can allow discovery of stellar companions up to separations equal to $a\sim75$ AU after 10 years since the discovery of the planet $a\sim36$ AU after 5 years)."1067" Then. considering the observed frequency of binaries in the solar surrounding with periods P>5-10° days (27%) and the case of coplanar orbits. after 5 years since the discovery of a sample of transiting planets 1.0%£0.2% transiting planet host-stars will have a probability Pj,>99% to present detectable (>50 sec) transit timing variations induced by stellar binarity. and 2.8%40.3% after 10 years."," Then, considering the observed frequency of binaries in the solar surrounding with periods $P>5\,\cdot\,10^3$ days $\%$ ) and the case of coplanar orbits, after 5 years since the discovery of a sample of transiting planets $1.0\%\pm0.2\%$ transiting planet host-stars will have a probability $P_{det}>99\%$ to present $detectable$ $>50$ sec) transit timing variations induced by stellar binarity, and $2.8\%\pm0.3\%$ after 10 years."1068 Considering the case of random inclinations the expected frequencies (1) are 0.6%40.1% and 1.7%+0.2% after 5 and 10 years respectively., Considering the case of random inclinations the expected frequencies $f_{det}$ ) are $0.6\%\pm0.1\%$ and $1.7\%\pm0.2\%$ after 5 and 10 years respectively.1069 These results üre summarized in Table 1.., These results are summarized in Table \ref{tab:prob_p1_p2}.1070 Our estimates can be considered a conservative lower limit. since we have excluded binaries with periodsP«5-10? days.," Our estimates can be considered a conservative lower limit, since we have excluded binaries with periods $P<5\,\cdot\,10^3$ days."1071 In this paper we have investigated how known transiting extrasolar planets can be used to constrain the frequency of multiple stellar systems among planet-host stars., In this paper we have investigated how known transiting extrasolar planets can be used to constrain the frequency of multiple stellar systems among planet-host stars.1072 The presence of a stellar companion in these systems is expected to induce transit timing variations of the transiting planets even once perturbing effects are neglected. due to the orbital revolution of the primary around the barycenter of the binary stellar system.," The presence of a stellar companion in these systems is expected to induce transit timing variations of the transiting planets even once perturbing effects are neglected, due to the orbital revolution of the primary around the barycenter of the binary stellar system."1073 If the frequency of binaries among planet-host stars is the same as determined in the solar neighborhood. after 5 years since the discovery of a sample of transiting planets 1.0%+0.2% of them have a probability >99% to present a transit timing variations >50 see induced by stellar binarity. and 2.8+0.3% after 10 years 1f the planetary and binary orbits are coplanar.," If the frequency of binaries among planet-host stars is the same as determined in the solar neighborhood, after 5 years since the discovery of a sample of transiting planets $1.0\%\pm0.2\%$ of them have a probability $>99\%$ to present a transit timing variations $>50$ sec induced by stellar binarity, and $2.8\pm0.3\%$ after 10 years if the planetary and binary orbits are coplanar."1074 Considering the case of random inclinations theprobabilities are0.6%+0.1% and 1.7%+0.2% after 5 and 10 years respectively., Considering the case of random inclinations theprobabilities are$0.6\%\pm0.1\%$ and $1.7\%\pm0.2\%$ after 5 and 10 years respectively.1075 Our results have been obtaied assuming a binary period P>5-10° (a>6 AU).," Our results have been obtained assuming a binary period $P>5\,\cdot\,10^3$ $a\gtrsim6$ AU)."1076 Moreover. we derived that we can expect to discover stellar companions of transiting planets host stars up to a maximum separations à.~75 AU," Moreover, we derived that we can expect to discover stellar companions of transiting planets host stars up to a maximum separations $a\sim75$ AU"1077another intermediate mass planet (Msini=31.2Mg).,"another intermediate mass planet $\msini=31.2\, \mearth$ )."1078 Its semi-major axis is 0.25 AU., Its semi-major axis is 0.25 AU.1079" With intermediate mass planets we here refer to planets with a mass larger than Neptune, but less than Saturn."," With intermediate mass planets we here refer to planets with a mass larger than Neptune, but less than Saturn."1080" They therefore fall in the significant mass gap of a factor 5 between Uranus (17.2 Mg) and Saturn (95.2Mg) which corresponds to an important difference in internal composition, too: Uranus and Neptune are ice giants which consist mainly of a heavy elements (iron, rocks and ices) and only about in mass of hydrogen and helium (e.g. Figueira et al. [2009))."," They therefore fall in the significant mass gap of a factor 5 between Uranus $17.2\,\mearth$ ) and Saturn $95.2\,\mearth$ ) which corresponds to an important difference in internal composition, too: Uranus and Neptune are ice giants which consist mainly of a heavy elements (iron, rocks and ices) and only about in mass of hydrogen and helium (e.g. Figueira et al. \cite{figueiraetal2009}) )."1081 Saturn and Jupiter are in contrast gas giants with only a small fraction of heavy elements., Saturn and Jupiter are in contrast gas giants with only a small fraction of heavy elements.1082" Instead, their massive gas envelopes account for most of the mass, namely for about 67 to for Saturn, and 87 to for Jupiter (Guillot[1999))."," Instead, their massive gas envelopes account for most of the mass, namely for about 67 to for Saturn, and 87 to for Jupiter (Guillot \cite{guillot1999}) )."1083" This difference in mass and composition is in turn understood in theoretical formation models based on the core accretion paradigm as the consequence of an important difference in the formation history of these two groups of planets: Jupiter and Saturn are thought to have undergone a phase of rapid ""runaway"" gas accretion, during which they accreted their large gaseous envelopes on a relatively short timescale, whereas the ice giants never underwent this phase (Pollack et al. [1996))."," This difference in mass and composition is in turn understood in theoretical formation models based on the core accretion paradigm as the consequence of an important difference in the formation history of these two groups of planets: Jupiter and Saturn are thought to have undergone a phase of rapid “runaway” gas accretion, during which they accreted their large gaseous envelopes on a relatively short timescale, whereas the ice giants never underwent this phase (Pollack et al. \cite{pollacketal1996}) )."1084" If the transformation phase from a Neptunian mass to a Jovian mass is very short compared to the lifetime of the protoplanetary disk, i.e. if the gas accretion rate onto the planet is very high, then it is very unlikely that the protoplanetary disk disappears just during the phase, which would obviously terminate the gas accretion and leave the planet at some intermediate mass."," If the transformation phase from a Neptunian mass to a Jovian mass is very short compared to the lifetime of the protoplanetary disk, i.e. if the gas accretion rate onto the planet is very high, then it is very unlikely that the protoplanetary disk disappears just during the phase, which would obviously terminate the gas accretion and leave the planet at some intermediate mass."1085" We thus see that in this scenario, the absence of a planet with a mass between Uranus and Saturn is not a surprise, but rather an expected outcome (Ida Lin 2004))."," We thus see that in this scenario, the absence of a planet with a mass between Uranus and Saturn is not a surprise, but rather an expected outcome (Ida Lin \cite{idalin2004}) )."1086" One can therefore say that the higher (smaller) the gas accretion rate in runaway, the smaller (higher) the expected relative frequency of intermediate mass planets compared to giant planets."," One can therefore say that the higher (smaller) the gas accretion rate in runaway, the smaller (higher) the expected relative frequency of intermediate mass planets compared to giant planets."1087 This correlation will manifest itself in the Observed detection rate of intermediate mass extrasolar planets., This correlation will manifest itself in the observed detection rate of intermediate mass extrasolar planets.1088 Indeed point detections made over the last years in particular with to a minimum of the observed planetary mass function at about 30-40 Μο (e.g. Bouchy et al. [2009))., Indeed point detections made over the last years in particular with to a minimum of the observed planetary mass function at about 30-40 $\mearth$ (e.g. Bouchy et al. \cite{bouchyetal2009}) ).1089" Such a minimum (if further confirmed) would be a very important observational finding, as it cannot be explained by an observational bias."," Such a minimum (if further confirmed) would be a very important observational finding, as it cannot be explained by an observational bias."1090 This is because the lower mass Neptunian and Super-Earth planets are more difficult to detect than intermediate mass planets., This is because the lower mass Neptunian and Super-Earth planets are more difficult to detect than intermediate mass planets.1091 It is interesting to note that a relative paucity of intermediate mass planets seems to exist also in transit searches (Hartman et al. 2009))., It is interesting to note that a relative paucity of intermediate mass planets seems to exist also in transit searches (Hartman et al. \cite{hartmanetal2009}) ).1092 Several physical effects influence the runaway gas accretion rate (Lubow et al. [[999};, Several physical effects influence the runaway gas accretion rate (Lubow et al. \cite{lubowetal1999};1093 Ida Lin 20045; Mordasini et al. 2009a:;, Ida Lin \cite{idalin2004}; Mordasini et al. \cite{mordasinietal2009a};1094 Lissauer et al. [2009)):, Lissauer et al. \cite{lissaueretal2009}) ):1095" Thermal pressure in the planetary envelope at lower planetary masses, the global evolution (dissipation) of the protoplanetary disk which reduces the amount of matter available for the planet, the rate of mass transport due to viscosity within the disk towards the planet as well as local phenomena like the formation of a gap."," Thermal pressure in the planetary envelope at lower planetary masses, the global evolution (dissipation) of the protoplanetary disk which reduces the amount of matter available for the planet, the rate of mass transport due to viscosity within the disk towards the planet as well as local phenomena like the formation of a gap."1096" While several planetary population synthesis simulations built on the core accretion paradigm share the common feature that they predict some depletion of intermediate mass planets relative to other planet types, they attribute in their underlying formation models (cf."," While several planetary population synthesis simulations built on the core accretion paradigm share the common feature that they predict some depletion of intermediate mass planets relative to other planet types, they attribute in their underlying formation models (cf."1097 Alibert et al. [2005)), Alibert et al. \cite{alibertetal2005}) )1098" distinct importance to the various effects mentioned above, so that there are clear differences in the degree of depletion: from an almost complete absence of such planets between 0.1-1 AU (Miguel Brunini [2009)), over a significant depletion (Ida Lin [2008b)) to a rather moderate one (about a factor 2-3 relative to Jovian planets) in Mordasini et al. [2009b))."," distinct importance to the various effects mentioned above, so that there are clear differences in the degree of depletion: from an almost complete absence of such planets between 0.1-1 AU (Miguel Brunini \cite{miguelbrunini2008,miguelbrunini2009}) ), over a significant depletion (Ida Lin \cite{idalin2004,idalin2008a,idalin2008b}) ) to a rather moderate one (about a factor 2-3 relative to Jovian planets) in Mordasini et al. \cite{mordasinietal2009a,mordasinietal2009b}) )."1099 The reason for the difference between the latter two models was discussed in details in Mordasini et al. (2009a))., The reason for the difference between the latter two models was discussed in details in Mordasini et al. \cite{mordasinietal2009a}) ).1100 It is related to different assumptions about the amount of gas present in the vicinity of the planet as discussed next., It is related to different assumptions about the amount of gas present in the vicinity of the planet as discussed next.1101 Here we illustrate this in Fig., Here we illustrate this in Fig.1102 [5] that shows the theoretically obtained mass distribution from Neptunian to Jovian planets using two different assumptions for the gas accretion rate in runaway., \ref{fig:msiniinter} that shows the theoretically obtained mass distribution from Neptunian to Jovian planets using two different assumptions for the gas accretion rate in runaway.1103 Other settings are similar as in Mordasini et al. (2009a))., Other settings are similar as in Mordasini et al. \cite{mordasinietal2009a}) ).1104" Only planets with a period less than 5 years are shown which are detectable for a RV instrument of | m/s precision, similar asHARPS."," Only planets with a period less than 5 years are shown which are detectable for a RV instrument of 1 m/s precision, similar as."1105" In one simulation (solid line), the gas accretion rate of the planet is limited by the accretion rate in the disk only if the mass of the planet is larger than the local gas isolation mass, calculated with the undisturbed gas surface density."," In one simulation (solid line), the gas accretion rate of the planet is limited by the accretion rate in the disk only if the mass of the planet is larger than the local gas isolation mass, calculated with the undisturbed gas surface density."1106 This is similar to the criterion used by Ida Lin (2004))., This is similar to the criterion used by Ida Lin \cite{idalin2004}) ).1107" It means that for masses between roughly 30 up to ~100Mg i.e. before the gas isolation mass is reached, high gas accretion rates of up to roughly 107? Mg/yr can occur, so that the transformation from a Neptunian to a Jovian planet only takes of order 104 years, much less than typical disk lifetimes."," It means that for masses between roughly 30 up to $\sim100\, \mearth$ i.e. before the gas isolation mass is reached, high gas accretion rates of up to roughly $10^{-2}\,\mearth$ /yr can occur, so that the transformation from a Neptunian to a Jovian planet only takes of order $10^{4}$ years, much less than typical disk lifetimes."1108 The underlying assumption is here that gas already inside the planet’s Hill sphere can be accreted independently of the inflow from further away in the disk., The underlying assumption is here that gas already inside the planet's Hill sphere can be accreted independently of the inflow from further away in the disk.1109 This inflow is in turn limited by the disk viscosity., This inflow is in turn limited by the disk viscosity.1110" In the other simulation (dotted line), the planetary gas accretion rate is limited by the accretion rate in the disk as soon as runaway starts."," In the other simulation (dotted line), the planetary gas accretion rate is limited by the accretion rate in the disk as soon as runaway starts."1111" As cores reach a mass large enough to trigger gas runaway accretion at a moment typically not much before the disk goes away (it is recalled that around about of FGK stars, there are no giant planets), the accretion rate in the disk has usually already fallen to quite low values at this moment, of order a few 107 to 1077 Maf/yr."," As cores reach a mass large enough to trigger gas runaway accretion at a moment typically not much before the disk goes away (it is recalled that around about of FGK stars, there are no giant planets), the accretion rate in the disk has usually already fallen to quite low values at this moment, of order a few $10^{-4}$ to $10^{-3}\,\mearth$ /yr."1112" This means that the transformation from a Neptunian to a Jovian planet takes now several 10? years, which is no more much shorter than the remaining disk life time."," This means that the transformation from a Neptunian to a Jovian planet takes now several $10^{5}$ years, which is no more much shorter than the remaining disk life time."1113" The underlying assumption is here that due to gap formation, the mass directly available to the planet is in fact small, in particular smaller than the gas isolation mass as calculated above, as (beginning) gap formation reduces the gas surface density around the planet."," The underlying assumption is here that due to gap formation, the mass directly available to the planet is in fact small, in particular smaller than the gas isolation mass as calculated above, as (beginning) gap formation reduces the gas surface density around the planet."1114 This corresponds to the setting in Mordasini et al. (2009a))., This corresponds to the setting in Mordasini et al. \cite{mordasinietal2009a}) ).1115darkeniig coellicients [for cool stars obtainecl from fits to model stelar atinospheres of late-type stars (Allard ischildt 1995: the fits to the moclel attnospleres are presente iu Froniug 1999b).,darkening coefficients for cool stars obtained from fits to model stellar atmospheres of late-type stars (Allard Hauschildt 1995; the fits to the model atmospheres are presented in Froning 1999b).1116 We fit three sets of models to the &0620—00 light curves., We fit three sets of models to the A0620–00 light curves.1117 In the first set we assunniec that ouly the hsar coutributes to the observed flux., In the first set we assumed that only the K star contributes to the observed flux.1118 For the second set we acded an accretion disk with a bright spot on its edge., For the second set we added an accretion disk with a bright spot on its edge.1119 The third set was a series of moclels to estimate tlie effect of a uou-varyiug source of dilutiug Παν ou the derived orbital inclinatious., The third set was a series of models to estimate the effect of a non-varying source of diluting flux on the derived orbital inclinations.1120 We mocleled the H light curves ouly: the J aud Ix data are of lower quality ancl do not warrant detailed modeling., We modeled the H light curves only; the J and K data are of lower quality and do not warrant detailed modeling.1121 We initially assumed that the Ix star is the ouy source of the observed flux., We initially assumed that the K star is the only source of the observed flux.1122 We assumed that the asyiumetry in the observed light curve is caused by extra flux added to he peak at © = 0.75 auc. to avoid this extra flux. we fit the moclels to he ligit curves oulv over phases from © = 0 to 0.51.," We assumed that the asymmetry in the observed light curve is caused by extra flux added to the peak at $\phi$ = 0.75 and, to avoid this extra flux, we fit the models to the light curves only over phases from $\phi$ = 0 to 0.51."1123 We generated mocel ligit curves of the Ix sta yovaryiug each of the parameters that affect the shape of the ellipsoidal modulation iu turn. aud fi the Iielit curves to the observed data by least squares.," We generated model light curves of the K star, varying each of the parameters that affect the shape of the ellipsoidal modulation in turn, and fit the light curves to the observed data by least squares."1124 The parameters that allect this mocel :we: tdje. orbital inclination. 7. the mass ratio. q. the temperature of the Ix star at the pole. 75. he quadratic Liub-darkeniug coefficieuts aud the gravity darkening coellicient. 9.," The parameters that affect this model are: the orbital inclination, $i$, the mass ratio, $q$, the temperature of the K star at the pole, $T_{2}$, the quadratic limb-darkening coefficients and the gravity darkening coefficient, $\beta$ ."1125 We varied the iicliation from i= 1° — SO* in d steps ancl the Ix star temperature from 75 = 1000 — 1500 Ix in τιnents o- 100 Ix. The mass ratio in 0620-00 is known to be gy = 0.067 (Marsh.Robinson&Woe1991)., We varied the inclination from $i$ = $\arcdeg$ – $\arcdeg$ in $\arcdeg$ steps and the K star temperature from $T_{2}$ = 4000 – 4500 K in increments of 100 K. The mass ratio in A0620–00 is known to be $q$ = 0.067 \citep{marsh1994}.1126. We calculated 110dels [for this mass ratio but also for q = 0.056. 0.083 and 0.10 to chee stle depeidence of our resuls On mass ratio.," We calculated models for this mass ratio but also for $q$ = 0.056, 0.083 and 0.10 to check the dependence of our results on mass ratio."1127 The limb darkering was mocleled using quadratic μι ib¢larkenii© coeíficieuts oXained from the model stellaratinospleres ciscussed above., The limb darkening was modeled using quadratic limb darkening coefficients obtained from the model stellaratmospheres discussed above.1128 We calcilated noctels [or gravity darseiiug coellictents of 0.05 and 0.08 (Sarua1980) axl we assumed that the Ix star {ills its Roche lobe., We calculated models for gravity darkening coefficients of 0.05 and 0.08 \citep{sarna1989} and we assumed that the K star fills its Roche lobe.1129 yieure 3. shows the best-fit mocles for each of the H-baud liguo curves., Figure \ref{fig_mod1} shows the best-fit models for each of the H-band light curves.1130 The parameters [or the mocels are givel in Table L. , The parameters for the models are given in Table \ref{tab_fits}. .1131For he 1996 December light curve. the best-fit model for gy = 0.067 las ali inclination of 7= LL:for he 1996 January light curve. 7=38° gives the best fit: and for the 1995 December lel curve. models with [3 — 15° give σος fits.," For the 1996 December light curve, the best-fit model for $q$ = 0.067 has an inclination of $i$ = $\arcdeg$ ;for the 1996 January light curve, $i = 38\arcdeg$ gives the best fit; and for the 1995 December light curve, models with $\arcdeg$ – $\arcdeg$ give good fits."1132 The models are largely llisensitve to variations 1n t1e Lass ralio and are ouly weakly clepeucent ou the temperature of the ]x star aud the value of the eravity darkening coefficient., The models are largely insensitive to variations in the mass ratio and are only weakly dependent on the temperature of the K star and the value of the gravity darkening coefficient.1133 The best fis were typically obtained lor Ts = IUJO I aud ο = 0.08., The best fits were typically obtained for $T_{2}$ = 4100 K and $\beta$ = 0.08.1134 lore important. Changes in these parameers had virtually no ellect ou the value of the best-fit inclination.," More important, changes in these parameters had virtually no effect on the value of the best-fit inclination."1135 For the 1996 January light curve. [or example. the reduced 4? of the best fits for each combination of the other parameters tested ange [ror χο LOG — L8. but the inclinations for these moclels vary only from 7 = 38° — 11°.," For the 1996 January light curve, for example, the reduced $\chi^2$ of the best fits for each combination of the other parameters tested range from $\chi^{2}_{\nu}$ = 1.06 – 1.58, but the inclinations for these models vary only from $i$ = $\arcdeg$ – $\arcdeg$."1136 Similarly. tle models calculated for the 1996 December light curve |ad 47 between2.62 and D+)3.01 and inclinations from =[5 — 207. while for the 1995 December light curve. VL = 113-118 and? = 13° —- 15”.," Similarly, the models calculated for the 1996 December light curve had $\chi^{2}_{\nu}$ between2.62 and 3.01 and inclinations from $i$ =$\arcdeg$ – $\arcdeg$, while for the 1995 December light curve, $\chi^{2}_{\nu}$ = 1.13 – 1.18 and $i$ = $\arcdeg$ – $\arcdeg$ ."1137"This simple addition procedure is valid because the foreground cloud is transparent and the outflow is seen at an angle, so that neither component blocks the view of the dense core.","This simple addition procedure is valid because the foreground cloud is transparent and the outflow is seen at an angle, so that neither component blocks the view of the dense core."1138" The modeled optical depths are 2.8 for the lline and 0.77 for?CO,, and the derived abundances are uncertain to ~35% which is the quadratic sum of from the line intensity and from the core mass estimate."," The modeled optical depths are 2.8 for the line and 0.77 for, and the derived abundances are uncertain to $\approx$ which is the quadratic sum of from the line intensity and from the core mass estimate."1139" The model underproduces the observed line strengths at large position offsets, which suggests that the density profile flattens out at large radii, as also indicated by the observed continuum brightness at the offset positions."," The model underproduces the observed line strengths at large position offsets, which suggests that the density profile flattens out at large radii, as also indicated by the observed continuum brightness at the offset positions."1140 Deriving the aand aabundances in the outflow and the foreground using radiative transfer models is not possible because the masses and ccolumn densities of these components are unknown., Deriving the and abundances in the outflow and the foreground using radiative transfer models is not possible because the masses and column densities of these components are unknown.1141" Instead we have used RADEX (?) to estimate their aand ccolumn densities,N(H20)/N(H»)."," Instead we have used RADEX \citep{radex} to estimate their and column densities,."1142". For the foreground cloud, we adopt Tkin= 10KK and n(H2)= 104ccm,, and for the outflow, we adopt Τις KK and n(H2)=ccm,, each."," For the foreground cloud, we adopt = K and = $^4$, and for the outflow, we adopt = K and =, ."1143". To estimate N(1?CO)) for the foreground cloud, we use the J=1-0 observations by ?.."," To estimate $N$ ) for the foreground cloud, we use the $J$ =1-0 observations by \citet{jakob}."1144 Table 2 summarizes our derived column densities and abundances of aand iin the various physical components of the DR21 region., Table \ref{t:abs} summarizes our derived column densities and abundances of and in the various physical components of the DR21 region.1145" At the low temperatures and high densities in the core, most iis likely frozen on grains, and the observed line may arise in a small region with a high aabundance."," At the low temperatures and high densities in the core, most is likely frozen on grains, and the observed line may arise in a small region with a high abundance."1146" The derived aabundance for the core is —4x lower than expected for the above values of the CO isotopic ratios and abundance, which suggests that even some CO is frozen out in the outer parts of the core."," The derived abundance for the core is $\sim$ $\times$ lower than expected for the above values of the CO isotopic ratios and abundance, which suggests that even some CO is frozen out in the outer parts of the core."1147" The density of the foreground cloud is too low for significant freeze-out, but with Ay «1.2 mmag, photodissociation is rapid for bbut not for3CO."," The density of the foreground cloud is too low for significant freeze-out, but with $A_V$ $\approx$ mag, photodissociation is rapid for but not for."1148". The high aabundance for the outflow is likely related to its temperature of ~200 KK, which is high enough to have rreleased from thedust grains by thermal evaporation, or possibly by shocks (?).."," The high abundance for the outflow is likely related to its temperature of $\sim$ K, which is high enough to have released from thedust grains by thermal evaporation, or possibly by shocks \citep{melnick}."1149Balsara switch (?)..,Balsara switch \citep{balsaraphd89}.1150 We used. variable timesteps controlled bv the Courant time with a Courant [actor of 0.4., We used variable timesteps controlled by the Courant time with a Courant factor of $0.4$.1151 The improved stability and error properties of nunotivate a full. re-examination of the standard. SPILL artificial viscosity., The improved stability and error properties of motivate a full re-examination of the standard SPH artificial viscosity.1152 This is bevond the scope of this present work., This is beyond the scope of this present work.1153 However. we note that the improved stability in nuneans that. particles better follow. characteristics of the How. while the gradients in the Balsara switeh will be less noisy.," However, we note that the improved stability in means that particles better follow characteristics of the flow, while the gradients in the Balsara switch will be less noisy."1154 Both of these effects. should. act to decrease the viscosity. in regions of steady Blow. (, Both of these effects should act to decrease the viscosity in regions of steady flow. (1155Note that all numerical schemes carry numerical viscosity. whether it is manifested through limited resolution or artificial shock-capturing viscosity.,"Note that all numerical schemes carry numerical viscosity, whether it is manifested through limited resolution or artificial shock-capturing viscosity."1156 Indeed. these viscous terms are vital for successfully modelling shocks.)," Indeed, these viscous terms are vital for successfully modelling shocks.)"1157 ln relsce:kh.. relsce:sock and relsce:blob.. we show that our rresults agree very well with analytic expectations. and with," In \\ref{sec:kh}, \\ref{sec:sod} and \\ref{sec:blob}, we show that our results agree very well with analytic expectations, and with"11581987).,.1159. During the course of such a collapse. the j)»roto-neuron star slainss [rom a radius of ~1Q0 kin to 10 kin.," During the course of such a collapse, the proto-neutron star shrinks from a radius of $\sim 100$ km to $\sim 10$ km."1160 We empinslze tha| the assumpion of the foregoing cliscussion that the kick is applied at a single locatio Lis simplislic (see. e.g.. 5pruit&Phinney(1998))): ia a real supernova. both linear aud augular nomentuli1 will be aceruulated by the proto-jieutron star throughout its formation at many clifle'ent locations.," We emphasize that the assumption of the foregoing discussion that the kick is applied at a single location is simplistic (see, e.g., \citet{Spruit1998}) ); in a real supernova, both linear and angular momentum will be accumulated by the proto-neutron star throughout its formation at many different locations."1161" Iu this more ""Ceilistie context. the constraiuts above ou offset distauces should be iuerpretect iusteacl as costraluts oh the offset scale at which the of the linear aud augular 101nentUL1 js accumulaed."," In this more realistic context, the constraints above on offset distances should be interpreted instead as constraints on the offset scale at which the of the linear and angular momentum is accumulated."1162 In)re precise coustraints will require detailed modeling of he hydrodyaamüc process of nonetUl ac‘Ctunulation int 109 slyernova 11at formed PSR , More precise constraints will require detailed modeling of the hydrodynamic process of momentum accumulation in the supernova that formed PSR J0737-3039B.1163Neverleless. it is Interesting ilabd he location of kicks inferred frorsuch a siiple mocle is COLsIStebL with kick «rigius in the UL sof he shrinkine proto-neutrou star duriug the supernova (see Figure 3 .," Nevertheless, it is interesting that the location of kicks inferred from such a simple model is consistent with kick origins in the bulk of the shrinking proto-neutron star during the supernova (see Figure \ref{fig:offset-pdf}) )."1164" Some recent SN moceliD>oO suggests that the processes that produce the kick ane hose tha inm»art rotation to tle ""eslin& neutron star produce incepeucent kicks aud splus. aud thereore there is little corelatidL )elW'eeu the kick maguLttde aid direction and t1e rotation imipartec (lothe post-SN compact oject (Wongwatliauarateta.2010:RauUsioual.2011)."," Some recent SN modeling suggests that the processes that produce the kick and those that impart rotation to the resulting neutron star produce independent kicks and spins, and therefore there is little correlation between the kick magnitude and direction and the rotation imparted to the post-SN compact object \citep{Wongwathanarat2010,Rantsiou2011}."1165. B his case he ollset leneth scae infered. above [rom he dyuauiical coustrainuts on the kick wouk 100 be relevan., In this case the offset length scale inferred above from the dynamical constraints on the kick would not be relevant.1166 We conclude tlal pulsar B's spin is actually linked to the torque inclucec w the plvsical mechanism: prexluciie the kick it must be offset from the center of mass of the collapsing lelrol star progenitor., We conclude that pulsar B's spin is actually linked to the torque induced by the physical mechanism producing the kick it must be offset from the center of mass of the collapsing neutron star progenitor.1167 Rega‘less of the specifics c ofthe collaose process. Lowever. the expected aligniueut of the spin of pulsar B's SN progenitor witli e pre-SN orbital angular momenttun aud the observed misaligniuent of pulsar B's spin and orbit at oreseut uniquely imply that pulsar B's spin is dominated by augular momenttun procced during the SN process. uot angular inomentum provided by the progenitor.," Regardless of the specifics of the collapse process, however, the expected alignment of the spin of pulsar B's SN progenitor with the pre-SN orbital angular momentum and the observed misalignment of pulsar B's spin and orbit at present uniquely imply that pulsar B's spin is dominated by angular momentum produced during the SN process, not angular momentum provided by the progenitor."1168 The realization of this empirical constraint on angular momentuum production in supernovae presente here is uuiquely eabled by the spin spin iuisaliguineut in the PSR JOT37-3039 system (Lyueetal.2008:Lyutikov&Thompson2005) and can be used to guide core-collapse simulations and te quest for the uuderstaudiug of compact object spius aud kicks.," The realization of this empirical constraint on angular momentum production in supernovae presented here is uniquely enabled by the spin spin misalignment in the PSR J0737-3039 system \citep{Lyne2004,Ferdman2008,Lyutikov2005}1169 and can be used to guide core-collapse simulations and the quest for the understanding of compact object spins and kicks."1170 We thank Tsiig Wal Woi© for providiug «ata ou the distributiou of allowed pulsar kicks from the current orbita constraints from Wongetal.(2010) (Figure 5)., We thank Tsing Wai Wong for providing data on the distribution of allowed pulsar kicks from the current orbital constraints from \citet{Wong2010} (Figure 5).1171 We also thank Luerid Stairs. Hans-Thomas Jatca. Aclain BWrows. aud Rodrigo Fernaucdez fo: helpful comments ou drafts of this manuscript.," We also thank Ingrid Stairs, Hans-Thomas Janka, Adam Burrows, and Rodrigo Fernandez for helpful comments on drafts of this manuscript."1172 WALIF aud Vis are partially supported by NSF erait AST-0008930., WMF and VK are partially supported by NSF grant AST-0908930.1173 Ids acknowledges support from a NASA summe: research: grant through the Illinois Space Grant NNGO2COS81IH. ML is supported by NASA grant INXO00AHAJTCG., KK acknowledges support from a NASA summer research grant through the Illinois Space Grant NNG05G381H. ML is supported by NASA grant NNX09AH37G.1174mid-IK. emission from the torus is anisotropic by a factor of <3 (Shietal.2005).,mid-IR emission from the torus is anisotropic by a factor of $\le 3$ \citep{srh05}.1175. We presentSpitzer observations of a sample of 42 FR IL radio galaxies ancl quasars selected [rom the 3CRR survey., We present observations of a sample of 42 FR II radio galaxies and quasars selected from the 3CRR survey.1176 The goals are (o search for mid-IR. emission [rom hidden quasar nuclei and (est (he ubiquity of the unification hvpothesis., The goals are to search for mid-IR emission from hidden quasar nuclei and test the ubiquity of the unification hypothesis.1177 TheSpilzer Inlrared Spectrograph (RS) combines the advantages of unprecedented sensitivity [rom 5-36.5 pam to measure (he mid-IB. continuum and spectral resolution to measure hieh ionization emission lines powered bv hidden AGNs., The Infrared Spectrograph (IRS) combines the advantages of unprecedented sensitivity from 5-36.5 $\mu$ m to measure the mid-IR continuum and spectral resolution to measure high ionization emission lines powered by hidden AGNs.1178 In the current paper. we present evidence lor hidden quasar nuclei based on mid-IR photometry extracted from the IRS spectra.," In the current paper, we present evidence for hidden quasar nuclei based on mid-IR photometry extracted from the IRS spectra."1179 We examine in detail (he spectra of the subset of 14 mic-IR luminous radio galaxies which appear to contain hidden quasar nuclei., We examine in detail the spectra of the subset of 14 mid-IR luminous radio galaxies which appear to contain hidden quasar nuclei.1180 Spectra of the quasars and mid-IR. weak radio galaxies ancl a statistical study of the complete sample will be presented in separate papers., Spectra of the quasars and mid-IR weak radio galaxies and a statistical study of the complete sample will be presented in separate papers.1181 We beein by selecting a well-defined. radio flux-limited ancl redshift-limited sample of 55 radio galaxies and quasars from the 3CRR catalog (Laing.Rilev.&Longair1983).," We begin by selecting a well-defined, radio flux-limited and redshift-limited sample of 55 radio galaxies and quasars from the 3CRR catalog \citep{lrl83}."1182. We include all 3CRB sources with FR Π radio morphology. a flix of S45:>16.4 at 178 MIIz. and a redshift of 2<1.," We include all 3CRR sources with FR II radio morphology, a flux of $S_{178}>16.4$ at 178 MHz, and a redshift of $z<1$."1183 The original οςRR catalog has a flux limit of 10 Jv at 178 MIIz. is restricted to northern declinations (9> 107). and has galactic latitude |b>10°.," The original 3CRR catalog has a flux limit of 10 Jy at 178 MHz, is restricted to northern declinations $\delta >10 \arcdeg$ ), and has galactic latitude $|b|>10\arcdeg$."1184 It is the canonical low-lrequency selected catalog of bright radio sources. has optical identifications and redshifts for all entries. and has been extensively observed in most wavebands.," It is the canonical low-frequency selected catalog of bright radio sources, has optical identifications and redshifts for all entries, and has been extensively observed in most wavebands."1185 We select only sources with FR IL radio morphology., We select only sources with FR II radio morphology.1186 We verify or update the FR Classification of all sources by inspection of the latest published radio maps., We verify or update the FR classification of all sources by inspection of the latest published radio maps.1187 Compact. sources (CSSs: 3C 48. 138. 147. 286. and 309.1) with radio major axis D«10 kpc (Fantietal.1985). ave excluded from (he sample because (hey may constitute a class of voung or frustrated radio sources.," Compact, steep-spectrum sources (CSSs: 3C 48, 138, 147, 286, and 309.1) with radio major axis $D<10$ kpc \citep{ffp85} are excluded from the sample because they may constitute a class of young or frustrated radio sources."1188" Here and throughout (his paper. we assume a cosmology with Jf)—70 km ! |f. Q4,=0.3. and O4=0.7."," Here and throughout this paper, we assume a cosmology with $H_0=70$ km $^{-1}$ $^{-1}$, $\Omega_\mathrm{m}=0.3$, and $\Omega_\Lambda=0.7$."1189 Size and morphology indicate tha (ος are not related to FR Is by orientation., Size and morphology indicate that CSSs are not related to FR IIs by orientation.1190 It is essential lor our unification studies Chat we select a sample based on isotropic radio lobe flux. ancl on optical or I. properties. so that it is unbiased by selection effects.," It is essential for our unification studies that we select a sample based on isotropic radio lobe flux, and on optical or IR properties, so that it is unbiased by orientation-dependent selection effects."1191 In particular. our sample includes no blazars.," In particular, our sample includes no blazars."1192 No sources make the flux limit only because of beamed emission from (he core of the radio jet., No sources make the flux limit only because of beamed emission from the core of the radio jet.1193 Our sample, Our sample1194As a first step. we rotated our galaxies (to align their major axes with the horizontal axis) using an iterative method until the position angle (PA) had been determined to an accuracy of approximately 0.05 degrees.,"As a first step, we rotated our galaxies (to align their major axes with the horizontal axis) using an iterative method until the position angle (PA) had been determined to an accuracy of approximately 0.05 degrees."1195 We used the PA from the ESO-LV Catalogue as a starting point; we then iteratively fitted a straight line to the central region of the galaxy., We used the PA from the ESO-LV Catalogue as a starting point; we then iteratively fitted a straight line to the central region of the galaxy.1196 Our results are shown in Fig., Our results are shown in Fig.1197 2 as isophote maps of our 20 galaxies., \ref{isofotas} as isophote maps of our 20 galaxies.1198 To calculate the warp curves. we fitted Gaussians to the vertical profiles of the galaxies in each filter.," To calculate the warp curves, we fitted Gaussians to the vertical profiles of the galaxies in each filter."1199" We only considered data with a signal-to-noise ratio greater than 3 (c = Yous+Onus Where oy, is the systematic error associated with the random background noise and σι the standard deviation). and with the FWHM of the peak smaller than 100 pixels."," We only considered data with a signal-to-noise ratio greater than 3 $\sigma$ = $\sqrt{\sigma^{2}_{\rm{sys}} + \sigma^{2}_{\rm{std}}}$, where $\sigma_{\rm{sys}}$ is the systematic error associated with the random background noise and $\sigma_{\rm{std}}$ the standard deviation), and with the FWHM of the peak smaller than 100 pixels."1200 The results of our analysis are presented in Fig. 3..," The results of our analysis are presented in Fig. \ref{alabeos},"

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