ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 For us letter we consider that accretion occurs over a range of radii within the corotation radius., For this letter we consider that accretion occurs over a range of radii within the corotation radius.3" Εις is equivalent to the ipproach taken previously by 27... ὃν, 2. and ? who have emonstrated that such an assumption reproduces observed μα»ectral line profiles ancl variability."," This is equivalent to the approach taken previously by \citet*{har94}, \citet*{muz01}, \citet*{sym05} and \citet{aze06} who have demonstrated that such an assumption reproduces observed spectral line profiles and variability."4 It should also be noted iu the accreting field geometries which we consider here are only snap-shots in time. and in reality will evolve due to the interaction with the disc.," It should also be noted that the accreting field geometries which we consider here are only snap-shots in time, and in reality will evolve due to the interaction with the disc."5 The accretion filling factors are 0.9% for the AB Dor-like field and 1.2% for the LO 1να- field. smaller than would be expected for accretion to a dipole. but consistent with observationally inferred. values (e.g. 7)).," The accretion filling factors are $0.9\%$ for the AB Dor-like field and $1.2\%$ for the LQ Hya-like field, smaller than would be expected for accretion to a dipole, but consistent with observationally inferred values (e.g. \citealt{val04}) )."6 We assume that material is supplied. by the disc. and accretes onto the star at a constant rate., We assume that material is supplied by the disc and accretes onto the star at a constant rate.7 For a. dipolar magnetic field. accretion occurs into two rings in opposite, For a dipolar magnetic field accretion occurs into two rings in opposite8There απο two objections to accepting the above arguments as disproof of the idea of flowing space.,There are two objections to accepting the above arguments as disproof of the idea of flowing space.9 One is that the derivatious here are cutirely Newtonian. aud so of course will uot show au effect which is relativistic iu origin: that problem will be addressed below.," One is that the derivations here are entirely Newtonian, and so of course will not show an effect which is relativistic in origin; that problem will be addressed below."10 The secoud is that an origin placed a finite distance away from the test particle is not appropriate. aud one should really consider the problem to be that of a particle eiven an initial peculiar velocity.," The second is that an origin placed a finite distance away from the test particle is not appropriate, and one should really consider the problem to be that of a particle given an initial peculiar velocity."11 In that case. the Dubble flow could be thought of as exerting a sort of viscous force.," In that case, the Hubble flow could be thought of as exerting a sort of viscous force."12 One would then not necessarily expect the expansion of space to act inmiediatelv and overwheliinglv., One would then not necessarily expect the expansion of space to act immediately and overwhelmingly.13 While a leaf placed in a river iuight become part of the flow at once. a ship has inertia. aud ouce the engines are stopped will continue to move with respect to the water for some time (as anvoue who has attempted fine maneuvering knows very well).," While a leaf placed in a river might become part of the flow at once, a ship has inertia, and once the engines are stopped will continue to move with respect to the water for some time (as anyone who has attempted fine maneuvering knows very well)."14 Aud it does appear. since peculiar velocity is know to decay with tine. that the free particle indeed joius the Hubble flow after passing through the origin.," And it does appear, since peculiar velocity is known to decay with time, that the free particle indeed joins the Hubble flow after passing through the origin."15 It is to the matter of joining the IIubble flow that we now turn., It is to the matter of joining the Hubble flow that we now turn.16 It is a general property of expauding universes that anv motion which departs from the general flow decreases with time. the familiar decay of peculiarvelocity.," It is a general property of expanding universes that any motion which departs from the general flow decreases with time, the familiar decay of peculiar."17 In the solutions above (and in the Appendix) it is readily seen that the time derivatives of the e» functions are all monotonically decreasing (with the exception of the overdense universe i some phases). aud iudeed decreasing more quickly than the ey (Iubble flow) fictions.," In the solutions above (and in the Appendix) it is readily seen that the time derivatives of the $c_2$ functions are all monotonically decreasing (with the exception of the overdense universe in some phases), and indeed decreasing more quickly than the $c_1$ (Hubble flow) functions."18 From this it seems clear that all particles nist eventually join the IHubble Sow as their velocity away from it vanishes., From this it seems clear that all particles must eventually join the Hubble flow as their velocity away from it vanishes.19" However. this is worth investigating quantitatively,"," However, this is worth investigating quantitatively."20 Considering the motion of a free particle iu a critical universe without a cosmological coustaut. its asvinptotic speed is The backerouud particle with this same asviuptotic speed has as its equation of notion The spatial distance between these two particles is thusunbounded.," Considering the motion of a free particle in a critical universe without a cosmological constant, its asymptotic speed is The background particle with this same asymptotic speed has as its equation of motion The spatial distance between these two particles is thus."21 Fox an uuderdeuse universe. the distance between the free particle aud the backerouud particle with the same asviuptotic speed (we might call it a ‘peculiar distance’) is more complicated. but as time becomes large approaches a constaut.," For an underdense universe, the distance between the free particle and the background particle with the same asymptotic speed (we might call it a `peculiar distance') is more complicated, but as time becomes large approaches a constant."22 That is. the free particle stavs at least this far from its correspouding backgound particle. even at infinite times.," That is, the free particle stays at least this far from its corresponding backgound particle, even at infinite times."23" A particle which keeps a distauce frou its ""proper place in the Uubble flow. even after an infinitely long time. caunot really be said to join it."," A particle which keeps a distance from its `proper' place in the Hubble flow, even after an infinitely long time, cannot really be said to join it."24reutralino signal obviously may chanec this result.,neutralino signal obviously may change this result.25 Actually we show here that in a Ον scenario. even. in case of a finite antiproton lifetime. we can produce uodels which give at the same time a very good fit of he existing data aud whose spectral features are oeculiar Chough to be distiuguished. from the standard jckerounud once measurements at higher energies will be available.," Actually we show here that in a clumpy scenario, even in case of a finite antiproton lifetime, we can produce models which give at the same time a very good fit of the existing data and whose spectral features are peculiar enough to be distinguished from the standard background once measurements at higher energies will be available."26 Iu Fie., In Fig.27 5 we compare three such models with he case of standard backerouud aud iufiuite rz., \ref{fig:spectau} we compare three such models with the case of standard background and infinite $\tau_{\bar{p}}$.28" Model 5 is à 51 GeV eate10-like neutralino whose autiproton flux iis been scaledby £6=19 aac for 7,=0.82 Myr."," Model 5 is a 51 GeV gaugino-like neutralino whose antiproton flux has been scaled by $f\,\delta = 49$ and for $\tau_{\bar{p}} = 0.82$ Myr."29" Also shown in the figure is the reduction in the background flux luneed In:"" cousidering rz,=0.82 Myr (dashed nue abelledby 5 3b).", Also shown in the figure is the reduction in the background flux induced by considering $\tau_{\bar{p}} = 0.82$ Myr (dashed line labelled by $5b$ ).30" Model7 is the heaviest model for which the oxedieted Buiis stillin excellent agreement with existiug data C.L. fit for i,=1017 GeV. fó=12:H and 7,=2.92 My»). while model 6 is some intermediate case (Gn.=188 GeV. fó=T8 aud 7,= Alvr)."," Model 7 is the heaviest model for which the predicted flux is still in excellent agreement with existing data C.L. fit for $m_{\chi} = 477$ GeV, $f\,\delta = 1.2 \cdot 10^4$ and $\tau_{\bar{p}} = 2.92$ Myr), while model 6 is some intermediate case $m_{\chi} = 188$ GeV, $f\,\delta = 78$ and $\tau_{\bar{p}} = 1.32$ Myr)."31 The trend is that for heavier ueutralinos there is a larger overproduction of autiprotous in the high energy ranec., The trend is that for heavier neutralinos there is a larger overproduction of antiprotons in the high energy range.32 Applving the largest possible rescalings consistent with -rav lneasurements we fud that a 56 CV ueutralino model gives a flux which is consistent with data at CLL. in case the autiproton lifetime is as low as Ty=045 Myr., Applying the largest possible rescalings consistent with $\gamma$ -ray measurements we find that a 56 GeV neutralino model gives a flux which is consistent with data at C.L. in case the antiproton lifetime is as low as $\tau_{\bar{p}} = 0.15$ Myr.33 The bound of Ceer Ixeuucdy (1998) as clearly violated., The bound of Geer Kennedy (1998) is clearly violated.34 Notice that we are comparing with a more aboundant data set than in that reference 97 data were not included there) and that we used our standard values for the parameters which define the diffusion model aud solar iiodulation., Notice that we are comparing with a more aboundant data set than in that reference 97 data were not included there) and that we used our standard values for the parameters which define the diffusion model and solar modulation.35" If uncertainties were included the lower bound we would ect with this method would probably be very close to the most stringent direct experimental bounud z,>0.05 Myr or lower.", If uncertainties were included the lower bound we would get with this method would probably be very close to the most stringent direct experimental bound $\tau_{\bar{p}} > 0.05$ Myr or lower.36 To conclude. we have shown that there is a chance of detecting ueutralino dark iwatter im upcoming nieasurenaents of the cosmic autiprotou flux at hieh energies.," To conclude, we have shown that there is a chance of detecting neutralino dark matter in upcoming measurements of the cosmic antiproton flux at high energies."37 The signatures we propose here are alternative to the signature of au exotic component low kinetic enereies. which secius not to be required by present data.," The signatures we propose here are alternative to the signature of an exotic component at low kinetic energies, which seems not to be required by present data."38" We have aso discussed the possibility that antiprotous have a finite lifetime and shown that the lini ou 7, which is possible to set ou the basis of cosmic rav measurements is comparable to those iu direct experiments.", We have also discussed the possibility that antiprotons have a finite lifetime and shown that the limit on $\tau_{\bar{p}}$ which is possible to set on the basis of cosmic ray measurements is comparable to those in direct experiments.39 I an eratefulrateful to LLars BerestBerestrouun audL JoakimJoal EdEdsyo for αμα useful discussions., I am grateful to Lars Bergströmm and Joakim Edsjö for many useful discussions.40 I thauk Paolo Coudolo for collaboration ou thenumerical. supersviuuuetry calculations., I thank Paolo Gondolo for collaboration on thenumerical supersymmetry calculations.41We processed 5.030 sessions of the permanent geodetic and astrometric VLBI program since 1979. totalling 7.285.312 group delay measurements at 84 GHz.,"We processed 5,030 sessions of the permanent geodetic and astrometric VLBI program since 1979, totalling 7,285,312 group delay measurements at 8.4 GHz."42 Radio source coordinates Were estimated once per session. together with Earth orientation parameters and station coordinates.," Radio source coordinates were estimated once per session, together with Earth orientation parameters and station coordinates."43 The off elevation angle was set to 5°., The cut-off elevation angle was set to $^{\circ}$.44 A priori zenith delays were determined from local pressure values (Saastamoinen 1972). which were then mapped to the elevation of the observation using the Vienna mapping functions (Bóhhm et al.," A priori zenith delays were determined from local pressure values (Saastamoinen 1972), which were then mapped to the elevation of the observation using the Vienna mapping functions (Böhhm et al."45 2006)., 2006).46 Zenith wet delays were estimated as a continuous piecewise linear function at 30-min intervals., Zenith wet delays were estimated as a continuous piecewise linear function at 30-min intervals.47 Troposphere gradients were estimated às 8-hr east and north piecewise functions at all stations except a set of 110 stations with short observational histories., Troposphere gradients were estimated as 8-hr east and north piecewise functions at all stations except a set of 110 stations with short observational histories.48 Statior heights were corrected for atmospheric pressure and oceanic tidal loading., Station heights were corrected for atmospheric pressure and oceanic tidal loading.49 The relevant loading quantities were deduced from surface pressure grids from the U. 5. NCEP/NCAR reanalysis project atmospheric global circulation model (Kalnay et al., The relevant loading quantities were deduced from surface pressure grids from the U. S. NCEP/NCAR reanalysis project atmospheric global circulation model (Kalnay et al.50 1996. Petrov Boy 2004) and from the FES 2004 ocean tide model (Lyard et al.," 1996, Petrov Boy 2004) and from the FES 2004 ocean tide model (Lyard et al."51 2004)., 2004).52 No-net rotation (NR) and translation constraints per session were applied to the positions of all stations. excluding Fort Davis (Texas). Pie Town (New Mexico). Fairbanks (Alaska). and the TIGO antenna at Concepeiónn. Chile because of strong non linear displacements (These two sites experienced post-seismic relaxation effects after large earthquakes on the Denali fault in 2003. and between Talca and Concepeiónn in early 2010).," No-net rotation (NNR) and translation constraints per session were applied to the positions of all stations, excluding Fort Davis (Texas), Pie Town (New Mexico), Fairbanks (Alaska), and the TIGO antenna at Concepciónn, Chile because of strong non linear displacements (These two sites experienced post-seismic relaxation effects after large earthquakes on the Denali fault in 2003, and between Talca and Concepciónn in early 2010)."53 A priori precession and nutation comply with the [AU 2000/2006 resolutions. which include the nutation model of Mathews et al. (," A priori precession and nutation comply with the IAU 2000/2006 resolutions, which include the nutation model of Mathews et al. ("542002). the improved precession model of Capitaine et al. (,"2002), the improved precession model of Capitaine et al. ("552003b). and the non rotating origin-based coordinate. transformation. between terrestrial and celestial coordinate systems (Capitaine et al.,"2003b), and the non rotating origin-based coordinate transformation between terrestrial and celestial coordinate systems (Capitaine et al."56 2003a)., 2003a).57 Usually. an NNR constraint is used to fix the ICRS axes.," Usually, an NNR constraint is used to fix the ICRS axes."58 However. Titov (2010) argues that application of a tight NNR constraint may wipe out all systemtic effects in the proper motion of reference radio sources.," However, Titov (2010) argues that application of a tight NNR constraint may wipe out all systemtic effects in the proper motion of reference radio sources."59 Therefore. we tied the celestial frame to the ICRF? using a loose NNR constraint uniformly applied for each session.," Therefore, we tied the celestial frame to the ICRF2 using a loose NNR constraint uniformly applied for each session."60 More details are discussed later in Section 3.2., More details are discussed later in Section 3.2.61 The calculations used the Cale 10.0/Solve 2010.05.21 geodetic VLBI analysis software package. which was developed and maintained at NASA Goddard Space Flight Center. and were carried out at the Paris Observatory IVS Analysis Center (Gontier et al.," The calculations used the Calc 10.0/Solve 2010.05.21 geodetic VLBI analysis software package, which was developed and maintained at NASA Goddard Space Flight Center, and were carried out at the Paris Observatory IVS Analysis Center (Gontier et al."62 2008)., 2008).63 Before 1990. the general deficiency of the VLBI networks. including the number of observed sources and observing antennas per session. makes the VLBI products less reliable (see. e.g.. Gontier et al.," Before 1990, the general deficiency of the VLBI networks, including the number of observed sources and observing antennas per session, makes the VLBI products less reliable (see, e.g., Gontier et al."64 2001. Malkin 2004. Feissel-Vernier et al.," 2001, Malkin 2004, Feissel-Vernier et al."65 2004. Lambert Gontier 2009 who reports interesting statistical results and remarks about the VLBI evolution over the past two decades).," 2004, Lambert Gontier 2009 who reports interesting statistical results and remarks about the VLBI evolution over the past two decades)."66 For this reason we removed data before 1990., For this reason we removed data before 1990.67 A treatment of the full data base over 1979-2010 1s nevertheless presented later for comparison., A treatment of the full data base over 1979–2010 is nevertheless presented later for comparison.68 In the coordinate time series. data points resulting from fewer than three reliable observations within à session. were removed. and outliers were eliminated so that the y is reasonably close to unity.," In the coordinate time series, data points resulting from fewer than three reliable observations within a session were removed, and outliers were eliminated so that the $\chi^2$ is reasonably close to unity."69 Then. proper motions were computed by weighted least-squares for time series containing at least ten points and longer than ten years.," Then, proper motions were computed by weighted least-squares for time series containing at least ten points and longer than ten years."70 Weights were taken as the inverse of the squared formal error., Weights were taken as the inverse of the squared formal error.71" A set of 39 sources showing significant non linear positional variations due to large-scale variations in their structure (including 3C84. 3C273B. 3C279, 3C345. 3C454.3. and 1C39.25) were isolated in the ICRF2 work and treated in such a manner that they did not perturb the geodetic solutions (Fey et al."," A set of 39 sources showing significant non linear positional variations due to large-scale variations in their structure (including 3C84, 3C273B, 3C279, 3C345, 3C454.3, and 4C39.25) were isolated in the ICRF2 work and treated in such a manner that they did not perturb the geodetic solutions (Fey et al."72 2009)., 2009).73 We removed these 39 sources from our data set., We removed these 39 sources from our data set.74 The final sample contains proper motions of 555 sources and is made available electronically., The final sample contains proper motions of 555 sources and is made available electronically.75 Figure | displays the distribution of sources and proper motion formal errors in declination., Figure \ref{fig00} displays the distribution of sources and proper motion formal errors in declination.76 Near the polar areas. the number of sources decreases proportionally to the cosine of the declination.," Near the polar areas, the number of sources decreases proportionally to the cosine of the declination."77 It also shows the nonuniformity of the sample and a lack of sources at declinations under —407., It also shows the nonuniformity of the sample and a lack of sources at declinations under $-40^{\circ}$.78" Figure 2 displays 45,cos0 versus a for 40 sources observed in more than 1.000 sessions (see Section 3.2 for details)."," Figure \ref{fig20} displays $\mu_{\alpha}\cos\delta$ versus $\alpha$ for 40 sources observed in more than 1,000 sessions (see Section 3.2 for details)."79 The apparent motions of these sources are estimated very accurately thanks to a large number of observations., The apparent motions of these sources are estimated very accurately thanks to a large number of observations.80 A systematic in sinc of magnitude less than 10 uas clearly shows up., A systematic in $\sin\alpha$ of magnitude less than 10 $\mu$ as clearly shows up.81 Tiny underlying aberrational drift is indicative even for a limited number of well-observed radio sources., Tiny underlying aberrational drift is indicative even for a limited number of well-observed radio sources.82 This section comprises our results of the dipole component estimation., This section comprises our results of the dipole component estimation.83 We start with a main solution including all 555 radio sources., We start with a main solution including all 555 radio sources.84 Then we consider different subsets of radio sources to verify the robustness of the main solution., Then we consider different subsets of radio sources to verify the robustness of the main solution.85 First. dipole and rotation coefficients. were fitted by weighted least-squares following Eqs. (," First, dipole and rotation coefficients were fitted by weighted least-squares following Eqs. ("863)}-(4) and (6)-(7) (Table .. column DR).,"3)–(4) and (6)–(7) (Table \ref{tab01}, column DR)."87 Reported errors are standard formal errors., Reported errors are standard formal errors.88 The fit produces correlations of ~O.4 between dj and r» and between d» and r|., The fit produces correlations of $\sim$ 0.4 between $d_1$ and $r_2$ and between $d_2$ and $r_1$.89 Figure 3 displays both the proper motions of the 555 sources and the estimated dipole component of the velocity field., Figure \ref{fig01} displays both the proper motions of the 555 sources and the estimated dipole component of the velocity field.90 Within error bars. the dipole amplitude and direction agree with predictions from measurements of the Galactic parameters.," Within error bars, the dipole amplitude and direction agree with predictions from measurements of the Galactic parameters."91 The corresponding, The corresponding92niaxiumui of 6045WOkkinss | atrτο ppc.,maximum of $v_{rot}=160$ $^{-1}$ at $r=375$ pc.93 The rotation curve (ο) aud the related principal frequencies derived in the epieyclie approximation (Q and 0-4/2) are represented iu Fig., The rotation curve $v_{rot}$ ) and the related principal frequencies derived in the epicyclic approximation $\Omega$ and $\Omega$ $\kappa$ /2) are represented in Fig.94 7 for the immer ppc., \ref{fig:omega} for the inner pc.95 A laree amount of eas is detected at velocities lower than deterinined by the ridge o‘terminal velocities shown in Fie. 6.., A large amount of gas is detected at velocities lower than determined by the ridge of terminal velocities shown in Fig. \ref{fig:posvel}.96 There Is an S-shaped feature in the p-v diaeranuue gonig across the vuanucal centre (denoted as SP in Fie. 6))., There is an S-shaped feature in the p-v diagramme going across the dynamical centre (denoted as SP in Fig. \ref{fig:posvel}) ).97 The characteristic pattern of SP sugeests the presence of a ioni-axisviunetrie distribution of inolecular gas 1i the form o clear nünirspiral aris. which extend froui the cenre to r~300 ppc.," The characteristic pattern of SP suggests the presence of a non-axisymmetric distribution of molecular gas in the form of nuclear mini-spiral arms, which extend from the centre to $r\sim 300$ pc."98" ALorcover. the ""figure-cielit xivttern of the p-v diagranune formed by SP and the curve of termiwl velocities is typical of a burriven eas fiow (Ixuijiken Merrifield 1995j)."," Moreover, the `figure-eight' pattern of the p-v diagramme formed by SP and the curve of terminal velocities is typical of a bar-driven gas flow (Kuijken Merrifield \cite{kuijken}) )."99" The existeice of a nucear stellar bar of kk)c diameter was iready establishec w Telesco Cezari 19923). base Ol Learifrared observations 1 ithe J. Is iux Thanes,"," The existence of a nuclear stellar bar of kpc diameter was already established by Telesco Gezari \cite{telesco9}) ), based on near-infrared observations in the J, K and I bands."100 Driven by alxw potential. eas clotds follow nonu sclfintersecting elipsoiclal orbits.," Driven by a bar potential, gas clouds follow non self-intersecting ellipsoidal orbits."101 TiC snajex axes of these orbits precess as a netion of radius. axl hence eud up delineating niviral wis. owlne o orbi crowding.," The major axes of these orbits precess as a function of radius, and hence end up delineating spiral arms, owing to orbit crowding."102 The precession of eas OFits is due to tI !dissipative nature of eas: uolecular Coud-oud collisions aud the implied viscosiv of tie pTOCeSS οιπο a smooth trasition between the har-criven + orbits (parallel to the bar major axis) towards ο orbits serpendicular to bar major axis). when we go across t1C Imer Lindblad Resouance (IER).," The precession of gas orbits is due to the dissipative nature of gas: molecular cloud-cloud collisions and the implied viscosity of the process cause a smooth transition between the bar-driven $x_1$ orbits (parallel to the bar major axis) towards $x_2$ orbits (perpendicular to bar major axis), when we go across the Inner Lindblad Resonance (ILR)."103" Therefore. the presence ) fa nuclear bar potential and a spiral eas response are ""PSituaatelv related."," Therefore, the presence of a nuclear bar potential and a spiral gas response are intimately related."104 The existence of two ILBs in the nucleus of is clearly sugecsted by our observations (see Fig. 7))., The existence of two ILRs in the nucleus of is clearly suggested by our observations (see Fig. \ref{fig:omega}) ).105identification.,identification.106" As can be seen from Figure 2,, however, also passively evolving or dusty galaxies at intermediate redshifts 2.5—4) can exhibit similarly red colors in ὅπου(z—Hi69."," As can be seen from Figure \ref{fig:colsel}, however, also passively evolving or dusty galaxies at intermediate redshifts $z\sim2.5-4$ ) can exhibit similarly red colors in $J_{125}-H_{160}$ ."107" While the requirement of optical non-detections removes the bulk of lower redshift contamination, certain intermediate redshift galaxies with evolved or dusty stellar populations can still be included due to the fact that the optical data does not reach deep enough, if at similar depth as the IR (see Fig. 3))."," While the requirement of optical non-detections removes the bulk of lower redshift contamination, certain intermediate redshift galaxies with evolved or dusty stellar populations can still be included due to the fact that the optical data does not reach deep enough, if at similar depth as the IR (see Fig. \ref{fig:maglim}) )."108 Deep Spitzer IRAC data provides a way to identify contaminating galaxies., Deep Spitzer IRAC data provides a way to identify contaminating galaxies.109" These are expected to exhibit very red Hiso—[3.6m] colors, which discriminates them from genuine z~10 candidates."," These are expected to exhibit very red $H_{160}-[3.6\micron]$ colors, which discriminates them from genuine $z\sim10$ candidates."110" To exclude possible low-redshift contamination, we thus use two steps to select z=9.5 galaxy candidates."," To exclude possible low-redshift contamination, we thus use two steps to select $z\gtrsim9.5$ galaxy candidates."111" For the first step, the primary criteria are based on HST data only: Additional to excluding objects that are detected in any band blueward of Jj25 at more than 20, we include a cut in the optical X2pt value of a galaxy (seee.g.Bouwensetal.2011b; 2011a)."," For the first step, the primary criteria are based on HST data only: Additional to excluding objects that are detected in any band blueward of $J_{125}$ at more than $2\sigma$, we include a cut in the optical $\chi^2_{opt}$ value of a galaxy \citep[see e.g.][]{Bouwens10c,Bouwens11}."112". This is computed from the 0725 radius aperture fluxes as =sign(fi)(fioi), where the sum runs over all X2ptthe bands5 available in the given data set blueward of J125, i.e. it includes all the available optical data as well as the NIR band Yos for the HUDF09 data, and Yoos in the ERS."," This is computed from the $0\farcs25$ radius aperture fluxes as $\chi^2_{opt} = \sum_i\mathrm{sign}(f_i)\left(f_i/\sigma_i\right)^2$, where the sum runs over all the bands available in the given data set blueward of $J_{125}$, i.e. it includes all the available optical data as well as the NIR band $Y_{105}$ for the HUDF09 data, and $Y_{098}$ in the ERS."113 The relatively large apertures were chosen in order sample >7096 of the light of point-like sources., The relatively large apertures were chosen in order sample $>70\%$ of the light of point-like sources.114" The limiting x2,, are derived from photometric scatter simulations.", The limiting $\chi^2_{cut}$ are derived from photometric scatter simulations.115" They are set to exclude the majority of interlopers which remain undetected at 2σ purely due to photometric noise, but not to cut a substantial fraction of galaxies with real zero flux in the optical bands."," They are set to exclude the majority of interlopers which remain undetected at $2\sigma$ purely due to photometric noise, but not to cut a substantial fraction of galaxies with real zero flux in the optical bands."116" The scatter simulations utilize all galaxies in our catalogs that are 1—3 mag above the completeness limit, applying photometric Gaussian noise from 1 mag fainter sources."," The scatter simulations utilize all galaxies in our catalogs that are $1-3$ mag above the completeness limit, applying photometric Gaussian noise from 1 mag fainter sources."117" From these simulations it is clear that contamination is mainly an issue at 0.75 mag above the completeness limits, but that ~60—80% of contaminants can be eliminated by using a limit of x2,,=2.8 or 2.4, for 5 filters or 4 filters, respectively."," From these simulations it is clear that contamination is mainly an issue at 0.75 mag above the completeness limits, but that $\sim60-80$ of contaminants can be eliminated by using a $\chi^2_{opt}$ limit of $\chi^2_{cut}=2.8$ or 2.4, for 5 filters or 4 filters, respectively."118"X2pt In the HUDFO09 data, the resulting number of expected contaminants due to photometric scatter is thus reduced from ~0.5 source per WFC3/IR field to 0.1 source."," In the HUDF09 data, the resulting number of expected contaminants due to photometric scatter is thus reduced from $\sim0.5$ source per WFC3/IR field to $\sim0.1$ source."119" On the other hand, the adopted x2,, limits do remove an additional ~20% of sources with real zero flux, simply due to Gaussian statistics."," On the other hand, the adopted $\chi^2_{cut}$ limits do remove an additional $\sim20$ of sources with real zero flux, simply due to Gaussian statistics."120 This reduction of the real galaxy sample is reflected in our subsequent analysis in the reduction of the selection volume., This reduction of the real galaxy sample is reflected in our subsequent analysis in the reduction of the selection volume.121" All galaxies passing the above selection criteria, using both the ACS and WFC3/IR data, are retained and analyzed individually."," All galaxies passing the above selection criteria, using both the ACS and WFC3/IR data, are retained and analyzed individually."122" These total to 17 sources with Hi60,AB in the range 23.6—28.8 mag; one source in the HUDF, none in the parallel HUDFO09 fields, three in the ERS and 8 and 5 in the CANDELS Deep and Wide, respectively Tables 2 and A4))."," These total to 17 sources with $H_{160,AB}$ in the range $23.6-28.8$ mag; one source in the HUDF, none in the parallel HUDF09 fields, three in the ERS and 8 and 5 in the CANDELS Deep and Wide, respectively (see Tables \ref{tab:phot} and \ref{tab:photContamin}) )."123" Interestingly, (seeonly one source previously reported galaxy with Heo~29 mag from (theBouwens et al."," Interestingly, only one source (the previously reported galaxy with $H_{160}\sim29$ mag from Bouwens et al."124" 2011) did pass our selection in the three deep HUDF09 fields,"," 2011) did pass our selection in the three deep HUDF09 fields,"125"annuli, and fits the data to flattened sinusoidal models as a function of position angle.","annuli, and fits the data to flattened sinusoidal models as a function of position angle."126" For the discrete velocity data (GCs), rotation and dispersion are fitted simultaneously through a maximum likelihood method (Figs."," For the discrete velocity data (GCs), rotation and dispersion are fitted simultaneously through a maximum likelihood method (Figs."127 2a-d)., 2a–d).128" 'The position angles and ellipticities of the rotation field and the sampling bins (PAkin, €kin) are part of the fit for the SKiMS data but are not well constrained for the discrete velocity data, which we assume follows the stellar isophotes (PA=43.5°, e=0.5)."," The position angles and ellipticities of the rotation field and the sampling bins $\rm{PA_{kin}}$ , $\rm{\epsilon_{kin}}$ ) are part of the fit for the SKiMS data but are not well constrained for the discrete velocity data, which we assume follows the stellar isophotes $^{\circ}$, $\epsilon$ =0.5)."129 Our results are insensitive to reasonable variations in these parameters., Our results are insensitive to reasonable variations in these parameters.130 Uncertainties are estimated via Monte Carlo fitting of mock data-sets., Uncertainties are estimated via Monte Carlo fitting of mock data-sets.131 The resulting rotation profiles for the different subcomponents are shown in Figs., The resulting rotation profiles for the different subcomponents are shown in Figs.132" 2e and 2f, where rolling fits with radius are used to capture the details of any radial kinematic transitions from the inner to outer bulge/halo)."," 2e and 2f, where rolling fits with radius are used to capture the details of any radial kinematic transitions (e.g from the inner to outer bulge/halo)."133" Within ~ 1.5 (e.gΠο the MRGC system rotates nearly as rapidly as the stellar bulge, supporting the coevolution of these two components, as also inferred from their similar ages and metallicities [2006).."," Within $\sim$ 1.5 $R_{\rm e}$ the MRGC system rotates nearly as rapidly as the stellar bulge, supporting the coevolution of these two components, as also inferred from their similar ages and metallicities \citep{2006MNRAS.367..815N}."134" At larger radii, this rotation decreases dramatically (see also Figs."," At larger radii, this rotation decreases dramatically (see also Figs."135" 2a,b)."," 2a,b)."136" The MPGCs have moderate rotation with a decline outside ~4R,.", The MPGCs have moderate rotation with a decline outside $\sim 4 R_{\rm e}$.137 An alternative rotation profile for the MRGCs is shown in Fig., An alternative rotation profile for the MRGCs is shown in Fig.138" 2g, after normalizing by the local velocity dispersion."," 2g, after normalizing by the local velocity dispersion."139" The photometric ellipticity profile is also plotted, showing a decrease with radius that parallels the rotational gradient."," The photometric ellipticity profile is also plotted, showing a decrease with radius that parallels the rotational gradient."140" The overall implication is for a bulge that has a high degree of rotational flattening in its central regions, while becoming rounder and dispersion-dominated in its outskirts."," The overall implication is for a bulge that has a high degree of rotational flattening in its central regions, while becoming rounder and dispersion-dominated in its outskirts."141" Having found kinematic transitions in both GC subpopulations, we look for analogous transitions in the radial metallicity profiles."," Having found kinematic transitions in both GC subpopulations, we look for analogous transitions in the radial metallicity profiles."142 First we summarize the overall color distribution of the GCs in Fig., First we summarize the overall color distribution of the GCs in Fig.143" 3a, which shows a classic bimodality."," 3a, which shows a classic bimodality."144" We will assume that this bimodality persists with increasing radius, but that the location of the color peaks may shift."," We will assume that this bimodality persists with increasing radius, but that the location of the color peaks may shift."145" At large radii we must cope with the contaminating effects of foreground stars, whose color distribution we also show in Fig."," At large radii we must cope with the contaminating effects of foreground stars, whose color distribution we also show in Fig."146" 3a, and use to construct Monte Carlo mock datasets to iteratively correct for the contaminant bias on the color peak locations."," 3a, and use to construct Monte Carlo mock datasets to iteratively correct for the contaminant bias on the color peak locations."147 Fig., Fig.148" 3c shows color versus radius, both for individual GC candidates and for the fitted peak locations."," 3c shows color versus radius, both for individual GC candidates and for the fitted peak locations."149" Both GC subpopulations have radially-decreasing colors, whichwe quantify as power-law color gradients with slopes of —0.05 and —0.07 mag per dex for MPGCs and MRGCs, respectively."," Both GC subpopulations have radially-decreasing colors, whichwe quantify as power-law color gradients with slopes of $-0.05$ and $-0.07$ mag per dex for MPGCs and MRGCs, respectively."150" Using our own empirical calibration to the (g—z) color used in ACS surveys (Pengal.[2006),, the gradients are —0.07 and —0.10 mag peret dex."," Using our own empirical calibration to the $(g-z)$ color used in ACS surveys \citep{2006ApJ...639...95P}, the gradients are $-0.07$ and $-0.10$ mag per dex."151" onverting to [Fe/H] metallicity 2006),, we estimate gradients of —0.38+0.06 and —0.17+0.04 dex per dex."," Converting to [Fe/H] metallicity \citep{2006ApJ...639...95P}, we estimate gradients of $-0.38\pm 0.06$ and $-0.17\pm 0.04$ dex per dex."152" To our knowledge, this is the first time that metallicity gradients in both GC subpopulations have been measured to large radii in any galaxy besides a few very massive ellipticals (see Section ??))."," To our knowledge, this is the first time that metallicity gradients in both GC subpopulations have been measured to large radii in any galaxy besides a few very massive ellipticals (see Section \ref{intro}) )."153 It is also one of the first cases of any galaxy type where joint rotation and metallicity gradients are observed in the halo (see also NGC 4697: 2009}; and NGC 4125: 3010).," It is also one of the first cases of any galaxy type where joint rotation and metallicity gradients are observed in the halo (see also NGC 4697: \citealt{2005ApJ...627..767M,2009ApJ...691..228M}; and NGC 4125: \citealt{2010A&A...516A...4P}) )."154 We now consider some possible implications of the rotation and metallicity gradients for NGC 3115’s assembly history., We now consider some possible implications of the rotation and metallicity gradients for NGC 3115's assembly history.155" The central bulge properties are generally consistent with a standard major merger picture, with the very high amount of rotation in this case indicative of a gas-rich merger with an uneven mass-ratio (og.⋅⋅"," The central bulge properties are generally consistent with a standard major merger picture, with the very high amount of rotation in this case indicative of a gas-rich merger with an uneven mass-ratio \citep[e.g.,][]{2005A&A...437...69B,2006MNRAS.372..839N}."156" Alternatively, the inner bulge might have formed via the inward migration of giant star forming clumps within a turbulent disk fed by cold streams from the cosmicweb at early epochs (e. al.||2009)).."," Alternatively, the inner bulge might have formed via the inward migration of giant star forming clumps within a turbulent disk fed by cold streams from the cosmicweb at early epochs \citep[e.g.,][]{1999ApJ...514...77N,2008ApJ...688...67E,2009ApJ...703..785D}. ."157" In either case, the exceptionally high inner-bulge rotation in NGC 3115 may require a residual thick disk component"," In either case, the exceptionally high inner-bulge rotation in NGC 3115 may require a residual thick disk component"158We used different methods for the ealeulation of the column density.,We used different methods for the calculation of the column density.159" In fact. for (he NU, (1.1) and (3-2) transitions the optical depth and the excitation temperature have been derived from the hyperline fitting whilst for the other lines these two crucial parameters cannot be derived. therefore we must assume the optical thin and LTE approximation."," In fact, for the (1-0), $_3$ (1,1) and (3-2) transitions the optical depth and the excitation temperature have been derived from the hyperfine fitting whilst for the other lines these two crucial parameters cannot be derived, therefore we must assume the optical thin and LTE approximation."160" For (1-0) and (3-2) the column density is given by the following formula (from Casellietal. (2002b))) valid lor optically thick transitions where Ar is the line width. ν is lrequeney of the observed transition. μι is the Einstein coellicient. g, are the statistical weight of the upper level. 7 is the optical depth.Zi, is the excitation temperature. Q is the partition function. £; is the energy of the lower level."," For (1-0) and (3-2) the column density is given by the following formula (from \cite{caselli02b}) ) valid for optically thick transitions where $\Delta v$ is the line width, $\nu$ is frequency of the observed transition, $A_{ul}$ is the Einstein coefficient, $g_u$ are the statistical weight of the upper level, $\tau$ is the optical depth,$T_{ex}$ is the excitation temperature, Q is the partition function, $E_l$ is the energy of the lower level."161 In parüeular. for the NII; (1.1) line we used an approximated formula derived by (1987).. also valid for optically thin lines," In particular, for the $_3$ (1,1) line we used an approximated formula derived by \cite{bachiller87}, also valid for optically thin lines"162lue ratio and we discuss their implications.,line ratio and we discuss their implications.163 Finally. we give a brief smunnuw in 5 5.," Finally, we give a brief summary in $\S$ 5."164" Wo assume the following ACDAM cosmology throughout the paper: | 1204, =0.3. and Q4=0.72007)."," We assume the following $\Lambda$ CDM cosmology throughout the paper: $H_0=70 \ $ $ \ $ $^{-1} \ $ $^{-1}$, $\Omega_M=0.3$, and $\Omega_\Lambda165=0.7$."166. Our sample is composed of 22 targets: 3 new spectra observed with VET-ISAAC (see Tab. 1)), Our sample is composed of 22 targets: 3 new spectra observed with VLT-ISAAC (see Tab. \ref{new_source}) )167 aud 19 sources from the literature (see Tab. 2)).," and 19 sources from the literature (see Tab. \ref{lit_spec}) ),"168 kindly provided by the respective authors., kindly provided by the respective authors.169 Ten of the literature sourceshave redshift L50«τς5.702002).. while the remaining 9 have 5.70<2«6.12 and z-band maguitudes apo20.9 GQuagnitudes are taken from the discovery papers).," Ten of the literature sourceshave redshift $4.50 <z<5.70$, while the remaining 9 have $5.70 < z < 6.43$ and z-band magnitudes $z'_{AB}<20.9$ (magnitudes are taken from the discovery papers)."170 We lave observed 3 SDSS QSOs with magnitudes 19.6«ig<20.7 and redshifts 6.05<2<6.08 (sec Tab. 13)., We have observed 3 SDSS QSOs with magnitudes $19.6<z'_{AB}<20.7$ and redshifts $<$ $<$ 6.08 (see Tab. \ref{new_source}) ).171 The two faintest oues have beeu selected fro the SDSS Stripe 82 aud extend the existing sample towards the faint end of the QSO luminosity function., The two faintest ones have been selected from the SDSS Stripe 82 and extend the existing sample towards the faint end of the QSO luminosity function.172 The observations were carried out with ISAAC ou Autu (VLT-UT1) in low resolution mode (LR). using the 10511051 Tawa Rockwell array of the Short Waveleneth arm.," The observations were carried out with ISAAC on Antu (VLT-UT1) in low resolution mode (LR), using the 1024x1024 Hawaii Rockwell array of the Short Wavelength arm."173 For cach QSO the lline and the cecomiplex were observed: given the redshift of the sources.tb these features fall in the IK band.," For each QSO the line and the complex were observed: given the redshift of the sources, these features fall in the K band."174" The selected slit had a width of D"" aud combined with the order selection filter it gives a spectral resolution A/AA~ 150."," The selected slit had a width of 1"" and combined with the order selection filter it gives a spectral resolution $\lambda/\Delta\lambda \sim$ 450."175 Table 1) stuumarizes the exposure time for each object., Table \ref{new_source} summarizes the exposure time for each object.176 For cach observation block (OB). sixteen frames of 115 seconds were taken following an ABBA dithering patteru Qvith larec offsets among the dithered positions: frou WwYo to 307).," For each observation block (OB), sixteen frames of 148 seconds were taken following an ABBA dithering pattern (with large offsets among the dithered positions: from 20” to 30”)."177 Further. μπα] random offsets within a box of LY to 18 were applied at cach dithered position im order to avoid pixebrelated artifacts (jitteriue).," Further, small random offsets within a box of 4” to 18” were applied at each dithered position in order to avoid pixel-related artifacts (jittering)."178" CaveEeeR, the faintuess of the sources. the observation setup wa: chosen so that a bright star was always in the slit. i order to allow a correct centering of the target."," Given the faintness of the sources, the observation setup was chosen so that a bright star was always in the slit, in order to allow a correct centering of the target."179 The QSO SDSS J2051-0005 has been discovered in the SDSS deep Stripe 82 bv(2008)., The QSO SDSS J2054-0005 has been discovered in the SDSS deep Stripe 82 by.180. Our new ISAAC spectymu confirms the weak-line nature of this source2008)., Our new ISAAC spectrum confirms the weak-line nature of this source.181. Given the iutriusic wealsucss of the eemission. we do not include this QSO in the following analysis.," Given the intrinsic weakness of the emission, we do not include this QSO in the following analysis."182 The ESO ISAAC pipeline produces waveleneth calibrated co-added 2-D spectra from the individual frames that were acquired curing cach OD., The ESO ISAAC pipeline produces wavelength calibrated co-added 2-D spectra from the individual frames that were acquired during each OB.183 Subsequenut reduction was carried out within ΤΗΔΕ., Subsequent reduction was carried out within IRAF.184 One-dimensional spectra were oxtracted using the task., One-dimensional spectra were extracted using the task.185 The tracing of the 1-D spectra was performed first on the bright stars in the slit., The tracing of the 1-D spectra was performed first on the bright stars in the slit.186 The resulting traciug fuuctious were then used for the extraction of the QSO spectra., The resulting tracing functions were then used for the extraction of the QSO spectra.187 Tudividual 1D spectra were corrected for the tellurie absorptions using the task., Individual 1D spectra were corrected for the telluric absorptions using the task.188 Usually the telluzic correction is performed by dividing the observed spectrum by the oue of a telluvic standard star observed shortly after the science target., Usually the telluric correction is performed by dividing the observed spectrum by the one of a telluric standard star observed shortly after the science target.189 This ratio is subsequently multiplied by the model atmosphere corresponding to the spectral type of the tellure standard aud scaled. to its observed Ik magnitude. iu order to recover the correct slope of the QSO spectriuu.," This ratio is subsequently multiplied by the model atmosphere corresponding to the spectral type of the telluric standard and scaled to its observed K magnitude, in order to recover the correct slope of the QSO spectrum."190 This operation allows us also to flux-calibrate the QSO spectrum., This operation allows us also to flux-calibrate the QSO spectrum.191 Iustead. of using the observed telluric standard stars. we cuploved the ESO sky absorption spectrum measured on the Paranal site at a nonünal aiiuass of 1.," Instead of using the observed telluric standard stars, we employed the ESO sky absorption spectrum measured on the Paranal site at a nominal airmass of 1."192 This choice was criven by two reasons: i} the spectral regions in the QSO spectrmm where τομής absorptions are lost severe are characterized by a verv low signal-to-noise. ic. insuffiicicut to properly correct for the actual detailed shape of the night sky. features: 1) the observed telluric standard stars were often characterized by a spectral aud lumnimnositv class for which accurate model atiuosphiere could not be computed: the derived QSO. continua slope would have hence been distorted by the stellar spectral shape.," This choice was driven by two reasons: i) the spectral regions in the QSO spectrum where telluric absorptions are most severe are characterized by a very low signal-to-noise, i.e. insufficient to properly correct for the actual detailed shape of the night sky features; ii) the observed telluric standard stars were often characterized by a spectral and luminosity class for which accurate model atmosphere could not be computed; the derived QSO continuum slope would have hence been distorted by the stellar spectral shape."193 We thus decided to focus ou the spectral regions with higher sgual where au accurate correction was possible: we assumed the template sky absorption spectrum scaled to the airmass of our QSO spectra and ignored the stability of the sky transparency.," We thus decided to focus on the spectral regions with higher signal, where an accurate correction was possible: we assumed the template sky absorption spectrum, scaled to the airmass of our QSO spectra and ignored the stability of the sky transparency."194 This way. we could preserve the intrinsic shape of the QSO continua.," This way, we could preserve the intrinsic shape of the QSO continua."195 The task corrects for the difference im aininass between the science aud calibration spectra via the Beer-Laibert-Bounener law., The task corrects for the difference in airmass between the science and calibration spectra via the Beer-Lambert-Bouguer law.196 The telluric absorptions are in any case well removed from the spectra of bright QSOs. while significant residuals are left in the spectra of faint QSOs at lower S/N. The relative flux calibration was obtained with thesensfunction.. aud tasks.," The telluric absorptions are in any case well removed from the spectra of bright QSOs, while significant residuals are left in the spectra of faint QSOs at lower S/N. The relative flux calibration was obtained with the, and tasks."197 The iustrument sensitivity. function was obtained from the observed tellure stars of luminosity class V (giants) and spectral class D. For these stars it is possible to compute reliable model atmospheres since effective temperature and surface gravity are well estimated., The instrument sensitivity function was obtained from the observed telluric stars of luminosity class V (giants) and spectral class B. For these stars it is possible to compute reliable model atmospheres since effective temperature and surface gravity are well estimated.198" The model atinosphieres were computed by interpolating the NIR spectral cuerey distributions available from πι feniperature ancl surtace-eravity in order to match the observed spectral type and magnitudes,", The model atmospheres were computed by interpolating the NIR spectral energy distributions available from in temperature and surface-gravity in order to match the observed spectral type and magnitudes.199 After the relative flux calibration. the individual 1D QSO spectra were averaged to form a single spectrum.," After the relative flux calibration, the individual 1D QSO spectra were averaged to form a single spectrum."200 The absolute flix calibration was performed scaling the observed spectra to match the QSOs Ἱνας magnitudes., The absolute flux calibration was performed scaling the observed spectra to match the QSOs K-band magnitudes.201 Since the IK-baud magnitude was not measured for the QSOs iu Tab. 1l.," Since the K-band magnitude was not measured for the QSOs in Tab. \ref{new_source},"202 we derived it from the SDSS QSO template scaled to the observed Jo aud ILband magnitude., we derived it from the SDSS QSO template scaled to the observed J and H-band magnitude.203 The whole reduction procedure was exteusivelv tested on some of the observed telluric standard stars used as reference targets., The whole reduction procedure was extensively tested on some of the observed telluric standard stars used as reference targets.204 The recovered spectra match the theoretical wnodel atinospheres typically within104.. even iu the spectral reeious iore affected by τοήΊο absorptions.," The recovered spectra match the theoretical model atmospheres typically within, even in the spectral regions more affected by telluric absorptions."205 The reduced spectra are shown in Fie. 3.L., The reduced spectra are shown in Fig. \ref{fig_new}.206 The QSOs taken from the literature are siumnUmavized in Tab. 2.., The QSOs taken from the literature are summarized in Tab. \ref{lit_spec}.207 We have collected a total of 22 spectra(19 different sources) which cover all the features of interest at sufficicutly hieh S/N to perform the spectral decomposition (sec Sec., We have collected a total of 22 spectra(19 different sources) which cover all the features of interest at sufficiently high S/N to perform the spectral decomposition (see Sec.208 L2. and Sec. 1.6))., \ref{sec_SN_mass} and Sec. \ref{sec_SN_met}) ).209 The literature saluple is composed of LO sources with :«5.70 aud 9 sources with 5.70<2 6.13., The literature sample is composed of 10 sources with $z<5.70$ and 9 sources with $5.70 < z < 6.43$ .210 The 10 QSOs with :« are selected. amongst the 13 sources published by (2002)..., The 10 QSOs with $z<5.70$ are selected amongst the 13 sources published by .211 These QSOs have redshifts, These QSOs have redshifts212"Here, we summarise the effects of various foreground cleaning methods on the potential power spectrum measurement.","Here, we summarise the effects of various foreground cleaning methods on the potential power spectrum measurement."213" Figure 10 shows the spherically averaged 21-cm power spectra for a foreground cleaning run performed on a simulated, noiseless Stokes J signal over a 32 MHz band, centred on 177.5 MHz (6.3<z 7.8) using 128 channels."," Figure \ref{fig:ps} shows the spherically averaged 21-cm power spectra for a foreground cleaning run performed on a simulated, noiseless Stokes $I$ signal over a 32 MHz band, centred on 177.5 MHz $6.3 \lsim\,\,z \lsim\,\,7.8$ ) using 128 channels."214" Instrumental sensitivity curves for MWA500 (dashed)) and MWAS5000 (dot-dashed)) and 1000 hr integrations are also shown, where it has been assumed that the total 32 MHz bandwidth has been split into four 8 MHz sub-bands."," Instrumental sensitivity curves for MWA500 ) and MWA5000 ) and 1000 hr integrations are also shown, where it has been assumed that the total 32 MHz bandwidth has been split into four 8 MHz sub-bands."215" It can be seen that the power spectrum of the cleaned signal agrees very well with the power spectrum of the model cosmic signal, once it has been apodised and further cleaned"," It can be seen that the power spectrum of the cleaned signal agrees very well with the power spectrum of the model cosmic signal, once it has been apodised and further cleaned"216significamtlv smaller than the previously evaluated (e.g. <2=~ 2.8 [or 16 bursts by Jakobssonetal. 2006)).,"significantly smaller than the previously evaluated (e.g., $<z>\,\sim$ 2.8 for 16 bursts by \citealp{jakobsson06}) )."217" The distributions of the BAT duration 75, both in the observer and vest frame are shown in Figure 1..", The distributions of the BAT duration $T_{90}$ both in the observer and rest frame are shown in Figure \ref{fig:t90}.218 The mean value of Z5) in observers frame is 80.4 s. consistent with that deduced in a sample of 237 bursts bv Sakamotoetal.(2008).," The mean value of $T_{90}$ in observer's frame is 80.4 s, consistent with that deduced in a sample of 237 bursts by \citet{sakamoto08}."219. With the 150 measured redshifts in our sample. we find (hat (he mean value of Loy in the rest [rame is 29.2 s. Shorter bursts increase significantly in munber. but no positive classification can be made.," With the 150 measured redshifts in our sample, we find that the mean value of $T_{90}$ in the rest frame is 29.2 s. Shorter bursts increase significantly in number, but no positive classification can be made."220" Even so. as indicated in Figure 1.. there is an apparent oddball. GRBO50509B. which has an extremely short curation Z5,~0.04 s in the rest frame."," Even so, as indicated in Figure \ref{fig:t90}, there is an apparent oddball, GRB050509B, which has an extremely short duration $T_{90}\sim 0.04$ s in the rest frame."221" To view the overall behaviors of GRBs in their rest frame. we take advantage of theSwift Burst Analyser. which provides the combined DAT-XRT light curves in the form of Πας density F,,(77) at LO keV in the observers rest frame (Evansetal.2010)."," To view the overall behaviors of GRBs in their rest frame, we take advantage of the Burst Analyser, which provides the combined BAT-XRT light curves in the form of flux density $F_\nu(\nu)$ at 10 keV in the observer's rest frame \citep{evans10}."222. With the measured redshift z. the isotropic spectral luminosity al a given photon [frequency vy (which is 10 keV in this work) in the rest frame can be calculated by where νι.= m(l-4zi)is the emitted photon energy in the rest. frame for the observed photonenergv mj amd /= fg/(lo:)is the time measured in the rest. frame related (o (he time /j measured in the observers frame.," With the measured redshift $z$, the isotropic spectral luminosity at a given photon frequency $\nu_0$ (which is 10 keV in this work) in the rest frame can be calculated by where $\nu_e=\nu_0(1+z)$ is the emitted photon energy in the rest frame for the observed photonenergy $\nu_0$ and $t=t_0/(1+z)$ is the time measured in the rest frame related to the time $t_0$ measured in the observer's frame."223" Here. a single power-law spectrum has been assumed F(t)xpl1HU, where P(/) is the time-dependent photon index. which is available bv the Burst Analyser (Evansetal.2010)."," Here, a single power-law spectrum has been assumed $F_\nu(\nu,t)\propto\nu^{1-\Gamma(t)}$, where $\Gamma(t)$ is the time-dependent photon index, which is available by the Burst Analyser \citep{evans10}."224". Throughout this work. the luminosity distance Dj(2) is caleulated by assuming the cosmological parameters 44,=7lkims!Mpe. t. Oy,=0.27. and O4=0.73."," Throughout this work, the luminosity distance $D_{\rm L}(z)$ is calculated by assuming the cosmological parameters $H_0=71\,{\rm km}\, {\rm s}^{-1}\, {\rm Mpc}^{-1}$ , $\Omega_M=0.27$, and $\Omega_\Lambda=0.73$."225 According to the error propagation rule. the uncertainty in (he isolropic spectral luminosity can be given by where σι: and gp are the uncertainties in (he measured fIux density and photon index. respectively. and the uncertainty in GRB redshift measurements is considered negligible.," According to the error propagation rule, the uncertainty in the isotropic spectral luminosity can be given by where $\sigma_{F_\nu}$ and $\sigma_\Gamma$ are the uncertainties in the measured flux density and photon index, respectively, and the uncertainty in GRB redshift measurements is considered negligible."226 In Figure 2.. the calculated. X-ray light curves from 0.01 s to 10* s after BAT trigger al 10 keV in the rest frame with the spectral evolution of 150 GRBs are plotted together.," In Figure \ref{fig:restframelc}, the calculated X-ray light curves from 0.01 s to $10^7$ s after BAT trigger at 10 keV in the rest frame with the spectral evolution of 150 GRBs are plotted together."227 Interestingly. some underlyingglobal features are revealed in both the light curves," Interestingly, some underlyingglobal features are revealed in both the light curves"228"black hole would accrete al a rate corresponding to a higher effective ambient temperature ol a stationary black hole such that: where{τους is (he actual ambient temperature and 77,,,—0 is Lhe ambient temperature (and corresponding sound speed) (that woulcl give (he proper accretion rate in the calculations of Park&Ricotti(2011). lor stationary black holes.","black hole would accrete at a rate corresponding to a higher effective ambient temperature of a stationary black hole such that: where$T_{ism,v>0}$ is the actual ambient temperature and $T_{ism,v=0}$ is the ambient temperature (and corresponding sound speed) that would give the proper accretion rate in the calculations of \citet{PR11} for stationary black holes."229" I. for a given. ambient temperature. the accretion rate for a moving black hole corresponds to accretion al effectively a higher ambient temperature for a stationary black hole. then the values of the radiative feedback parameter derived by Park&Ricotti(2011) for a given ambient temperature must also be scaled to that same. higher effective temperature or. from Equation 15:: where <A,4429 is (lie radiative feedback efficiency parameter for the same value of Thom but for a stationary black hole [rom Park&Dicotti(2011)."," If, for a given ambient temperature, the accretion rate for a moving black hole corresponds to accretion at effectively a higher ambient temperature for a stationary black hole, then the values of the radiative feedback parameter derived by \citet{PR11} for a given ambient temperature must also be scaled to that same, higher effective temperature or, from Equation \ref{lambda_rad}: where $<\lambda_{rad}>_0$ is the radiative feedback efficiency parameter for the same value of $T_{ism}$, but for a stationary black hole from \cite{PR11}."230". We can thus write for a black hole moving with a velocitv. v: The results of (Park&Ricotti2011) for a given ambient temperature are thus equivalent (o those of a somewhat higher ""effec(ive ambient temperature for a moving black hole."," We can thus write for a black hole moving with a velocity, v: The results of \citep{PR11} for a given ambient temperature are thus equivalent to those of a somewhat higher “effective"" ambient temperature for a moving black hole."231 This correctionfactor is quantitatively important. but not qualitatively important for eSe.," This correctionfactor is quantitatively important, but not qualitatively important for $v \lesssim c_{s}$."232 The steep dependence on this [actor may become significant if the black holes move supersonically through the ambient medium., The steep dependence on this factor may become significant if the black holes move supersonically through the ambient medium.233 With Equations 10.. 18.. 13. and 14. we can write: where the factor. F. is: Or," With Equations \ref{bondi}, \ref{mdot1}, , \ref{rad1} and \ref{rad2}, , we can write: where the factor, F, is: or"234"void volume function of the L model more or less catches up those in the N models, as is evident in the upper left panel of Fig. 2..","void volume function of the L model more or less catches up those in the N models, as is evident in the upper left panel of Fig. \ref{vvf}."235" We also note that at a—1.0 there are less small voids in the N4 model than in the L and other N models, which is likely because of the fact that small voids have been used up to merge to form bigger ones."," We also note that at $a=1.0$ there are less small voids in the N4 model than in the L and other N models, which is likely because of the fact that small voids have been used up to merge to form bigger ones."236" For the C models, the suppress of the fifth force means that the clustering of matter and growth of voids are less affected by it*."," For the C models, the suppress of the fifth force means that the clustering of matter and growth of voids are less affected by it."237". This is easily seen in the a=0.5 case (Fig. 2,,"," This is easily seen in the $a=0.5$ case (Fig. \ref{vvf},"238" lower right panel), which shows that the void volume functions for the C models do not deviate much from that for the L model (one might appreciate the effect of the suppress in the fifth force by considering that the ratio between the magnitudes of fifth force and gravity is 24? if the former is not suppressed, and y~O(0.1) for N models while y~O(1) for C models)."," lower right panel), which shows that the void volume functions for the C models do not deviate much from that for the L model (one might appreciate the effect of the suppress in the fifth force by considering that the ratio between the magnitudes of fifth force and gravity is $2\gamma^2$ if the former is not suppressed, and $\gamma\sim\mathcal{O}(0.1)$ for N models while $\gamma\sim\mathcal{O}(1)$ for C models)."239" We could also have an examination of the void filling factor, defined as the fraction of total space that is filled by voids which are either bigger or smaller than V."," We could also have an examination of the void filling factor, defined as the fraction of total space that is filled by voids which are either bigger or smaller than $V$."240" Because our algorithm leaves the very small voids undetected, we choose to show the former, and the results are given in Fig. 3.."," Because our algorithm leaves the very small voids undetected, we choose to show the former, and the results are given in Fig. \ref{ff}."241 It turns out that this plot shows more clearly the effects of the scalar coupling., It turns out that this plot shows more clearly the effects of the scalar coupling.242" As our first example, for the L model at a=0.5 (Fig. 3,,"," As our first example, for the L model at $a=0.5$ (Fig. \ref{ff},"243" lower left panel), we notice that only of the total space is filled by voids larger than 35)ὃ Mpc’, in contrast to more than and for the models N3 and N4 respectively."," lower left panel), we notice that only of the total space is filled by voids larger than $35h^{-3}$ $^3$, in contrast to more than and for the models N3 and N4 respectively."244" At a=1.0, as a result of void growth and mergers, the numbers for these three models are changed to25%,, and respectively."," At $a=1.0$, as a result of void growth and mergers, the numbers for these three models are changed to, and respectively."245" In both cases, the scalar field coupling dramatically changes the total volume of void"," In both cases, the scalar field coupling dramatically changes the total volume of void"246that the PCI chip with its strong backeround light gradient was the only area where | was a significant function of radius: on the WE chips a single function. (11) could be used.,that the PC1 chip with its strong background light gradient was the only area where $f$ was a significant function of radius; on the WF chips a single function $f(m)$ could be used.247 The instrumental magnitudes were converted to (he Johnson-Cousins J svstem with the stancard transformations for ΕΣΕΤ found in Holtzmanetal.(1995)., The instrumental magnitudes were converted to the Johnson-Cousins $I$ system with the standard transformations for $F814W$ found in \citet{Hol95}.248". For the four individual BCGs. Galactic exlinclion corrections of A,= 0.16. 0.10. 0.11 and 0.06 respectively have been adopted (.1, is given by A,=1.532Byy. where the Eg\ values are obtained Dor each ealaxv [rom NASA/IPAC Extragalactie Database (NED))."," For the four individual BCGs, Galactic extinction corrections of $A_I =\;$ 0.16, 0.10, 0.11 and 0.06 respectively have been adopted $A_I$ is given by $A_I=1.82 \; E_{B-V}$, where the $E_{B-V}$ values are obtained for each galaxy from NASA/IPAC Extragalactic Database (NED))."249 To emplov the transformation equations and also to step back and forth between { and the (more normally used) V magnitude scale for globular clusters. we have simply assumed a color index of 1.10.1. typical of moderately metal-rich globular clusters in giant E galaxies (e.g. (1999))).," To employ the transformation equations and also to step back and forth between $I$ and the (more normally used) $V$ magnitude scale for globular clusters, we have simply assumed a color index of $(V-I)_0 = 1.1 \pm 0.1$ , typical of moderately metal-rich globular clusters in giant E galaxies (e.g. \citet{Kun99}) )."250 The intrinsic range in (V.— Z)s. folded through the trausformation equations. will not introduce uncertainties larger than £0.03 in the calibration of J.," The intrinsic range in $(V-I)_0$ , folded through the transformation equations, will not introduce uncertainties larger than $\pm 0.03$ in the calibration of $I$."251 The assumed (V—η value is (he mean value representative of most other gE galaxies. ancl if the GCSs were entirely metal-rich or metal-poor. (he error introduced by (he assumption would be at most 0.1 mae.," The assumed $(V-I)_0$ value is the mean value representative of most other gE galaxies, and if the GCSs were entirely metal-rich or metal-poor, the error introduced by the assumption would be at most 0.1 mag."252 The projected number density σ of detected objects around each galaxy plainly reveals an extensive GCS concentrated around (he galaxy center in each case., The projected number density $\sigma$ of detected objects around each galaxy plainly reveals an extensive GCS concentrated around the galaxy center in each case.253 The profile is reasonably well represented by a simple power-law form o(r)=eut(r)+OngephTig. Where σε 1s the background number density of starlike objects (mostly faint. small galaxies which passed through the image classification routines. plus a few foreground Galactic stus).," The profile is reasonably well represented by a simple power-law form $ \sigma(r) = \sigma_{cl}(r) + \sigma_{bg} = a254\;r^b + \sigma_{bg} $, where $\sigma_{bg}$ is the background number density of starlike objects (mostly faint, small galaxies which passed through the image classification routines, plus a few foreground Galactic stars)."255 To obtain the profile parameters of the GCS for each galaxy. we subdivided the WEDPC? fields into annuli 50 pixels wide. centered on the DCGs.," To obtain the profile parameters of the GCS for each galaxy, we subdivided the WFPC2 fields into annuli 50 pixels wide, centered on the BCGs."256 The number densitv of objects was then calculated down to a cutoff magnitude at which incompleteness corrections were still small., The number density of objects was then calculated down to a cutoff magnitude at which incompleteness corrections were still small.257 The projected number density is then just &.=N/A. where N is the number of detected objects within each annulus and ο is the area of that annulus which falls within the ΝΕΟΣ boundaries (minus (he small masked-out areas).," The projected number density is then just $\sigma = N/A$, where $N$ is the number of detected objects within each annulus and $A$ is the area of that annulus which falls within the WFPC2 boundaries (minus the small masked-out areas)."258 Completeness corrections. (hough small. were explicitly accounted for.," Completeness corrections, though small, were explicitly accounted for."259" Finally. the background density 05, on each of the four fields was delined as the mean of the outermost eight annuli. which fall on the outskirts of the WF chips."," Finally, the background density $\sigma_{bg}$ on each of the four fields was defined as the mean of the outermost eight annuli, which fall on the outskirts of the WF chips."260" This corresponds to a radial distance greater than 105"" (or about 30 kpe) from the centers of the DCGs.", This corresponds to a radial distance greater than $105''$ (or about 30 kpc) from the centers of the BCGs.261 Although the GCSs probably extend at trace amounts farther out than this boundary. the directly observed o(r) curves (Figure 1)) have plainlyalmost leveled off there. indicating that we are already including the main portion of the GCS.," Although the GCSs probably extend at trace amounts farther out than this boundary, the directly observed $\sigma(r)$ curves (Figure \ref{rad_plot}) ) have plainlyalmost leveled off there, indicating that we are already including the main portion of the GCS."262The estimates of £45. Dy; and mi obtained in equations (4). (5) for the shell model are valid in (his case as well. (,"The estimates of $R_{15}$ , $D_{15}$ and $n_{15}$ obtained in equations (4), (5) for the shell model are valid in this case as well. ("263Recall that ej; refers to the projected radius of the emitüng region. and that 7 refers to the optical depth parallel to the line of sieht. not perpendicular to the inclined. funnel surface.),"Recall that $R_{15}$ refers to the projected radius of the emitting region, and that $\tau$ refers to the optical depth parallel to the line of sight, not perpendicular to the inclined funnel surface.)"264 In this model there is no longer any clilficeult with having a very small value of d/D., In this model there is no longer any difficult with having a very small value of $d/D$.265 The ionizing photons trom the GRB impinge on the inner surface of the funnel aud (he line photons are also emitted from (he same surface., The ionizing photons from the GRB impinge on the inner surface of the funnel and the line photons are also emitted from the same surface.266 Thus. the radiating region can be an arbitrarily thin laver (unlike in the case of the shell model).," Thus, the radiating region can be an arbitrarily thin layer (unlike in the case of the shell model)."267 There is. however. a problem with the amount of mass required in the model.," There is, however, a problem with the amount of mass required in the model."268 Let us assume (hat (he funnel is carved oul of a roughly quasi-spherical external medium (11651 probably the supernova ejecta)., Let us assume that the funnel is carved out of a roughly quasi-spherical external medium (most probably the supernova ejecta).269 Given (he electron densitv. à». ancl the radius of the sphere D. we estimate the external mass to be For 8H~1. the mass is unacceptable large.," Given the electron density $n_e$ and the radius of the sphere $D$, we estimate the external mass to be For $\theta_{-1}\sim1$, the mass is unacceptable large."270oS Even if we take ϐH~10. 1.e.. no beamineg.5 ihe mass is still much too large.," Even if we take $\theta_{-1}\sim10$, i.e., no beaming, the mass is still much too large."2715 As in the case of the shell model. we do not have much freedom in choosing the other parameters.," As in the case of the shell model, we do not have much freedom in choosing the other parameters."272 One wav lo decrease the mass requirement in the funnel model is to enhance (he density in the funnel wall relative to (he rest of the ejecta., One way to decrease the mass requirement in the funnel model is to enhance the density in the funnel wall relative to the rest of the ejecta.273 For example. one could imagine that initiallv there was no funnel. and (hat it was the GRB itself that pushed the material aside to form the Tunnel.," For example, one could imagine that initially there was no funnel, and that it was the GRB itself that pushed the material aside to form the funnel."274 It is conceivable that the material pushed aside could have piled up on the walls. giving an enhanced density there.," It is conceivable that the material pushed aside could have piled up on the walls, giving an enhanced density there."275 It is not clear that this mechanism can produce the orders of magnitude density enhancement needed (o reduce (he mass estimate to a reasonable value.Model:, It is not clear that this mechanism can produce the orders of magnitude density enhancement needed to reduce the mass estimate to a reasonable value.:276 In this model. the photoionization occurs indirectly.," In this model, the photoionization occurs indirectly."277 Photons from the GRB travel out to a screen. are scattered. and then irradiate the eas. which is located in the surface lavers of (he supernova ejecta.," Photons from the GRB travel out to a screen, are scattered, and then irradiate the line-emitting gas, which is located in the surface layers of the supernova ejecta."278 Because the eeomelry is very cdillerent [rom (he previous(wo moclels. there are (wo important changes.," Because the geometry is very different from the previoustwo models, there are two important changes."279 First. the time delay of 104 s between the GRB and the line emission measures (he distance io the scattering screen (>10!! em) but not the size of the line-emitting ejecta.," First, the time delay of $10^4$ s between the GRB and the line emission measures the distance to the scattering screen $> 10^{14}$ cm) but not the size of the line-emitting ejecta."280 Therefore. (he estimates given in equation (4) are not valid.," Therefore, the estimates given in equation (4) are not valid."281" Instead. if we assume Chat the ejecta have expanded αἱ speed O.le3.4 for 1047...) s. we estimate that R~Dc1077/1,413,4 em."," Instead, if we assume that the ejecta have expanded at speed $0.1c\beta_{*,-1}$ for $10^4\txg$ s, we estimate that $R\sim282D\sim 10^{13.5}\txg\beta_{*,-1}$ cm."283" Second. the irradiation occurs for a (ime ~10,4; s ratherthan 107 s. and so we expect lo~ 10011."," Second, the irradiation occurs for a time $\sim10^4\txg$ s ratherthan $10^2$ s, and so we expect $t_2\sim100\txg$ ."284 Both changes help to ease some of the constraints., Both changes help to ease some of the constraints.285Galactic plane as measured with extragalactic radio sources is where 1607 7 £10 per cent rans.,"Galactic plane as measured with extragalactic radio sources is where =1607 $^{-2}$ $\pm$ 10 per cent r.m.s.,"286 462.17 (Cleggοἱal.1986)., $^\circ$ \citep{clegg92}.287.. WT5N is al the galactic longitude 81.97. and [rom (7) one has A544 7.," W75N is at the galactic longitude $^\circ$, and from (7) one has =–544 $^{-2}$."288 At the OIL frequency (A=I8 em) the Faraday rotation is x(0.13)?— 17.6 radizc5.4 rad if one assumes maximum uncertainty 30 per cent., At the OH frequency $\lambda=18$ cm) the Faraday rotation is $\times$ $^2$ =– 17.6 $\pm$ 5.4 rad if one assumes maximum uncertainty 30 per cent.289 This is the total Faraday rotation throughout the ealactic disk: W75N is at (he distance of 2 kpe. which is probably about a half of the total effective cistance. and the Faraday rotation to W75N can be a [actor of 2 lower. or 9 racdzE2.7 rad.," This is the total Faraday rotation throughout the galactic disk; W75N is at the distance of 2 kpc, which is probably about a half of the total effective distance, and the Faraday rotation to W75N can be a factor of 2 lower, or –9 $\pm$ 2.7 rad."290 This is a large rotation. about 3 full turns. and a correction for the Faraday rotation to the position angle of the linear polarization could be quite uncertain.," This is a large rotation, about 3 full turns, and a correction for the Faraday rotation to the position angle of the linear polarization could be quite uncertain."291 Therefore il is not posible to determine the direction of (he magnetic field in OIL maser spots., Therefore it is not posible to determine the direction of the magnetic field in OH maser spots.292" Physical parameters of the maser spots can be estimated [rom maser models. which require gas density nj,-10'cm 7. kinetic temperature LOO IX. dust. temperature. 150. Is. and OL abundance 10? (GrayandField1995)."," 	Physical parameters of the maser spots can be estimated from maser models, which require gas density $_{H_2}$ $^7$ $^{-3}$, kinetic temperature 100 K, dust temperature 150 K, and OH abundance $^{-5}$ \citep*{gray95}."293. Such parameters can provide inversion of 1665 MlIz OII transition in a model with FIR line overlap and a velocity gradient of about 0.025 ΑΙ. (CravanclField 1995).., Such parameters can provide inversion of 1665 MHz OH transition in a model with FIR line overlap and a velocity gradient of about 0.025 $^{-1}$ /A.U. \citep*{gray95}. .294 With the magnetic field strength of 10 miligauss the model of GravandField(1995) provides 100 per cent elliptically polarized c components. wilh 7 components suppressed. in agreement wilh results of this paper for W75N. The size of maser spots LO A.U. and molecular hydrogen density 10*em. ? correspond to the mass of maser spots of 2x10. ‘AL.. which is less than the mass of the Earth.," With the magnetic field strength of 10 milligauss the model of \citet*{gray95} provides 100 per cent elliptically polarized $\sigma$ –components, with $\pi$ –components suppressed, in agreement with results of this paper for W75N. The size of maser spots 10 A.U. and molecular hydrogen density $^7$ $^{-3}$ correspond to the mass of maser spots of $\times10^{-7}$ $_{\odot}$, which is less than the mass of the Earth."295 If the maser spots are discrete physical objects dense cold gas condensations surrounded a low density medium — thev should be confined by the external pressure., If the maser spots are discrete physical objects – dense cold gas condensations surrounded a low density medium – they should be confined by the external pressure.296 A gas condensation with the density L0'em* and temperature LOO IX. can be in pressure equilibrium with the gas of density 10?em.* and temperature 107 IX. Llowever. the magnetic pressure in the maser spots. with the magnetic field strength of 10 milligauss. is an order of magnitude higher. ancl can be compensated by (πριΙου or ram pressure of the hot medium (see discussion in a paper by Reidetal. (1987))).," A gas condensation with the density $^{7}$ $^{-3}$ and temperature 100 K can be in pressure equilibrium with the gas of density $^{5}$ $^{-3}$ and temperature $^4$ K. However, the magnetic pressure in the maser spots, with the magnetic field strength of 10 milligauss, is an order of magnitude higher, and can be compensated by turbulent or ram pressure of the hot medium (see discussion in a paper by \citet{reid87}) )."297 Another model of maser spots proposed for Class II methanol masers (Slvshetal.1999). assumed (hat the maser spots are extended gaseous envelopes of solid icv planets orbiting around O. B-stars. outside their HII regions.," Another model of maser spots proposed for Class II methanol masers \citep{slysh99} assumed that the maser spots are extended gaseous envelopes of solid icy planets orbiting around O, B-stars, outside their HII regions."298 OL molecules as well as methanol molecules are continuosly. supplied to the envelope by evaporation of ice from the surface of the planets., OH molecules as well as methanol molecules are continuosly supplied to the envelope by evaporation of ice from the surface of the planets.299 In W75N the ultracompact III region VLAI (Fig., In W75N the ultracompact HII region VLA1 (Fig.300 5) marks the position of the central star with luminosity 14x 10°L. (Mooreetal.1991). which corresponds to the main sequence O9.star with mass 20M..., 5) marks the position of the central star with luminosity $\times10^{5}$ $_{\odot}$ \citep{moore91} which corresponds to the main sequence O9–star with mass $_{\odot}$.301 The largest distance from VLAÀI to a maser spot is about 1000 mas. or 2000 astronomical units.," The largest distance from VLA1 to a maser spot is about 1000 mas, or 2000 astronomical units."302 Atthis distance from a 20M. star the orbital velocity is 3 | which is consistent with the observed. velocity range of maser, Atthis distance from a $_{\odot}$ star the orbital velocity is 3 $^{-1}$ which is consistent with the observed velocity range of maser303cannot be excluded.,cannot be excluded.304 The above described observations and further follow-up observations will be used to derive a coherent timing solution for PSR J1952+2630., The above described observations and further follow-up observations will be used to derive a coherent timing solution for PSR J1952+2630.305 This will provide a more precisely measured sky position. orbital parameters. and values for the orbital eccentricity and the intrinsic spin-down of the pulsar.," This will provide a more precisely measured sky position, orbital parameters, and values for the orbital eccentricity and the intrinsic spin-down of the pulsar."306 This should enable a detailed description of this binary system and constrain its possible formation., This should enable a detailed description of this binary system and constrain its possible formation.307 A precise position would also enable searches for counterparts in X-ray. infrared. and optical wavelengths. although the large distance makes detections challenging.," A precise position would also enable searches for counterparts in X-ray, infrared, and optical wavelengths, although the large distance makes detections challenging."308 Furthermore. detection of Shapiro delay could be possible with further timing observations for high orbital inclinations.," Furthermore, detection of Shapiro delay could be possible with further timing observations for high orbital inclinations."309 Given the already high minimum companion mass derived in this Letter. even a non-detection of the Shapiro delay could provide interesting. more stringent limits on the companion mass and its nature.," Given the already high minimum companion mass derived in this Letter, even a non-detection of the Shapiro delay could provide interesting, more stringent limits on the companion mass and its nature."310 This pulsar is the second pulsar discovered by the global distributed volunteer computing project Einstein?Home (?).., This pulsar is the second pulsar discovered by the global distributed volunteer computing project EinsteinHome \citep{2010Sci...329.1305K}.311 This further demonstrates the value of volunteer computing for discoveries 1n astronomy and other data-driven science., This further demonstrates the value of volunteer computing for discoveries in astronomy and other data-driven science.312 We thank the Einstein@Home volunteers. who made this discovery possible.," We thank the EinsteinHome volunteers, who made this discovery possible."313 The Einstein?Home users whose computers detected the pulsar with the highest significance are VVitaliy. Shiryaev (Moscow. Russia) and Stacey Eastham (Darwen. UK).," The EinsteinHome users whose computers detected the pulsar with the highest significance are Vitaliy Shiryaev (Moscow, Russia) and Stacey Eastham (Darwen, UK)."314 This work was supported by CFI. CIFAR. FORNT. MPG. NAIC. NRAO. NSERC. NSF. NWO. and STFC.," This work was supported by CFI, CIFAR, FQRNT, MPG, NAIC, NRAO, NSERC, NSF, NWO, and STFC."315 Arecibo is operated by the National Astronomy and lonosphere Center under a cooperative agreement with the NSF., Arecibo is operated by the National Astronomy and Ionosphere Center under a cooperative agreement with the NSF.316 This work was supported by NSF grant AST 0807151 to Cornell University., This work was supported by NSF grant AST 0807151 to Cornell University.317 Pulsar research at UBC ts supported by an NSERC Discovery Grant and by the CFI., Pulsar research at UBC is supported by an NSERC Discovery Grant and by the CFI.318 UWM and U. C. Berkeley acknowledge support by NSF grant 0555655., UWM and U. C. Berkeley acknowledge support by NSF grant 0555655.319 ggratefully acknowledges the support of the Max Planck Society., gratefully acknowledges the support of the Max Planck Society.320 acknowledges support from NSF grant AST-0806942., acknowledges support from NSF grant AST-0806942.321(see Dhattacharva&vandenHeuvel 1991 [or a review).,(see \citeauthor{bha91} 1991 for a review).322 The mechanism for this remains unclear. wilh suggestions including decay of crustal fields due to heating (Blonclin&Freese 1986).. burial of the field (Disnovatvi-Ikogan&Ixomberg1974:Romani1990.1993). and decay of core fields due to flux tube expulsion from the superfhiud interior (Srinivasanetal. 1990))).," The mechanism for this remains unclear, with suggestions including decay of crustal fields due to heating \citep{blo86}, burial of the field \citep{bis74,rom90,rom93} and decay of core fields due to flux tube expulsion from the superfluid interior \citep{sri90}) )."323 Ikonar&Bhattacharva(1997). suggested that rapid ohmic decay in the accretion heated crust occurs., \cite{kon97} suggested that rapid ohmic decay in the accretion heated crust occurs.324 On the one hand. the heating reduces the electrical conductivity ancl consequently the ohimic decay (me-scale induces a faster decay of the field.," On the one hand, the heating reduces the electrical conductivity and consequently the ohmic decay time-scale induces a faster decay of the field."325 Ou the other hand. the deposition of matter on top of the crust pushes the original current carrving lavers into deeper aud denser regions where (he higher conductivity slows down the decay (Ixonar&Bhattacharva1997).," On the other hand, the deposition of matter on top of the crust pushes the original current carrying layers into deeper and denser regions where the higher conductivity slows down the decay \citep{kon97}."326. In a class of reeveling models in which the magnetic field decrease is a finction only of the amount aecreted onto the neutron star. it was shown that no model of this class is consistent wilh all available data (Wijers1997).," In a class of recycling models in which the magnetic field decrease is a function only of the amount accreted onto the neutron star, it was shown that no model of this class is consistent with all available data \citep{wij97}."327. The detection of coherent X-ray pulsations with a millisecond period in a handful of LMXBDs (Lamb&Yu2005). is often used in support of the idea of accretion-induced [field decay (Wijnands&vanderIxlis1993)., The detection of coherent X-ray pulsations with a millisecond period in a handful of LMXBs \citep{lam05} is often used in support of the idea of accretion-induced field decay \citep{wij98}.328. However. whether this is evidence simply [or field submersion and spin-up during the accretion disk phase. or for field decay ancl spin-up. remains {ο be established (Ferrario&Wickramasinghe2007).," However, whether this is evidence simply for field submersion and spin-up during the accretion disk phase, or for field decay and spin-up, remains to be established \citep{fer07}."329. We conclude Irom the above. that it is not clear if low field ~105 G field MSPs can be produced in LMXDs.," We conclude from the above, that it is not clear if low field $\sim 10^{8}$ G field MSPs can be produced in LMXBs."330 We pointed out that the field could simply be submerged by (he accreted matter and would then re-emeree later on when accretion stops., We pointed out that the field could simply be submerged by the accreted matter and would then re-emerge later on when accretion stops.331 Zhangetal.(2009) show that there is no evidence for field restructuring and/or decay in accreting magnetic white dwarls., \cite{zha09} show that there is no evidence for field restructuring and/or decay in accreting magnetic white dwarfs.332 This fact was already. known for the isolated magnetic white cwarls., This fact was already known for the isolated magnetic white dwarfs.333 As far as we know. this could apply also to neutron stars.," As far as we know, this could apply also to neutron stars."334ionizing background.,ionizing background.335 In $3 and $4. we present the PDFs and power spectra of various UVBs. respectively.," In \ref{sec:PDF} and \ref{sec:ps} we present the PDFs and power spectra of various UVBs, respectively."336 In $5.. we study the impact of a spatially varying UVB on the Lya forest.," In \ref{sec:forest}, we study the impact of a spatially varying UVB on the $\alpha$ forest."337 Finally. in $6... we summarize our findings and offer conclusions.," Finally, in \ref{sec:conc}, we summarize our findings and offer conclusions."338" We quote all quantities in comoving units. with the exception of flux. and we denote proper units with a prefix ""p."," We quote all quantities in comoving units, with the exception of flux, and we denote proper units with a prefix 'p'."339 We adopt the background cosmological parameters (OQ. O1. Qn. n. as. Ho = (0.72. 0.28. 0.046. 0.96. 0.82. 70 km + . matching the five-year results of the satellite (22).," We adopt the background cosmological parameters $\Omega_\Lambda$, $\Omega_{\rm M}$, $\Omega_b$, $n$, $\sigma_8$ , $H_0$ ) = (0.72, 0.28, 0.046, 0.96, 0.82, 70 km $^{-1}$ $^{-1}$ ), matching the five–year results of the satellite \citep{Dunkley08, Komatsu08}."340 We generate our ionizing flux fields following the procedure described in ?.., We generate our ionizing flux fields following the procedure described in \citet{MD08}.341 We briefly outline the procedure below., We briefly outline the procedure below.342 We begin with halo fields at >=5 and 2=6 generated with the semi-numerical simulationDexM?*.. which has been shown to reproduce the correet number density and clustering properties of halos well into the quasi-linear and non-linear regimes (?:: Fig.," We begin with halo fields at $z=5$ and $z=6$ generated with the semi-numerical simulation, which has been shown to reproduce the correct number density and clustering properties of halos well into the quasi-linear and non-linear regimes \citealt{MF07}; Fig."343 | in 2: Mesinger et al., 1 in \citealt{Dijkstra08}; Mesinger et al.344 in preparation)., in preparation).345 Our simulation boxes are 150Mpe on a side. with a 150/1800 Mpe halo grid cell size.," Our simulation boxes are 150Mpc on a side, with a 150/1800 Mpc halo grid cell size."346 The velocity fields used to perturb the halo field were generated on a lower resolution 900% erid: thus our final halo field resolution is Ar= 150/900 = 0.17 Mpc., The velocity fields used to perturb the halo field were generated on a lower resolution $^3$ grid; thus our final halo field resolution is $\Delta x =$ 150/900 = 0.17 Mpc.347 For each halo field. we create a corresponding UV flux field on a 150° erid (spatial resolution of | Mpc). by performing a halo mass/r7 weighted sum.," For each halo field, we create a corresponding UV flux field on a $^3$ grid (spatial resolution of 1 Mpc), by performing a halo $r^2$ weighted sum."348 Specitically. we compute the flux of ionizing photons (in units of ionizing photons s.+ 7) with where x is the location of the cell of interest. AJ; is the total halo miss. xj is its location. and the factorof (1|z)? converts the factor [x xi[ from comoving into properunits.," Specifically, we compute the flux of ionizing photons (in units of ionizing photons $^{-1}$ $^{-2}$ ) with where ${\bf x}$ is the location of the cell of interest, $M_i$ is the total halo mass, ${\bf x_i}$ is its location, and the factor of $(1+z)^2$ converts the factor $|{\bf x} - {\bf x_i}|^2$ from comoving into properunits."349 We also include a duty parameter through the random variable .[N;(D€). which has a value of | with likelihood DC or O with likelihood 1.DC.," We also include a duty parameter through the random variable $X_i(\DC)$, which has a value of 1 with likelihood $\DC$ or 0 with likelihood $1-\DC$."350" Finally. the factor ej, in eg. CL) "," Finally, the factor $\epsilon_{\rm ion}$ in eq. \ref{eq:sum}) )"351denotes the rate at which ionizing photons are released into the IGM by a dark matter halo per unit mass., denotes the rate at which ionizing photons are released into the IGM by a dark matter halo per unit mass.352 Unless stated otherwise. we assume a fiducial value of Gaz)2A385107/DCEOulul|M.petias which provides a good fit to the observed luminosity functions of Lya emitting galaxies (LAEs) (22222)... and the .=6 Lyman Break galaxies (LBGs) (?)..," Unless stated otherwise, we assume a fiducial value of $\epsilon_{\rm ion}(z) = 3.8 \times 10^{58}/\DC \left[\frac{\Omega_b}{\Omega_{\rm M}} \frac{1}{t_H(z)}\right]\hs\frac{{\rm photons}}{M_{\odot}\hs {\rm s}}.$, which provides a good fit to the observed luminosity functions of $\alpha$ emitting galaxies (LAEs) \citep{Shimasaku06, Kashikawa06, DWH07, SLE07, McQuinn07LAE}, and the $z=6$ Lyman Break galaxies (LBGs) \citep{Bouwens06}."353 However. as we are mostly concerned with the shape of the flux PDF. where applicable we present results in units of the mean flux or intensity. J/0./5.," However, as we are mostly concerned with the shape of the flux PDF, where applicable we present results in units of the mean flux or intensity, $J/\langle J\rangle$."354 In constructing the UVB. we also assume a minimum halo mass able to host stars. A4. and explore mainly two different values for this parameter: 1.6.5107 and 1.4.10° AZ...," In constructing the UVB, we also assume a minimum halo mass able to host stars, $\Mmin$, and explore mainly two different values for this parameter: $1.6\times10^8$ and $1.4\times10^9$ $\Msun$."355 The former value corresponds to a virial temperature of 10 K at these redshifts. and represents the regime of ineffective feedback on atomically cooled halos during reionization (22)..," The former value corresponds to a virial temperature of $^4$ K at these redshifts, and represents the regime of ineffective feedback on atomically cooled halos during reionization \citep{MD08, OGT08}."356 The later value represents the regime of very inefficient star formation inside such small halos. perhaps due to strong radiative and/or mechanical feedback (e.g. 2??? ," The later value represents the regime of very inefficient star formation inside such small halos, perhaps due to strong radiative and/or mechanical feedback (e.g. \citealt{Yepes97, Scannapieco06, PS08}) )."357"The exact value of the higher Mi, was also motivated by the fact that the number density of halos with mass greater than 14.10° AZ. is roughly ten times smaller than the number density of halos with masses greater than 1.6«107 AZ. at 2=5. thus allowing us to compare results at fixed source number density with values of DC= 1.0 and 0.1 in the two models. respectively."," The exact value of the higher $\Mmin$ was also motivated by the fact that the number density of halos with mass greater than $1.4\times10^9$ $\Msun$ is roughly ten times smaller than the number density of halos with masses greater than $1.6\times10^8$ $\Msun$ at $z\approx5$, thus allowing us to compare results at fixed source number density with values of $\DC=$ 1.0 and 0.1 in the two models, respectively."358 Additionally. we generate a flux Ποιά corresponding to the >=5.71 source field from the cosmological hydrodynamical simulation presented in?) (their +=6 reionization model).," Additionally, we generate a flux field corresponding to the $z=5.71$ source field from the cosmological hydrodynamical simulation presented in \citet{TCL08} (their $z=6$ reionization model)."359 This simulation is fixed-grid. 143 Mpe on a side. and includes wescriptions for modeling dark matter. baryons and ionizing shotons (for details see ? and 2).," This simulation is fixed-grid, 143 Mpc on a side, and includes prescriptions for modeling dark matter, baryons and ionizing photons (for details see \citealt{TC07} and \citealt{TCL08}) )."360 The density field was calculated on grid of 0.19 Mpe cells. which resolves the Jeans length in he mean density. ionized intergalactic medium (IGM) by a factor of few. and then smoothed to a cell size of 0.74 Mpe.," The density field was calculated on grid of 0.19 Mpc cells, which resolves the Jeans length in the mean density, ionized intergalactic medium (IGM) by a factor of few, and then smoothed to a cell size of 0.74 Mpc."361 Each wlo’s instantaneous star-formation rate (SFR) is proportional to the instantaneous gas accretion rate. and enters our eq. CL) ," Each halo's instantaneous star-formation rate (SFR) is proportional to the instantaneous gas accretion rate, and enters our eq. \ref{eq:sum}) )"362in place of ye halo's mass. AZ;. with the normalization adjusted accordingly so as to match the mean value of the UVB.," in place of the halo's mass, $M_i$, with the normalization adjusted accordingly so as to match the mean value of the UVB."363 We make use of this flux field. combined with the +=5.71 gas density fields from 18 hydro-simulation. to generate more accurate mean Lya forest flux decrement statistics in S5...," We make use of this flux field, combined with the $z=5.71$ gas density fields from the hydro-simulation, to generate more accurate mean $\alpha$ forest flux decrement statistics in \ref{sec:forest}."364 At these redshifts and scales. our semi-numerically generated density fields somewhat over-predict ye rare voids which dominate the this statistic (Mesinger et al..," At these redshifts and scales, our semi-numerically generated density fields somewhat over-predict the rare voids which dominate the this statistic (Mesinger et al.,"365 in preparation)., in preparation).366 Furthermore. the ray-tracing algorithm from the numerical simulation over-predicts the fluctuations in the flux field. due to an insufficient number of ray splittings (Trac 2008. private communication).," Furthermore, the ray-tracing algorithm from the numerical simulation over-predicts the fluctuations in the flux field, due to an insufficient number of ray splittings (Trac 2008, private communication)."367 This problem only becomes severe following reionization. but since this is the epoch westudy here. we use this aybrid preseription in some of the results below (i.e. we do not use he radiative transfer field from the simulation).," This problem only becomes severe following reionization, but since this is the epoch westudy here, we use this hybrid prescription in some of the results below (i.e. we do not use the radiative transfer field from the simulation)."368 In Fig. l..," In Fig. \ref{fig:pics},"369 we present a 0.74 Mpe thick slice through the density and flux fields based on this 2=5.71 simulation output rom ?.., we present a 0.74 Mpc thick slice through the density and flux fields based on this $z=5.71$ simulation output from \citet{TCL08}.370 The flux fields were calculated assuming Ayu)=30 Mpe. and DC=|.," The flux fields were calculated assuming $\lmfp=30$ Mpc, and $\DC=1$."371 Even for such a moderately high choice of Aj. we can qualitatively see that there is significant inhomogeneity in the flux fields.," Even for such a moderately high choice of $\lmfp$, we can qualitatively see that there is significant inhomogeneity in the flux fields."372 Furthermore. it is evident that the flux and density fields are highly correlated. an issue we will return to in S5..," Furthermore, it is evident that the flux and density fields are highly correlated, an issue we will return to in \ref{sec:forest}."373 We complement these numerical techniques with a relatively simple analytic model., We complement these numerical techniques with a relatively simple analytic model.374 While this approach cannot incorporate all of the physics provided by the semi-numeric and numeric models. we will find it to be helpfulin elucidating the physies of the background radiation field.," While this approach cannot incorporate all of the physics provided by the semi-numeric and numeric models, we will find it to be helpfulin elucidating the physics of the background radiation field."375 As in the remainder of this paper. we will focus on computing two statistical descriptions of the flux field: the PDF and the power spectrum.," As in the remainder of this paper, we will focus on computing two statistical descriptions of the flux field: the PDF and the power spectrum."376 The PDF of the radiation background ./. normalized to its mean value. 0/5. can be computed exactly for randomly. distributed sources. provided that we assume a constant attenuation. length," The PDF of the radiation background $J$ , normalized to its mean value $\langle J\rangle$ , can be computed exactly for randomly distributed sources, provided that we assume a constant attenuation length"377remain eravitationally bouud to the low mass star.,remain gravitationally bound to the low mass star.378 The conditions for photou-pressure blow-out of erains iu the disk around the low mass star can be derived as follows., The conditions for photon-pressure blow-out of grains in the disk around the low mass star can be derived as follows.379 Let 3 be the ratio of photon to eravitational force on a grain in the vicinity of the O star by itself., Let $\beta$ be the ratio of photon to gravitational force on a grain in the vicinity of the O star by itself.380 Photon-pressure-driven blow-out occurs for jJ > 0.5., Photon-pressure-driven blow-out occurs for $\beta$ $>$ 0.5.381 The equivalent ratio of forces for à grain near the low mass star must include its gravity. resulting in a reduction of J approximately iu proportion to the ratio of eravitational forces from the two stars at the position of the erain.," The equivalent ratio of forces for a grain near the low mass star must include its gravity, resulting in a reduction of $\beta$ approximately in proportion to the ratio of gravitational forces from the two stars at the position of the grain."382 It is casily shown that the modified ratio has a value of 0.5 at a distance frou the low mass (mass Af ) star given by where dis its distance from the lieh mass star (mass My)., It is easily shown that the modified ratio has a value of 0.5 at a distance from the low mass (mass $M_2$ ) star given by where $d$ is its distance from the high mass star (mass $M_1$ ).383 The term + is of order 1 aud depends on where the erain is du its orbit., The term $\gamma$ is of order 1 and depends on where the grain is in its orbit.384 Lamy&Perrin(1997) have determined values of JJ for eras of various compositions and sizes around a uiuuber of stars., \citet{Lamy97} have determined values of $\beta$ for grains of various compositions and sizes around a number of stars.385 The values are roughly similar for a given grain size: erain composition enters as a secondary parameter., The values are roughly similar for a given grain size; grain composition enters as a secondary parameter.386 We will use typical values in the following discussion., We will use typical values in the following discussion.387 Taking the values for the 09.5 V star ¢ Oph. grains of radius between 0.01 aud 1/22 have ye 1000 - 10000.," Taking the values for the O9.5 V star $\zeta$ Oph, grains of radius between 0.01 and 1 $\mu$ m have $\beta \approx$ 1000 - 10000."388 We take a typical distance between our high aud low mass stars to be d = 0.2 pe. aud also asstune LO M for the high mass star aud 0.6ML.. for he low mass one.," We take a typical distance between our high and low mass stars to be $d$ = 0.2 pc, and also assume 40 ${\rm M_{\sun}}$ for the high mass star and ${\rm M_{\sun}}$ for the low mass one."389 We then find that the O-star photon xessure on such s1nall eras becomes dominant at a distance of z 110 AU from the low mass star., We then find that the O-star photon pressure on such small grains becomes dominant at a distance of $\approx$ 110 AU from the low mass star.390 Simall eyaius well inside this radius will be held iu orbit bv he gravitational field of the low-mass star., Small grains well inside this radius will be held in orbit by the gravitational field of the low-mass star.391 For larger eras. ο drops rapidly: typically it is < 50 for 10 μπι radius erains.," For larger grains, $\beta$ drops rapidly; typically it is $<$ 50 for 10 $\mu$ m radius grains."392" The corresponding distauce for them where ohoton pressure from the O star overcomes the gravity of the low-mass one is z 500 AU. that is. outside the youndary of a typical protoplanetary disk ίοιο,,Andrews&Willizuuus 2005)."," The corresponding distance for them where photon pressure from the O star overcomes the gravity of the low-mass one is $\approx$ 500 AU, that is, outside the boundary of a typical protoplanetary disk \citep[e.g.,][]{Andr05}."393. Therefore. the larger erains are likely ο remain bound to the low nass star throughout the disk. where they will coutiuue to be ground. down by collisions.," Therefore, the larger grains are likely to remain bound to the low mass star throughout the disk, where they will continue to be ground down by collisions."394" That is. the erains making up the ""conet tail” probably originate frou collisional cascades in the outer reeious of the circumstellar disk."," That is, the grains making up the “comet tail” probably originate from collisional cascades in the outer regions of the circumstellar disk."395 These regious lave previously been cleared of gas by photoevaporation aud the remaining solid particles are settling toward the disk mud-plain aud starting to assemble iuto larger bodies (see Throop&Bally(2005) for details). leacling to ecucration of these eraims idu a vigorous episode of collisional cascades.," These regions have previously been cleared of gas by photoevaporation and the remaining solid particles are settling toward the disk mid-plain and starting to assemble into larger bodies (see \citet{Thro05} for details), leading to generation of these grains in a vigorous episode of collisional cascades."396 IHTowever. eiveu typical sizes of protoplanctary disks. this process is likev to be effective only up to a WAN erain size of al)out a nuücron radius.," However, given typical sizes of protoplanetary disks, this process is likely to be effective only up to a maximum grain size of about a micron radius."397 The tails are conrposed. primarily ¢of erains of & O.OL gan to z 1 jaa du radius (where the ower limit is set by the model fits to their surface brighttess profiles: Balog ct al., The tails are composed primarily of grains of $\approx$ 0.01 $\mu$ m to $\approx$ 1 $\mu$ m in radius (where the lower limit is set by the model fits to their surface brightness profiles; Balog et al.398 2006)., 2006).399 We have modified the upper mass limits derived by Balogotal.(2006) to reflect the approximate eraiu size Iunit of l yan. This approximate nuit results from the teudeney of larecr grains to remai1 eravitationally bound to the low-luass star., We have modified the upper mass limits derived by \citet{Balo06} to reflect the approximate grain size limit of 1 $\mu$ m. This approximate limit results from the tendency of larger grains to remain gravitationally bound to the low-mass star.400 From the mass loss rates and typical disk masses iu Section Ld. we can estimate an approximate timescale of 10—109 vrs for which this phenomenon is visible.," From the mass loss rates and typical disk masses in Section 4.1, we can estimate an approximate timescale of $10^5 - 10^6$ yrs for which this phenomenon is visible."401 Therefore. the phenomenon might be quite common. however we see ouly three cases in our GTO survey of about 20 O stars (there are additional cases around three O-stars in the Wh region. Koenig et al.," Therefore, the phenomenon might be quite common, however we see only three cases in our GTO survey of about 20 O stars (there are additional cases around three O-stars in the W5 region, Koenig et al."402 in preparation) sugecsting that this is probably a short lived rather rare phenomenon (oulv about L/L of the observed O-stars have cometary structures in their ucighborhood)., in preparation) suggesting that this is probably a short lived rather rare phenomenon (only about 1/4 of the observed O-stars have cometary structures in their neighborhood).403 We present IIST/NICMIOS Pao images and IRS spectra of cometary structures detected in Spitzer/MIPS 2| yan images., We present HST/NICMOS $\alpha$ images and IRS spectra of cometary structures detected in /MIPS 24 $\mu$ m images.404 We estimate an upper lint to the amount of gas iu the comets’ disk and tail and find that the eas-to-dust mass ratio is much lower than the value observed iu the ISAL: the tails are esseutially eas free., We estimate an upper limit to the amount of gas in the comets' disk and tail and find that the gas-to-dust mass ratio is much lower than the value observed in the ISM: the tails are essentially gas free.405 Using this new observation we are able to coustrain the flow velocity aud thus the mass loss rate., Using this new observation we are able to constrain the flow velocity and thus the mass loss rate.406 The new iiass oss rates allow us to estimate the timescale ou which the shenomenon occurs (10106 vr)., The new mass loss rates allow us to estimate the timescale on which the phenomenon occurs $10^5 - 10^6$ yr).407 The short timescale avors photocvaporation models predicting quick removal of gas from the outer parts of the disk., The short timescale favors photoevaporation models predicting quick removal of gas from the outer parts of the disk.408 The “comet tails” are produced. from the outer regions of the disks. where areor erains collide at au elevated rate generating second eeneration dust.," The “comet tails” are produced from the outer regions of the disks, where larger grains collide at an elevated rate generating second generation dust."409 These suall erains are then ejected due o photon pressure from the nearby O-star., These small grains are then ejected due to photon pressure from the nearby O-star.410 The SED of the sources shows excess cussion between 3 aud 5 ju. in agreement with the IRS low resolution spectra.," The SED of the sources shows excess emission between 3 and 8 $\mu$ m, in agreement with the IRS low resolution spectra."411 This cussion mdicates that there is an iuncr disk that survives the plotoevaporation process for longer thu LO? years. as predicted by the photoevaporation models.," This emission indicates that there is an inner disk that survives the photoevaporation process for longer than $10^5$ years, as predicted by the photoevaporation models."412 The authors thank the anouvinous referce for cohunents and sugeestious which improved the paper., The authors thank the anonymous referee for comments and suggestions which improved the paper.413 We also thank Robert Iiug for providing the VLT data for the source iu NGC 2211., We also thank Robert King for providing the VLT data for the source in NGC 2244.414 This work is based on, This work is based on415fiial gaaxy would be minimized if simall fragments/cdwarl galaxies were absett from tlie galaxys envlrornent.,final galaxy would be minimized if small fragments/dwarf galaxies were absent from the galaxy's environment.416 Note. though. hat sucha situation is less likely for a ceutalcD galaxy.," Note, though, that such a situation is less likely for a central cD galaxy."417 Late mergers a'e expected to p'ovide a significant population of metaI-poor globular clusters originating iu the »ogenior galaxies. predomiuautly meta—rich systenis «:'ould be explaiued if the progenitors were κ.OrLIC but οuster-poor (see also Gebjardt Ixissle-Patig 1999).," Late mergers are expected to provide a significant population of metal–poor globular clusters originating in the progenitor galaxies, predominantly metal–rich systems could be explained if the progenitors were gas–rich but cluster–poor (see also Gebhardt Kissler-Patig 1999)."418 StrippitD>oO Call lesilt in the yrelerential re10val of metal-poor cluste‘s if the metal-j»oor population is more extended than the netal-‘ich one., Stripping can result in the preferential removal of metal–poor clusters if the metal–poor population is more extended than the metal–rich one.419 siuce stripping will act 1lore efficiently ou the more exended population. it could 'esult i systels hat are biased towards veh metallicitses.," Since stripping will act more efficiently on the more extended population, it could result in systems that are biased towards high metallicities."420 Iu this context. it is interesting to note hat the halo lel oL IC 1051 appears trucated beyonc  30 kpc (Jorgeusenetal.1992).," In this context, it is interesting to note that the halo light of IC 4051 appears truncated beyond $\sim$ 30 kpc \citep{jor92}."421. Preferentlally metal-rich globular €Uster svstenuis can be accommodated by adjusting the existiug scena‘ios., Preferentially metal–rich globular cluster systems can be accommodated by adjusting the existing scenarios.422" An exclusively metal-rich elobular clister population. should oue be discovered. would. probabv be most easily accomxlated in a siiele collapse mocel — au ""in situ galaxy formation sceiario with rapid star—lormaion preceding cluster formation."," An exclusively metal–rich globular cluster population, should one be discovered, would probably be most easily accommodated in a single collapse model – an “in situ” galaxy formation scenario with rapid star–formation preceding cluster formation."423 We have shown that the color distribution of globular clusters in NCC 3311 is normal [or bright elliptical galaxies., We have shown that the color distribution of globular clusters in NGC 3311 is normal for bright elliptical galaxies.424 It. is bi-inodal with peaks at V—/~0.91 aud. 1.09. correspoudiug to metallicity peaks at around [Fe/H~—1.D nali —(0.75 (the precise values being dependent. on the choice of the conversio1 ‘elation between color aud metallicity).," It is bi–modal with peaks at $V-I\sim0.91$ and $1.09$, corresponding to metallicity peaks at around $\sim -1.5$ and $-0.75$ (the precise values being dependent on the choice of the conversion relation between color and metallicity)."425 This range of metallicities is uormal for bright elliptical gaaxles. a result. whicἩ contradicts au earlier claim that NCC 3311 might host an extremely metal-ricl elobular cluster syseinn.," This range of metallicities is normal for bright elliptical galaxies, a result which contradicts an earlier claim that NGC 3311 might host an extremely metal–rich globular cluster system."426 We suggest that i is Worth revisitiug the globular cluster system of NGC 3923 whose globular cluster system was reported to be very re lontje basis of observations made ou the same run which produced the NCC 331:| results., We suggest that it is worth revisiting the globular cluster system of NGC 3923 whose globular cluster system was reported to be very red on the basis of observations made on the same run which produced the NGC 3311 results.427 Although the evidence orexclusively wetal-rich globular cluster systems has been weakenec. there are still cases of galaxies whose elobular cluster systems may lack a significant metal-poor component.," Although the evidence for metal–rich globular cluster systems has been weakened, there are still cases of galaxies whose globular cluster systems may lack a significant metal–poor component."428 We have briely discussed the implications of such systems for our wucderstaucing of elobular cluster aud galaxy formation and have concluded that. with some adjustiueuts. they cau be explained uuder existitD>& sCeLaLlos.," We have briefly discussed the implications of such systems for our understanding of globular cluster and galaxy formation and have concluded that, with some adjustments, they can be explained under existing scenarios."429 We thank Johu Huch‘a for lis help and useful suggestions aud Duucan Forbes. Carl Cuiulimair and Wen Freeman for their contributions.," We thank John Huchra for his help and useful suggestions and Duncan Forbes, Carl Grillmair and Ken Freeman for their contributions."430 This work was supported by HST eraut. GO.0655 National Science Foundation grant number AST0000732 aud Faculty Research fuus from the University of Califoruia. Santa Cruz.," This work was supported by HST grant GO.06554.01-95A, National Science Foundation grant number AST9900732 and Faculty Research funds from the University of California, Santa Cruz."431 The solution to this equation is (for more details see 1908) with the normalization N adjusted to give |T|?=1., The solution to this equation is (for more details see JS08) with the normalization $\cal N$ adjusted to give ${|T|}^2=1$.432" Apparently the thus constructed nulling weights depend on which (£4,£5,5) combination is considered and with respect to which cosmological parameter we optimize the information content."," Apparently the thus constructed nulling weights depend on which $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ combination is considered and with respect to which cosmological parameter we optimize the information content."433" In this paper the default cosmological parameter to optimize is O4, and we choose for each (i,j) combination the (£4,£2,£3) combination which maximizes FG»."," In this paper the default cosmological parameter to optimize is $\Omega_{\rm m}$, and we choose for each $(i,j)$ combination the $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ combination which maximizes $F_{\rm o}^{(ij)}$."434" However one needs to be aware that this serves only as a clear choice of a (δι,£5,€3) combination and is not necessarily the best in terms of information preservation considering all angular frequency bins and all cosmological parameters."," However one needs to be aware that this serves only as a clear choice of a $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ combination and is not necessarily the best in terms of information preservation considering all angular frequency bins and all cosmological parameters."435" To show which triangle shapes and sizes contain more information, we plot F? against the (21,&,£5) triangle shape and size for four typical (i,j) combinations in 1."," To show which triangle shapes and sizes contain more information, we plot $F_{\rm o}^{(ij)}$ against the $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ triangle shape and size for four typical $(i,j)$ combinations in $\,$."436" In the left panel, the nulled information FGD contained in different triangles with a common shortest side length £;=171 is plotted against o, which is the angle opposite to £j."," In the left panel, the nulled information $F_{\rm o}^{(ij)}$ contained in different triangles with a common shortest side length $\bar{\ell}_1=171$ is plotted against $\alpha$, which is the angle opposite to $\bar{\ell}_1$."437" Due to our logarithmic binning in angular frequency, only eight (61,£5,£5) combinations with £j,=171 can form triangles."," Due to our logarithmic binning in angular frequency, only eight $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ combinations with $\bar{\ell}_1=171$ can form triangles."438 One sees that the more elongated triangles (small α) contain much more Fisher information than the almost equilateral triangles (large a)., One sees that the more elongated triangles (small $\alpha$ ) contain much more Fisher information than the almost equilateral triangles (large $\alpha$ ).439" The small separation between the 3rd and the 4th points from the left is caused by the degeneracy of different triangle shapes with respect to o, e.g. two equal and very long side lengths can result in the same value of α as two shorter side lengths with a length difference close to the length of the shortest side length."," The small separation between the 3rd and the 4th points from the left is caused by the degeneracy of different triangle shapes with respect to $\alpha$, e.g. two equal and very long side lengths can result in the same value of $\alpha$ as two shorter side lengths with a length difference close to the length of the shortest side length."440" The right panel shows the distribution of the Fisher information contained in one (£i,£5,(4) bin over the triangle size."," The right panel shows the distribution of the Fisher information contained in one $(\bar{\ell}_1, \bar{\ell}_2, \bar{\ell}_3)$ bin over the triangle size."441" When the redshift in consideration is higher, the peak of the information distribution moves to higher angular frequencies."," When the redshift in consideration is higher, the peak of the information distribution moves to higher angular frequencies."442 The figure suggests that most information comes from high redshifts and small angular scales., The figure suggests that most information comes from high redshifts and small angular scales.443" To explore the sensitivity of nulling weights on the choice of the cosmological parameter, we construct seven sets of weight functions, each optimizing the information content in terms of one parameter."," To explore the sensitivity of nulling weights on the choice of the cosmological parameter, we construct seven sets of weight functions, each optimizing the information content in terms of one parameter."444" For all (i,7) combinations we find that the nulling weights are not very sensitive to the choice of parameter."," For all $(i,j)$ combinations we find that the nulling weights are not very sensitive to the choice of parameter."445" As an example, the weights for (i,j)=(1,2) are shown in 2."," As an example, the weights for $(i,j) = (1,2)$ are shown in $\,$."446" This result is rather surprising at first sight, since for different parameters the distribution of information (contained in the bispectrum) over redshift bins is quite different."," This result is rather surprising at first sight, since for different parameters the distribution of information (contained in the bispectrum) over redshift bins is quite different."447" However, such insensitivity suggests that the shapes of nulling weights are already strongly constrained under our construction scheme."," However, such insensitivity suggests that the shapes of nulling weights are already strongly constrained under our construction scheme."448" One constraint is, evidently, the nulling condition."," One constraint is, evidently, the nulling condition."449" Moreover, considering the fact that we optimize the nulling weights for each (i,j) combination with respect to the information content they preserve, we have already required the shapes of these first order nulling weights to be as smooth as possible."," Moreover, considering the fact that we optimize the nulling weights for each $(i,j)$ combination with respect to the information content they preserve, we have already required the shapes of these first order nulling weights to be as smooth as possible."450" The fact that these two conditions have already imposed strong constraints on the nulling weights also suggests that nulling weights can be robustly and efficiently constructed, i.e. it is not critical to construct the “best” nulling weights."," The fact that these two conditions have already imposed strong constraints on the nulling weights also suggests that nulling weights can be robustly and efficiently constructed, i.e. it is not critical to construct the “best” nulling weights."451" What the nulling technique “nulls” is the GGI signal Baar, so the GGI/GGG ratio is the most direct quantification of its performance."," What the nulling technique “nulls” is the GGI signal $B_{\rm GGI}$, so the GGI/GGG ratio is the most direct quantification of its performance."452 We plot the modeled GGI and GGG bispectra before and after nulling in Fig.3.," We plot the modeled GGI and GGG bispectra before and after nulling in $\,$."453 The original GGI signal is shown in the left panels by dashed lines., The original GGI signal is shown in the left panels by dashed lines.454 For comparison the GGG signals are shown as solid curves., For comparison the GGG signals are shown as solid curves.455 The results are shown for equilateral triangle configurations for the convenience of presenting., The results are shown for equilateral triangle configurations for the convenience of presenting.456" One sees that when the redshift bin number j and/or k increase, the changes in GGG and GGI signals are different, which shows the expected different redshift dependence."," One sees that when the redshift bin number $j$ and/or $k$ increase, the changes in GGG and GGI signals are different, which shows the expected different redshift dependence."457 For all redshift bin combinations the GGI signal is modeled to be subdominant to the GGG signal., For all redshift bin combinations the GGI signal is modeled to be subdominant to the GGG signal.458" In the nulled measures shown in the right panels, the GGI/GGG ratio is suppressed by a factor of 10 over all angular scales, which reflects the success of the nulling technique."," In the nulled measures shown in the right panels, the GGI/GGG ratio is suppressed by a factor of 10 over all angular scales, which reflects the success of the nulling technique."459" We further evaluate the performance of the nulling technique by looking at the constraining power of cosmic shear bispectrum tomography on cosmological parameters, as well as the biases caused by the GGI systematics before and after nulling."," We further evaluate the performance of the nulling technique by looking at the constraining power of cosmic shear bispectrum tomography on cosmological parameters, as well as the biases caused by the GGI systematics before and after nulling."460White dwarfs are the final remnants of low- and intermediate-mass stars.,White dwarfs are the final remnants of low- and intermediate-mass stars.461 About of main-sequence stars will end their evolutionary pathways as white dwarfs and. hence. the study of the white dwarf population provides details about the late stages of the life of the vast majority of stars.," About of main-sequence stars will end their evolutionary pathways as white dwarfs and, hence, the study of the white dwarf population provides details about the late stages of the life of the vast majority of stars."462 Since white dwarfs are long-lived objects. they also constitute useful objects to study the structure and evolution of our Galaxy (Liebert et al.," Since white dwarfs are long-lived objects, they also constitute useful objects to study the structure and evolution of our Galaxy (Liebert et al."463 2005a: Isern et al., 2005a; Isern et al.464 2001)., 2001).465 For instance. the initial- mass relationship (IFMR). which connects the properties of a white dwarf with those of its main-sequence progenitor. is of paramount importance for different aspects in. modern astrophysics.," For instance, the initial-final mass relationship (IFMR), which connects the properties of a white dwarf with those of its main-sequence progenitor, is of paramount importance for different aspects in modern astrophysics."466 It 1s required as an input for determining. the ages of globular clusters and their distances. for studying the chemical evolution of galaxies. and also to understand the properties of the Galactic population of white dwarfs.," It is required as an input for determining the ages of globular clusters and their distances, for studying the chemical evolution of galaxies, and also to understand the properties of the Galactic population of white dwarfs."467 Despite its relevance. this relationship is still poorly constrained. both from the theoretical and the observational points of view.," Despite its relevance, this relationship is still poorly constrained, both from the theoretical and the observational points of view."468 The first attempt to empirically determine the initial-final mass relationship was undertaken by Weidemann(1977).. who also provides a recent review on this subject (Weidemann 2000).," The first attempt to empirically determine the initial-final mass relationship was undertaken by \cite{wei77}, who also provides a recent review on this subject (Weidemann 2000)."469 It is still not clear how this function depends on the mass and metallicity of the progenitor. its angular momentum. or the presence of a strong magnetic field.," It is still not clear how this function depends on the mass and metallicity of the progenitor, its angular momentum, or the presence of a strong magnetic field."470 The total age of a white dwarf can be expressed as the sum of its cooling time and the main-sequence lifetime of its progenitor., The total age of a white dwarf can be expressed as the sum of its cooling time and the main-sequence lifetime of its progenitor.471 The latter depends on the metallicity of the progenitor of the white dwarf. but it cannot be determined from observations of single white dwarfs.," The latter depends on the metallicity of the progenitor of the white dwarf, but it cannot be determined from observations of single white dwarfs."472 This is because white dwarfs have such strong surface gravities that gravitational settling operates very efficiently in their atmospheres. and any information about their progenitors (e.g. metallicity) is lost in the very early evolutionary stages of the cooling track.," This is because white dwarfs have such strong surface gravities that gravitational settling operates very efficiently in their atmospheres, and any information about their progenitors (e.g. metallicity) is lost in the very early evolutionary stages of the cooling track."473 Moreover. the evolution during the AGB phase of the progenitors is essential in determining the size and composition of the atmospheres of the resulting white dwarfs. since the burning processes that take place in H and He shells determine their respective thicknesses and their detailed chemical compositions. which are crucial ingredients for determining the evolutionary cooling times.," Moreover, the evolution during the AGB phase of the progenitors is essential in determining the size and composition of the atmospheres of the resulting white dwarfs, since the burning processes that take place in H and He shells determine their respective thicknesses and their detailed chemical compositions, which are crucial ingredients for determining the evolutionary cooling times."474 A promising approach to circumvent the problem. and also to directly test the initial-final mass relationship. is to study white dwarfs for which external constraints are available.," A promising approach to circumvent the problem, and also to directly test the initial-final mass relationship, is to study white dwarfs for which external constraints are available."475 This is the case of white dwarfs in open and globular clusters (Ferrario et al., This is the case of white dwarfs in open and globular clusters (Ferrario et al.476 2005. Dobbie et al.," 2005, Dobbie et al."477 2006) or in non-interacting binaries. for Instance. common proper motion pairs (Wegner 1973. Oswalt et al.," 2006) or in non-interacting binaries, for instance, common proper motion pairs (Wegner 1973, Oswalt et al."478 1988)., 1988).479 Focusing on the latter. it is sound to assume that the members of a common proper motion par were born simultaneously and with the same chemical composition.," Focusing on the latter, it is sound to assume that the members of a common proper motion pair were born simultaneously and with the same chemical composition."480 Since the components are well separated (100 to 1000 AU). mass exchange between them ts unlikely and it can be considered that they have evolved as isolated stars.," Since the components are well separated (100 to 1000 AU), mass exchange between them is unlikely and it can be considered that they have evolved as isolated stars."481 Thus. important information of the white dwarf. such as its total age or the metallicity of the progenitor. can be inferred from the study of the companion.," Thus, important information of the white dwarf, such as its total age or the metallicity of the progenitor, can be inferred from the study of the companion."482 In particular. if the companion ts an EF. G or K type star the metallicity can be derived with high accuracy from detailed spectral analysis.," In particular, if the companion is an F, G or K type star the metallicity can be derived with high accuracy from detailed spectral analysis."483 On the other hand. the age can be obtained using different methods.," On the other hand, the age can be obtained using different methods."484 In particular. we will use stellar isochrones when the star is moderately evolved. or the X-ray luminosity if the star is very close to the ZAMS.," In particular, we will use stellar isochrones when the star is moderately evolved, or the X-ray luminosity if the star is very close to the ZAMS."485 The purpose of this work is to present our spectroscopic analysis of both members of some common proper motion pars containing a white dwarf. and the semi-empirical intial-final mass relationship that we have derived from this study.," The purpose of this work is to present our spectroscopic analysis of both members of some common proper motion pairs containing a white dwarf, and the semi-empirical intial-final mass relationship that we have derived from this study."486 The paper is organized as follows., The paper is organized as follows.487 In $2 we present the observations done so far and describe the data reduction., In 2 we present the observations done so far and describe the data reduction.488 Section 3 1s devoted to discuss the classification and the analysis of the observed white dwarfs. whereas in $4 we present the analysis of the companions.," Section 3 is devoted to discuss the classification and the analysis of the observed white dwarfs, whereas in 4 we present the analysis of the companions."489 This is followed by $5 where we present our main results and finally in $6 we elaborate our conclusions., This is followed by 5 where we present our main results and finally in 6 we elaborate our conclusions.490radius. (hen we need to correct for only enclosing part of the cooling gas in the aperture.,"radius, then we need to correct for only enclosing part of the cooling gas in the aperture."491 As described in BAIL. this is accomplished by multiplving the total My by the fraction of the X-ray luminosity projected into the aperture.," As described in BMI, this is accomplished by multiplying the total ${\dot{M}}_X$ by the fraction of the X-ray luminosity projected into the aperture."492 Here we use the value of Mx for the distributed model (q=1). corrected to the size of the aperture.," Here we use the value of ${\dot{M}}_X$ for the distributed model ), corrected to the size of the aperture."493 I we had used the model with no mass drop-out (y=). the median predicted value of Ma hardly changes. although for the most X-ray Iuminous galaxies in (he sample (often the largest). the predicted My could be a factor of two higher.," If we had used the model with no mass drop-out ), the median predicted value of ${\dot{M}}_X$ hardly changes, although for the most X-ray luminous galaxies in the sample (often the largest), the predicted ${\dot{M}}_X$ could be a factor of two higher."494 The various values of AM and other relevant derived quantities. including the N-vav huninositw. £y. and the X-ray. temperature. Ly are eiven in Table 3.," The various values of ${\dot{M}}$ and other relevant derived quantities, including the X-ray luminosity, $L_X$, and the X-ray temperature, $T_X$ are given in Table 3."495 There is a complete set of X-ray fluxes and luminosities for (his sample. so we beein bv comparing these to the OVI data.," There is a complete set of X-ray fluxes and luminosities for this sample, so we begin by comparing these to the OVI data."496 The nominal prediction was that there would be a connection between M and the detection of OVI emission. but we examined other relationships as well.," The nominal prediction was that there would be a connection between ${\dot{M}}_X$ and the detection of OVI emission, but we examined other relationships as well."497 Some of those relationships investigated were between OVI and Ly. Ly/Ty. and Tx. where no strong correlations were found and no strong relationships were predicted.," Some of those relationships investigated were between OVI and $L_X$, $L_X/T_X$, and $T_X$, where no strong correlations were found and no strong relationships were predicted."498 However. there seems {ο be a correlation between the OVI and My. as seen in Fieure 25.," However, there seems to be a correlation between the OVI and ${\dot{M}}_X$, as seen in Figure 25."499 We see that none of the six galaxies with the lowest values of Ma. have any OVI emission. vet (he significance of this correlation is difficult to quantilv.," We see that none of the six galaxies with the lowest values of ${\dot{M}}_X$ have any OVI emission, yet the significance of this correlation is difficult to quantify."500 We would like to use the Ixaplan-Meier estimator lor the analysis of censored data Nelson 19386).. but an underlving assumption is that the data are censored randomly.," We would like to use the Kaplan-Meier estimator for the analysis of censored data \citep{isobe86}, but an underlying assumption is that the data are censored randomly."501 Here. the upper limits (censored data) are not randomly. distributed. but preferentially occur at the lower values of M.," Here, the upper limits (censored data) are not randomly distributed, but preferentially occur at the lower values of ${\dot{M}}_X$."502" Also. we have introduced ""possible"" detections. which is difficult to incorporate in statistical schemes."," Also, we have introduced “possible” detections, which is difficult to incorporate in statistical schemes."503 Given these challenges. we can divide the sample into three bins by Ma. using our scoring for detections. possible detections. and upper limits.," Given these challenges, we can divide the sample into three bins by ${\dot{M}}_X$, using our scoring for detections, possible detections, and upper limits."504 We find that for galaxies with the lowest M. values. 1/8 have OVI. while 4.5/8 of the highest My objects have OVI (and 3.5/8 of the intermediate objects have OVI).," We find that for galaxies with the lowest ${\dot{M}}_X$ values, 1/8 have OVI, while 4.5/8 of the highest ${\dot{M}}_X$ objects have OVI (and 3.5/8 of the intermediate objects have OVI)."505 Using Poisson statistics. the joint probability that of the lowest My objects. 1/8 (or lewer) have OVI while 4.5/8 (or more) of the highest Mx objects have OVI would oceur by chance of the time confidence level).," Using Poisson statistics, the joint probability that of the lowest ${\dot{M}}_X$ objects, 1/8 (or fewer) have OVI while 4.5/8 (or more) of the highest ${\dot{M}}_X$ objects have OVI would occur by chance of the time confidence level)."506 The correlation is probably a bit stronger than this value since none of the lowest six My object have either an OVI detection or possible detection., The correlation is probably a bit stronger than this value since none of the lowest six ${\dot{M}}_X$ object have either an OVI detection or possible detection.507 For a Ixendall's T lest or a Spearman's p test for the whole sample (treating upper limits and detections equally). the significance improves (ο confidence. with the higher significance if the very poor upper limit object. NGC 3585 is eliminated from the sample.," For a Kendall's $\tau$ test or a Spearman's $\rho$ test for the whole sample (treating upper limits and detections equally), the significance improves to confidence, with the higher significance if the very poor upper limit object, NGC 3585 is eliminated from the sample."508 Conservatively. we conclude that the correlation exists at the confidence level when using the," Conservatively, we conclude that the correlation exists at the confidence level when using the"509Even on the long timescales probed by our time-averaged monitoring data. the flux-dependent behaviour of the broad iron line is complex.,"Even on the long timescales probed by our time-averaged monitoring data, the flux-dependent behaviour of the broad iron line is complex."510 For these data. the line flux increases more-or-less proportionally with the continuum flux. resulting. in a roughly constant equivalent width over a decade range of flux.," For these data, the line flux increases more-or-less proportionally with the continuum flux, resulting in a roughly constant equivalent width over a decade range of flux."511 The December 1996 long-look data also show line flux increasing with continuum flux. although the relation is not directly proportionate. so that the equivalent width decreases with flux.," The December 1996 long-look data also show line flux increasing with continuum flux, although the relation is not directly proportionate, so that the equivalent width decreases with flux."512 In fact. the line fluxes measured in December 1996 are consistently larger than the corresponding fluxes measured from the long-term monitoring data.," In fact, the line fluxes measured in December 1996 are consistently larger than the corresponding fluxes measured from the long-term monitoring data."513 The anti-correlation of line equivalent width and continuum flux in December 1996 timescales is also in contrast to the result of (Wangetal.1999).. who report a positive correlation during an observation in 1994.," The anti-correlation of line equivalent width and continuum flux in December 1996 timescales is also in contrast to the result of \cite {Wang}, who report a positive correlation during an observation in 1994."514 The discrepancy between the iron line behaviour in the December 1996 and long-term monitoring data might be explained if there is additional short-term variability in the iron line which is not simply related to the continuum flux., The discrepancy between the iron line behaviour in the December 1996 and long-term monitoring data might be explained if there is additional short-term variability in the iron line which is not simply related to the continuum flux.515 For example. Vaughan Edelson (2001) show that in the Seyfert | MCG-6-30-15. the broad iron line flux varies significantly but independently of short term continuum variations.," For example, Vaughan Edelson (2001) show that in the Seyfert 1 MCG-6-30-15, the broad iron line flux varies significantly but independently of short term continuum variations."516 One possibility is that the iron line flux tracks the long-term variations in the continuum flux Qvhich are being probed to some extent with the long-term monitoring data). but responds only weakly to the short-term variations whieh are observed during the December 1996 long-look observation.," One possibility is that the iron line flux tracks the long-term variations in the continuum flux (which are being probed to some extent with the long-term monitoring data), but responds only weakly to the short-term variations which are observed during the December 1996 long-look observation."517 A number of theoretical papers have been written to explain why the iron line flux may not vary linearly with the continuum flux (teg Mattetal.1993.. Nayakshin&Kazanas 2002.. Ballantyne&Ross 20023). often involving ionised discs but these models have so far been largely untroubled by data.," A number of theoretical papers have been written to explain why the iron line flux may not vary linearly with the continuum flux (eg \ncite{Matt}, \ncite{Nayakshin}, \ncite{Ballantyne}) ), often involving ionised discs but these models have so far been largely untroubled by data."518 We are aquiring more long-term monitoring data. sampling a broader range of long-term flux. variations. to determine whether the iron line does follow the continuum on long timescales and to provide some constraints for theoretical models.," We are aquiring more long-term monitoring data, sampling a broader range of long-term flux variations, to determine whether the iron line does follow the continuum on long timescales and to provide some constraints for theoretical models."519 Our observations of NGC 40531 do not support the correlation between the photon index and the reflected fraction. # in the model as reported from spectroscopy of a sample of Seyfert galaxies (Zdziarski.Lu, Our observations of NGC 4051 do not support the correlation between the photon index and the reflected fraction $R$ in the model as reported from spectroscopy of a sample of Seyfert galaxies \cite{Zdziarski} .520bifiski&Smith, From Fig.5211999)... From Fig. + itis obvious that the reflected fraction remains below /?=1 even for the softest states of the source., \ref{dec_cont} it is obvious that the reflected fraction remains below $R=1$ even for the softest states of the source.522 There is also no evidence for this correlation in theAXE spectra of NGC 5506 (Lamer.Ut-tley&M'Hardy 2000)., There is also no evidence for this correlation in the spectra of NGC 5506 \cite{Lamer2000}.523.. We therefore suggest that the reported correlation does not apply to the variations of photon index and reflected fraction in a given object., We therefore suggest that the reported correlation does not apply to the variations of photon index and reflected fraction in a given object.524 During our monitoring campaign the primary continuum 2-10 keV photon index. L'. varies strongly with photon flux from 1.60 at the lowest flux levels to 2.35 at the highest fluxes.," During our monitoring campaign the primary continuum 2-10 keV photon index, $\Gamma$, varies strongly with photon flux from 1.60 at the lowest flux levels to 2.35 at the highest fluxes."525 The same correlation is also observed on the shorter time-scales of the December 1996 long look., The same correlation is also observed on the shorter time-scales of the December 1996 long look.526 Flux-L correlations have been observed before in NGC 4051 (Matsuokaetal.1987). and other Seyfert galaxies teg NGC 4151. Perolaetal.19869) but. as in the present paper. it has only recently been possible to disentangle the effects of reflection and variations of the primary X-ray spectrum (teg Lamer.Uttley&M'Hardy2000.. Chiangetal.2000.. Leeetal.20001).," $\Gamma$ correlations have been observed before in NGC 4051 \cite{Matsuoka} and other Seyfert galaxies (eg NGC 4151, \ncite{Perola}) ) but, as in the present paper, it has only recently been possible to disentangle the effects of reflection and variations of the primary X-ray spectrum (eg \ncite{Lamer2000}, \ncite{Chiang}, \ncite{Lee2000}) )."527 The slope-luminosity correlation is. often. explained by stronger cooling of the accretion disk corona during episodes of high thermal seed photon flux from the accretion disk itself (eg Pietrini&Krolik1995.. Malzac&Jourdain2000)).," The slope-luminosity correlation is often explained by stronger cooling of the accretion disk corona during episodes of high thermal seed photon flux from the accretion disk itself (eg \ncite{Pietrini}, \ncite{Malzac}) )."528 Haardt.Maraschi&Ghisellini(1997) have calculated luminosity — spectral index relations in Compton cooled accretion disk coronae., \scite{Haardt} have calculated luminosity – spectral index relations in Compton cooled accretion disk coronae.529 For a compact. pair dominated corona they predict spectral index variations of AL~0.3 for luminosity variations by more than a factor of 20.," For a compact, pair dominated corona they predict spectral index variations of $\Delta\Gamma\sim 0.3$ for luminosity variations by more than a factor of 20."530 However the variations seen here exceed their predictions and imply. in their scenario. a non-pair dominated corona.," However the variations seen here exceed their predictions and imply, in their scenario, a non-pair dominated corona."531 In certain regimes this model predicts a positve correlation of 2-10 keV flux and spectral hardness. which is not observed in NGC 4051.," In certain regimes this model predicts a positve correlation of 2-10 keV flux and spectral hardness, which is not observed in NGC 4051."532" Pietrint&Krolik(1995) point out that the spectral index depends almost solely on the ratio of seed photon compactness ἐς and hot plasma heating rate compactness ἐν with a=1L60./0,):17.", \scite{Pietrini} point out that the spectral index depends almost solely on the ratio of seed photon compactness $l_s$ and hot plasma heating rate compactness $l_h$ with $\alpha=1.6(l_s/l_h)^{1/4}$.533 The observed spectral indices in NGC 4051 then correspond to ἐς{δε=0.02..0.4., The observed spectral indices in NGC 4051 then correspond to $l_s/l_h=0.02..0.4$.534 Examination of fig 7 shows that the change of spectral index with flux is not linear., Examination of fig \ref{lineflux} shows that the change of spectral index with flux is not linear.535 The rate of increase of spectral index with flux is very rapid at low fluxes but decreases at the highest fluxes where the spectral index approaches an asymptotic level., The rate of increase of spectral index with flux is very rapid at low fluxes but decreases at the highest fluxes where the spectral index approaches an asymptotic level.536" This saturation of the ""spectral index/flux relationship has been known for some time: eg the saturation was clearly visible in our early RXTE monitoring observations of MCG-6-30-15 and was reported by M'Hardy.Papadakis&Uttley(1998) where it was suggested that the relationship might derive from the combination of a constant spectrum hard component. and a steeper spectrum variable component."," This saturation of the `spectral index/flux' relationship has been known for some time; eg the saturation was clearly visible in our early RXTE monitoring observations of MCG-6-30-15 and was reported by \scite{mch98} where it was suggested that the relationship might derive from the combination of a constant spectrum hard component, and a steeper spectrum variable component."537 Saturation of the spectral index/flux relationship was again reported in MCG-6-30-15 by Shihetal.(2002) from a long ASCA observation., Saturation of the spectral index/flux relationship was again reported in MCG-6-30-15 by \scite{Shih} from a long ASCA observation.538 In MCG-6-30-15 both the long term RXTE monitoring and short term ASCA observations agree that the saturation level of the spectral index is ~2.1 (see fig 7 of MHardy.Uttley1998 and tig 8 of Shihetal. 2002))., In MCG-6-30-15 both the long term RXTE monitoring and short term ASCA observations agree that the saturation level of the spectral index is $\sim2.1$ (see fig 7 of \ncite{mch98} and fig 8 of \ncite{Shih}) ).539 However the monitoring observations cover a wider flux and spectral range and show variation of the spectral index between 1.65 and 2.05 whereas the continuous ASCA observation only shows an index variation between 1.9 and 2.1., However the monitoring observations cover a wider flux and spectral range and show variation of the spectral index between 1.65 and 2.05 whereas the continuous ASCA observation only shows an index variation between 1.9 and 2.1.540 Similarly in NGC 4051 (fig 7)) we see that the monitoring observations cover a wider flux and spectral range than the December 1996 long look and. as with MCG-6-30- the resultant time-averaged spectral index/flux relationship is smoother.," Similarly in NGC 4051 (fig \ref{lineflux}) ) we see that the monitoring observations cover a wider flux and spectral range than the December 1996 long look and, as with MCG-6-30-15, the resultant time-averaged spectral index/flux relationship is smoother."541 We note. however. that although the lowest spectral index so far measured in the RATE monitoring observations is about the same in both NGC 4051 and MCG-6-30-15. the saturation level is ~ 2.4in NGC 4051 compared to ~2.1 in MCG-6-30-15.," We note, however, that although the lowest spectral index so far measured in the RXTE monitoring observations is about the same in both NGC 4051 and MCG-6-30-15, the saturation level is $\sim2.4$ in NGC 4051 compared to $\sim2.1$ in MCG-6-30-15."542 In M'Hardy.Papadakis&Uttley(1998). we suggested that the torus might be the source of the possible hard. constant. component.," In \scite{mch98} we suggested that the torus might be the source of the possible hard, constant, component."543 However the very hard spectral component found in the May 1998 very low state. which represents an upper limit to the torus contribution. was removed before producing the spectral index/flux. relationship (fig 7).," However the very hard spectral component found in the May 1998 very low state, which represents an upper limit to the torus contribution, was removed before producing the spectral index/flux relationship (fig \ref{lineflux}) )."544 Thus. if we wish to retain the two-component spectral model. we require a different location for," Thus, if we wish to retain the two-component spectral model, we require a different location for"545ssysteni would have a combined mass at least. 18% larger than Mes. but might not be a SN In progenitor if the conrpanion turus out to be an ONe WD instead of a CO WD (Carcia-Derroetal.1997).,"system would have a combined mass at least $43\%$ larger than $M_{Ch}$, but might not be a SN Ia progenitor if the companion turns out to be an ONe WD instead of a CO WD \citep{garcia-berro97:stars}."546.. Considering all the evidence. the most likely possibility bv far is that the companion of bbe the stellar remnant of| a supernova explosion. either a NS or a DII.," Considering all the evidence, the most likely possibility by far is that the companion of be the stellar remnant of a supernova explosion, either a NS or a BH."547 We note that our estimated. distance range places this object closer to the Solar System thu any other mown NS (Posseltetal.2007)., We note that our estimated distance range places this object closer to the Solar System than any other known NS \citep{posselt07:M7}.548. The puzzle of the uature of the companion of ccannot be solved with the observations that we preseut in this paper. and must be the subject of future work.," The puzzle of the nature of the companion of cannot be solved with the observations that we present in this paper, and must be the subject of future work."549 It is. however. uterestiug to speculate about the possibilities.," It is, however, interesting to speculate about the possibilities."550 For our best-fit value of A4 (0.92NL... Alp is below 1.86 (thelargestmasuredmassforaNS.seeLattimerM...&Prakash2007:Niceetal.2008) for PIT. which has a raudoia likelihood of 39%.," For our best-fit value of $M_{A}$ $0.92\,\mathrm{M_{\odot}}$ ), $M_{B}$ is below $1.86\,\mathrm{M_{\odot}}$ \citep[the largest551masured mass for a NS, see][]{lattimer07:NS_EOS,nice08:no_massive_NS} for $i \geq 67^{\circ}$, which has a random likelihood of $39\%$."552 Iu this case. the mass of the WD and the circularity of the orbit would place the system in the class of interiiediate-nass WDINS binaries (secTable1inStairs2001).," In this case, the mass of the WD and the circularity of the orbit would place the system in the class of `intermediate-mass' WD+NS binaries \citep[see Table 1 in][]{stairs04:pulsars_binary_systems}."553 Curent models for binary stellar evolution predict that hese svstenis undergo unstable mass transfer during a conimion-envelope phase (Stairs2001).. which results iu he NS becoming a nülkdv recvcled millisecond. pulsar (MSDP).," Current models for binary stellar evolution predict that these systems undergo unstable mass transfer during a common-envelope phase \citep{stairs04:pulsars_binary_systems}, which results in the NS becoming a mildly recycled millisecond pulsar (MSP)."554 This scenario would explain the high inferred uass for the companion of 5128.. well above he average for a NS. as a byproduct of the mass trausfer yroeess (Lattimer&Prakash2007).," This scenario would explain the high inferred mass for the companion of , well above the average for a NS, as a byproduct of the mass transfer process \citep{lattimer07:NS_EOS}."555". Df the NS is iudecd an MSP. we expect the magnetic field to be low (~10? Cass). the pulsar lifetime to be large (1 Cor). aud the opening augle to be wide (possiblymanytensofdegrees,seePhinney&I&ulkarui 1991)..."," If the NS is indeed an MSP, we expect the magnetic field to be low $\sim10^{9}$ Gauss), the pulsar lifetime to be large $\sim1$ Gyr), and the opening angle to be wide \citep[possibly many tens of degrees, see][]{phinney94:binary_and_ms_pulsars}."556 Under these couditious. the prospects for detecting such a nearby. AISP are good. although we note that the cooling time for the WD is also of the order of Cyr (see Section 3.3)). aud the MSP nuelt have lost a large part of its maenuetic field.," Under these conditions, the prospects for detecting such a nearby MSP are good, although we note that the cooling time for the WD is also of the order of Gyr (see Section \ref{subsec:spectrum}) ), and the MSP might have lost a large part of its magnetic field."557" For Af,=0.92M. and inclination angles below 67° (6114( rvaudom Likelihood). the companion would probably be a stellar mass DIT."," For $M_{A}=0.92\,\mathrm{M_{\odot}}$ and inclination angles below $67^{\circ}$ $61\%$ random likelihood), the companion would probably be a stellar mass BH."558 In this case. we would not expect αν sind of direct emission frou it.," In this case, we would not expect any kind of direct emission from it."559 No radio or N-rav source appears at this location iu any of the indexed. astronomical catalogs. but this does rot preclude the existence of a faint counterpart to the companion of5128.," No radio or X-ray source appears at this location in any of the indexed astronomical catalogs, but this does not preclude the existence of a faint counterpart to the companion of."560. This part of the sky iis onlv been shallowly surveved for MSPs in the racio., This part of the sky has only been shallowly surveyed for MSPs in the radio.561 The most stringent lanits are probably from the Crecn Bank 110 ft 350 MIITIZ survey. which lad a nominal sensitivity of 12-15 ταν (Saveretal.1997)..," The most stringent limits are probably from the Green Bank 140 ft 350 MHz survey, which had a nominal sensitivity of 12-15 mJy \citep{sayer97:GBT_northern_sky_pulsar_survey}."562" In the ravs. this location has never been observed with orNewton, aud a faint nearby source could have casily escaped detection by the All-Sky Survey."," In the X-rays, this location has never been observed with or, and a faint nearby source could have easily escaped detection by the All-Sky Survey."563 A systematic cross-correlation between stellar sources 1u SDSS (ποιος all the WDs from E06) aud the catalogues was performed by Agüerosetal.(2009).. who ound no counterpart to15128.," A systematic cross-correlation between stellar sources in SDSS (including all the WDs from E06) and the catalogues was performed by \citet{agueros09:ROSAT_SDSS_Stars}, who found no counterpart to."564. One interesting implication of the companion of ος a NS or DII is that the system should have received some kind of kick from the SN explosion., One interesting implication of the companion of being a NS or BH is that the system should have received some kind of kick from the SN explosion.565 Iu xinciple. the spatial velocity of cca be determuned by the temporal average of the jon-eravitational Doppler shifts iu the spectra. which eives the radial component. and the proper motion iieasured by SDSS (jp=04491228vr.1). which gives the colponcut on the plane of the sky.," In principle, the spatial velocity of can be determined by the temporal average of the non-gravitational Doppler shifts in the spectrum, which gives the radial component, and the proper motion measured by SDSS $\mu=0.049\,\mathrm{mas\,yr^{-1}}$ ), which gives the component on the plane of the sky."566 Taking our distance estimate. the proper motion translates iuto a transverse velocity of 11aus1," Taking our distance estimate, the proper motion translates into a transverse velocity of $11^{+2}_{-4}\,\mathrm{km\,s^{-1}}$."567 The radial coniponent is more difficult to estimate. because the coustaut conponent to the RV curve (24.=289+L6knis| in our fit. soe Section 3.2)) is the combination of the true radial velocity and the eravitational redshift of the WD.," The radial component is more difficult to estimate, because the constant component to the RV curve $\gamma_{A}=-28.9\pm4.6\,\mathrm{km\,s^{-1}}$ in our fit, see Section \ref{subsec:orbit}) ) is the combination of the true radial velocity and the gravitational redshift of the WD."568 The value of the eravitational redshift depends on the WD amass. which we cannot measure with accuracy (see Section 3.3)). and the WD radius. which is also model depeudenut.," The value of the gravitational redshift depends on the WD mass, which we cannot measure with accuracy (see Section \ref{subsec:spectrum}) ), and the WD radius, which is also model dependent."569 A detailed estimate of the eravitational redshitt for lis outside the scope of this work. but for a massive ~LAL. WD. we expect it to be of the order of οας+ (Weener&Reid1991).. which would require a radial velocity avound 120kus|.," A detailed estimate of the gravitational redshift for is outside the scope of this work, but for a massive $\sim1\,\mathrm{M_{\odot}}$ WD, we expect it to be of the order of $90\,\mathrm{km\,s^{-1}}$ \citep{wegner91:gravitational_redshift_WDs}, which would require a radial velocity around $-120\,\mathrm{km\,s^{-1}}$."570 The total spatial velocity would then be ~120128.1. mostly in the radial direction. which is comparable to the measured kicks for pulsars im binary svstems (Waneetal.2006)..," The total spatial velocity would then be $\sim120\,\mathrm{km\,s^{-1}}$, mostly in the radial direction, which is comparable to the measured kicks for pulsars in binary systems \citep{wang06:NS_kicks}."571 Regardless of what the nature of the companion to tturus out to be. the svsteii is clearly very interesting from απ astroplivsica ut oof view. and new observations should vield exciting results in the near future.," Regardless of what the nature of the companion to turns out to be, the system is clearly very interesting from an astrophysical point of view, and new observations should yield exciting results in the near future."572" From the measured orbital parameters. the separation of thecomponents must be small. with values of the semimajor axis lareer than 0.02 AU only for Ax17"" (0.0003 AU at ;/=607)."," From the measured orbital parameters, the separation of thecomponents must be small, with values of the semimajor axis larger than $0.02$ AU only for $i \leq 17^\circ$ $0.0093$ AU at $i = 60^\circ$ )."573 Tf pulsations from a NS companion were detected. this could allow for a sjenificaut measurement of the Shapiro delay. as in PSR (715 (vauStratenetal. 2001)..," If pulsations from a NS companion were detected, this could allow for a significant measurement of the Shapiro delay, as in PSR $-$ 4715 \citep{vanstraten01:Shapiro_delay}. ."574 The mereine time Paterge of the ssvsteni due to GW radiation is <511!M Aber with the uncertainty im this upper hut arising frou the uneertainty in the cetermination of A4.," The merging time $t_{Merge}$ of the system due to GW radiation is $\leq575511^{+342}_{-141}$ Myr, with the uncertainty in this upper limit arising from the uncertainty in the determination of $M_{A}$."576 For the canonical inclination angle ¢=GOP. fij44; bocones 267|e Ab.," For the canonical inclination angle $i = 60^\circ$, $t_{Merge}$ becomes $267^{+165}_{-70}$ Myr."577 We note that only three WD|NS binaries with fangge<fnanue Were kuown previously 2001:Immetal.200 D.. aud lias a shorter period aud taferge than anv of them.," We note that only three WD+NS binaries with $t_{Merge} < t_{Hubble}$ were known previously \citep{stairs04:pulsars_binary_systems,kim04:Ns_WD_Mergers}, and has a shorter period and $t_{Merge}$ than any of them."578 The discovery of a fourth object iu this class should prompt a revision of the estimated WD|NS mereecr rates and their expected contribution to the CAV background (simet200 1).., The discovery of a fourth object in this class should prompt a revision of the estimated WD+NS merger rates and their expected contribution to the GW background \citep{kim04:Ns_WD_Mergers}.579 The ultimate fate of His unclear., The ultimate fate of is unclear.580 Kinectal.(2007) lave proposed the nereine of massive WDs onto NSs as a scenario for he origin of lone-duration eanunaray bursts without an accolupanving SN like CRB 060611 (GalYametal.2006:DellaValleetal. 2006).. but theoretical simulations or the final accretion pliase in this kind of eveuts lave τος been performed vet.," \citet{king07:WD+NS_GRBs} have proposed the merging of massive WDs onto NSs as a scenario for the origin of long-duration gamma-ray bursts without an accompanying SN like GRB 060614 \citep[][]{gal-yam06:GRB060614,dellavalle06:GRB060614}, but theoretical simulations for the final accretion phase in this kind of events have not been performed yet."581 The coalescence of a DII with a nassive WD may also lead to a eanunaray burst (Frver 1999).. but the observational signature depends again onthe details of the final accretion phase. which arexoorlv uuderstood.," The coalescence of a BH with a massive WD may also lead to a gamma-ray burst \citep{fryer99:WD_BH_Mergers}, , but the observational signature depends again onthe details of the final accretion phase, which arepoorly understood."582 More exotic astroplivsical trausieuts. such as mw be observed by preseut (PTF.Rauetal. and fortheonuüug (ee.LSST.Ivezicetal.2008) svioptie survevs. cannot be discarded as outcomes.," More exotic astrophysical transients, such as may be observed by present \citep[PTF,][]{rau09:PTF_science} and forthcoming \citep[e.g. LSST,][]{ivezic08:LSST} synoptic surveys, cannot be discarded as outcomes."583"example, if [και=30% then for A1689, the observed (central) value r/ryir=0.072 for the effective Einstein radius can be compared with the predicted value (at the same projected surface density equal to the critical lensing density) of r/rvir=0.064 (with adiabatic compression in the satellites), and the previous pure dark-matter prediction (i.e., without adiabatic compression) of r/rvir= 0.045.","example, if $\fs=30\%$ then for A1689, the observed (central) value $r/\rv = 0.072$ for the effective Einstein radius can be compared with the predicted value (at the same projected surface density equal to the critical lensing density) of $r/\rv = 0.064$ (with adiabatic compression in the satellites), and the previous pure dark-matter prediction (i.e., without adiabatic compression) of $r/\rv = 0.045$ ."584" For A1703, the observed r/ryvir=0.070 can be compared with the theoretical r/ryir=0.070 (with adiabatic compression) and r/ryir=0.052 (without)."," For A1703, the observed $r/\rv = 0.070$ can be compared with the theoretical $r/\rv = 0.070$ (with adiabatic compression) and $r/\rv = 0.052$ (without)."585" For Cl0024-17 the corresponding numbers are r/ryir=0.100 compared to r/rvir=0.081 (with) and r/ryir=0.064 (without); for RXJ1347, r/rvir=0.109 compared to r/rvir=0.100 (with) and r/ryir=0.086 (without)."," For Cl0024-17 the corresponding numbers are $r/\rv = 0.100$ compared to $r/\rv = 0.081$ (with) and $r/\rv = 0.064$ (without); for RXJ1347, $r/\rv = 0.109$ compared to $r/\rv = 0.100$ (with) and $r/\rv586= 0.086$ (without)."587" While the observed clusters still have slightly high Einstein radii compared to the typical expected cluster profile, the theoretical scatter in cyir together with the observational errors make the theoretical and observational predictions consistent with each other."," While the observed clusters still have slightly high Einstein radii compared to the typical expected cluster profile, the theoretical scatter in $\cv$ together with the observational errors make the theoretical and observational predictions consistent with each other."588" As we have shown, our results depend only weakly on feat, aS long as it is within a reasonable range."," As we have shown, our results depend only weakly on $\fs$, as long as it is within a reasonable range."589 The results also depend slightly on other assumed properties of the satellites., The results also depend slightly on other assumed properties of the satellites.590" We illustrate this for A1689, fixing fa.=30% and adjusting the host concentration accordingly in each case."," We illustrate this for A1689, fixing $\fs=30\%$ and adjusting the host concentration accordingly in each case."591" We find that lowering the satellite cyir to 3 at Zsat=2 decreases the predicted r/ryir by 14%, while raising cvi; to 5 increases it by 12%."," We find that lowering the satellite $\cv$ to 3 at $\zs=2$ decreases the predicted $r/\rv$ by $14\%$ , while raising $\cv$ to 5 increases it by $12\%$."592" Assuming cyir=3 at Zsat=3 raises the predicted r/rvir by 3%, while cvi=5 at Zar=1 lowers it by 1396, all compared to our standard case of Cvir=4 at z;4=2."," Assuming $\cv=3$ at $\zs=3$ raises the predicted $r/\rv$ by $3\%$, while $\cv=5$ at $\zs=1$ lowers it by $13\%$, all compared to our standard case of $\cv=4$ at $\zs=2$."593" Finally, if we assume that only 50% (rather than 10096) of the baryons in the galactic satellites cooled and condensed before their halos were stripped, i.e., in equation (5)) we use half the cosmic fraction for f», then the predicted r/ryi; is reduced by 9% for A1689 and feat=30%."," Finally, if we assume that only $50\%$ (rather than $100\%$ ) of the baryons in the galactic satellites cooled and condensed before their halos were stripped, i.e., in equation \ref{eq:Mf}) ) we use half the cosmic fraction for $f_b$, then the predicted $r/\rv$ is reduced by $9\%$ for A1689 and $\fs=30\%$."594 We have demonstrated that dark matter compression due to baryonic cooling inside galaxy halos can in turn lead to tidal stripping of these galactic halos closer to the center of the galaxy cluster in which they reside., We have demonstrated that dark matter compression due to baryonic cooling inside galaxy halos can in turn lead to tidal stripping of these galactic halos closer to the center of the galaxy cluster in which they reside.595" Even if the baryons are later redistributed within the cluster by feedback, a substantial effect remains due to the early adiabatic compression."," Even if the baryons are later redistributed within the cluster by feedback, a substantial effect remains due to the early adiabatic compression."596 This effect can explain the high central mass concentration of clusters in lensing observations., This effect can explain the high central mass concentration of clusters in lensing observations.597" Our scenario, in which only the inner 10-2096 of the virial radius is significantly modified, is consistent with weak lensing measurements at larger radii that find low cluster halo concentrations (e.g.,Mandelbaumetal.2008).."," Our scenario, in which only the inner $10$ $20\%$ of the virial radius is significantly modified, is consistent with weak lensing measurements at larger radii that find low cluster halo concentrations \citep[e.g.,][]{Hirata}."598 We have adopted a number of simplifying approximations inshowing the existence of the effect., We have adopted a number of simplifying approximations inshowing the existence of the effect.599 Hydrodynamical simulations that avoid overcooling of the baryons at the cluster core are necessary in order to testour proposed mechanism in quantitative detail., Hydrodynamical simulations that avoid overcooling of the baryons at the cluster core are necessary in order to testour proposed mechanism in quantitative detail.600" Nevertheless, our simplified treatment has demonstrated the general"," Nevertheless, our simplified treatment has demonstrated the general"601"The effect of this is considered in the Appendix, where it is shown that the G run has a systematically lower concentration compared to the HACC run over the same mass scale by about (Figure A11,, left panel).","The effect of this is considered in the Appendix, where it is shown that the G run has a systematically lower concentration compared to the HACC run over the same mass scale by about (Figure \ref{fig:cmsys}, left panel)."602" To compensate for this minor underestimate, we rescale concentrations obtained from the wCDM runs by a factor of 1.05, checking for correctness by comparing against the fit obtained for the reference cosmology."," To compensate for this minor underestimate, we rescale concentrations obtained from the $w$ CDM runs by a factor of $1.05$, checking for correctness by comparing against the fit obtained for the reference cosmology."603 Figure 5 shows the variation of the c—v relation with respect to the best-fit WMAPS cosmology., Figure \ref{fig:meancwcdm} shows the variation of the $c-\nu$ relation with respect to the best-fit WMAP5 cosmology.604 The mean c—v relation variesby about -- over the currently allowed wCDM cosmological parameter range., The mean $c-\nu$ relation variesby about $\pm$ over the currently allowed $w$ CDM cosmological parameter range.605" Note that v already accounts for some of the cosmology dependence of the c—M relation, so a part of the variation is actually hidden."," Note that $\nu$ already accounts for some of the cosmology dependence of the $c-M$ relation, so a part of the variation is actually hidden."606" Since we have already found that expressing c as a function of v explains the redshift evolution of NFW halo profiles, we illustrate the cosmology dependence using the c—v relation in place of the c—M relation."," Since we have already found that expressing $c$ as a function of $\nu$ explains the redshift evolution of NFW halo profiles, we illustrate the cosmology dependence using the $c-\nu$ relation in place of the $c-M$ relation."607 Table 3 shows the approximate difference between the (corrected) mean c—M relation seen in each of the wCDM runs compared to the mean c—M relation obtained for the reference ACDM cosmology., Table \ref{table_wcdm} shows the approximate difference between the (corrected) mean $c-M$ relation seen in each of the $w$ CDM runs compared to the mean $c-M$ relation obtained for the reference $\Lambda$ CDM cosmology.608" Note that although most of the variation in the c—M relation is in the overall amplitude, the slope also changes for some of the models (e.g., M003, M012)."," Note that although most of the variation in the $c-M$ relation is in the overall amplitude, the slope also changes for some of the models (e.g., M003, M012)."609" Interestingly, we find that some of the models show no variation compared to the reference, although these models differ across the range of cosmological parameters."," Interestingly, we find that some of the models show no variation compared to the reference, although these models differ across the range of cosmological parameters."610" For example, M014 and ΜΟΙ7 both have lower og compared to the reference model, but show essentially no variation — parameters other than og are clearly also active."," For example, M014 and M017 both have lower $\sigma_8$ compared to the reference model, but show essentially no variation – parameters other than $\sigma_8$ are clearly also active."611" The standard deviation of the concentration distribution, on the other hand, changes in the same way as the mean, leaving the ratio σε/ο almost universal (Dolagetal. 2004).."," The standard deviation of the concentration distribution, on the other hand, changes in the same way as the mean, leaving the ratio $\sigma_c/c$ almost universal \citep{dolag04}. ."612" Figure 6 shows that the o,/c varies by <5% over the range of wCDM cosmologies.", Figure \ref{fig:sigcwcdm} shows that the $\sigma_c/c$ varies by $< 5\%$ over the range of $w$ CDM cosmologies.613 Semianalytical ‘toy models’ based on Press-Schechter arguments, Semianalytical `toy models' based on Press-Schechter arguments614mainlv X-ray selected. favoring selection of objects wilh SEDs peaked at higher Irequencies. whereas the high luminosity objects were mainly radio selected. (hus favoring the selection ol objects with SEDs peaked at lower [requencies (Anton&Browne2005).,"mainly X-ray selected, favoring selection of objects with SEDs peaked at higher frequencies, whereas the high luminosity objects were mainly radio selected, thus favoring the selection of objects with SEDs peaked at lower frequencies \citep{ant05}."615". The correlation between radio luminosity and gamma-ray dominance (which in our analvsis would be characterized by 0,,) . originally seen by Fossatie£a£.(1998) is also seen here. after more than doubling the sample size."," The correlation between radio luminosity and gamma-ray dominance (which in our analysis would be characterized by $\alpha_{og}$ ), originally seen by \citet{fos98} is also seen here, after more than doubling the sample size."616" Here. because of the way the broad band indices ave defined. à lower value of à, corresponds to greater ganuna-ray dominance."," Here, because of the way the broad band indices are defined, a lower value of $\alpha_{og}$ corresponds to greater gamma-ray dominance."617 This correlation is further evidence for a blazar sequence defined by. relative importance of the ECS model at higher himinosities., This correlation is further evidence for a blazar sequence defined by relative importance of the ECS model at higher luminosities.618" We have also examined a possible correlation between all parameters individually ancl the gamma-ray variability (Nolane£af.2003) and find only a weak anti-correlation between eanmma-ray spectral index. o4. and variability (as parameterized by Nolanetαἱ,(2003):: see Table 1)."," We have also examined a possible correlation between all parameters individually and the gamma-ray variability \citep{nol03} and find only a weak anti-correlation between gamma-ray spectral index, $\alpha_{\gamma}$ , and variability (as parameterized by \citet{nol03}; see Table 1)."619 This max be related (o a spectral hysteresis effect observed by [or a sample of 26 particularly bright and well observed blazars., This may be related to a spectral hysteresis effect observed by \citet{nan07} for a sample of 26 particularly bright and well observed blazars.620" These authors find that during gamma rav. [lares the spectral index first tends to flatten with increasing Πας, and (hen returns (o a steeper index as the flare ends."," These authors find that during gamma ray flares the spectral index first tends to flatten with increasing flux, and then returns to a steeper index as the flare ends."621 Thus. it is plausible that if a blazar is more highlv variable there will be an increasing chance of observing it while its spectral index is flatter (han lor other blazars that are not @amima ray variables.," Thus, it is plausible that if a blazar is more highly variable there will be an increasing chance of observing it while its spectral index is flatter than for other blazars that are not gamma ray variables."622 Alter taking into account statistical tests and Monte Carlo analvsis. we findthe following:," After taking into account statistical tests and Monte Carlo analysis, we findthe following:"623The discovery of N-rayv (7) iux later optical atorelows (?) of eanuma-ray bursts led to identification of ezaunia-rav burst host ealaxies (?)..,The discovery of X-ray \cite{1997IAUC.6576....1C} and later optical afterglows \cite{1997IAUC.6584....1G} of gamma-ray bursts led to identification of gamma-ray burst host galaxies \cite{1997IAUC.6588....1G}.624 Such rost galaxies We‘re expectec in the cosmological model. since LOS of he theories identified CRBs with some. perhaps ex101110 stages of stell evolution.," Such host galaxies were expected in the cosmological model, since most of the theories identified GRBs with some, perhaps extreme, stages of stellar evolution."625 Ilowever. various plivsica inodels of GREs gave ciffercut xedictioas regardiic the location of GRBs with respect to the host galaxies.," However, various physical models of GRBs gave different predictions regarding the location of GRBs with respect to the host galaxies."626 T ithe frauework of the collapsar iuocl. and im general all nodels tha relate GRBs to the final stages of ονο]ιtion of nasslive stars (7).. one expecs that GRBs are Ua in the star formuneOo reeious.," In the framework of the collapsar model, and in general all models that relate GRBs to the final stages of evolution of massive stars \cite{Colgate}, one expects that GRBs are found in the star forming regions."627Oo In the secoud class of iioels where bursts are associated with nergers of conipact object biniaries stich associations are χο obvious., In the second class of models where bursts are associated with mergers of compact object binaries such associations are not obvious.628 Compact object πα. may live for quite a loic time before they merece. aud given the possibilitv that they could have high velocities they may travel away from the place they were formed.," Compact object binaries may live for quite a long time before they merge, and given the possibility that they could have high velocities they may travel away from the place they were formed."629 The distribution of compac object niergers can be found using stellar popuation codes., The distribution of compact object mergers can be found using stellar population codes.630" The main problem of this approaclois that such codes contain a number of poorly kuown ]xwainueters, which nav affect the results."," The main problem of this approach is that such codes contain a number of poorly known parameters, which may affect the results."631 Oue of the most duportant paraneter is the lick velocity a newly born οςmupact object receives at birth., One of the most important parameter is the kick velocity a newly born compact object receives at birth.632 The distribution of double neutron star svstenis around galaxies has beei calculated. for a few types of ealaxies by. Blooni etal. C2).., The distribution of double neutron star systems around galaxies has been calculated for a few types of galaxies by Bloom etal. \cite*{1999MNRAS.305..763B}.633 Bulilk et al., Bulik et al.634 (7) calculated such. distributin for the case of a massive galaxy like the Afillky Way. aud for the case of eniptyv space. sine four different kick velocity distributions.," \cite*{BBZ99}635 calculated such distribution for the case of a massive galaxy like the Milky Way, and for the case of empty space, using four different kick velocity distributions."636 These studies considered oulv ünaries coutaiine neutrou stars., These studies considered only binaries containing neutron stars.637 We used an asstuuption (?:7) +hat all supernovae lead to formation of a LAL. neutron star.," We used an assumption \cite{BBZ99,BBZ99-Rome} that all supernovae lead to formation of a $1.4\,M_\odot$ neutron star."638 However. unergers of binaries containiug a black hoe are now more favored for CRBs.," However, mergers of binaries containing a black hole are now more favored for GRBs."639 One reaπο1 for this is enereeties. GRDB9901223 had an equivale itdsOropic enocrev release of 10?! eres.," One reason for this is energetics, GRB990123 had an equivalent isotropic energy release of $10^{54}$ ergs."640 The energetic requinenienuts eo dow iwheu considering that the relativistic outfiow fro the central fireball is not isotropic mut beamect., The energetic requirements go down when considering that the relativistic outflow from the central fireball is not isotropic but beamed.641 Ou the other haud not all kinetic cucrey oeji the ouflow can be converted iuto gala ravs. and ierefore he ererectic requireineut for the GRB central eneiue will eo 1p.," On the other hand not all kinetic energy in the outflow can be converted into gamma rays, and therefore the energetic requirement for the GRB central engine will go up."642 Therefore it is inportaut to investigate je. distributioji around galaxies of binarics coutainiug lack holeμα, Therefore it is important to investigate the distribution around galaxies of binaries containing black holes.643 Iu this paper we exteud the results of (2).. to iuclude 1C Case ο COMict object binaries contaimineg black holes.," In this paper we extend the results of \cite{BBZ99}, to include the case of compact object binaries containing black holes."644 Tn section 2 woelescribe the model for population svutliesis and galactic potential used in this paper. in section 3 we xeseut the restIts. ancl we stuumarize this work in section 1.," In section 2 we describe the model for population synthesis and galactic potential used in this paper, in section 3 we present the results, and we summarize this work in section 4."645 We use the population svuthesis code described iu detail in Delezvüsski Bulik (?).., We use the population synthesis code described in detail in Belczyńsski Bulik \cite*{BB1998}. .646 Within this model we assume, Within this model we assume647strong silicatefeatures. while large compact aggregates of (he same composition ancl mass show weaker silicatefeatures (IXolokolova.etal.,"strong silicate–features, while large compact aggregates of the same composition and mass show weaker silicate–features \citep{Kolokolova}."6482007).. Furthermore. the degree of linear polarization [or aggregates mav slightly varv with their porosity and/or volatile content (INimuraetal.2006).," Furthermore, the degree of linear polarization for aggregates may slightly vary with their porosity and/or volatile content \citep{Kimura06}."649. In this study. we examine how well the standard model can explain the available observational data of the DI event by calculating color temperature. silicate.feature strength. and polarization lor aggregates in the framework of the standard model.," In this study, we examine how well the standard model can explain the available observational data of the DI event by calculating color temperature, silicate–feature strength, and polarization for aggregates in the framework of the standard model."650 We then discuss the advantage of the standard model over the ΑΕΠ model even when comparing with the results of Sugitaetal.(2005) who analvzed the same STO data as Ixadonoetal.(2007) (hereafter. we refer Kadonoetal.(2007) ancl Sugitaetal.(2005). as to STO-Ix and 5. respectivelv).," We then discuss the advantage of the standard model over the KFH model even when comparing with the results of \citet{STO-S} who analyzed the same STO data as \citet{STO-K}651 (hereafter, we refer \citet{STO-K} and \citet{STO-S} as to STO-K and STO-S, respectively)."652 We also discuss the groundless interpretations of the long retention age of the T1 primordial surface ancl the thickness of its dust mantle based on comprehensive observational and theoretical considerations., We also discuss the groundless interpretations of the long retention age of the T1 primordial surface and the thickness of its dust mantle based on comprehensive observational and theoretical considerations.653 We calculate the color temperature το. the silicatefeature strength Si. and (he degree of linear polarization J) for the standard model as follows.," We calculate the color temperature $T_{\rm c}$, the silicate–feature strength $S_{\rm si}$, and the degree of linear polarization $P_{\rm l}$ for the standard model as follows."654 In (he standard model. large compact aggregates exist in a dust mantle and large fluffy. aggregates are embedded: with volatiles below the dust mantle (Ixolokolovaetal.," In the standard model, large compact aggregates exist in a dust mantle and large fluffy aggregates are embedded with volatiles below the dust mantle \citep{Kolokolova}."6552007).. To describe largeaggregates. we use fractal clusters of 27? identical monomers with a radius of 0.1jan (Ixolokolovaοἱal.," To describe largeaggregates, we use fractal clusters of $2^{22}$ identical monomers with a radius of $0.1~\micron$ \citep{Kolokolova}."6562007).. Each monomer is assumed to have a concentric structure of an organic refractory outer laver. a lorsteritic inner laver. ancl an amorphous silicate core (Yamamotoetal.2007).," Each monomer is assumed to have a concentric structure of an organic refractory outer layer, a forsteritic inner layer, and an amorphous silicate core \citep{YamaT}."657. The compact aggregates are asstuned {ο have a fractal dimension D=2.5. which corresponds (o the value when the maximum compression of [hilly aggregates is attained (Wadaοἱal.2007).," The compact aggregates are assumed to have a fractal dimension $D=2.5$, which corresponds to the value when the maximum compression of fluffy aggregates is attained \citep{Wada}."658. Our calculation using the superposition T-matrix method (TAAL) showed that the geometric albedo of the compact aggregates is ~0.04 al visible wavelengths (seeINimuraetal.2003).. and this value is identical to the albedo derived from an observation of the Tl surface CALearn et al.," Our calculation using the superposition T-matrix method (TMM) showed that the geometric albedo of the compact aggregates is $\sim 0.04$ at visible wavelengths \citep[see][]{Kimura}, and this value is identical to the albedo derived from an observation of the T1 surface (A'Hearn et al."659 2005)., 2005).660 We also calculate ¢~0.3 (the ratio of solar radiation pressure (o solar gravitv) for (he compact aggregates. using Mie (heorv with the optical constants deduced. from the \laxwell-Garnett mixing rule (AIG) (seeMukaietal.1992): This is also consistent. with 2» [or high-velocity ejecta within the DI ejecta plume 2005)..," We also calculate $\beta \sim 0.3$ (the ratio of solar radiation pressure to solar gravity) for the compact aggregates, using Mie theory with the optical constants deduced from the Maxwell-Garnett mixing rule (MG) \citep[see][]{Mukai}: This is also consistent with $\beta$ for high-velocity ejecta within the DI ejecta plume \citep[e.g.,][]{Meech05}. ."661 Similarly. let large (hilly aggregates embecdcded with volatile materials below the clust mantle be presented by fractal clusters with D=1.9 (Alukaietal. 1992)..," Similarly, let large fluffy aggregates embedded with volatile materials below the dust mantle be presented by fractal clusters with $D=1.9$ \citep{Mukai}. ."662"where e,, (of nominal value 5%)) is a parametrisation of the total mass energy raclialed in the coalescence. and the reduced mass is fe=maqma/M.","where $\epsilon_m$ (of nominal value ) is a parametrisation of the total mass energy radiated in the coalescence, and the reduced mass is $\mu=m_1m_2/M$."663 The reduction [actor (44//M)? is unity for equal masses and gives the correct scaling law in (he test particle limit ji«AM., The reduction factor $(4\mu/M)^2$ is unity for equal masses and gives the correct scaling law in the test particle limit $\mu \ll M$.664" The Gequeney of gravitational radiation in (he in-spiral phase is well-defined as a [unetion of time. and increases monotonically,"," The frequency of gravitational radiation in the in-spiral phase is well-defined as a function of time, and increases monotonically."665 All the energy emitted in the in-spiral phase is al Irequencies less (han /;., All the energy emitted in the in-spiral phase is at frequencies less than $f_i$.666 We assume that the spectrum of radiation in the merger phase is confined to the lrequency regime [>f;., We assume that the spectrum of radiation in the merger phase is confined to the frequency regime $f >f_i$.667 As we discuss below. we cleline the end of the merger phase (ο occur when the waveform can be described by the /=m2 quasi-normal mode signal of a Ixerr black hole.," As we discuss below, we define the end of the merger phase to occur when the waveform can be described by the $l=m=2$ quasi-normal mode signal of a Kerr black hole."668" The quasi-normal ringing frequency f, gives an approximate upper-bound [ον the frequencies carrving substantial power during the merger (Flanagan llughes 1998).", The quasi-normal ringing frequency $f_q$ gives an approximate upper-bound for the frequencies carrying substantial power during the merger (Flanagan Hughes 1998).669 where F(a)=120.63(1—a)! and a is the dimensionless spin parameter of the black hole (Echeverria 1989)., where $F(a)=1-0.63 (1-a)^{3/10}$ and $a$ is the dimensionless spin parameter of the black hole (Echeverria 1989).670 Though (he energy spectrum could have some features related to the dvnamical instabilities (Zluge οἱ al., Though the energy spectrum could have some features related to the dynamical instabilities (Zhuge et al.671 1994: Dimmelmeier et al., 1994; Dimmelmeier et al.672" 2002). we assume the simplest flat spectrum with the following amplitude. Using eq (2)) and an approximation/,—f;~ J. the characteristic gravitational wave amplitude is given bv Ilere the AZ. ji and e; ave those appropriate to the end of the merger phase."," 2002), we assume the simplest flat spectrum with the following amplitude, Using eq \ref{eq:hc}) ) and an approximation$f_q-f_i \sim f_q$ , the characteristic gravitational wave amplitude is given by Here the $M$, $\mu$ and $\epsilon_m$ are those appropriate to the end of the merger phase."673 If dynamical instabiliües develop in the rotating core or in (he rotating massive disk during the merger phase. the deformed core/cdisk could radiate strong gravitational waves in a narrow frequency band.," If dynamical instabilities develop in the rotating core or in the rotating massive disk during the merger phase, the deformed core/disk could radiate strong gravitational waves in a narrow frequency band."674 The deformation may be considered. in ils simplest form. as either two blobs or a bar.," The deformation may be considered, in its simplest form, as either two blobs or a bar."675 Using in either case a formula appropriate for a rotating bar (e.g. Frver. llolz IIughes 2002). we can estimate the amplitude of the corresponding exavitational wave enussion.," Using in either case a formula appropriate for a rotating bar (e.g. Fryer, Holz Hughes 2002), we can estimate the amplitude of the corresponding gravitational wave emission."676 Considering a bar of mass mm and length 2r which rotates with angular frequency aw. (hemean strain is given by," Considering a bar of mass $m$ and length $2r$ which rotates with angular frequency $\omega$ , themean strain is given by"677outer regions.,outer regions.678 The resulting evolution leads to star formation that is systematically less racially extended as time goes on., The resulting evolution leads to star formation that is systematically less radially extended as time goes on.679 We can verily that the age gradient. observed. in. our simulations is established at formation rather than through stellar migration., We can verify that the age gradient observed in our simulations is established at formation rather than through stellar migration.680 Figures 7. and 8 show the radius where stars formed as a function of time., Figures \ref{fig:kmsrformtform} and \ref{fig:bfrformtform} show the radius where stars formed as a function of time.681 Ehe overall shape of Figures 7 and S closely resembles Figures 5. and 6:: both sets of figures show shrinking of the radial envelope within which stars form., The overall shape of Figures \ref{fig:kmsrformtform} and \ref{fig:bfrformtform} closely resembles Figures \ref{fig:kmsrtform} and \ref{fig:bfrtform}; both sets of figures show shrinking of the radial envelope within which stars form.682 Formation location is therefore à primary factor in determining the final age gradient. rather than the degree of stellar migration.," Formation location is therefore a primary factor in determining the final age gradient, rather than the degree of stellar migration."683 One aspect of Figures 5. and 6 that dillers greatly from Figures 7. and S is that many stars finish the simulation bevond where they form., One aspect of Figures \ref{fig:kmsrtform} and \ref{fig:bfrtform} that differs greatly from Figures \ref{fig:kmsrformtform} and \ref{fig:bfrformtform} is that many stars finish the simulation beyond where they form.684 This outward movement is investigated in more detail in the next section., This outward movement is investigated in more detail in the next section.685 Each of the models shows two distinct spatial regions., Each of the models shows two distinct spatial regions.686 At small radii the star formation is abundant and nearly continuous.," At small radii, the star formation is abundant and nearly continuous."687 At large radii. there is less star formation and what there is is episodic.," At large radii, there is less star formation and what there is is episodic."688 During these episodes of elevated: star. formation. gas temporarily reaches star forming densities in the outer region of the disk. typically in spiral patterns.," During these episodes of elevated star formation, gas temporarily reaches star forming densities in the outer region of the disk, typically in spiral patterns."689 Once the stars have formed. the supernova feedback disperses the dense gas and shuts down star formation.," Once the stars have formed, the supernova feedback disperses the dense gas and shuts down star formation."690 1n models with lower stellar mass. the cpisocic bursts are scattered. at regular. time intervals spacecl 300. Myr apart.," In models with lower stellar mass, the episodic bursts are scattered at regular time intervals spaced 300 Myr apart."691 This timescale is similar to the star formation interval reported in ?.. and is related to the free fall time at the center ob halos.," This timescale is similar to the star formation interval reported in \citet{Stinson07}, and is related to the free fall time at the center of halos."692 Higher mass galaxies also show similar episocic star formation. but only outside the inner stable star forming disk.," Higher mass galaxies also show similar episodic star formation, but only outside the inner stable star forming disk."693 At large radii. the episodic star formation has a longer timescale of ~1 Gyr due to the lower characteristic densities associated with longer free fall times.," At large radii, the episodic star formation has a longer timescale of $\sim 1$ Gyr due to the lower characteristic densities associated with longer free fall times."694 This episodie star formation can be explained with a delay cilferential equation as shown hy ?.., This episodic star formation can be explained with a delay differential equation as shown by \citet{Quillen2008}.695 lor a More direct οςmparisc211 with observations. Figure 9 plots model CMDs showing the expected stellar populations in several radial bins.," For a more direct comparison with observations, Figure \ref{fig:cmds} plots model CMDs showing the expected stellar populations in several radial bins."696 To create the CAIDs. we used StarISLE version 1.1 (?) to populate a set of isochrones (?) based on a user-supplied star-formation history. (SELL) set by the mass. age. metallicity (-2.3 « pefl] «-1 after 13.5 Gyr). and. position of the simulated star particles.," To create the CMDs, we used StarFISH version 1.1 \citep{harris01} to populate a set of isochrones \citep{Girardi02} based on a user-supplied star-formation history (SFH) set by the mass, age, metallicity (-2.3 $<$ [Fe/H] $<$ -1 after 13.5 Gyr), and position of the simulated star particles."697 To mimic observing conditions. we adopted artificial star tests from the ACS Nearby Galaxy Survey (7). assuming that the simulated. galaxy is located at a distance of 500 kpe.," To mimic observing conditions, we adopted artificial star tests from the ACS Nearby Galaxy Survey \citep{Dalcanton2008} assuming that the simulated galaxy is located at a distance of 500 kpc."698 The CALDs of the simulations show a prominent voung main sequence in the central 2 kpe that disappears at large radii., The CMDs of the simulations show a prominent young main sequence in the central 2 kpc that disappears at large radii.699 The absence of the voung main sequence outside 2 kpe agrees with the age gradient shown in Figure 4.., The absence of the young main sequence outside 2 kpc agrees with the age gradient shown in Figure \ref{fig:ageprofile}.700 Figure 9 also shows that any main sequence stars in the halo are more than 1 magnitude fainter than the red clump and horizontal branch. such that only the deepest. photometry in the nearest galaxies would. show evidence of this intermediate age population.," Figure \ref{fig:cmds} also shows that any main sequence stars in the halo are more than 1 magnitude fainter than the red clump and horizontal branch, such that only the deepest photometry in the nearest galaxies would show evidence of this intermediate age population."701 The age gradient apparent in Figure 9. is consistent with the deep CALDs presented in Figure 2 of ? suggesting that isolated models can produce realistic stellar structure., The age gradient apparent in Figure \ref{fig:cmds} is consistent with the deep CMDs presented in Figure 2 of \citet{hidalgo03} suggesting that isolated models can produce realistic stellar structure.702 While the location of star formation largely determines the age eracient and stellar structure of the galaxies. there are some features that are due to the motion of stars after they form.," While the location of star formation largely determines the age gradient and stellar structure of the galaxies, there are some features that are due to the motion of stars after they form."703 Figures 10 and 11 highlight these cdillerences hy showing how far the stars have migrated. racially., Figures \ref{fig:kmsmigration} and \ref{fig:bfmigration} highlight these differences by showing how far the stars have migrated radially.704" The mean migration (dashed. line) in the constan f, 0.1 models shown in Figure 10. shows little deviation from zero (solid line). although many individual particles move a large distance."," The mean migration (dashed line) in the constant $f_b$ =0.1 models shown in Figure \ref{fig:kmsmigration} shows little deviation from zero (solid line), although many individual particles move a large distance."705 For the lowest mass halos. most of this movement is in the outward direction and is concurren with star formation episodes.," For the lowest mass halos, most of this movement is in the outward direction and is concurrent with star formation episodes."706 To explore this behavior. Figure 12. shows the stellar apocenters as a function of their velocity at the time of formation.," To explore this behavior, Figure \ref{fig:velsf} shows the stellar apocenters as a function of their velocity at the time of formation."707 Compared to the vertical line at zero velocity. stars to the right formed from gas that was moving outwarels. while stars to the left formed as the gas was collapsing.," Compared to the vertical line at zero velocity, stars to the right formed from gas that was moving outwards, while stars to the left formed as the gas was collapsing."708 There is a noticeable trend for stars with larger apocenters to. form with larger outward. velocities. particularly for the lowest mass 15 kms + halo.," There is a noticeable trend for stars with larger apocenters to form with larger outward velocities, particularly for the lowest mass 15 km $^{-1}$ halo."709— Physically. 1ese halo stars formed [rom outward [owing gas that was vecelerated: by supernovae. blastwaves and. shocks against infalling gas.," Physically, these halo stars formed from outward flowing gas that was accelerated by supernovae blastwaves and shocks against infalling gas."710 Phe shock produces high densities and triggers μαar formation., The shock produces high densities and triggers star formation.711 These stars are thus launched. on racial bits that create an extended stellar halo and a positively 4sewed velocity distribution., These stars are thus launched on radial orbits that create an extended stellar halo and a positively skewed velocity distribution.712 Conversely. few stars. form uring the infall phase curing which shocks are weaker or ibsent.," Conversely, few stars form during the infall phase during which shocks are weaker or absent."713 No stars are observed to be ereateck with sullicient nergv to become unbound from the chvarl galaxy., No stars are observed to be created with sufficient energy to become unbound from the dwarf galaxy.714 Once the warf interacts with other satellites. some of these stars may become unbound.," Once the dwarf interacts with other satellites, some of these stars may become unbound."715 Figure 12. suggests that shocks and instabilities in the SN-driven wind may be necessary for the formation of the most extended halo stars seen in dwarls., Figure \ref{fig:velsf} suggests that shocks and instabilities in the SN-driven wind may be necessary for the formation of the most extended halo stars seen in dwarfs.716 To verily that the large radius. high initial velocity stars are indeed due to feedback. we have rerun the 15 km simulation. with feedback turned. oll.," To verify that the large radius, high initial velocity stars are indeed due to feedback, we have rerun the 15 km $^{-1}$ simulation with feedback turned off."717 In. this simulation. no extended: halo forms.," In this simulation, no extended halo forms."718 Instead. the distribution of apocenters is centered at zero radial velocity. consisting of disk stars formed from gas moving in stable circular orbits.," Instead, the distribution of apocenters is centered at zero radial velocity, consisting of disk stars formed from gas moving in stable circular orbits."719 The signature of halo formation due to. supernova feedback is much reduced in higher mass galaxies. which form long-lived stable stelar disks.," The signature of halo formation due to supernova feedback is much reduced in higher mass galaxies, which form long-lived stable stellar disks."720 In these cases. disk stars form on circular orbits anc appear near the zero initial racial velocity line in Figure 12..," In these cases, disk stars form on circular orbits and appear near the zero initial radial velocity line in Figure \ref{fig:velsf}."721 ‘These stars may migrate outwards through disk instabilities (?).. but are unlikely to swell into a 3-dimensional halo.," These stars may migrate outwards through disk instabilities \citep{Roskar2008}, but are unlikely to swell into a 3-dimensional halo."722 ligure 12 also sugecs sthat in the lowest mass galaxies. one expects a kinematic sienature where the outermost halo stars are all preferentially on radial orbits.," Figure \ref{fig:velsf} also suggests that in the lowest mass galaxies, one expects a kinematic signature where the outermost halo stars are all preferentially on radial orbits."723 However. these stars can be strongly inlueneecl by tidal ellects. possibly making this signature cillicult to detect observationallv.," However, these stars can be strongly influenced by tidal effects, possibly making this signature difficult to detect observationally."724" Unlike in the f, nxxdels. the low f, models show an ollset between the mean initial and final radii as shown in Figure 11.."," Unlike in the $f_b$ models, the low $f_b$ models show an offset between the mean initial and final radii as shown in Figure \ref{fig:bfmigration}."725 This olfset develops because while the halo potential well is deep enoteh that eas cools onto a disk. the gas is pressure supported and not dense enough to form stars outside the central region.," This offset develops because while the halo potential well is deep enough that gas cools onto a disk, the gas is pressure supported and not dense enough to form stars outside the central region."726 ‘Thus. in these low fi models. the eas disk is significantly more massive than the stellar clisk ancl dominates the disk dynamics. as shown in Figure J DE ," Thus, in these low $f_b$ models, the gas disk is significantly more massive than the stellar disk and dominates the disk dynamics, as shown in Figure \ref{fig:gasstarmass}. ."727The stellar disk then meanders in response to the dominant influence of the gas., The stellar disk then meanders in response to the dominant influence of the gas.728 As the stellar disk. meanders. Figure 1l shows that it scatters stars into the halo.," As the stellar disk meanders, Figure \ref{fig:bfmigration} shows that it scatters stars into the halo."729 Stars that, Stars that730The spectruui of V132. Aur at primary uiuima is due ouly to the secoudary star. anc subtracting it from the spectrum at secondary eclipse would provide the spectrum of the primary component.,"The spectrum of V432 Aur at primary minimum is due only to the secondary star, and subtracting it from the spectrum at secondary eclipse would provide the spectrum of the primary component."731 Unfortunately. during the scheduled spectroscopic observiug rus at the telescope. V132 Aur never passed through the primary eclipse. and therefore it has not been possible to obtain the isolated spectra of the individual components.," Unfortunately, during the scheduled spectroscopic observing runs at the telescope, V432 Aur never passed through the primary eclipse, and therefore it has not been possible to obtain the isolated spectra of the individual components."732 To the aiu of performing an atmospheric analysis of the two components of V132 Aur we have therefore focused on the spectra of the highest S/N around phases 0.25 and 0.75 (cf., To the aim of performing an atmospheric analysis of the two components of V432 Aur we have therefore focused on the spectra of the highest S/N around phases 0.25 and 0.75 (cf.733 Table 2). where the velocity separation of the components is asian aud the same lines from the two components are unbleded.," Table 2), where the velocity separation of the components is maximum and the same lines from the two components are unbleded."734 Working siuultauncously at phase 0.25 and 0.75 allows also to eet rid of line superposition aud stronely reimforces the robustuess of the overall ft., Working simultaneously at phase 0.25 and 0.75 allows also to get rid of line superposition and strongly reinforces the robustness of the overall fit.735 To derive the basic stellar parameters (τω. logg.o ‘Z| and Voor) we have performed a u best match analvsis of the high S/N observed spectra around phases 0.25 and 0.75 (spectrum 392323. 39239. 39605. 39291 and 39607 in Table 1) against the extensive erid of svuthetie Iuruczs spectra computed by Muni et al. (," To derive the basic stellar parameters $T_{\rm eff}$ , $\log g$, $_\odot$ ] and $V_{\rm rot}$ ) we have performed a $\chi^2$ best match analysis of the high S/N observed spectra around phases 0.25 and 0.75 (spectrum 39233, 39239, 39605, 39294 and 39607 in Table 1) against the extensive grid of synthetic Kurucz's spectra computed by Munari et al. ("7362003) for the same resolution of the Asiago Echelle spectrograph (R=20 0000) over the 5500 iranee (thus fully covering the wavelength rauge recorded or VI32 Aur),2003) for the same resolution of the Asiago Echelle spectrograph $R$ 000) over the $-$ 500 range (thus fully covering the wavelength range recorded for V432 Aur).737 The Muuar et al. (, The Munari et al. (738"2003) svuthetic atlas covers the range 355002Tigx 5500. 0loggyx:5.0. 5x |Z/Z.] € (with solar relative abundances or metals) and O<Vu,«500 kin |.","2003) synthetic atlas covers the range $\leq T_{\rm eff} \leq$ 500, $\leq \log g739\leq$ 5.0, $\leq$ $_\odot$ ] $\leq -2.5$ (with solar relative abundances for metals) and $\leq V_{\rm rot}\leq$ 500 km $^{-1}$."740 The grid steps around the values appropriate to V132 Aur are 250 K iu eniperature. 0.5 dex in gravity. 0.5 dex iu iictallicitv aud 10 km 1 in rotation.," The grid steps around the values appropriate to V432 Aur are 250 K in temperature, 0.5 dex in gravity, 0.5 dex in metallicity and 10 km $^{-1}$ in rotation."741 To proceed independently frou: the orbital solution. ↿∖⊳∖↿⋮↧↓⋅⇉⋪↓⊔↾↓∖∥∣⋡⇂⋖⋅↓⊐⋡∪⊔↿⇂⊓⋅↓⋅∢⋅⋡∖⋪⊓⇂⋯∐⋡∖∩⇂⋅↿↓↕≺⋅↓−⋯⇂⋪⋯⇂∖⇁≺⋅⇂⋯∙⋪↓↿⋪⊓⋅⊳∖⋜⋯≺↧↑↕⋯↴∖↴↑∪↿↴∖↴↸∖↑∐↸∖⋜↕↕⊔∪↴∖↴↻↕∐∖↥⋅∐⊳⋜⋯⋜↧↕⋅↖↽↴∖↴↕↴∖↴∪↖⇁⋜↧∐≼↧⋜↧↑↸∖↕↑∙ ' : | ∙ ⇂∎↓⋅∪⊔↓⇂↓↕⋖⋅⊓↓⋅∣⋡⋪⊔⋮↧↓⊳∖∪↓⋯⋪↓∪⊔↿∖↓⊲∖⋪↓⋏∙≟⊔↓⋅⋖⋅↓⋜⋯∠⇂↾↓∖⋜↧∣⋡↓⋖⋅⇉⊐⊳↾↓∖↓↕⋖⋅↓≻⋖⋅↓⋰↓⋯⇂↖↖⊽↸∖ iad to first define the region where the true 1iniumun of: tle 47 D4:distribution: has to be looked for.+ to avoid: beiug: trapped iuto uuplivsical local minima.," To proceed independently from the orbital solution, and thus to use the atmospheric analysis to validate it, we had to first define the region where the true minimum of the $\chi^2$ distribution has to be looked for, to avoid being trapped into unphysical local minima."742 To accomplish this. we have isolated from the spectrum secured at secondary eclipse (#33968 Lin Table 1) the wavelength range 8180STIO ({dominated by diagnostic Call triplet aud Pascheu series lines).," To accomplish this, we have isolated from the spectrum secured at secondary eclipse 39684 in Table 1) the wavelength range 8480--8740 (dominated by diagnostic CaII triplet and Paschen series lines)."743 We have classified it against the spectral atlases of Afbunari and Tomasella (1999) aud Aarrese et al. (, We have classified it against the spectral atlases of Munari and Tomasella (1999) and Marrese et al. (7442003) obtained over the same wavelength range with the same oeistrumental set-up for the Asiago Echelle spectrograph.,2003) obtained over the same wavelength range with the same instrumental set-up for the Asiago Echelle spectrograph.745 We also checked the consistency of the classification east the svuthetic atlas of Muni and Castelli (2000) that covers this wavelength region at the same resolving power of our observatious., We also checked the consistency of the classification against the synthetic atlas of Munari and Castelli (2000) that covers this wavelength region at the same resolving power of our observations.746 Haviug determined im 76500 Is the colmbined temperature of the two stars. we have imoved to the analytical 47 test by fixing to 55007500 the temperature boundary and to 100 kin 1 he rotational velocities aud inposimeg uo restriction on gravity auc metallicity.," Having determined in $\sim$ 6500 K the combined temperature of the two stars, we have moved to the analytical $\chi^2$ test by fixing to 5500–7500 the temperature boundary and to $\leq$ 100 km $^{-1}$ the rotational velocities and imposing no restriction on gravity and metallicity."747 The analysis has been limited to the same sx⋅ Echelle orders adopted for⋅⋅ the derivation. of. radial velocities (cf, The analysis has been limited to the same six Echelle orders adopted for the derivation of radial velocities (cf.748 sect., sect.749 2.3). ready to iuclude additional orders if uot satisfactorily converging on the first six.," 2.3), ready to include additional orders if not satisfactorily converging on the first six."750 This has uot been necessary given the close sinübluity of the results among the 6 selected orders., This has not been necessary given the close similarity of the results among the 6 selected orders.751 The results of the atmospheric analysis are given in Table Ll where a comparison is provided with the pariuneters common to the orbital solution.," The results of the atmospheric analysis are given in Table 4, where a comparison is provided with the parameters common to the orbital solution."752 Given the absolute independence of the two methods (carried. out mdepeudeutlv by two distinct sub-eroups of the authors of thispaper). the remarkable correspondence mutually reinforce the confidence d the orbital solution. and atmospheric analysis.," Given the absolute independence of the two methods (carried out independently by two distinct sub-groups of the authors of thispaper), the remarkable correspondence mutually reinforce the confidence in the orbital solution and atmospheric analysis."753 Two sinall and diagnostic samples, Two small and diagnostic samples754 , 755Baldwin. Phillips “Terlevich 1981).,"Baldwin, Phillips Terlevich 1981)."756 In. all four knots. woe find that the optical line ratios are most naturally explained. with photoionization by a hard. continuum.," In all four knots, we find that the optical line ratios are most naturally explained with photoionization by a hard continuum."757" In particular. the ΝΟOLLI] ratio (Pable 1) is consistent with photoionization and A,.=O0 mae (also Humphrey ct al."," In particular, the [NeIII]/[OIII] ratio (Table 1) is consistent with photoionization and $_{v}$ =0 mag (also Humphrey et al."758" 2008a). but is substantially lower than predicted: by fast shock models (νο, 0.2-1: Dopita Sutherland 1996)."," 2008a), but is substantially lower than predicted by fast shock models (i.e., 0.2-1: Dopita Sutherland 1996)."759 The /LE3 ratio is also consistent with photoionization by a hard continuum (e.g. Robinson et al., The $\beta$ ratio is also consistent with photoionization by a hard continuum (e.g. Robinson et al.760 LOST)., 1987).761 The detection of NeV] in apertures B and C also supports this idea., The detection of [NeV] in apertures B and C also supports this idea.762 This conclusion is in agreement with that of Taniguchi ct al. (, This conclusion is in agreement with that of Taniguchi et al. (7632001).,2001).764 It is. however. important to recognise that the line ratios themselves do not allow us to determine the source of the hard ionizing continuum (ie. the AGN. the metagalactic background. radiation or radiative shocks).," It is, however, important to recognise that the line ratios themselves do not allow us to determine the source of the hard ionizing continuum (i.e., the AGN, the metagalactic background radiation or radiative shocks)."765 Extensive imaging in à variety of wavebands has revealed spectacular filamentary or bubble-like features roughly along the racio axis of AIRC 0406-244 (Rush et al., Extensive imaging in a variety of wavebands has revealed spectacular filamentary or bubble-like features roughly along the radio axis of MRC 0406-244 (Rush et al.766 1997: Pentericci et al., 1997; Pentericci et al.767 2001: see Figure 2)., 2001; see Figure 2).768 While imaging alone provides insullicient information to elucidate the nature of. these morphological features. spectroscopic data can be used to assess what fraction of this spatially extended. emission is emitted by warm ionized gas.," While imaging alone provides insufficient information to elucidate the nature of these morphological features, spectroscopic data can be used to assess what fraction of this spatially extended emission is emitted by warm ionized gas."769 For cach of our spectroscopic apertures. we first caleulated. the tux expected. due. to nebular continuum within the spectral range of the FIGOW filter.," For each of our spectroscopic apertures, we first calculated the flux expected due to nebular continuum within the spectral range of the F160W filter."770 This we have estimated. using the 112 flux and. the nebular continuum emission coellicients in. Aller (1984). assuming that the eas has a temperature of 15.000. Ix. We then summed. the expected nebular coninuum [lux with the measured. ας of line. emission. and οςmipared this against the total flux measured. within the spectral range corresponding to the wavelength range of the relevant broad-band filter.," This we have estimated using the $\beta$ flux and the nebular continuum emission coefficients in Aller (1984), assuming that the gas has a temperature of 15,000 K. We then summed the expected nebular continuum flux with the measured flux of line emission, and compared this against the total flux measured within the spectral range corresponding to the wavelength range of the relevant broad-band filter."771 In the spatially extended: figure-of-cight structure (apertures A.C ancl DJ. we find that the overall II-band fux is dominated by nebular emission (i.e. 2:90 per cent).," In the spatially extended figure-of-eight structure (apertures A, C and D), we find that the overall H-band flux is dominated by nebular emission (i.e. $\ge$ 90 per cent)."772 While Rush ct al. (, While Rush et al. (7731997) suggested that the igure-ol-eight structure has a tidal origin. we prefer the interpretation of MeCarthy ct al. (,"1997) suggested that the figure-of-eight structure has a tidal origin, we prefer the interpretation of McCarthy et al. ("7741999) and ‘Taniguchi et al. (,1999) and Taniguchi et al. (7752001). lo..1 this structure is ambient ISM that has been swept up into a bubble by a galactic superwind.,"2001), i.e., this structure is ambient ISM that has been swept up into a bubble by a galactic superwind."776 Our reasoning for this is as follows., Our reasoning for this is as follows.777 In the L-band image. the bright. emission knots are arranged into two ellipses. one on cach side of the brightest emission peak. which we assume marks the position of the galactic nucleus.," In the H-band image, the bright emission knots are arranged into two ellipses, one on each side of the brightest emission peak, which we assume marks the position of the galactic nucleus."778 Such morphological features can be naturally explained as eas that has been swept up from tae ISAT of AIRC 0406-244 by a galactic wind. anc which has since cooled at the interface between the wir and the external medium.," Such morphological features can be naturally explained as gas that has been swept up from the ISM of MRC 0406-244 by a galactic wind, and which has since cooled at the interface between the wind and the external medium."779 Indeed. as remarked by MeCarthy et al. (," Indeed, as remarked by McCarthy et al. ("7801999). these features are reminiscent of the winel-blown xibbles associated with some ultra-Iuminous infrarec ealaxies.,"1999), these features are reminiscent of the wind-blown bubbles associated with some ultra-luminous infrared galaxies."781 In. addition. at the centre of cach of the emission ellipses there is a region of essentially zero Hux.," In addition, at the centre of each of the emission ellipses there is a region of essentially zero flux."782 This can be naturally explained if the wind has been relatively elficien ab sweeping up. or at shock-heating. the ambient ISM. caving little or no cool gas in its wake.," This can be naturally explained if the wind has been relatively efficient at sweeping up, or at shock-heating, the ambient ISM, leaving little or no cool gas in its wake."783 Outside of the xibbles. a significant ας of emission is detected. which can » explained as emission [rom ISM across which the bubble las not vet. expanded.," Outside of the bubbles, a significant flux of emission is detected, which can be explained as emission from ISM across which the bubble has not yet expanded."784 In 53.1. we analysed ratios between emission. lines. conclucing hat the gas in the super bubbles is ionized wimarily by a hard continuum.," In $\S3.1$, we analysed ratios between emission lines, concluding that the gas in the super bubbles is ionized primarily by a hard continuum."785 As mentioned in 983.1. the ine ratios themselves have not allowed: us to discriminate tween three potential sources of hard ionizing continuum emission.," As mentioned in $\S$ 3.1, the line ratios themselves have not allowed us to discriminate between three potential sources of hard ionizing continuum emission."786 In this context. the detailed. morphology. of the super bubble can provide a useful test. to. discriminate tween photoionization by the anisotrpic radiation field of he AGN. and ionization by shocks or bv the isotropic ealactic background radiation (ALIBI).," In this context, the detailed morphology of the super bubble can provide a useful test to discriminate between photoionization by the anisotrpic radiation field of the AGN, and ionization by shocks or by the isotropic meta-galactic background radiation (MBR)."787 Let us suppose that each superbubble in MIRC'0406-244 is à spherical bubble of cool gas. which is optically thick to ionizing radiation. which has a uniform cistribution of material. and which is irraciated by an AGN located at some," Let us suppose that each superbubble in MRC0406-244 is a spherical bubble of cool gas, which is optically thick to ionizing radiation, which has a uniform distribution of material, and which is irradiated by an AGN located at some"7881.7 ffrom Fig. 2..,1.7 from Fig. \ref{Fig:HR_Sb}.789 The situation. looks quite different indeed when considering intermediate-mass K giants (upper panel of Fig. 7))., The situation looks quite different indeed when considering intermediate-mass K giants (upper panel of Fig. \ref{Fig:elogP_cluster_mass}) ).790 If the sample of K giant binaries were only made out of low-mass stars. the correcting factor would then be 7/9 = 0.78 instead of 83/126 = 0.66 for the intermediate-mass K giants (close to the value 0.72 derived above for the total sample of K giants).," If the sample of K giant binaries were only made out of low-mass stars, the correcting factor would then be 7/9 = 0.78 instead of 83/126 = 0.66 for the intermediate-mass K giants (close to the value 0.72 derived above for the total sample of K giants)."791 The true correction factor should thus be the average of these two values. weighted by the (unknown) fraction of low-mass with respect to intermediate-mass stars among the sample of field K-giant binaries.," The true correction factor should thus be the average of these two values, weighted by the (unknown) fraction of low-mass with respect to intermediate-mass stars among the sample of field K-giant binaries."792 This true correction factor should be in the range 0.6-0.8 (irrespective of the exact weighing function). which is still too much to account for the factor of 2—3 difference between the binary. frequencies among K and M etants.," This true correction factor should be in the range 0.6–0.8 (irrespective of the exact weighing function), which is still too much to account for the factor of 2–3 difference between the binary frequencies among K and M giants."793 In the end. it is the difficulty of finding spectroscopic. binaries among M giants due to their larger radial-velocity jitter that must be invoked to account for the binaries missing among M giants.," In the end, it is the difficulty of finding spectroscopic binaries among M giants due to their larger radial-velocity jitter that must be invoked to account for the binaries missing among M giants."794 In this paper we have derived the frequency of spectroscopic binaries among field M giants. for the first time based on an extensive sample. 771 M giants in total.," In this paper we have derived the frequency of spectroscopic binaries among field M giants, for the first time based on an extensive sample, 771 M giants in total."795 The frequency obtained.6.3%.. is much lower than for K giants (ranging from 14 to31%.. depending on the samples considered).," The frequency obtained, is much lower than for K giants (ranging from 14 to, depending on the samples considered)."796 However. the binary frequency derived from this large sample of M giants is not very meaningful. since there are important observational difficulties with detecting M giant binaries. because on average they have smaller orbital velocity amplitudes (longer periods) and larger intrinsic velocity jitter than K giants.," However, the binary frequency derived from this large sample of M giants is not very meaningful, since there are important observational difficulties with detecting M giant binaries, because on average they have smaller orbital velocity amplitudes (longer periods) and larger intrinsic velocity jitter than K giants."797 This effect is especially important for samples with only a fewmeasurements per star. as is the case with this sample of 771 M giants.," This effect is especially important for samples with only a fewmeasurements per star, as is the case with this sample of 771 M giants."798 A higher binary frequency was obtained in a smaller M giant sample with more radial-velocity measurements per object: confirmed plus possible binaries., A higher binary frequency was obtained in a smaller M giant sample with more radial-velocity measurements per object: confirmed plus possible binaries.799 Even though this frequency is close to the lower bound of the binary frequency among K giants. we have shown that they cannot be directly compared because they were obtained under different observing conditions (mostly the number of observations per object).," Even though this frequency is close to the lower bound of the binary frequency among K giants, we have shown that they cannot be directly compared because they were obtained under different observing conditions (mostly the number of observations per object)."800 When one tries to compare the binary frequencies for samples of K and M giants under similar observing conditions. the binary frequency among M giants remains lower than that among K giants by a factor of about 2.," When one tries to compare the binary frequencies for samples of K and M giants under similar observing conditions, the binary frequency among M giants remains lower than that among K giants by a factor of about 2."801 Part of this difference may stem from the destruction of shorter-period K giant binaries due to catastrophic RLOF leading to common envelope. before they can become M giant binaries.," Part of this difference may stem from the destruction of shorter-period K giant binaries due to catastrophic RLOF leading to common envelope, before they can become M giant binaries."802 However. how important this effect is depends on the mass distribution in the considered populatioαυ," However, how important this effect is depends on the mass distribution in the considered population."803] The lack of Μ-σιαπί binaries with orbital periods shorter than 160 d. as compared to K-giant binaries that have orbital periods as short as 4 d. has been clearly demonstrated from our extensive set of orbital elements.," The lack of M-giant binaries with orbital periods shorter than 160 d, as compared to K-giant binaries that have orbital periods as short as 4 d, has been clearly demonstrated from our extensive set of orbital elements."804 We have found that the CORAVEL line-width parameter b 1s better correlated with the stellar radius than with either luminosity or effective temperature separately., We have found that the CORAVEL line-width parameter $Sb$ is better correlated with the stellar radius than with either luminosity or effective temperature separately.805 This allowed identification of the R — Sh relation outliers HD 190658 and HD 219654 as fast rotators. possibly due to the binary companion influence: indeed. HD 190658 and HD 219654 turn out to be known as spectroscopic binaries.," This allowed identification of the $R$ – $Sb$ relation outliers HD 190658 and HD 219654 as fast rotators, possibly due to the binary companion influence; indeed, HD 190658 and HD 219654 turn out to be known as spectroscopic binaries."806 A third candidate was rejected. because its large οὐ ," A third candidate was rejected, because its large $Sb$ "807A fundamental aspect of our knowledge of the Universe concerns the existence of non-baryonie dark matter (DM). believed to represents about of the total energy budget of the Universe.,"A fundamental aspect of our knowledge of the Universe concerns the existence of non-baryonic dark matter (DM), believed to represents about of the total energy budget of the Universe."808 Signatures of DM are found in galaxies and cluster of galaxies due to its dynamical effects and gravitational lensing., Signatures of DM are found in galaxies and cluster of galaxies due to its dynamical effects and gravitational lensing.809 While DM appears in wildly variable quantities and distributions among different types of objects. it appears to exhibit systematic (but not yet understood) behaviours (c.f..," While DM appears in wildly variable quantities and distributions among different types of objects, it appears to exhibit systematic (but not yet understood) behaviours (c.f.,"810 recent findings by Gentileetal.2009 and. by Donato et al.," recent findings by \cite{gentile09}811 and by Donato et al."812 2009)., 2009).813 The most remarkable (e.g.. Binney 2004)) being that DM is needed to reconcile the observations with the expectations of Newtonian dynamics when and only when the acceleration of gravity goes below a value. 1.2x107 em s (Begemanetal. 1991)).," The most remarkable (e.g., \cite{binney04}) ) being that DM is needed to reconcile the observations with the expectations of Newtonian dynamics when and only when the acceleration of gravity goes below a value, $a_0 \sim 1.2 \times 10^{-8}$ cm $^{-2}$ \cite{begeman91}) )."814 This fact led to suggest that Newtonian dynamic might not be applicable below this acceleration and the most successful proposal of this. type. known as MOND (Milgrom 1983)). nicely explains the rotation curves of spiral galaxies and many other dynamical properties of galaxies without the needs for DM (MeGaugh&deBlock1998:; Mortlock&Turner 2001:: see Sanders&MeGaugh2002 for a review).," This fact led to suggest that Newtonian dynamic might not be applicable below this acceleration and the most successful proposal of this type, known as MOND \cite{milgrom83}) ), nicely explains the rotation curves of spiral galaxies and many other dynamical properties of galaxies without the needs for DM \cite{mcgaugh98}; \cite{mortlock01}; see \cite{sanders02} for a review)."815 Such an hypothesis. however. has so many serious consequences for the standard physics (e.g. Milgrom2009. for a review). that as many test as possible should be carried out to verify it.," Such an hypothesis, however, has so many serious consequences for the standard physics (e.g. \cite{milgrom09} for a review), that as many test as possible should be carried out to verify it."816 Experiments can be made in the laboratory or studying astrophysical systems where DM ts absent., Experiments can be made in the laboratory or studying astrophysical systems where DM is absent.817 A first pioneering study along this line focused on the dynamics of the external regions of globular clusters. the largest virialized structures that do not contain significant amount of dark matter.," A first pioneering study along this line focused on the dynamics of the external regions of globular clusters, the largest virialized structures that do not contain significant amount of dark matter."818 Measurement of the velocity dispersion in the outskirt of c Centauri (Scarpa.MarconiandGilmozzi 2003A..B). showed a clear flattening of the velocity dispersion profile starting. at the radius where the cluster's internal acceleratior of gravity I8 ~do. With no evidence of the expected Keplerian falloff.," Measurement of the velocity dispersion in the outskirt of $\omega$ Centauri \cite{scarpa03A}, ,B), showed a clear flattening of the velocity dispersion profile starting at the radius where the cluster's internal acceleration of gravity is $\sim a_0$, with no evidence of the expected Keplerian falloff."819 The very same behaviour observed in elliptical galaxies and explained invoking the presence of large amounts of dark matter., The very same behaviour observed in elliptical galaxies and explained invoking the presence of large amounts of dark matter.820 This result was then extended to other 6 globular clusters (Scarpa. Marconi. and Gilmozzi 2003 A.B: Scarpa. Marconi. and Gilmozzi 2004 A.B: Scarpa et al.," This result was then extended to other 6 globular clusters (Scarpa, Marconi, and Gilmozzi 2003 A,B; Scarpa, Marconi, and Gilmozzi 2004 A,B; Scarpa et al."821 2007 A.B: Scarpa et al.," 2007 A,B; Scarpa et al."822 2010). showing the behaviour seen in «o Cen is not a peculiar property of this cluster.," 2010), showing the behaviour seen in $\omega$ Cen is not a peculiar property of this cluster."823 Given the relevance of this result. the dynamics of w Cen was recently carefully reconsidered by Sollima et al. (," Given the relevance of this result, the dynamics of $\omega$ Cen was recently carefully reconsidered by Sollima et al. ("8242009. S09 hereafter).,"2009, S09 hereafter)."825 Based on the analysis of a new large dataset of radial velocities measurements. it was claimed that the velocity dispersion decreases monotonically with radius. in agreement with Newtonian prediction. and in clear contrast with the claim by Searpa Marconi and Gilmozzi (2003B: SMG hereafter).," Based on the analysis of a new large dataset of radial velocities measurements, it was claimed that the velocity dispersion decreases monotonically with radius, in agreement with Newtonian prediction, and in clear contrast with the claim by Scarpa Marconi and Gilmozzi (2003B; SMG hereafter)."826 We reconsider here the data presented by S09., We reconsider here the data presented by S09.827 The reanalysis ts done joining both the S09 and SMG radial velocities data to create a larger sample., The reanalysis is done joining both the S09 and SMG radial velocities data to create a larger sample.828 In both works. the selection criteria for cluster members identification was basically the same (a selection based on position 11 the color magnitude diagram combined with a cut in radial velocity). ensuring that the whole dataset is homogeneous.," In both works, the selection criteria for cluster members identification was basically the same (a selection based on position in the color magnitude diagram combined with a cut in radial velocity), ensuring that the whole dataset is homogeneous."829 Located at 6.4 kpe from the galactic center (Harris 1996)). co Cen is the most massive and luminous globular cluster of the Milky Way.," Located at 6.4 kpc from the galactic center \cite{harris96}) ), $\omega$ Cen is the most massive and luminous globular cluster of the Milky Way."830 It is sufficiently massive to contain more than one stellar population. as indicated by helium abundance variation (Norris 2004; Piotto et al.," It is sufficiently massive to contain more than one stellar population, as indicated by helium abundance variation (Norris 2004; Piotto et al."831 2005) and its peculiar position in the size-luminosity plane (MackeyandvanderBergh 2005))., 2005) and its peculiar position in the size-luminosity plane \cite{mackey05}) ).832 It has been also argued that co Cen could be not a genuine globular cluster but the nuclear remnant of a dwarf galaxy that merged in the past with the Milky Way (e.g.. Bellazzinietal. 2008)).," It has been also argued that $\omega$ Cen could be not a genuine globular cluster but the nuclear remnant of a dwarf galaxy that merged in the past with the Milky Way (e.g., \cite{bellazzini08}) )."833 Among the earliest dynamical studies of ω Cen relevant to this work. Meylan Mayor (1986) discussed theradial," Among the earliest dynamical studies of $\omega$ Cen relevant to this work, Meylan Mayor (1986) discussed theradial"834 , 835introduced by the longitudinal density structuring 1n. two-dimensional thread models. solutions to Equations (5))-(9)) have been first obtained by using a two-dimensional (2D) density distribution with Linear=L.,"introduced by the longitudinal density structuring in two-dimensional thread models, solutions to Equations \ref{first}) \ref{last}) ) have been first obtained by using a two-dimensional (2D) density distribution with $L_{\rm thread}=L$."836 The purpose of these numerical experiments has been to check the correct behaviour of the code by reproducing the results obtained by Arreguietal. (2008)., The purpose of these numerical experiments has been to check the correct behaviour of the code by reproducing the results obtained by \citet{Arregui08thread}.837". In addition. we have also considered the magnetic Reynolds number. AR,ραη. Which should not affect the computed damping times. in the limit of large Reynolds numbers."," In addition, we have also considered the magnetic Reynolds number, $R_{\rm m}=v_{\rm Af}a/\eta$, which should not affect the computed damping times, in the limit of large Reynolds numbers."838 Figure 3. displays the obtained results., Figure \ref{check} displays the obtained results.839 The damping time of resonantly danped kink waves is independent of the nagnetic Reynolds number. as long as this quantity is large enough for resonance absorption to be the operating damping nechanism.," The damping time of resonantly damped kink waves is independent of the magnetic Reynolds number, as long as this quantity is large enough for resonance absorption to be the operating damping mechanism."840 This regine (see the plateau regions) is obtained for different values of Ry when different transitional layers are considered., This regime (see the plateau regions) is obtained for different values of $R_{\rm m}$ when different transitional layers are considered.841 Figure 3aa shows a perfect agreement between the ID results and the current computations using the 2D code., Figure \ref{check}a a shows a perfect agreement between the 1D results and the current computations using the 2D code.842 The perfect correspondence between 1D and 2D computations, The perfect correspondence between 1D and 2D computations843are the Mossetal.(2010). DEMO models: models 19-24 are the axisvmmetric analytic models.,are the \citet{M10} DEHO models; models 19-24 are the axisymmetric analytic models.844 For each model. the P and PA for each star in each Monte-Carlo stellar sample was calculated.," For each model, the P and PA for each star in each Monte-Carlo stellar sample was calculated."845 The mean PA and its dispersion were ealeulated for one 10x10' field of view ab each sky grid direction for each model., The mean PA and its dispersion were calculated for one $10\arcmin \times 10\arcmin$ field of view at each sky grid direction for each model.846 Cumulative distribution functions (CDFs) of the polarizations from each model were tabulated., Cumulative distribution functions (CDFs) of the polarizations from each model were tabulated.847 The polarization CDFs are shown in Fig., The polarization CDFs are shown in Fig.848" 2 and the quartile values of the cumulative P/12,,, distributions are presented in Table 1..", \ref{CDFs} and the quartile values of the cumulative $P/P_{max}$ distributions are presented in Table \ref{model_table}.849" For example. for model 1. of the P/12,,, values ave smaller (han 0.0038."," For example, for model 1, of the $P/P_{max}$ values are smaller than 0.0038."850 The variation of simulated starlight polarization across (he sky. centered on the Galactic anti-center. for a subset of the S0. AO. DELO. and axisvaumetric analytic magnetic field models ave shown in Figs.," The variation of simulated starlight polarization across the sky, centered on the Galactic anti-center, for a subset of the S0, A0, DEHO, and axisymmetric analytic magnetic field models are shown in Figs."851 3. through 26.., \ref{predic_S0} through \ref{R24_SMB}.852 Predictions towards the Galactic center. and in particular in the Galactic plane. are likely less reliable than the Outer Galaxy ancl hieher Galactic latitudes because of the effects of star formation ancl supernovae which may interact with the large-scale magnetic field.," Predictions towards the Galactic center, and in particular in the Galactic plane, are likely less reliable than the Outer Galaxy and higher Galactic latitudes because of the effects of star formation and supernovae which may interact with the large-scale magnetic field."853 The S0 (Fig. 3)).," The S0 (Fig. \ref{predic_S0}) ),"854 DEIIO (Fig. 15)).," DEHO (Fig. \ref{predic_DEHO}) ),"855 and axisvmmetrie analvtie (Fig. 23)), and axisymmetric analytic (Fig. \ref{predic_AA}) )856 magnetic field models all show similar structures. wilh magnetic field orientations in the disk parallel to the Galactic plane ancl magnetic nulls near Galactic longitudes 90° and 2707. which are along the toroidal magnetic component.," magnetic field models all show similar structures, with magnetic field orientations in the disk parallel to the Galactic plane and magnetic nulls near Galactic longitudes $90\degr$ and $270\degr$, which are along the toroidal magnetic component."857 This is significantly different from the predictions ol the AO models shown in Fig., This is significantly different from the predictions of the A0 models shown in Fig.858 9 where the lield is predominantly perpendicular to the Galactic plane and the magnetic nulls occur al high Galactic latitudes., \ref{predic_A0} where the field is predominantly perpendicular to the Galactic plane and the magnetic nulls occur at high Galactic latitudes.859 Toward the 6=0 and 1807 directions. many of the model predictions become degenerate.," Toward the $\ell=0\degr$ and $180\degr$ directions, many of the model predictions become degenerate."860 In. particular. the three S0. three DEIIO. and 0.=0° analytic models make identical predictions toward these directions al all Galactic Iatitudes.," In particular, the three S0, three DEHO, and $\theta=0\degr$ analytic models make identical predictions toward these directions at all Galactic latitudes."861 The other, The other862f(a)= constant.),$f(a)=$ constant.)863" Thus, defining and with the (flatness) constraint we may rewrite the Friedmann equation as We have here normalized the expansion factor so that απ)=1. which we assume throughout this paper."," Thus, defining and with the (flatness) constraint we may rewrite the Friedmann equation as We have here normalized the expansion factor so that $a(t_0)=1$, which we assume throughout this paper."864" Introducing the cosmological redshitt z, where we can re-arrange this equation to read so that That is where we have also defined the function g(z)=f/(1 zY—I. and <(f,) is the redshift of light reaching us at £j, but emitted at cosmic time ¢,."," Introducing the cosmological redshift $z$, where we can re-arrange this equation to read so that That is where we have also defined the function $g(z)\equiv f/(1+z)^2-1$, and $z(t_e)$ is the redshift of light reaching us at $t_0$ , but emitted at cosmic time $t_e$."865" In this expression, we have used the equality Ry=c/H, which is valid in a flat (A= 0) cosmology."," In this expression, we have used the equality $R_{\rm h}=c/H$, which is valid in a flat $k=0$ ) cosmology."866" Other than this flat condition, Equation (25) is identical to that obtained in the concordance model, subject to the density in Equation (16)."," Other than this flat condition, Equation (25) is identical to that obtained in the concordance model, subject to the density in Equation (16)."867" If we now put f,—0 and c(t.)>oo, then clearly Our proposed form of the gravitational (1.c., Hubble) radius in Equation (8) leads to the equality Αμ)= ομ."," If we now put $t_e\rightarrow 0$ and $z(t_e)\rightarrow \infty$, then clearly Our proposed form of the gravitational (i.e., Hubble) radius in Equation (8) leads to the equality $R_{\rm h}(t_0)=ct_0$ ."868 Therefore; any cosmological model consistent withthe Weyl Postulate and the," Therefore, any cosmological model consistent withthe Weyl Postulate and the"869Period doubling. a phenomenon often observed in dynamical systems. was also found in stellar models and actual »ulsating variables in the last decades.,"Period doubling, a phenomenon often observed in dynamical systems, was also found in stellar models and actual pulsating variables in the last decades."870 Period doubling (PD) means that the observed. quantity of the svstem alternates rclween a high. and low amplitude: evele., Period doubling (PD) means that the observed quantity of the system alternates between a high and low amplitude cycle.871 Dynamical systenis as the simple Rosssler oscillator are usually capable of period doubling bifurcation. and through a series. of rifurcations called the can evolve to chaotic behaviour.," Dynamical systems as the simple Rösssler oscillator are usually capable of period doubling bifurcation, and through a series of bifurcations called the can evolve to chaotic behaviour."872 The very definition of RY Tauri variables was originally the alternation of deep and shallow minima. a clear sign of PD Prestonetal. 1963)).," The very definition of RV Tauri variables was originally the alternation of deep and shallow minima, a clear sign of PD \citealt{preston63}) )."873 “Phe phenomenon was reproduced by Fokin(1994). in radiative stellar models for RV Tauri stars as well as by Saitou.Takeuti&Tanaka(1989) in one-zone stellar models., The phenomenon was reproduced by \citet{fokin94} in radiative stellar models for RV Tauri stars as well as by \citet*{saitou89} in one-zone stellar models.874 Models of Wo Vir variables (Buehler&IxovácsLOST) are also capable of bifurcation cascade towards chaos. and chaotic pulsations wereindeed identified in semiregular stars (Buehleretal.1996.. Ixolláth 1998... Buehler.Ixolláth.&Caclmus 2004)).," Models of W Vir variables \citep{bk87} are also capable of bifurcation cascade towards chaos, and chaotic pulsations wereindeed identified in semiregular stars \citealt{bksm96}, \citealt{kbsm98}, \citealt*{bkc04}) )."875. Period doubling was reported by Wiss&Szatmáry(2002). in the Mira star It Cvg., Period doubling was reported by \citet{kiss02} in the Mira star R Cyg.876 Models of classical pulsators also showed oonmisng results., Models of classical pulsators also showed promising results.877 Moskalik.&Buehler(1990) searched or half-integer resonances both in Cepheicl ancl RR Lyrac models and. indeed found PD in the former case (sec also Duchler&Aloskalik 1990))., \citet{mb90} searched for half-integer resonances both in Cepheid and RR Lyrae models and indeed found PD in the former case (see also \citealt{bm90}) ).878 Α 2:3. resonance between the 'uncdamental mode and the first overtone was also iclentified as a root cause., A 2:3 resonance between the fundamental mode and the first overtone was also identified as a root cause.879 But neither period. doubling nor suitable resonances were found in Rit Lyrae stars between the tundamental or first overtone and any higher modes up to he fourth overtone., But neither period doubling nor suitable resonances were found in RR Lyrae stars between the fundamental or first overtone and any higher modes up to the fourth overtone.880 Aikawa(2001) also reported PD in very ong-period. radiative Cepheid mocdels but. did not identify any underlying resonance.," \citet{aikawa01} also reported PD in very long-period, radiative Cepheid models but did not identify any underlying resonance."881 The signs of possible PD are the hall-integer frequencies (MILES) in the Fourier spectrum. in the form of (27|1)/2fy with respect to the fy main periodicity. sometimes. called as subharmonics.," The signs of possible PD are the half-integer frequencies (HIFs) in the Fourier spectrum, in the form of $ (2n+1)/2\, f_0 $ with respect to the $f_0$ main periodicity, sometimes called as subharmonics."882 Similar peaks were found in some variable white dwarfs as well. suggesting PD in 11251|489 (Cioupiletal.1988). and even signs of four-period. behaviour in €G191-16 (Vauclairetal. 1989)..," Similar peaks were found in some variable white dwarfs as well, suggesting PD in PG1351+489 \citep{goupil88} and even signs of four-period behaviour in G191-16 \citep{vauc89}. ."883 But because exact values usually differ slightlv from hall-inte@er values. they could," But because exact values usually differ slightly from half-integer values, they could"884origin is distinctly clouded.,origin is distinctly clouded.885 In the solar svstem. copper owes ~30% of its abundance to the s-process weak (uassive star) and main (low mass star) components combined 1993).. with the remainder likely a product of some combination of SNe Type Ia and Type II.," In the solar system, copper owes $\sim$ of its abundance to the s-process weak (massive star) and main (low mass star) components combined \citep{Matteucci1993}, with the remainder likely a product of some combination of SNe Type Ia and Type II."886" Both the weak and main s-process are “secondary processes in (he sense that a ""secondary"" element requires a Fe (or Fe-peak) seed nucleus to build upon.", Both the weak and main s-process are “secondary” processes in the sense that a “secondary” element requires a Fe (or Fe-peak) seed nucleus to build upon.887 Such an element is expected to decrease in abundance (relative to Fe) as the stars overall iron content decreases (because the production of the element depends upon the availability of seed Fe nuclei)., Such an element is expected to decrease in abundance (relative to Fe) as the star's overall iron content decreases (because the production of the element depends upon the availability of seed Fe nuclei).888 In contrast. a “primary” process does not require a Fe seed nucleus. aud so the abundance of a primary element (relative to iron) should be constant wilh overall metallicity.," In contrast, a “primary” process does not require a Fe seed nucleus, and so the abundance of a primary element (relative to iron) should be constant with overall metallicity."889 While [Cu/Fe] does decrease with metallicity in the manner of a secondary element. (he s-process contribution to copper is sharply restricted.," While [Cu/Fe] does decrease with metallicity in the manner of a secondary element, the s-process contribution to copper is sharply restricted."890 Increasing the s-process component of copper in s-process nucleosvnthesis models overproduces other light s-process elements (like Sr). and models of chemical evolution [ail to reflect the observed. solar composition (Matteuccietal.1993:Daraffe&Takahashi1993).," Increasing the s-process component of copper in s-process nucleosynthesis models overproduces other light s-process elements (like Sr), and models of chemical evolution fail to reflect the observed solar composition \citep{Matteucci1993,Baraffe1993}."891. The massive star (weak s-process) contribution appears to be an order of magnitude too small to reproduce the Cu content of either the Sun or metal-poor stars. and it must be augmented by explosive nucleosvithesis.," The massive star (weak s-process) contribution appears to be an order of magnitude too small to reproduce the Cu content of either the Sun or metal-poor stars, and it must be augmented by explosive nucleosynthesis."892 It is unclear whether the additional copper in field stars is produced mainly bv Type Ia supernovae (Matteuccietal.1993:Baralle&Takahashi1993) or bv Type 11 supernovae (Iimnmesetal.1995).," It is unclear whether the additional copper in field stars is produced mainly by Type Ia supernovae \citep{Matteucci1993, Baraffe1993} or by Type II supernovae \citep{Timmes1995}."893. Cohen(1978.1979. 1980)s early observations of copper abundances in giant stars of several globular clusters of varving metallicity (—2.4X [Fe/H] < —0.4) vielded somewhat ambiguous resulls.," 	 \citet{Cohen1978, Cohen1979, Cohen1980}' 's early observations of copper abundances in giant stars of several globular clusters of varying metallicity $-2.4 \leq$ [Fe/H] $\leq -0.4$ ) yielded somewhat ambiguous results."894 Cohen found evidence that copper is deficient. wilh respect to iron and that this deficiency may become larger at lower metallicities., Cohen found evidence that copper is deficient with respect to iron and that this deficiency may become larger at lower metallicities.895 ILowever. uncertainties in the derived abundances related to hyperfine broadening in the Cu line made a clearer delineation of any such trend impossible (Cohen1980).," However, uncertainties in the derived abundances related to hyperfine broadening in the Cu line made a clearer delineation of any such trend impossible \citep{Cohen1980}."896. Early studies such as these showed no real abundance trends in globular cluster iron-peak elements. copper included (e.e..," Early studies such as these showed no real abundance trends in globular cluster iron-peak elements, copper included (e.g.,"897much as a factor~2 in planet-star radius for an Earth mass moon with ες=0.01) from the parent star than (hie classical. stellar insolation. zone alone.,"much as a factor$\sim 2$ in planet-star radius for an Earth mass moon with $e_s=0.01$ ) from the parent star than the classical, stellar insolation, zone alone."898" llowever. in addition to the major caveat that the ""energv balance"" represented. by Equation 5 is extremely. approximate. it should be noted that the actual levels of tidal heating required in these zones can be as high as ~10° erg ! ? - which is a [actor ~10* larger than the surface heat flow estimated for Io."," However, in addition to the major caveat that the “energy balance” represented by Equation 5 is extremely approximate, it should be noted that the actual levels of tidal heating required in these zones can be as high as $\sim 10^6$ erg $^{-1}$ $^{-2}$ - which is a factor $\sim 10^3$ larger than the surface heat flow estimated for Io."899 It is an interesting question whether or not this would create such an unstable surface environment that habitability would be compronmised., It is an interesting question whether or not this would create such an unstable surface environment that habitability would be compromised.900 More modest heating levels are required {ο boost temperatures which are already close to temperate., More modest heating levels are required to boost temperatures which are already close to temperate.901 The outer stable orbit (Equation 1) provides a natural scaling for a moon svstem architecture., The outer stable orbit (Equation 1) provides a natural scaling for a moon system architecture.902" A given hypothetical moon must have a semi-major axis of 3aeelop (assuming small eccentricities). where 3<1 and sae’>q""""*,"," A given hypothetical moon must have a semi-major axis of $\beta a_s^{outer}$ (assuming small eccentricities), where $\beta \leq 1$ and $\beta a_s^{outer}\geq a_s^{inner}$."903" We can also write this in terms of tidal heat flow: Thus. for a given exoplanet in (he sample used here we can estimate the (ime-averaged ITj vequired to attain a given 774, (Equation 5). and for an assumed set of moon properties. such as mass. orbital eccentricity. density. rigidity. dissipation. aud atmosphliere we can then evaluate the requirecl orbital radius of the moon aoulcrE"," We can also write this in terms of tidal heat flow: Thus, for a given exoplanet in the sample used here we can estimate the time-averaged $H_T$ required to attain a given $T_{eq}$ (Equation 5), and for an assumed set of moon properties, such as mass, orbital eccentricity, density, rigidity, dissipation, and atmosphere we can then evaluate the required orbital radius of the moon $\beta a_s^{outer}$."904 We have applied (his calculation to the subset of known exoplanets considered above., We have applied this calculation to the subset of known exoplanets considered above.905 Figure T summarizes the results for the example of a 0.1 moon (the minim mass moon likely capable of retaining a terrestrial-tvpe atmosphere) with an orbital eccentricity ol ος=0.01. density p;= 3g oE. and albedo. rigidity and dissipation commensurate wilh (hat estimated for Europa.," Figure 7 summarizes the results for the example of a $0.1$ $_{\oplus}$ moon (the minimum mass moon likely capable of retaining a terrestrial-type atmosphere) with an orbital eccentricity of $e_s=0.01$, density $\rho_s=3$ g $^{-3}$, and albedo, rigidity and dissipation commensurate with that estimated for Europa."906 We have assumed an atmosphere with e=0.62., We have assumed an atmosphere with $\epsilon=0.62$.907" To attain a moon surface temperature 7;,=272A then (as described above) 46 of the exoplanets would require a moon to have My>0 to boost the stellar insolation. lor Ti,=3134Y. then 10 of the exoplanets would require a moon to have fy>0."," To attain a moon surface temperature $T_{eq}=273 K$ then (as described above) 46 of the exoplanets would require a moon to have $H_T>0$ to boost the stellar insolation, for $T_{eq}=373 K$ then 70 of the exoplanets would require a moon to have $H_T>0$."908 In Figure 7 we plot the distribution of both the absolute orbital semianajor axis required [for such moons aud (he ratio of this orbital axis to the inner (a) stable orbital semi-major axis., In Figure 7 we plot the distribution of both the absolute orbital semi-major axis required for such moons and the ratio of this orbital axis to the $a_s^{inner}$ ) stable orbital semi-major axis.909 The latter plot confirms that such moons would reside comfortably outside of the inner Roche limit., The latter plot confirms that such moons would reside comfortably outside of the inner Roche limit.91019343--2026. is a high mass protostellar object candidate located al a kinematic distance (measured using NIL; line velocities) of 4.2 kpe. with a FIR luninosity of 2.7x10! L. (Molinarietal.1996).,"19343+2026, is a high mass protostellar object candidate located at a kinematic distance (measured using $_3$ line velocities) of 4.2 kpc, with a FIR luminosity of $\times$ $^4$ $_{\odot}$ \citep{mol96}."911. This cluster is detected in the GLIMIPSE survey (cluster 24 of Mercer 2005))., This cluster is detected in the GLIMPSE survey (cluster 24 of \citealt*{mcm05}) ).912 Section 2 describes (he detailed observations., Section 2 describes the detailed observations.913 The photometric. spectroscopic and SED modelling results are also presented im (his section.," The photometric, spectroscopic and SED modelling results are also presented in this section."914 The implications of the results are discussed in Section 3., The implications of the results are discussed in Section 3.915 We then summarize our conclusions in Section 4., We then summarize our conclusions in Section 4.916 NIB. photometric imaging observations were made al (he 3.8 m United Kingdom luilrarecl Telescope (UAIRT) with the [αοαν imager UFTI (Rocheetal.2002)., NIR photometric imaging observations were made at the 3.8 m United Kingdom Infrared Telescope (UKIRT) with the facility imager UFTI \citep{roche02}.917. UFTI houses a ILAWAII-I 1024 x 1024 pixel array., UFTI houses a HAWAII-1 1024 $\times$ 1024 pixel array.918" The UFTI plate scale of 0.091"" &ives an available field of view (FOV) of ~90"".", The UFTI plate scale of $\arcsec$ gives an available field of view (FOV) of $\sim 90\arcsec$.919" Photometric observations through J (A=1.25pam... AA=0.16 jn). IH (À=1.64yon.. XA=0.29 jam) and WA=2.20jam.. AA=0.34 jm) broad-band filters were obtained for the IRAS source during the night of 26"" June 2002."," Photometric observations through $J$ $\lambda=1.25$, $\Delta\lambda=0.16$ ), $H$ $\lambda=1.64$, $\Delta\lambda=0.29$ ) and $K$ $\lambda=2.20$, $\Delta\lambda=0.34$ ) broad-band filters were obtained for the IRAS source during the night of $^{th}$ June 2002."920 An integration time of GO sec was used in each of the J. ££ and A band filters: averaging Ποιοί exposures vielded a total exposure Gime of 540 seconds in each band.," An integration time of 60 sec was used in each of the $J$, $H$ and $K$ band filters; averaging jittered exposures yielded a total exposure time of 540 seconds in each band."921" The mean seeing measured was 0.6"" in the A-band images.", The mean seeing measured was $\arcsec$ in the $K$ -band images.922" A nine point (2x3) jittered observing sequence was executed to obtain data that provided final mosaics with a total FOV of —115""x115"".", A nine point $\times$ 3) jittered observing sequence was executed to obtain data that provided final mosaics with a total FOV of $\sim$ $\arcsec\times$ $\arcsec$.923 We note that the signal-to-noise ratio at the edges of these mosaics is lower than within the central area., We note that the signal-to-noise ratio at the edges of these mosaics is lower than within the central area.924 Standard data reduction techniques involving dark subtraction aud mecdian-sky-fLTat-fielding ol the jittered object Iraaes were applied., Standard data reduction techniques involving dark subtraction and median-sky-flat-fielding of the jittered object frames were applied.925 The A-band image of IRAS 193434-2026 is shown in Fig., The $K$ -band image of IRAS 19343+2026 is shown in Fig.926 1l. with an overlay ofSpitzer MIPS 24 jm contours.," 1, with an overlay of MIPS 24 $\mu$ m contours."927 Subsequent (ο our observations. (his region was recently covered by (he UAIRT Infrared Deep Sky Survey (UIXIDS5: Lawrenceοἱal. 2007)) Galactic Plane Survey (GPS).," Subsequent to our observations, this region was recently covered by the UKIRT Infrared Deep Sky Survey (UKIDSS; \citealt*[]{lawrence07}) ) Galactic Plane Survey (GPS)."928 The GPS is an ambitioussurvey of the Northern Galactic plane (Lucasetal.2003)., The GPS is an ambitioussurvey of the Northern Galactic plane \citep{lucas08}.929. The aim of the survey is (o map 1800 square degrees of the plane (|b| « 5 deg) in J. H/. and K to a depth ol J ~ 20.0. HI. — 19.1. ἐν  19.0 at sub-arcsecond resolution.," The aim of the survey is to map 1800 square degrees of the plane $|$ $|$ $<$ 5 deg) in $J$, $H$, and $K$ to a depth of $J$ $\sim$ 20.0, $H$ $\sim$ 19.1, $K$ $\sim$ 19.0 at sub-arcsecond resolution."930 UIXKIDSS emplovs the Wide Field Camera (WFECAXM: Casalietal. 2007)) at UIXIRE., UKIDSS employs the Wide Field Camera (WFCAM; \citealt*[]{cas07}) ) at UKIRT.931 WECAM contains four Rockwell Hawaii-II (11οςΤο 2048x2048 pixel) arravs spaced by in the focal plane., WFCAM contains four Rockwell Hawaii-II (HgCdTe 2048x2048 pixel) arrays spaced by in the focal plane.932 With a pixel scale of the field. of view of each array is 13., With a pixel scale of the field of view of each array is .9337... All UIXIDSS. survey data are reduced. by the, All UKIDSS survey data are reduced by the934principle no natural mechanism in the setup that might break spherical symmetry.,principle no natural mechanism in the setup that might break spherical symmetry.935" So, what was the point in considering different-from-spherical enhangoon shells?"," So, what was the point in considering different-from-spherical enhançoon shells?"936 Consider the case of two monopoles., Consider the case of two monopoles.937 The first brane is placed at the origin so it is one monopole., The first brane is placed at the origin so it is one monopole.938" Now try to bring another brane to wrap on it, another brane on the same footing not a probe."," Now try to bring another brane to wrap on it, another brane on the same footing not a probe."939 The six dimensional supergravity ansatz for the metric of this brane setup is clearly axially symmetric., The six dimensional supergravity ansatz for the metric of this brane setup is clearly axially symmetric.940" So, what is the point in considering that the branes finally melt out in a spherical enhancoon shell?"," So, what is the point in considering that the branes finally melt out in a spherical enhançoon shell?"941 The above argument suggests that there is in principle no especial preference for spherical shells., The above argument suggests that there is in principle no especial preference for spherical shells.942" Moreover, that a non spherical geometry for the enhangoon locus is naturally favoured."," Moreover, that a non spherical geometry for the enhançoon locus is naturally favoured."943" However, although suggesting, the argument is inconclusive."," However, although suggesting, the argument is inconclusive."944 Sphericalness may seem capricious but still within the brane picture whereas it is in a two or higher charge configuration from the field theory point of view., Sphericalness may seem capricious but still within the brane picture whereas it is in a two or higher charge configuration from the field theory point of view.945 A satisfactory dynamical mechanism should rule out spherical enhangoon shells., A satisfactory dynamical mechanism should rule out spherical enhançoon shells.946" Fortunately, there is an elegant way out the spherically symmetric puzzle which involves phenomena beyond SUGRA."," Fortunately, there is an elegant way out the spherically symmetric puzzle which involves phenomena beyond SUGRA."947 Recall that the enhancoon locus is the place where the probe brane (and also the gravitating one as seen by the probe) becomes tensionless., Recall that the enhançoon locus is the place where the probe brane (and also the gravitating one as seen by the probe) becomes tensionless.948 At large distance from r=0 the probe was pointlike in the six noncompact dimensions., At large distance from $r=0$ the probe was pointlike in the six noncompact dimensions.949" However as the distance to the center approaches r,. the brane seems to ""emerge"" and smear on a four dimensional sphere over the noncompact dimensions.", However as the distance to the center approaches $r_e$ the brane seems to “emerge” and smear on a four dimensional sphere over the noncompact dimensions.950" Some lines above, it was explained that the geometry of the enhancoon need not to be spherical but a large family of shapes are both available an consistent with the enhancoon mechanism [14,26].."," Some lines above, it was explained that the geometry of the enhançoon need not to be spherical but a large family of shapes are both available an consistent with the enhançoon mechanism \cite{AM, DJJ}."951 This was a comforting way out of the necessity of spherical symmetry but left us with an uneasy family of arbitrary shapes which are all on equal footing with the sphere., This was a comforting way out of the necessity of spherical symmetry but left us with an uneasy family of arbitrary shapes which are all on equal footing with the sphere.952 Why are all valid?, Why are all valid?953 Is there any mechanism that rules out any of them?, Is there any mechanism that rules out any of them?954 It seems a little naive though to think of such geometries at a region where the branes are blown in new (noncompact) dimensions and spacetime seem not to behave ordinarily., It seems a little naive though to think of such geometries at a region where the branes are blown in new (noncompact) dimensions and spacetime seem not to behave ordinarily.955" Indeed, it is believed that the correct description of the geometry near the enhangcoon locus is a fuzzy sphere [27].."," Indeed, it is believed that the correct description of the geometry near the enhançoon locus is a fuzzy sphere \cite{M}."956" Fuzzy spheres and more generally, non-commutative geometries break unavoidably spherical symmetry."," Fuzzy spheres and more generally, non-commutative geometries break unavoidably spherical symmetry."957" Moreover, Riemannian geometry and the concept of manifold are no longer valid in this context."," Moreover, Riemannian geometry and the concept of manifold are no longer valid in this context."958" It is remarkable, besides, that the fuzzy geometry appears for .N>1 and so the charge-1 monopole does not get affected and recovers spherical symmetry as expected."," It is remarkable, besides, that the fuzzy geometry appears for $N>1$ and so the charge-1 monopole does not get affected and recovers spherical symmetry as expected."959" We will indicate within the next sections the relation between the brane setup, the monopoles and the fuzzy geometry."," We will indicate within the next sections the relation between the brane setup, the monopoles and the fuzzy geometry."960"In Figure 3. we compare the relation between Y ray luminosity and massweighted. temperature. Z5,4. for the WR and LR runs with gravitational heating (GIL) only.","In Figure \ref{fi:lt_res} we compare the relation between $X$ –ray luminosity and mass–weighted temperature, $T_{mw}$, for the HR and LR runs with gravitational heating (GH) only."961 Increasing the resolution has a nonnegligible elect on the estimated Vo ray. luminosity. while. as expected. has only a marginal effect on the massweighted temperature.," Increasing the resolution has a non–negligible effect on the estimated $X$ –ray luminosity, while, as expected, has only a marginal effect on the mass–weighted temperature."962 The £x value for the Εν run of the Virgo cluster is ~25% higher han for the LI run. the dilference increasing to 50% for he Fornax group and to 10054 for the Lickson group.," The $L_X$ value for the HR run of the Virgo cluster is $\sim 25\%$ higher than for the LR run, the difference increasing to $\sim 50\%$ for the Fornax group and to $\sim 100\%$ for the Hickson group."963 We also verified that decreasing the softening by a factor two or the Fornax and Hickson runs. while keeping the mass resolution fixed. increases Ly by aboutπι Lus showing hat the adopted spatial resolution is adequate to resolve all he structures which are responsible for the No ray emission.," We also verified that decreasing the softening by a factor two for the Fornax and Hickson runs, while keeping the mass resolution fixed, increases $L_X$ by about, thus showing that the adopted spatial resolution is adequate to resolve all the structures which are responsible for the $X$ –ray emission."964 In fact. at the highest resolution of our runs. the number of massive subhalos (ic. with circular velocity larger than enth of that of the main halo) is likely to have converged (Ghigna et al.," In fact, at the highest resolution of our runs, the number of massive subhalos (i.e., with circular velocity larger than one--tenth of that of the main halo) is likely to have converged (Ghigna et al."965 2000)., 2000).966 We expect spatial resolution to be even loss of an issue when considering simulations with extra wating (see below)., We expect spatial resolution to be even less of an issue when considering simulations with extra heating (see below).967 Indeed. in this case the σας is put on a ügher adiabat and. therefore. does not follow the smallscale potential wells which are characterized by a virial emperature of a few tens of keV. We stress that the ΕΙ runs predict a Ly 1) relation which agrees well. both in slope and normalization. with the prediction of the sealing model of eq.(2)) for a pure bremsstrahlung cooling function.," Indeed, in this case the gas is put on a higher adiabat and, therefore, does not follow the small--scale potential wells which are characterized by a virial temperature of a few tens of keV. We stress that the HR runs predict a $L_X$ $T$ relation which agrees well, both in slope and normalization, with the prediction of the scaling model of \ref{eq:lth}) ) for a pure bremsstrahlung cooling function."968 We consider this agreement as a convincing indication that the Lt runs provide a correct description of the gas distribution. which requires using at least c5«103 particles within the virial racius.," We consider this agreement as a convincing indication that the HR runs provide a correct description of the gas distribution, which requires using at least $\simeq 5\times 10^4$ particles within the virial radius."969 This result is in agreement with that from resolution studies involving the collisionless component only (e.g. Moore et al., This result is in agreement with that from resolution studies involving the collisionless component only (e.g. Moore et al.970 1998). which established the minimum number of particles required to correctly model the central profile of dark matter halos.," 1998), which established the minimum number of particles required to correctly model the central profile of dark matter halos."971 Fig., Fig.972 3. also highlights a common problem of simulations of large cosmological volumes atresolution: as resolution becomes worse for smaller. masses. Lx becomes underestimated.," \ref{fi:lt_res} also highlights a common problem of simulations of large cosmological volumes at: as resolution becomes worse for smaller masses, $L_X$ becomes underestimated."973 For instance. taking the same mass and spatial resolution for the Virgo. Fornax and Lickson runs would produce spurious steepening of the relation.," For instance, taking the same mass and spatial resolution for the Virgo, Fornax and Hickson runs would produce spurious steepening of the relation."974 Finally. while a pure bremsstrahlung emissivity is à σου approximation at Z22 keV. the effect. of line emission is shown to become important for smaller/cooler systems.," Finally, while a pure bremsstrahlung emissivity is a good approximation at $T\magcir 2$ keV, the effect of line emission is shown to become important for smaller/cooler systems."975 Lhe net effect of properly accounting for them is that of Hlattening the relation. thus further increasing the ciserepaney with respect to observations at the scale of poor clusters and. groups.," The net effect of properly accounting for them is that of flattening the relation, thus further increasing the discrepancy with respect to observations at the scale of poor clusters and groups."976This is illustrated iu the iuset pancl in Fie.,This is illustrated in the inset panel in Fig.977 H4. (owervieht).," \ref{maxdisk}978 (lower-right)."979 Using the aremuents m Zwaanctal.(1995) we fud that Sy(Y)? needs to be coustaut for galaxies to obev a TF relation independent of surface brightucss., Using the arguments in \citet{zwaan_tf95} we find that $\Sigma_0(\Upsilon)^2$ needs to be constant for galaxies to obey a TF relation independent of surface brightness.980 If all galaxies were truly maxiuun disk Gu the seuse that Vuax(disk) 2Vinax (observed). oue could replace this by the requirement that ου(ΓιHas needs to be constant.," If all galaxies were truly maximum disk (in the sense that $V_{\rm max}$ (disk) $\simeq V_{\rm max}$ (observed), one could replace this by the requirement that $\Sigma_0(\Upsilon_\star)_{\rm max}^2$ needs to be constant."981 The lower-right pancl shows that this is not the case: at fixed Vas there is a substantial scatter which would translate in D mag scatter in TF., The lower-right panel shows that this is not the case: at fixed $V_{\rm max}$ there is a substantial scatter which would translate in $\sim 5$ mag scatter in TF.982 Clearly the observed scatter is ος snidler. and this shows the clear need for au additional mass component to make TE work.," Clearly the observed scatter is much smaller, and this shows the clear need for an additional mass component to make TF work."983 Iu other words. naxiuun disk for ealaxies aud TF are incompatible.," In other words, maximum disk for galaxies and TF are incompatible."984 The imost mnuportaut conclusion frou this work is that the huge majority of the ligh-resolution rotation curves presented here prefer the pseudo-isothermal core-donünated halo model., The most important conclusion from this work is that the large majority of the high-resolution rotation curves presented here prefer the pseudo-isothermal core-dominated halo model.985 For a small nuuber of galaxies neither the pseudo-isothermal nor the NFW models are an adequate description of the data., For a small number of galaxies neither the pseudo-isothermal nor the NFW models are an adequate description of the data.986 This should not come as a surprise as the true DAL distribution is likely to be more complex than the models presented here., This should not come as a surprise as the true DM distribution is likely to be more complex than the models presented here.987 Nevertheless. the general trend is that for alinost all galaxies discussed here the relative quality of the fits using the pseudo-isothermal model is better than those for the NFW 1nocdel.," Nevertheless, the general trend is that for almost all galaxies discussed here the relative quality of the fits using the pseudo-isothermal model is better than those for the NFW model."988 For a siuiall umuer of galaxies the NEW model provides a eood fit. but geucrally the couceutrations derived from the observed rotation curves are lower than predicted by the simulations.," For a small numer of galaxies the NFW model provides a good fit, but generally the concentrations derived from the observed rotation curves are lower than predicted by the simulations."989 This is hard to fix: the most likely effect which iav alter the initial cosmological NEW halo is adiabatic contraction. but this has the effect of making the fal (observed) halo couceutrated. so one would have to start off with (cosimologically relevant) halos that are even concentrated.," This is hard to fix: the most likely effect which may alter the initial cosmological NFW halo is adiabatic contraction, but this has the effect of making the final (observed) halo concentrated, so one would have to start off with (cosmologically relevant) halos that are even concentrated."990" It is worrving that for one or two extreme cases he difference between ""CDAML does work” or ""CDM doesu't work” depends on subtle differences in data. data iudlius. or analysis."," It is worrying that for one or two extreme cases the difference between “CDM does work” or “CDM doesn't work” depends on subtle differences in data, data handling, or analysis."991 Figure 7 illustrates that opposite clainis can sometimes be made from the same data., Figure \ref{nfwswaters} illustrates that opposite claims can sometimes be made from the same data.992 Heuce we re-iterate the need for the highest quality data of a aree saniple. in order to ΙΙχο these effects.," Hence we re-iterate the need for the highest quality data of a large sample, in order to minimize these effects."993 We refer to deBloketal.(2001) where it is shown hat data prescuted here are consistent with a core-douninated model: the good NEW fits that are fouud for a number of LSB ealaxics can be attributed to resolution effects., We refer to \citet{paper3} where it is shown that data presented here are consistent with a core-dominated model; the good NFW fits that are found for a number of LSB galaxies can be attributed to resolution effects.994 We stmunarize our results as follows., We summarize our results as follows.995 We thank Roclof Bottema aud Rob Swaters for their helpful comments on carly drafts of this paper., We thank Roelof Bottema and Rob Swaters for their helpful comments on early drafts of this paper.996 We thas the anonvinous referee for a thorough examination of the data., We thank the anonymous referee for a thorough examination of the data.997 The work of SSM is supported iu part by NSF eraut. AST9901663., The work of SSM is supported in part by NSF grant AST9901663.998 This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Tustitute of Techuologv. uncer contract with the National Acronauties and Space Adiinistration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."999 This rescarch has made use of NASA's Astrophysics Data System Abstract Service., This research has made use of NASA's Astrophysics Data System Abstract Service.1000has led to the notion of the condensation rank of topological spaces. see [11]. and also cf.,"has led to the notion of the condensation rank of topological spaces, see \cite{mg} and also c.f."1001 [ο for the classical rank due to Cantor-Dendixson., \cite{CP} for the classical rank due to Cantor-Bendixson.1002 The measure-preserving property of the condensation derivative ancl its iteration up to the condensation rank of the space provides a sufficient tool to prove a modified version of the Aleksandrov's theorem: a set with finite and positive measure in any regular non-atomic Borel measure space contains a perlect sel whose measure is positive., The measure-preserving property of the condensation derivative and its iteration up to the condensation rank of the space provides a sufficient tool to prove a modified version of the Aleksandrov's theorem: a set with finite and positive measure in any regular non-atomic Borel measure space contains a perfect set whose measure is positive.1003" It is also shown that lor any ordinal number. sav a. there exists an appropriate totallv imperlect IHauscdorff topological space whose condensation rank is a. In the case of non-limit ordinal numbers. the space can be a totally imperfect ""compact space. see [L1].."," It is also shown that for any ordinal number, say $\alpha,$ there exists an appropriate totally imperfect Hausdorff topological space whose condensation rank is $\alpha.$ In the case of non-limit ordinal numbers, the space can be a totally imperfect “compact"" space, see \cite{mg}."1004 Some relevant discussions and applications in locally compact groups are also presented in |12].., Some relevant discussions and applications in locally compact groups are also presented in \cite{mgph}.1005 However. there has not been any discussion on or result of the condensation rank of Danach spaces.," However, there has not been any discussion on or result of the condensation rank of Banach spaces."1006 In the next section we define (he condensation derivative and the condensation rank., In the next section we define the condensation derivative and the condensation rank.1007 We also provide some necessary preliminary results from |LI.12]..," We also provide some necessary preliminary results from \cite{mg, mgph}."1008 In section ??.. we present our main result. that is. (he condensation rank of any infinite dimensional injective Banach space is equal to or greater than the first uncountable ordinal number.," In section \ref{inject}, we present our main result, that is, the condensation rank of any infinite dimensional injective Banach space is equal to or greater than the first uncountable ordinal number."1009 Let X be a Hausdorff topological space., Let $X$ be a Hausdorff topological space.1010 A Borel derivative on 25 is a Borel map which is monotone on the closed subsets of X. i.e.D(N)&NK for any closed set Iv.," A Borel derivative on $2^X$ is a Borel map $D:2^X\rightarrow 2^X$ which is monotone on the closed subsets of $X$, i.e.,$D(K)\subseteq K$ for any closed set $K$."1011" For a Borel derivative D:22—25 and an ordinal number o. the a-th iterated derivative D:25—25 is defined inductively as follows: D'(N)=N. D(NK)=DID""(K)) and D(N)=f),.,,DUx). For limit ordinal number a."," For a Borel derivative $D:2^X\rightarrow 2^X$ and an ordinal number $\alpha$ , the $\alpha$ -th iterated derivative $D^\alpha: 2^X\rightarrow 2^X$ is defined inductively as follows: $D^0(K)=K$, $D^{\alpha+1}(K)=D(D^\alpha(K))$ and $D^\alpha(K)=\bigcap_{\beta\prec\alpha} D^\beta(K)$, for limit ordinal number $\alpha$."1012 Each D is à Borel map. c.f [5].. where ihe Borel complexity of the iterations is investigated.," Each $D^\alpha$ is a Borel map, c.f \cite{CM83}, where the Borel complexity of the iterations is investigated."1013 An illuminating presentation of Borel derivatives is given by IXechris [13].., An illuminating presentation of Borel derivatives is given by Kechris \cite{Ke94}.1014 A point p€X is called a condensation point of ACNX if any neighborhood of p contains an uncountable number of points Irom sl., A point $p\in X$ is called a condensation point of $A\subseteq X$ if any neighborhood of $p$ contains an uncountable number of points from $A$.1015 We reler to the set of all condensation points of 4 as the condensation derived set (CDS) of A and denote CD for the condensation derivative. (hat is à set valued funcüon. which maps any set (o its CDs.," We refer to the set of all condensation points of $A$ as the condensation derived set $\rm CDS$ ) of $A$ and denote $\rm CD$ for the condensation derivative, that is a set valued function, which maps any set to its $\rm CDS$."1016 Note that the maximal perfect subset of the closure of a set is called its perfect. kernel., Note that the maximal perfect subset of the closure of a set is called its perfect kernel.1017 Denote ayB as the least ordinal number among5 those which satisfy Lemma 2.1.., Denote $\alpha_A$ as the least ordinal number among those which satisfy Lemma \ref{ordinal}. .1018 Then. ↽↽ bU ⋅ ↕∫≼↽⊲↻∪⋖↜⇀∣↥∏⊳∩≜↓↕⋟∖⊽≀↧↪∖⊽⊔⋅↕≺∢⊔⋡∖⇁≼⇂≼↲⋟∖⊽≺∢≼↲∐≼∐∐≸≟≺∢↥⋯↴↕∐⋅⋅⋅," Then, $\{{\rm CD}^\alpha(A)\}^{\alpha_A}_{\alpha=1}$ is a strictly descendingchain."1019 The idea of the CR defined here is similar. but not identical idea. to the idea of classical rank due Cantor-Dendixson. cL. [6]..," The idea of the ${\rm CR}$ defined here is similar, but not identical idea, to the idea of classical rank due Cantor-Bendixson, c.f. \cite{CP}. ."1020is luminous closer to than to of the time.,is luminous closer to than to of the time.1021 For the duty evele to fully reach a track with either a large (25% or more) radiative efficiency or a steep decline (towards £L) appears to be required.," For the duty cycle to fully reach, a track with either a large $\sim 25\%$ or more) radiative efficiency or a steep decline (towards $t^{-10}$ ) appears to be required."1022" The implication is that SAIBLE evolution might not involve ""Hickering"". i.e. interspersed periods of luminous accretion and. quiescence on à LO’ vear timescale (?).. Ó]n"," The implication is that SMBH evolution might not involve “flickering”, i.e., interspersed periods of luminous accretion and quiescence on a $\sim 10^7$ year timescale \citep{Hatziminaoglou2001}."1023nstead. a simpler picture suggests itself: (1) the SALDLIL is seeded: (2) the SMDII grows until it enters the SDSS catalog at the low-mass end and at the Iddington uminositv: (3) the SMBLIID accretes as a Type 1 quasar for 12 Gyr while the Eddington ratio declines sharply: and (4) the SIBLE permanently ceases its rapid. unobscured. unminous accretion. with various possible post-turnolf states including dquiescence. Sevferts. ancl Type 2 quasars.," Instead, a simpler picture suggests itself: (1) the SMBH is seeded; (2) the SMBH grows until it enters the SDSS catalog at the low-mass end and at the Eddington luminosity; (3) the SMBH accretes as a Type 1 quasar for $1-2$ Gyr while the Eddington ratio declines sharply; and (4) the SMBH permanently ceases its rapid, unobscured, luminous accretion, with various possible post-turnoff states including quiescence, Seyferts, and Type 2 quasars."1024 “Phis λογο also requires a low value of 5. corresponding to decline proportional to f° or sleeper.," This picture also requires a low value of $\kappa$ , corresponding to decline proportional to $t^{-6}$ or steeper."1025 The possibility ofa short. 1.2 €ivr quasar lifetime is xurticularlv intriguing in light of two additional results from he AJ.L plane.," The possibility of a short, $1-2$ Gyr quasar lifetime is particularly intriguing in light of two additional results from the $M-L$ plane."1026 The characteristic luminosity for quasar acecretion at fixed mass and redshift requires that accretion rates be svnchronizec to within 1.2 Gwe (7)..., The characteristic luminosity for quasar accretion at fixed mass and redshift requires that accretion rates be synchronized to within $1-2$ Gyr \citep{Steinhardt2010b}.1027 Also. quasar urnoll is svnchronized. depending upon the mass. to within 1753 Gor for Aledf10AJ. (2).," Also, quasar turnoff is synchronized, depending upon the mass, to within $0.75-3$ Gyr for $M_BH > 10^9 M_\odot$ \citep{Steinhardt2010b}."1028 Perhaps the similarity of these three synchronization timescales could be explaine » à svnchronization in the times with which quasars turn ancl follow a common track in the ALL plane combine with short lifetimes., Perhaps the similarity of these three synchronization timescales could be explained by a synchronization in the times with which quasars turn and follow a common track in the $M-L$ plane combined with short lifetimes.1029 Ashort-lived Pype 1 quasar phase might seem to violate the Soltan argument because the SMDII spends most. of its time in another state., A short-lived Type 1 quasar phase might seem to violate the Soltan argument because the SMBH spends most of its time in another state.1030 The Soltan argument shows tha most of the total quasar mass in the universe was accretec Iuminously in Twpe 1 quasar states., The Soltan argument shows that most of the total quasar mass in the universe was accreted luminously in Type 1 quasar states.1031 However. the Soltan argument only places a limit on the last 2-32 e-foldings of mass growth. during which most of the mass is added.," However, the Soltan argument only places a limit on the last 2-3 e-foldings of mass growth, during which most of the mass is added."1032 Prior to these last e-foldings. we do not know how much of the erowth takes place through luminous accretion.," Prior to these last e-foldings, we do not know how much of the growth takes place through luminous accretion."1033 Even these short tracks with decline between //U and £I grow the SALDII by 1-1.4 dex. i.e. 2-3 e-foldings. so such solutions are allowed.," Even these short tracks with decline between $t^{-6}$ and $t^{-10}$ grow the SMBH by 1-1.4 dex, i.e., 2-3 e-foldings, so such solutions are allowed."1034 As shown in Figures Figures 11.12.. 15.. ancl 16.. it is possible that one scaling law for feedback. with universal parameters might be able to describe the evolution of all quasars at all initial masses and times.," As shown in Figures Figures \ref{fig:allpl4panel}, \ref{fig:allexp4panel}, \ref{fig:allbothMgoverlap}, and \ref{fig:allintbothMgoverlap}, it is possible that one scaling law for feedback with universal parameters might be able to describe the evolution of all quasars at all initial masses and times."1035 The existence of a characteristic luminosity at. each combination of mass and redshift’ is insullicient to require such a uniformity among the evolution of individual quasars., The existence of a characteristic luminosity at each combination of mass and redshift is insufficient to require such a uniformity among the evolution of individual quasars.1036 However. (1) the synchronization between quasars at fixed mass. (2) the narrow Luminosity range at fixed mass and redshift. and (3) the sharp peak in number density at a single. characteristic luminosity at fixecl mass and redshift (ef 2))," However, (1) the synchronization between quasars at fixed mass, (2) the narrow luminosity range at fixed mass and redshift, and (3) the sharp peak in number density at a single, characteristic luminosity at fixed mass and redshift (cf. \citet{Steinhardt2010b}) )"1037 make such a moclel intriguing., make such a model intriguing.1038 In this paper. we have investigated tracks for SMDII accretion histories and have shown that the quasar mass-liuminosity plane constrains these mocdels remarkably precisely.," In this paper, we have investigated tracks for SMBH accretion histories and have shown that the quasar mass-luminosity plane constrains these models remarkably precisely."1039 The most intriguing result is Chat we can rule out models in which the SMDILI aceretion rate is proportional to the matter number clensity » in the universe., The most intriguing result is that we can rule out models in which the SMBH accretion rate is proportional to the matter number density $n$ in the universe.1040 Even models in which the accretion rate is proportional to n7 are not allowed without a combination of higher-than-expected racliative cllicteney and a SALBLL that spends most of its time quiescent. rather than in a quasar state.," Even models in which the accretion rate is proportional to $n^2$ are not allowed without a combination of higher-than-expected radiative efficiency and a SMBH that spends most of its time quiescent, rather than in a quasar state."1041 This paper is an intermediate. phenomenological step hat has produced. constraints that seem to be required of heoretical models for SAIBLL growth., This paper is an intermediate phenomenological step that has produced constraints that seem to be required of theoretical models for SMBH growth.1042 The next step is to oduce a physical model of fuelling ancl feedback. leading ο quasar tracks satisving these constraints., The next step is to produce a physical model of fuelling and feedback leading to quasar tracks satisfying these constraints.1043 The authors would like to thank Lars Hernaquist. and orm Murrav for valuable comments., The authors would like to thank Lars Hernquist and Norm Murray for valuable comments.1044 This work was supported in part by Chandra:grant number C07-8136A. This work was supported by World. Premier International tesearch Center Initiative (WPL Initiative). MIZNXT. Japan.," This work was supported in part by Chandragrant number G07-8136A. This work was supported by World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan."1045interact with material within 2 scale heights from their orbit (2)..,interact with material within 2 scale heights from their orbit \citep{bate03}.1046 This is the maximum distance that is considered in. the inset of Fig. 3.., This is the maximum distance that is considered in the inset of Fig. \ref{fig3}.1047 For a Keplerian disc. à gap forms with a half width that is approximately 2H. and in the inset of Fig.," For a Keplerian disc, a gap forms with a half width that is approximately $2H$, and in the inset of Fig."1048 5. all three curves fall below 0.1 for f=1., \ref{fig3} all three curves fall below $0.1$ for $f=1$.1049 For f=0.8. there is still a density depression around the orbit of the planet. but for lower values of f there is no clear gap near r=ry. it has shifted inwards enough so that the planet is basically embedded in the dise again.," For $f=0.8$, there is still a density depression around the orbit of the planet, but for lower values of $f$ there is no clear gap near $r=\rp$, it has shifted inwards enough so that the planet is basically embedded in the disc again."1050 However. since for the lowest values of f there are no more resonances located close to the planet. the Type I torque will not be fully restored.," However, since for the lowest values of $f$ there are no more resonances located close to the planet, the Type I torque will not be fully restored."1051 We find that typically Mf=0.6/I(£1)2 for gap-opening planets., We find that typically $\Gamma(f=0.6)/\Gamma(f=1)\approx 2$ for gap-opening planets.1052 The influence of the head wind (see Eq. 1)), The influence of the head wind (see Eq. \ref{eqhead}) )1053 will be strong. however. since it does not rely on resonances.," will be strong, however, since it does not rely on resonances."1054 We comment that a opening planet is still tidally locked to the gap. just as in the Keplerian Type II migration case.," We comment that a gap-opening planet is still tidally locked to the gap, just as in the Keplerian Type II migration case."1055 The possibility of a high-mass planet fully embedded in the disc may have some important consequences for gas accretion., The possibility of a high-mass planet fully embedded in the disc may have some important consequences for gas accretion.1056 In a Keplerian disc. accretion drops by an order of magnitude when a gap is formed (?)..," In a Keplerian disc, accretion drops by an order of magnitude when a gap is formed \citep{dangelo3D}."1057 In a subkeplerian disc. there a significant amount of mass remains near the orbit of the planet. making accretion potentially very efficient.," In a subkeplerian disc, there a significant amount of mass remains near the orbit of the planet, making accretion potentially very efficient."1058 It is not clear. however. if the planet is able to accept material that has such a high relative velocity.," It is not clear, however, if the planet is able to accept material that has such a high relative velocity."1059 The results for a more massive planet of g=0.001 (1 My around a | Μι star) are very similar to those presented in Fig. 3.., The results for a more massive planet of $q=0.001$ (1 $\mj$ around a 1 $\msun$ star) are very similar to those presented in Fig. \ref{fig3}.1060 The gap shifts inward. and for f<0.7 the planet is located on the outer edge of its own density depression.," The gap shifts inward, and for $f<0.7$ the planet is located on the outer edge of its own density depression."1061 For smaller planets. which only open up a shallow density depression for f=1. remain fully embedded for f£«1. as shown in Fig. 4..," For smaller planets, which only open up a shallow density depression for $f=1$, remain fully embedded for $f<1$, as shown in Fig. \ref{fig4}."1062 While for f=Ll. aq=107+ planet decreases the surface density around its orbit by a factor 0.7. for f«0.7 there is no more evidence for any density depression.," While for $f=1$, a $q=10^{-4}$ planet decreases the surface density around its orbit by a factor $0.7$, for $f<0.7$ there is no more evidence for any density depression."1063 For this lower planet mass. the important resonances become too weak to affect the surface density in strongly subkeplerian dises.," For this lower planet mass, the important resonances become too weak to affect the surface density in strongly subkeplerian discs."1064 We have presented hydrodynamical simulations. of. planets embedded in subkeplerian dises., We have presented hydrodynamical simulations of planets embedded in subkeplerian discs.1065 They represent the first step towards modelling fully 3D magnetised dises., They represent the first step towards modelling fully 3D magnetised discs.1066 It is known for Keplerian disces that magnetic fields can have a strong impact on planet migration: regular fields introduce magnetic resonances (?).. while magnetic turbulence introduces stochastic migration (?2)..," It is known for Keplerian discs that magnetic fields can have a strong impact on planet migration: regular fields introduce magnetic resonances \citep{terquem03}, while magnetic turbulence introduces stochastic migration \citep{nelson04,adamsbloch09}."1067 It remains to be seen what impact a magnetic field configuration that gives rise to subkeplerian dises can have on the simple hydrodynamic picture presented here., It remains to be seen what impact a magnetic field configuration that gives rise to subkeplerian discs can have on the simple hydrodynamic picture presented here.1068 We have worked in the isothermal limit. but. since corotation torques only play à minor role in strongly subkeplerian dises. results for more realistic dises should be similar.," We have worked in the isothermal limit, but since corotation torques only play a minor role in strongly subkeplerian discs, results for more realistic discs should be similar."1069 The two-dimensional approximation. in combination with a gravitational softening parameter of order /. gives similar results to fully three-dimensional simulations. again at least as far as the Lindblad torque is concerned (?)..," The two-dimensional approximation, in combination with a gravitational softening parameter of order $h$, gives similar results to fully three-dimensional simulations, again at least as far as the Lindblad torque is concerned \citep{drag}."1070 We have found that the results do not depend strongly on the initial surface density profile., We have found that the results do not depend strongly on the initial surface density profile.1071 We have considered migration of low-mass planets (the Type I regime). finding that there is a strong dependence on f.," We have considered migration of low-mass planets (the Type I regime), finding that there is a strong dependence on $f$."1072 For |-fs/. the planet feels almost the full one-sided Lindblad torque. which ts a factor 1// stronger than the classical Type I torque.," For $1-f \approx h$, the planet feels almost the full one-sided Lindblad torque, which is a factor $1/h$ stronger than the classical Type I torque."1073 Such a dise would be very hazardous to low-mass planets. since inward migration is sped up by more than an order of magnitude compared to Keplerian dises.," Such a disc would be very hazardous to low-mass planets, since inward migration is sped up by more than an order of magnitude compared to Keplerian discs."1074 For |—-f>hr. the torque decreases because the resonances the planet interacts with become weaker.," For $1-f >h$, the torque decreases because the resonances the planet interacts with become weaker."1075 The dependence of the torque on /r and f is quite complicated. and it is not easy to say for a disc of given. f and /r whether the Type I torque will be stronger of weaker than the head-wind torque.," The dependence of the torque on $h$ and $f$ is quite complicated, and it is not easy to say for a disc of given $f$ and $h$ whether the Type I torque will be stronger of weaker than the head-wind torque."1076 Gap formation proceeds similar to that in. Keplerian discs., Gap formation proceeds similar to that in Keplerian discs.1077 However. because of the inward shift of the important resonances. the gap will be located inside the planet’s orbit.," However, because of the inward shift of the important resonances, the gap will be located inside the planet's orbit."1078 This then leaves the planet on the outside of its ow gap., This then leaves the planet on the outside of its own gap.1079 For strongly subkeplerian dises. a gap-opening planet can become fully embedded again.," For strongly subkeplerian discs, a gap-opening planet can become fully embedded again."1080 Since the torque-generating resonances are located far away. this does not restore the full Type I torque.," Since the torque-generating resonances are located far away, this does not restore the full Type I torque."1081 Accretion time scales could be very short. 1f the planet is able to accept the available matter.," Accretion time scales could be very short, if the planet is able to accept the available matter."1082of candidate YSOs in the IC1396N. elobule.,of candidate YSOs in the IC1396W globule.1083 We distinguish ive eroups of objects., We distinguish five groups of objects.1084 For all five groups we acdcditionallv require the objects to be point. sources and to be spatially ocated. within the area of the globule (<9 distance fron Ηλ 21246|5748): 1n Table 1 we list the coordinates and photometry for the candidates. selected xwed on these four criteria: their positions are overplotted in Fig. 1.., For all five groups we additionally require the objects to be point sources and to be spatially located within the area of the globule $<9'$ distance from IRAS 21246+5743): In Table \ref{t2} we list the coordinates and photometry for the candidates selected based on these four criteria; their positions are overplotted in Fig. \ref{f100}.1085 In total. 15 candidates are identified from colour. variability. ancl spatial position.," In total, 15 candidates are identified from colour, variability, and spatial position."1086 ALL require spectroscopic confirmation. in particular the eroups (iii)-(v).," All require spectroscopic confirmation, in particular the groups (iii)-(v)."1087 Five of them do not show colour excess and could. thus be WETS in 1C1396W. In addition. the elobule is known to harbour the likely Class 0 source LRAS 21246|57435 (?)..," Five of them do not show colour excess and could thus be WTTS in IC1396W. In addition, the globule is known to harbour the likely Class 0 source IRAS 21246+5743 \citep{2003MNRAS.346..163F}."1088 This can be contrasted with the 31 red objects identified in the shallow survey by 2.., This can be contrasted with the 31 red objects identified in the shallow survey by \citet{2003A&A...407..207F}.1089 A pure colour selection. as used by 2.. clearly overestimates he number of YSOs in the cloud.," A pure colour selection, as used by \citet{2003A&A...407..207F}, clearly overestimates the number of YSOs in the cloud."1090 The depth and completeness. of our. sample is not trivial to determine. because we used a variety of cilferent indicators.," The depth and completeness of our sample is not trivial to determine, because we used a variety of different indicators."1091 Our photometry database requires the objects to have uncertainties <0.2 mmage in Ix- and J-band. which elfectively limits the survey ο.)«19 mmag (Sect. 3.1)).," Our photometry database requires the objects to have uncertainties $<0.2$ mag in K- and J-band, which effectively limits the survey to $J<19$ mag (Sect. \ref{ccd}) )."1092" Assuming a maximum extinction of shy= 10mmage and a distance of ppc. this corresponds to AL,~ 7."," Assuming a maximum extinction of $A_V=10$ mag and a distance of pc, this corresponds to $M_J \sim 7$ ."1093 With the MMsyr track by ? this vields a mass limit 0.05AZ... which is applicable to categories. (i). (it) and (ui) above.," With the Myr track by \citet{1998A&A...337..403B} this yields a mass limit $\sim 0.05\,M_{\odot}$, which is applicable to categories (i), (ii) and (iii) above."1094 Since all objects identified in category (iv) have amplitudes < O.lmmag. the depth for this sample is likely to be mmaeg lower. corresponding to ~0.2AZ..," Since all objects identified in category (iv) have amplitudes $<0.1$ mag, the depth for this sample is likely to be mag lower, corresponding to $\sim 0.2\,M_{\odot}$."1095 Thus. the survey covers the regime around the peak in the EME.," Thus, the survey covers the regime around the peak in the IMF."1096 The sample of candidate. YSOs in the categories. (iii)- may contain à substantial [fraction of contaminating background objects., The sample of candidate YSOs in the categories (iii)-(v) may contain a substantial fraction of contaminating background objects.1097 On the other hand our method will miss certain tvpes of sources (see discussion in 3))., On the other hand our method will miss certain types of sources (see discussion in \ref{ident}) ).1098 Lt is sensitive to CTTS and variable WETS. which are the major fractions of objects in regions at MMwyr.," It is sensitive to CTTS and variable WTTS, which are the major fractions of objects in regions at Myr."1099 Fherefore we do not expect that the incompleteness alfects our most relevant result: The number of candidate YSOs in this globule is found to be low. probably less than 10.," Therefore we do not expect that the incompleteness affects our most relevant result: The number of candidate YSOs in this globule is found to be low, probably less than 10."1100 ICL39GW is one of the largest anc most. massive cloucls in the 1€1396 region., IC1396W is one of the largest and most massive clouds in the IC1396 region.1101 ? determined. a e@lobule mass of 400-550AZ. [rom near-infrared extinction. maps. assuming a distance of ppc.," \citet{2005A&A...432..575F}1102 determined a globule mass of $\,M_{\odot}$ from near-infrared extinction maps, assuming a distance of pc."1103 Combined: with the low number of YSOs this indicates low star formation elliciencv., Combined with the low number of YSOs this indicates low star formation efficiency.1104 Before we can quantify this more accurately. however. we re-assess the distance of 11396W. A standard: way of probing distances to dense clouds is by separating foreground ancl background objects using the near-infrared colour and comparing the number of foreground objects with predictions from stellar population models.," Before we can quantify this more accurately, however, we re-assess the distance of IC1396W. A standard way of probing distances to dense clouds is by separating foreground and background objects using the near-infrared colour and comparing the number of foreground objects with predictions from stellar population models."1105 In 1€1396NV this procedure is only applicable to the innermost part of the cloud. where the extinction is strong enough to block the background objects.," In IC1396W this procedure is only applicable to the innermost part of the cloud, where the extinction is strong enough to block the background objects."1106 Within 1.5. from the LRAS source there are 8S objects with unredcdened colours of JÁINo1.0. which agrees well with the typical colours of objects outside the cloud.," Within 1.5' from the IRAS source there are 8 objects with unreddened colours of $J-K \sim 1.0$, which agrees well with the typical colours of objects outside the cloud."1107 The other 14 objects in this area are clearly reddened with J9A>1.5., The other 14 objects in this area are clearly reddened with $J-K>1.5$.1108 We used the Galaxy model from the Besancon to simulate a catalogue of objects over an area of ssedeg in the direction of LCL396\W matching the dynamic range of our catalogue (9.8<A«18.7 mmag)., We used the Galaxy model from the Besancon to simulate a catalogue of objects over an area of sqdeg in the direction of IC1396W matching the dynamic range of our catalogue $9.3<K<18.7$ mag).1109 This vields 4798 objects with distance «750 ppc. which corresponds to 9 objects for an area of &£7 with R=L5'.," This yields 4798 objects with distance $<750$ pc, which corresponds to 9 objects for an area of $\pi R^2$ with $R=1.5'$."1110 Thus. the model predicts a number of foreground objects that is remarkably close to the actual number of foreground objects.," Thus, the model predicts a number of foreground objects that is remarkably close to the actual number of foreground objects."1111 This indicates that the assumption of d=τοῦ ppe isplausible., This indicates that the assumption of $d=750$ pc isplausible.1112 Xssuming Poissonian errors. the distance is unlikely to be below d=600 ppc.," Assuming Poissonian errors, the distance is unlikely to be below $d=600$ pc."1113 Thus the estimates for the cloud. mass reported above, Thus the estimates for the cloud mass reported above1114wwith implications for WL surveys.,with implications for WL surveys.1115" If iis to be used for shape analysis, the sampling rate must be such that the PSF minor-axis FWHM is sampled by >2.8 pixels; otherwise, ffails to converge to a unique shape refnativeconvergence))."," If is to be used for shape analysis, the sampling rate must be such that the PSF minor-axis FWHM is sampled by $\ge2.8$ pixels; otherwise, fails to converge to a unique shape \\ref{nativeconvergence}) )."1116" Similarly, the PSF-smeared objects which produce useful shapes must have a minimum resolution that depends on the significance (or vice-versa) in order to converge to a measurement of the pre-seeing galaxy shape."," Similarly, the PSF-smeared objects which produce useful shapes must have a minimum resolution that depends on the significance (or vice-versa) in order to converge to a measurement of the pre-seeing galaxy shape."1117" If we require shape measurements of all values of e to be successful, then rp>1.8 is required at v=20, or ry>1.2 at v>40 refdeconvconvergence))."," If we require shape measurements of all values of $e$ to be successful, then $r_b\geq1.8$ is required at $\nu=20$, or $r_b\geq1.2$ at $\nu\geq40$ \\ref{deconvconvergence}) )."1118" This implies that improvements in WL statistical accuracy are rapidly limited by the resolution of the galaxies: although going deep in the exposure increases the S/N of the detected objects and reduces the scatter and systematic in the shear estimate, the required S/N grows very rapidly for poorly resolved images, and it becomes difficult to produce additional useful WL data for poorly resolved galaxies."," This implies that improvements in WL statistical accuracy are rapidly limited by the resolution of the galaxies: although going deep in the exposure increases the $S/N$ of the detected objects and reduces the scatter and systematic in the shear estimate, the required $S/N$ grows very rapidly for poorly resolved images, and it becomes difficult to produce additional useful WL data for poorly resolved galaxies."1119" We have also found that larger, well resolved galaxies (rp25) exhibit a problem with the deconvolution method when the PSF is an Airy function, due to the poor approximation of the Airy-functions wings by the GL expansion."," We have also found that larger, well resolved galaxies $r_b\simgeq\,5$ ) exhibit a problem with the deconvolution method when the PSF is an Airy function, due to the poor approximation of the Airy-functions wings by the GL expansion."1120 A simple remedy is available: smooth the PSF and galaxy images before measuring both., A simple remedy is available: smooth the PSF and galaxy images before measuring both.1121" Another solution is to choose an alternative to Gauss-Laguerre as the PSF decomposition basis functions, one that is better suited to describe an Airy function."," Another solution is to choose an alternative to Gauss-Laguerre as the PSF decomposition basis functions, one that is better suited to describe an Airy function."1122" However, a non-GL decomposition will make the deconvolution process excessively complex; another solution is to apodize the telescope to suppress the Airy wings."," However, a non-GL decomposition will make the deconvolution process excessively complex; another solution is to apodize the telescope to suppress the Airy wings."1123" Kuijken(2006) offers an excellent description of the difference between various WL techniques; we refer the reader to this paper for the difference between, for example, the KSB method and Kuijken's method, which is applicable to the difference bewteen KSB method and the EGL method."," \citet{Kuijken} offers an excellent description of the difference between various WL techniques; we refer the reader to this paper for the difference between, for example, the KSB method and Kuijken's method, which is applicable to the difference bewteen KSB method and the EGL method."1124 Our EGL method is similar to Kuijken's polar shapelet method (Kuijken2006)., Our EGL method is similar to Kuijken's polar shapelet method \citep{Kuijken}.1125". What the methods have in common are: the deconvolution of the PSF, which in principle allows for any PSF effects to be removed; forward fitting, which allows error propagation, and hence an error estimate to the measured shape; and the definition of shape as shear, which has a well-defined shear transformation."," What the methods have in common are: the deconvolution of the PSF, which in principle allows for any PSF effects to be removed; forward fitting, which allows error propagation, and hence an error estimate to the measured shape; and the definition of shape as shear, which has a well-defined shear transformation."1126 All of these features contribute to a better shear accuracy., All of these features contribute to a better shear accuracy.1127 The differences between EGL and Kuijken's methods are subtle., The differences between EGL and Kuijken's methods are subtle.1128 The first difference is that Kuijen method works the deconvolution and shearing in shapelet-coefficient space., The first difference is that Kuijken method works the deconvolution and shearing in shapelet-coefficient space.1129 Our EGL method determines the shape-as-shear by iteratively fitting within the pixel space., Our EGL method determines the shape-as-shear by iteratively fitting within the pixel space.1130" Secondly, Kuijken's method obtains the shape using a shear transformation valid to first-order in e, which can be off by up to at e~0.9 (g= 0.6)."," Secondly, Kuijken's method obtains the shape using a shear transformation valid to first-order in $e$, which can be off by up to at $e\simeq0.9$ $g=0.6$ )."1131" Our method uses basis functions that are elliptical, the shear transformation is valid to all orders, which allows for the shape to be measured accurately for any e."," Our method uses basis functions that are elliptical, the shear transformation is valid to all orders, which allows for the shape to be measured accurately for any $e$."1132" 'The third difference is that, in Kuijken's method, only the m—0 terms are used to describe the galaxy."," The third difference is that, in Kuijken's method, only the $m=0$ terms are used to describe the galaxy."1133" This allows for less coefficients needed, and hence is efficient."," This allows for less coefficients needed, and hence is efficient."1134" In comparison, our methodobtains the full set of coefficients to the specified"," In comparison, our methodobtains the full set of coefficients to the specified"1135(Ptak et al. 2006..,"(Ptak et al. \citeyear{ptak06},"1136 Zakamska et al. 20045)., Zakamska et al. \citeyear{zak04}) ).1137 They show a wide range of X-ray luminosities and obscuring column densities., They show a wide range of X-ray luminosities and obscuring column densities.1138 About 40 objects in their sample were detected with IRAS and have infrared luminosities among the most luminous quasars at similar redshift., About 40 objects in their sample were detected with IRAS and have infrared luminosities among the most luminous quasars at similar redshift.1139 The host galaxies are ellipticals. although with irregular morphologies. and the nuclear optical emission is highly polarized (Zakamska et al. 2006)).," The host galaxies are ellipticals, although with irregular morphologies, and the nuclear optical emission is highly polarized (Zakamska et al. \citeyear{zak06}) )."1140 The detection rate in radio (~10%.. Zakamska et al. 2004::," The detection rate in radio $\sim$ , Zakamska et al. \citeyear{zak04};"1141 Vir Lal Ho 2007)) is consistent with that of other AGN types., Vir Lal Ho \citeyear{vir07}) ) is consistent with that of other AGN types.1142 Spectroscopic studies of type 2 quasars have focused so far on identifying emission lines. measuring some basic parameters (line luminosities. redshift. line widths) and searching for correlations among them and with other observables (equivalent widths. color magnitudes. radio luminosities. ete).," Spectroscopic studies of type 2 quasars have focused so far on identifying emission lines, measuring some basic parameters (line luminosities, redshift, line widths) and searching for correlations among them and with other observables (equivalent widths, color magnitudes, radio luminosities, etc)."1143 All this is critical to classify the objects. to investigate the nature of the powering mechanism. the obscuring structure and. ultimately. test the validity of the unification models (e.g. Reyes et al. 2008)).," All this is critical to classify the objects, to investigate the nature of the powering mechanism, the obscuring structure and, ultimately, test the validity of the unification models (e.g. Reyes et al. \citeyear{rey08}) )."1144 However. little work has been done to characterize the gaseous and ionization properties of type 2 quasars.," However, little work has been done to characterize the gaseous and ionization properties of type 2 quasars."1145 Given this lack of knowledge. the primary goal of the work presented here is to use the emission line information to study the excitation mechanism. the physical conditions and ionization properties of the gas.," Given this lack of knowledge, the primary goal of the work presented here is to use the emission line information to study the excitation mechanism, the physical conditions and ionization properties of the gas."1146 Detailed spectroscopic studies of other active galaxy types. such as narrow line radio galaxies (NLRGs. e.g. Robinson et al. 1987))," Detailed spectroscopic studies of other active galaxy types, such as narrow line radio galaxies (NLRGs, e.g. Robinson et al. \citeyear{rob87}) )"1147 have allowed during the last few decades the characterization of the chemical abundances of the gas. the ionization mechanism. the physical and kinematic properties. etc.," have allowed during the last few decades the characterization of the chemical abundances of the gas, the ionization mechanism, the physical and kinematic properties, etc."1148 Similar work must be done for type 2 quasars., Similar work must be done for type 2 quasars.1149 These studies will ultimately provide valuable information about the formation process ofthe host galaxy. the star forming and chemical enrichment histories and the origin of the nuclear activity (e.g. Tadhunter. Fosbury Quinn 1989.. n et al. 2005..," These studies will ultimately provide valuable information about the formation process of the host galaxy, the star forming and chemical enrichment histories and the origin of the nuclear activity (e.g. Tadhunter, Fosbury Quinn \citeyear{tadh89}, n et al. \citeyear{vm05},"1150 Humphrey et al. 2008))., Humphrey et al. \citeyear{hum08}) ).1151 An important disadvantage of NLRG over type 2 quasar studies is that the radio activity. via jet-gas interactions. imprints important distortions on the properties of the ionized gas making it difficult to investigate the intrinsic properties of the host galaxy and environment. as well as the chemical abundances and the nature of the excitation mechanism (e.g. Tadhunter 20023).," An important disadvantage of NLRG over type 2 quasar studies is that the radio activity, via jet-gas interactions, imprints important distortions on the properties of the ionized gas making it difficult to investigate the intrinsic properties of the host galaxy and environment, as well as the chemical abundances and the nature of the excitation mechanism (e.g. Tadhunter \citeyear{tadh02}) )."1152 The study of radio-quiet type 2 quasars should not suffer from such effects., The study of radio-quiet type 2 quasars should not suffer from such effects.1153" Throughout this paper we assume O4= 0.73. ©,,, = 0.27 and ily =7i kms ! Μρο|l"," Throughout this paper we assume $\Omega_{\Lambda} =$ 0.73, $\Omega_{m}$ = 0.27 and $H_{0}$ = 71 km $^{-1}$ $^{-1}$."1154 The objects studied in this paper are a sub-sample of candidate type 2 quasars from the Sloan Digital Sky Survey (SDSS) selected by Zakamskaetal.(2003) in the redshift range 0.37-2 550.8 (~ 145 objects).," The objects studied in this paper are a sub-sample of candidate type 2 quasars from the Sloan Digital Sky Survey (SDSS) selected by \cite{zak03}1155 in the redshift range $\la z \la$ 0.8 $\sim$ 145 objects)."1156 The selection criteria applied by Zakamskaetal.(2003) are listed below (notice that not necessarily all criteria apply to all objects., The selection criteria applied by \cite{zak03} are listed below (notice that not necessarily all criteria apply to all objects.1157 See Zakamska et al., See Zakamska et al.1158 2003. for more detailed The sub-sample studied here contains 50 type 2 quasars. which have been selected to span the full z range and the full [ΟΠΗ luminosity range (223. 107 L. ) of the original type 2 quasar sample.," \citeyear{zak03} for more detailed The sub-sample studied here contains 50 type 2 quasars, which have been selected to span the full $z$ range and the full [OIII] luminosity range $\ge$ $\times$ $^8$ $_ {\odot}$ ) of the original type 2 quasar sample."1159 We did not set constraints on the line equivalent widths. although the selection criteria on the original sample do contain an EW criterion. as stated above.," We did not set constraints on the line equivalent widths, although the selection criteria on the original sample do contain an EW criterion, as stated above."1160 The spectra were corrected for Galactic extinction., The spectra were corrected for Galactic extinction.1161 A galaxy template spectrum was subtracted for objects with low line EWs. to correct for possible underlying stellar and interstellar absorption (Zhang. Dultzin-Hacyan Wang 2007).," A galaxy template spectrum was subtracted for objects with low line EWs, to correct for possible underlying stellar and interstellar absorption (Zhang, Dultzin-Hacyan Wang \citeyear{zh07}) )."1162 This procedure was necessary for a small fraction of objects (~ only)., This procedure was necessary for a small fraction of objects $\sim$ only).1163 The spectra were not corrected for internal dust reddening because such correction was not possible for all objects., The spectra were not corrected for internal dust reddening because such correction was not possible for all objects.1164 Tt should not affect our conclusions since we will compare our results with previous works on other AGN types (Seyfert 2s. radio galaxies). in which no internal extinction correction was applied either.," It should not affect our conclusions since we will compare our results with previous works on other AGN types (Seyfert 2s, radio galaxies), in which no internal extinction correction was applied either."1165" We investigate in this section the dominant ionizing mechanism of the optical line emitting gas in the type 2 quasar sub-sample: AGN vs. stellar photoionization,", We investigate in this section the dominant ionizing mechanism of the optical line emitting gas in the type 2 quasar sub-sample: AGN vs. stellar photoionization.1166 We will ignore shocks in ourdiscussion as an alternative ionization mechanism (e.g. Dopita Sutherland 906)., We will ignore shocks in ourdiscussion as an alternative ionization mechanism (e.g. Dopita Sutherland \citeyear{dop96}) ).1167 In our sub-sample. 44 out of the 47 objects for which radio information is available are radio quiet (Le. Liaw. «107 ere 1 3 Land therefore. shocks induced by the radio structures are not expected to play a significant role in the ionization of the gas (e.g. Clark et al. 1998..," In our sub-sample, 44 out of the 47 objects for which radio information is available are radio quiet (i.e. $L_{1.4 GHz}<$ $^{31}$ erg $^{-1}$ $^{-1}$ $^{-1}$ ) and therefore, shocks induced by the radio structures are not expected to play a significant role in the ionization of the gas (e.g. Clark et al. \citeyear{clark98},"1168 n et al. 1999))., n et al. \citeyear{vm99}) ).1169 In the vast majority of the type 2 quasars in the sample considered here. the bulk of the line profiles is characterized by rather quiescent kinematies. rather than perturbed. as one would expect if shocks were present.," In the vast majority of the type 2 quasars in the sample considered here, the bulk of the line profiles is characterized by rather quiescent kinematics, rather than perturbed, as one would expect if shocks were present."1170 We show in Fig., We show in Fig.1171 | several diagnostic diagrams involving optica emission lines in which we plot the location of the SDSS type 2 quasar sub-sample (green. solid circles and blue solid triangles).," 1 several diagnostic diagrams involving optical emission lines in which we plot the location of the SDSS type 2 quasar sub-sample (green, solid circles and blue solid triangles)."1172 For comparison. we plot in the same diagrams the locus of HIT galaxies from the catalogue of Terlevichetal.(1991). (magenta small symbols).," For comparison, we plot in the same diagrams the locus of HII galaxies from the catalogue of \cite{ter91} (magenta small symbols)."1173 These are absent in diagrams involving the Hell and [NeV] because such lines are rarely detected in this objec class., These are absent in diagrams involving the HeII and [NeV] because such lines are rarely detected in this object class.1174 The solid lines represent the standard “Usequenceof photoionization models. buillwiththemallipurposccodeAl APL ," The solid lines represent the standard $U$ sequence of photoionization models, built with the multipurpose code MAPPINGS Ic (Binette, Dopita Tuohy \citeyear{bin85}; Ferruit et al. \citeyear{fer97}) )"1175G5xv UY ," that reproduces some of the main properties of the emission line spectra of narrow line radio galaxies at different redshifts (e.g. Robinson et al. \citeyear{rob87}, ,"1176, Humphrey et al. \citeyear{hum08}) ).1177with a eut off energy of 50 keV. The clouds are," The ionizing continuum isa power law of index $\alpha$ =1.5 $F_{\nu} \propto \nu^{-\alpha}$ ), with a cut off energy of 50 keV. The clouds are"1178ereater than that of the narrow component.,greater than that of the narrow component.1179 No strong trends can be seen with either date of observation or orbital phase. with the exception of the equivalent width of the narrow component. which has been generally decreasing since 1976 (except for the extremely low value in April 1978).," No strong trends can be seen with either date of observation or orbital phase, with the exception of the equivalent width of the narrow component, which has been generally decreasing since 1976 (except for the extremely low value in April 1978)."1180 Infrared. spectra were obtained one orbit after the optical spectra., Infrared spectra were obtained one orbit after the optical spectra.1181 The mean spectrum is shown in Fig. 3.., The mean spectrum is shown in Fig. \ref{fig:IRspectrum}.1182 Bright emission lines of Br aand A2.058 wavere visible in the spectrum., Bright emission lines of Br $\gamma$ and $\lambda$ were visible in the spectrum.1183 TFhese two lines are generally observed in the Ix-band. spectra of low- and high-mass X-rav binaries 1999))) and may arise in accretion disces or larger emitting regions. such as dense winds fron the mass donors and/or ejecta from the binaries.," These two lines are generally observed in the K-band spectra of low- and high-mass X-ray binaries ) and may arise in accretion discs or larger emitting regions, such as dense winds from the mass donors and/or ejecta from the binaries."1184 Phere also seems to be a hint of asymmetry in the blue wing of the infrared lines. so we again fit. gaussians to the line profile.," There also seems to be a hint of asymmetry in the blue wing of the infrared lines, so we again fit gaussians to the line profile."1185 We modelled cach emission line as the sum of two gaussians. constraining the widths anc velocities of the blue anc red components to be the same for both lines.," We modelled each emission line as the sum of two gaussians, constraining the widths and velocities of the blue and red components to be the same for both lines."1186 Phe results are shown in Table 3.., The results are shown in Table \ref{tab:IRspec}.1187 Both lines are well fit bv one gaussian at |450lans.P. a velocity. very similar to the narrow red-shifted: optical component. plus a bluc-shifted component at a velocity of —2000kms+.," Both lines are well fit by one gaussian at $+450\kms$, a velocity very similar to the narrow red-shifted optical component, plus a blue-shifted component at a velocity of $\sim -2000\kms$."1188 Initially. it was suggested that Cir X-1 is a high mass binary Consisting of a compact star (cither a neutron star or black hole) and an OB supergiant companion star 1980))).," Initially, it was suggested that Cir X-1 is a high mass binary consisting of a compact star (either a neutron star or black hole) and an OB supergiant companion star )."1189 Its X-ray properties imply a very eccentric binary orbit., Its X-ray properties imply a very eccentric binary orbit.1190 Phe 16.6-clay modulation in the X-ray luminosity is due to orbital variations in the mass accretion of the compact star., The 16.6-day modulation in the X-ray luminosity is due to orbital variations in the mass accretion of the compact star.1191 Phe X-ray. ancl radio bursts occur when the compact star encounters the dense stellar wind from the supereiant1980).. or when tidal mass transfer is induced. during the periastron passage1987).," The X-ray and radio bursts occur when the compact star encounters the dense stellar wind from the supergiant, or when tidal mass transfer is induced during the periastron passage."1192. The Type E X-ray bursts that were discovered in a brief epoch in the mid-1980s indicated that the compact star is a neutron star., The Type I X-ray bursts that were discovered in a brief epoch in the mid-1980s indicated that the compact star is a neutron star.1193 Recent observations citesnp91.gla94)). indicate that the companion star is unlikely to be a massive supergiant star., Recent observations ) indicate that the companion star is unlikely to be a massive supergiant star.1194 No Ένρο I bursts have vet been reported in ΗΝΕΣ observations., No Type I bursts have yet been reported in RXTE observations.1195 Llere. we propose a low-mass binary model. where the system consists of a neutron star orbiting around a subeiant companion star of about 3 to ((Figs. 4: ," Here, we propose a low-mass binary model, where the system consists of a neutron star orbiting around a subgiant companion star of about 3 to (Fig. \ref{fig:model}; ;"1196see also Tennant Wu ‘Phe orbital eccentricity is 0.7 0.9., see also Tennant Wu The orbital eccentricity is $\sim 0.7$ –0.9.1197 During the periastron passage. the companion star overfills its Roche-lobe. causing a transfer of mass at à super-IExldington rate onto the neutron star.," During the periastron passage, the companion star overfills its Roche-lobe, causing a transfer of mass at a super-Eddington rate onto the neutron star."1198 As Cir N-1 is a X-ray. burster. the magnetic field. of the neutron star is relatively weak. and so the aceretion [low is probably quasi-spherical during the periastron passage.," As Cir X-1 is a X-ray burster, the magnetic field of the neutron star is relatively weak, and so the accretion flow is probably quasi-spherical during the periastron passage."1199 Because of the Large radiative pressure of emission. [rom the neutron star. there must also be a strong (anisotropic) matter outflow.," Because of the large radiative pressure of emission from the neutron star, there must also be a strong (anisotropic) matter outflow."1200 The inllow/outllow geometry may be related to the larger scale jets observed. from. the system. with symmetrical collimated outllows along some preferred axis.," The inflow/outflow geometry may be related to the larger scale jets observed from the system, with symmetrical collimated outflows along some preferred axis."

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