source,target
" The only exception was G337.826+0.0 for which we calculated the integrated intensity by determining the area within the FHWM of the '?CO emission, as this line-of-sight shows complex velocity structure."," The only exception was G337.826+0.0 for which we calculated the integrated intensity by determining the area within the FHWM of the $^{13}$ CO emission, as this line–of–sight shows complex velocity structure."
 Based on the !?CO line parameters we identify 58 π]] components associated with dense molecular gas., Based on the $^{13}$ CO line parameters we identify 58 ] components associated with dense molecular gas.
" All of them also show ""CO emission while 12 show C!*O emission.", All of them also show $^{12}$ CO emission while 12 show $^{18}$ O emission.
 The remaining diffuse atomic and/or diffuse molecular u]]-emitting clouds that do not have ?CO counterparts are discussed by ? and ?..," The remaining diffuse atomic and/or diffuse molecular ]–emitting clouds that do not have $^{13}$ CO counterparts are discussed by \citet{Langer2010}
 and \citet{Velusamy2010}."
" In the left panel of Figure 2,, we summarize the observed characteristics by plotting the u]]CO and [Cnu]]/?CO integrated intensity ratios for the identified components as a function of the u]] integrated intensity."," In the left panel of Figure \ref{fig:results_pdr}, we summarize the observed characteristics by plotting the $^{12}$ CO and $^{13}$ CO integrated intensity ratios for the identified components as a function of the ] integrated intensity."
 The ratios are calculated from integrated intensities in units of ss!, The ratios are calculated from integrated intensities in units of $^{-1}$.
 The mean value and standard deviation are 0.29 and 0.6 for the n]]/?CO integrated intensity ratio and 1.75 and 2.54 for n]]/?CO., The mean value and standard deviation are 0.29 and 0.6 for the $^{12}$ CO integrated intensity ratio and 1.75 and 2.54 for $^{13}$ CO.
 The ratios vary over 2 orders of magnitude suggesting a wide range of physical conditions in our sample., The ratios vary over 2 orders of magnitude suggesting a wide range of physical conditions in our sample.
 We use the 1]]/? CO and i]]/?CO integrated intensity ratios to constrain the physical conditions of the line-emitting gas., We use the $^{12}$ CO and $^{13}$ CO integrated intensity ratios to constrain the physical conditions of the line–emitting gas.
" The '*CO emission, which becomes optically thick quickly after a modest fraction of the gas-phase carbon is converted to CO, is not very sensitive to the FUV radiation field, as the temperature at the C*/C°/CO transition layer is also insensitive to this quantity (??).."," The $^{12}$ CO emission, which becomes optically thick quickly after a modest fraction of the gas–phase carbon is converted to CO, is not very sensitive to the FUV radiation field, as the temperature at the $^+$ $^0$ /CO transition layer is also insensitive to this quantity \citep{Wolfire1989,
Kaufman99}. ."
" Therefore, the u]]/"" CO ratio is determined by the column density of C* and the temperature at the surface of the PDR, which are in turn dependent on the FUV radiation field and H5 density."," Therefore, the $^{12}$ CO ratio is determined by the column density of $^+$ and the temperature at the surface of the PDR, which are in turn dependent on the FUV radiation field and $_2$ density."
 The n]]/?CO ratio is proportional to the ratio between the Ct and CO column densities., The $^{13}$ COratio is proportional to the ratio between the $^+$ and $^{13}$ CO column densities.
" It therefore gives, provided that extra constraints on the total column of material are available and that there are no significant variations of the FUV field within the beam, a constraint on the location of the C*/C°/CO transition layer which in turn depends on the strength of the FUV field and H» density."," It therefore gives, provided that extra constraints on the total column of material are available and that there are no significant variations of the FUV field within the beam, a constraint on the location of the $^+$ $^0$ /CO transition layer which in turn depends on the strength of the FUV field and $_2$ density."
 We compare the observed 1]]/2ΟΟ and [Cu]]/CO integrated intensity ratios with the results of a PDR model grid in order to constrain physical conditions of the u]]-emitting clouds., We compare the observed $^{12}$ CO and $^{13}$ CO integrated intensity ratios with the results of a PDR model grid in order to constrain physical conditions of the ]–emitting clouds.
 The model grid was calculated using the KOSMA-r PDR model (??) which is available online?.," The model grid was calculated using the $\tau$ PDR model \citep{Stoerzer96,Roellig06} which is available ."
". The model provides a self-consistent solution of the chemistry and thermal balance of a spherical cloud, with a truncated density profile, which is illuminated isotropically by a FUV radiation field."," The model provides a self–consistent solution of the chemistry and thermal balance of a spherical cloud, with a truncated density profile, which is illuminated isotropically by a FUV radiation field."
" The density distribution has the form, n(r)2n,(r/r.)|? for 02r. <r€re and a constant density of n(r)= n5(0.2)-!? in the central region of the cloud (r« 0.2r,)."," The density distribution has the form, $n(r)$ $n_{s}(r/r_c)^{-1.5}$ for $r_c$ $\leq$ $ r \leq r_c$ and a constant density of $n(r)=n_s$ $^{-1.5}$ in the central region of the cloud $r <
0.2r_c$ )."
 Here γε is the cloud radius and n; is the density at the cloud surface., Here $r_c$ is the cloud radius and $n_{s}$ is the density at the cloud surface.
" Note that with a power-law index of 1.5, the average density of the clump is about twice the density at the cloud surface."," Note that with a power--law index of 1.5, the average density of the clump is about twice the density at the cloud surface."
 The line intensities are calculated using a non-LTEradiativetransfer code by ?.., The line intensities are calculated using a non–LTEradiativetransfer code by \cite{Gierens92}. .
" Each model is characterized by the clump mass, the density at the cloud surface, and strength of the FUV field."," Each model is characterized by the clump mass, the density at the cloud surface, and strength of the FUV field."
" The clump mass ranges from 107? to 100 ΜΜΕ, the density at the cloud surface from 10? to 106 ccm?, and the strengthof the FUV field from Xo=1 to 10° (in units of the 2 field?))."," The clump mass ranges from $10^{-2}$ to $100$ $_\odot$, the density at the cloud surface from $^3$ to $^6$ $^{-3}$, and the strengthof the FUV field from $\chi_{\rm
0}=1$ to $^6$ (in units of the \citealt{Draine78} )."
 We do not use and, We do not use and
The search for molecular lines at high redshifts ἐς21) offers one of the richest and most exciting avenues to follow in the study of galaxy formation and evolution (Solomon Vanden Bout 2005).,The search for molecular lines at high redshifts $z\gs 1$ ) offers one of the richest and most exciting avenues to follow in the study of galaxy formation and evolution (Solomon Vanden Bout 2005).
" The detections of rotational transitions of CO in extremely distant quasars (Carilli et 22002: Walter et 22003). high-z radio galaxies (Papadopoulos οἱ 22000: De Brenek et 22005) and submillimeter (submnm) selected. galaxies (Fraver et 11998. 1999: Neri οἱ 22003). have established that the most luminous ancl extreme objects in the early Universe harbour vast amounts of molecular gas (~LOM1M, | Neri οἱ 22003: Greve οἱ 22005) and have large dynamical masses (~LOYDAL. Genzel et 22003: Tacconi et 22006)."," The detections of rotational transitions of CO in extremely distant quasars (Carilli et 2002; Walter et 2003), high-z radio galaxies (Papadopoulos et 2000; De Breuck et 2005) and submillimeter (submm) selected galaxies (Frayer et 1998, 1999; Neri et 2003), have established that the most luminous and extreme objects in the early Universe harbour vast amounts of molecular gas $\sim 10^{10-11}\,\Msolar$ – Neri et 2003; Greve et 2005) and have large dynamical masses $\sim 10^{11}\,\Msolar$ – Genzel et 2003; Tacconi et 2006)."
 Modelling of the bulk physical conditions of the molecular gas in hieh-z QSOs have even been possible in cases where several CO lines have been detected eel 22005)., Modelling of the bulk physical conditions of the molecular gas in $z$ QSOs have even been possible in cases where several CO lines have been detected et 2005).
 The success of CO observations is primarily due to its much larger abundance compared to other species as well asils easily excitable low-/ (CO J4-1—JJ« 3) rotational lines. whieh makes it an excellent tracer of the total amount of 7) molecular gas.," The success of CO observations is primarily due to its much larger abundance compared to other species as well asits easily excitable $J$ (CO $J+1\longrightarrow J, J<3$ ) rotational lines, which makes it an excellent tracer of the total amount of $n(\mbox{H}_2) \sim 10^{2-3}$ $^{-3}$ ) molecular gas."
 However. the diffuse gas only serves as a reservoir olpolential fuel available for star formation. and is not actively involved in forming stars.," However, the diffuse gas only serves as a reservoir of fuel available for star formation, and is not actively involved in forming stars."
 In our own Galaxy. the sites of active star formation coincide with the dense cores of giant molecular clouds (G\ICs).," In our own Galaxy, the sites of active star formation coincide with the dense cores of giant molecular clouds (GMCs)."
 This dense gas phase (n(IIs)2I0 7?) is best traced using molecules with high critical densities such as HCN., This dense gas phase $n(\mbox{H}_2) \gs 10^{4}$ $^{-3}$ ) is best traced using molecules with high critical densities such as HCN.
" This basic picture also appears {ο hold in nearby galaxies. as suggested by IICN(1—0) survevs of local (2z 0.1) Luminous Infrared Galaxies (LIRGs Ly,~10 EL.) and Ultra Luminous Infrared Galaxies (ULIRGs LycLOL. ). which found a remarkably tight correlation between IB. luminosity and ICN line luminosity (Solomon et 11992: Gao Solomon 2004a.b)."," This basic picture also appears to hold in nearby galaxies, as suggested by $(1-0)$ surveys of local $z\ls 0.1$ ) Luminous Infrared Galaxies (LIRGs – $L_{\mbox{\tiny{IR}}}\sim 10^{11}\,\Lsolar$ ) and Ultra Luminous Infrared Galaxies (ULIRGs – $L_{\mbox{\tiny{IR}}}\sim 10^{12}\,\Lsolar$ ), which found a remarkably tight correlation between IR luminosity and HCN line luminosity (Solomon et 1992; Gao Solomon 2004a,b)."
 Recently. this relation was shown to extend all the way clown to single GMCs. thus extending over 7-8 orders of magnitude in IR. luminosity (Wu et 22005).," Recently, this relation was shown to extend all the way down to single GMCs, thus extending over 7-8 orders of magnitude in IR luminosity (Wu et 2005)."
 This was interpreted as suggesting that the (rue star formation elliciency (the star formation rate per unit mass of dense gas) is the same in these widely different svstems., This was interpreted as suggesting that the true star formation efficiency the star formation rate per unit mass of gas) is the same in these widely different systems.
behind the detonation can have significant amounts of cucdothermic reactions. violating the assumptions of the CJ structure.,"behind the detonation can have significant amounts of endothermic reactions, violating the assumptions of the CJ structure."
 Because in the case of a pathological detonation some fraction of the reactions powering the detonation are decoupled from the shock. calculating the speed of a pathological detonation requires detailed integration of the detonation structure. rather than simply using jump conditions.," Because in the case of a pathological detonation some fraction of the reactions powering the detonation are decoupled from the shock, calculating the speed of a pathological detonation requires detailed integration of the detonation structure, rather than simply using jump conditions."
 For either kind of detonation. oue can estimate where most burning occurs with slock speed D. Gvhich. even in non-CJ case. can be estimated with CJ speed) aud 7: /;=Dr; is the position behind the shock at which the induction zone chads and rapid burning takes place.," For either kind of detonation, one can estimate where most burning occurs with shock speed $D$ (which, even in non-CJ case, can be estimated with CJ speed) and $\tau_i$; $l_i = D \tau_i$ is the position behind the shock at which the induction zone ends and rapid burning takes place."
" Iu the case of a detonation iuto a very low-density. cold material. the material immediately behind the shock will still not burn significantly uutil a leugth of time equal to the ignition time passes. resulting in a ""quare wave detonation."," In the case of a detonation into a very low-density, cold material, the material immediately behind the shock will still not burn significantly until a length of time equal to the ignition time passes, resulting in a `square wave' detonation."
 For conditious relevant to near the core of a white dwarf. however. the post-shock fluid will typically have temperatures on order 75225 that is. temperatures which are already. near the aNd temperature which will be obtaiued by burning.," For conditions relevant to near the core of a white dwarf, however, the post-shock fluid will typically have temperatures on order $T_9 \approx 5$ – that is, temperatures which are already near the maximum temperature which will be obtained by burning."
 Even in these cases. this estimate of 7; provides a good measure of the thickness of the detonation structure behind the shock.as is shown in Fig & where it naeasures the position of maxima burniug.," Even in these cases, this estimate of $l_i$ provides a good measure of the thickness of the detonation structure behind the shock,as is shown in Fig \ref{fig:detstructure} where it measures the position of maximum burning."
 One mechanisin for ignition of a detonation by a hotspot is au initial rapid input of enerey which leads to a Sedov blast wave: the material shocked by the outgoing spherical blast ignites. aud as the outeoiug wave slows. a steady outeoing detonation results.," One mechanism for ignition of a detonation by a hotspot is an initial rapid input of energy which leads to a Sedov blast wave; the material shocked by the outgoing spherical blast ignites, and as the outgoing wave slows, a steady outgoing detonation results."
 A steady detonation cannot form until the outgoing shock wave speed drops to the detonation velocity detouation. or else the energy released. by reactions behind the blast wave will not be able to cateli up! to the outgoing shock to drive it.," A steady detonation cannot form until the outgoing shock wave speed drops to the detonation velocity detonation, or else the energy released by reactions behind the blast wave will not be able to `catch up' to the outgoing shock to drive it."
 Ou the other Iud. if the shock drops sienificautly below the detonation speed before ignition takes place. not enough material will be burning per nuit time to sustain a detonation wave.," On the other hand, if the shock drops significantly below the detonation speed before ignition takes place, not enough material will be burning per unit time to sustain a detonation wave."
 We denote the position of the shock when it reaches the detonation speed. froii above as Rp., We denote the position of the shock when it reaches the detonation speed from above as $R_{D}$.
 Naively. the condition for successful detonation ignition would be that Rpz Hf. ," Naively, the condition for successful detonation ignition would be that $R_{D} \gtrsim l_{i}$ "
The angle approximation Z5δρ is au estimate to b/p=2sin(Z/2). a Πιτ of a bascline so short that it does uot matter whether it is measured along a straight line (as drawn in refPoint Rho.ps)) or alone the circular perimeter.,"The angle approximation $Z\approx b/\rho$ is an estimate to $b/\rho =2\sin (Z/2)$, a limit of a baseline so short that it does not matter whether it is measured along a straight line (as drawn in \\ref{PointRho.ps}) ) or along the circular perimeter."
 The relative error in this approNimationu ds zc Z/2|. oy z10H for the example of b=100 ll. Pointingοπήπιο is a functionality of the individual telescopes: the existence of à nonzero pointing difference is absorbed in the telescope operation. aud any delay originating from there is to first order recovered in the tracking.," The relative error in this approximation is $\approx Z^2/24$ , or $\approx 10^{-11}$ for the example of $b=100$ m. Pointing/guiding is a functionality of the individual telescopes: the existence of a nonzero pointing difference is absorbed in the telescope operation, and any delay originating from there is to first order recovered in the tracking."
" We hereby explicitly discard auv “separation” term of Gubler and Tytler enmünent from their consideration of suele telescopes(ὃν, and acknowledge that both telescopes are ""1uispoinutiug"" at the same time."," We hereby explicitly discard any “separation” term of Gubler and Tytler eminent from their consideration of single telescopes, and acknowledge that both telescopes are “mispointing” at the same time."
 This cross-eved geometry iudirectly trausforis into a contibution to the delay that may boe understoodand calculated to lowest orderon the basis of what is said in refsec flatearth.. but has no parallel with sinele telescopes as long as their diameters are much sanaller than the earth radius: The example of 3.2 arcsec frou above translates into an additional anele of refraction of AR~Azg(1|tan?το)(η1)As(nyl)eZ(ng—z650 pas at 2600 ni above sea level. which is the first derivative of (1)) zy.," This cross-eyed geometry indirectly transforms into a contribution to the delay that may be understood—and calculated to lowest order—on the basis of what is said in \\ref{sec.flatearth}, but has no parallel with single telescopes as long as their diameters are much smaller than the earth radius: The example of $3.2$ arcsec from above translates into an additional angle of refraction of $\Delta R\approx \Delta z_0(1+\tan^2z_0)(n_0-1)>\Delta z_0(n_0-1)\approx
Z (n_0-1)\approx 650$ $\mu$ as at 2600 m above sea level, which is the first derivative of \ref{eq.Rofnflat}) ) $z_0$."
 Dropping the term tau?zy hore niens this is a lower estimate in the Init of stars at the zenith., Dropping the term $\tan^2z_0$ here means this is a lower estimate in the limit of stars at the zenith.
" The effect on the delay could be uuderstood in the ""staudard model of delay line correction.” where the two ravs of the star that will eventually ut the two telescopes ""eeucrate a phase difference in the vacuna (undone later ou in the delay. line uunel) as they hit the top laver of the earth atinosphiere with a path difference D rofDelModl.ps))."," The effect on the delay could be understood in the “standard model of delay line correction,” where the two rays of the star that will eventually hit the two telescopes “generate” a phase difference in the vacuum (undone later on in the delay line tunnel) as they hit the top layer of the earth atmosphere with a path difference $D$ \\ref{DelModl.ps}) )."
 The curvature correction means hat this top laver “bends back” a bit more for he telescope further away from) the star. which slightly increases the anele of incidence on the atimosplicre.," The curvature correction means that this top layer “bends back” a bit more for the telescope further away from the star, which slightly increases the angle of incidence on the atmosphere."
 refBase.ps shows that the ouly obvious definition of the baseline leneth b is on carth., \\ref{Base.ps} shows that the only obvious definition of the baseline length $b$ is on earth.
 Three alternatives have been narked withquestion marks m the figure: lu a formal way. I define an effective baseline leugth 5 above the atimosphere via," Three alternatives have been marked withquestion marks in the figure: In a formal way, I define an effective baseline length $b^*$ above the atmosphere via"
responsible for the radio enission. rather than tle amount of soft X-ray cussion. provided they account for the suppressed radio cussion du the soft state.,"responsible for the radio emission, rather than the amount of soft X-ray emission, provided they account for the suppressed radio emission in the soft state."
 Such a hypothesis ueath explains the behavior of Cre N-3. which is quite bright iu all the three energy ranges and hence the observational uncertaiutv is quite low.," Such a hypothesis neatly explains the behavior of Cyg X-3, which is quite bright in all the three energy ranges and hence the observational uncertainty is quite low."
 It cau also be noticed from Table 1 that the most significant correlation for (νο XN-3 is between the radio enüssiou and the ratio of hard. N-rav. flux to soft N-ravs., It can also be noticed from Table 1 that the most significant correlation for Cyg X-3 is between the radio emission and the ratio of hard X-ray flux to soft X-rays.
 Though such au explanation is not very clear in other sources. all the available observations are consisteut with this.," Though such an explanation is not very clear in other sources, all the available observations are consistent with this."
 Since the radio ciission too shows mereasing Clissio—-c from ‘off state to low-hard state (iu GN 339-1 aud VLO Cre). correlated. behavior m low-hard state (the above two sources and Cre X-1). aud high radio emission in au intermediate state very close to the high state (n CRS 19151105 and (νο XN-3) and suppressed radio enissio- in the high state (Cvg N-3. CRS 19151105. (νο N-1& CX 339-1). we speculate that the soft N-rav intensity deteriuues the spectral shape aud the accretion disc condition iu these sources. which iu turn determines the αλλο of radio enussion (in the X-ray quiescent hd state of these sources).," Since the radio emission too shows increasing emission from `off' state to low-hard state (in GX 339-4 and V404 Cyg), correlated behavior in low-hard state (the above two sources and Cyg X-1), and high radio emission in an intermediate state very close to the high state (in GRS 1915+105 and Cyg X-3) and suppressed radio emission in the high state (Cyg X-3, GRS 1915+105, Cyg X-1 GX 339-4), we speculate that the soft X-ray intensity determines the spectral shape and the accretion disc condition in these sources, which in turn determines the amount of radio emission (in the X-ray quiescent hard state of these sources)."
 The trausition of the svsteuis (Cre νο GRS 19151105) iuto flaring state απ their correspouding behavior in the radio as well as X-ray bands Is an issue not discussed here., The transition of the systems (Cyg X-3 GRS 1915+105) into flaring state and their corresponding behavior in the radio as well as X-ray bands is an issue not discussed here.
 Though there are models describing the accretion disk enüssion (Zdziarski=2000)— or jet cimission (ALarkotfetal. 2003). there are very few inodels which self solve the accretion aud. ejection phenomena seen in black hole sources.," Though there are models describing the accretion disk emission \citep{zdz00} or jet emission \citep{mar03}, there are very few models which self-consistently solve the accretion and ejection phenomena seen in black hole sources."
 Since our ficdines sugeest a close connection between these two phenomena. we attempt below to qualitatively explain the N-rayv radio association using one such model: the Two Conmponeut Advective Flow (TCAF) model of Chakrabarti(1996).," Since our findings suggest a close connection between these two phenomena, we attempt below to qualitatively explain the X-ray radio association using one such model: the Two Component Advective Flow (TCAF) model of \citet{cha96}."
. According to this model. the Compton scattered N-ravs iu a black hole source originates froun a region close to the compact object. coufined within the Ceutrifugal Boundary Laver (CENBOL).," According to this model, the Compton scattered X-rays in a black hole source originates from a region close to the compact object, confined within the Centrifugal Boundary Layer (CENBOL)."
 The N-vay spectral shape iu various ‘states’ of the source essentially depends on the location of the CENBOL and a detailed. description can be found in Chakrabarti&Titarchul1905} anc Ebisawaetal. (1996)., The X-ray spectral shape in various `states' of the source essentially depends on the location of the CENBOL and a detailed description can be found in \citet{cha95} and \citet{ebi96}.
. At low accretion rates. the CENBOL is far away from the compact object aud the N-vay spectrum is dominated by a thermalCompton spectra. originating roni the hieh temperature region within the CENBOL.," At low accretion rates, the CENBOL is far away from the compact object and the X-ray spectrum is dominated by a thermal-Compton spectrum, originating from the high temperature region within the CENBOL."
 Iu the transition state. the CENBOL comes closer to he compact object aud the CENBOL can sometimes eive rise to radial shocks. causing iuteuse quasi-periodic oscillations. as seen in GRS 1915|105.," In the transition state, the CENBOL comes closer to the compact object and the CENBOL can sometimes give rise to radial shocks, causing intense quasi-periodic oscillations, as seen in GRS 1915+105."
 Tn the high state. he increased accretion rate produces copious photons i- he accretion dise which cool the Compton region. giving rise to very intense disk blackbody cussion along with ilk inotiou Comptonization (a power-law in hard N-aravs with a photon index of —2.5).," In the high state, the increased accretion rate produces copious photons in the accretion disc which cool the Compton region, giving rise to very intense disk blackbody emission along with bulk motion Comptonization (a power-law in hard X-rays with a photon index of $\sim$ 2.5)."
 At some critical accretio- rates. the state transitions could be oscillatory as seen in GRS 19151105 (Chakrabarti&Mauickiun|20001," At some critical accretion rates, the state transitions could be oscillatory as seen in GRS 1915+105 \citep{cha00}."
", The behavior of TCAF disks aud the outflow has been stucied in detail in Das&Chakrabarti(1998).", The behavior of TCAF disks and the outflow has been studied in detail in \citet{das98}.
. The outflow rate is found to be a monotonic function of the conrpressiou ratio. BR. of the eas at the shock region.," The outflow rate is found to be a monotonic function of the compression ratio, R, of the gas at the shock region."
 Iu this scenario. at low accretion rates. the CENBOL is far away from the compact object. and a weal shock can forma with low compression ratio. eiviug low aud steady outflow.," In this scenario, at low accretion rates, the CENBOL is far away from the compact object, and a weak shock can form with low compression ratio, giving low and steady outflow."
 If this outflow gives rise to radio emission. one can expect a relation between the radio enission aud the N-ray cussion.," If this outflow gives rise to radio emission, one can expect a relation between the radio emission and the X-ray emission."
 In this state (off state to low-hard state). an ducreased accretion rate increases the overall amount of energv available to the Couptouizine region and heuce inercasing the XN-rav cuissionu.," In this state (off state to low-hard state), an increased accretion rate increases the overall amount of energy available to the Comptonizing region and hence increasing the X-ray emission."
 The CENBOL location would be pushed imward. increasing the compression ratio (aud hence inereasiug the radio Cluission) and also can increase the temperature aud optical depth of the Comptouizing region. thus eiving rise to a pivoting behavior at hard N-vavs (50. 90 keV) as seen in Cve X-1 and GN 339-1.," The CENBOL location would be pushed inward, increasing the compression ratio (and hence increasing the radio emission) and also can increase the temperature and optical depth of the Comptonizing region, thus giving rise to a pivoting behavior at hard X-rays (50 – 90 keV) as seen in Cyg X-1 and GX 339-4."
 At increased accretion rate. the CENBOL can come closer to the compact region. giving the spectral aud radio properties as seen im GRS 1915|105 and Cye N-3.," At increased accretion rate, the CENBOL can come closer to the compact region, giving the spectral and radio properties as seen in GRS 1915+105 and Cyg X-3."
 For a eiven accretion rate the compression ratio. after reaching a critical value (with the shock region comune correspondinely closer to the eveut horizou). causes the source to transit iuto the high-soft state state. for which the radio enuüssiou is progressively suppressed (Chakrabarti1999).," For a given accretion rate the compression ratio, after reaching a critical value (with the shock region coming correspondingly closer to the event horizon), causes the source to transit into the high-soft state state, for which the radio emission is progressively suppressed \citep{cha99}."
. This model qualitatively explains all the observed. X-ray spectral aud radio properties of Galactic black hole sources preseuted here., This model qualitatively explains all the observed X-ray spectral and radio properties of Galactic black hole sources presented here.
 A complete understaundiug of the accretion-ejection physics in Galactic icroquasars demands proper interpretation and modeling of all the varied states of X-rav and radio enuüssion. euconipassing the various flaring and steady enüssious covering all ranges of time scales.," A complete understanding of the accretion-ejection physics in Galactic microquasars demands proper interpretation and modeling of all the varied states of X-ray and radio emission, encompassing the various flaring and steady emissions covering all ranges of time scales."
 Iu this paper we have taken a first step in this direction by atteniptiug to understand the loug term variation of the (non flaring) radio enüssion associated with the X-rav chuission in the steady hard aud soft states;, In this paper we have taken a first step in this direction by attempting to understand the long term variation of the (non flaring) radio emission associated with the X-ray emission in the steady hard and soft states.
 We have analyzed the (quasi) smmltaueous observations on CRS 19151105 and (νο X-1 using the RVTE-ASAL CCRO-BATSE and GBI data and made a detailed study of correlation between radio aud X-ray fluxes.," We have analyzed the (quasi) simultaneous observations on GRS 1915+105 and Cyg X-1 using the -ASM, }-BATSE and GBI data and made a detailed study of correlation between radio and X-ray fluxes."
 Dased on this analvsis along with discussion of earlier published results ou Galactic niücroquasars we find that:, Based on this analysis along with discussion of earlier published results on Galactic microquasars we find that:
AMpgy.,$M_{BH}$.
 Black hole virialm 1iasses may be estimatedB as AfpyxRe.Di where Rois B the broad lineB regionB size.B and eds the velocity dispersion of the eas emiting the broad. emission lines.," Black hole virial masses may be estimated as $M_{BH} \propto R v^2$ , where $R$ is the broad line region size, and $v$ is the velocity dispersion of the gas emitting the broad emission lines."
 A correlation has been found between the luminosity of à source aud tie size of it’s broad line region 2005).," A correlation has been found between the luminosity of a source and the size of it's broad line region \citep[the $R$--$L$ relationship, e.g.,][]{kaspi05}."
. One can then exploit this τεJatiolsup. and use the broad line FWTAL as an estimate for 0. obtaining virial mass estimates MpXLee? (¢NOOO.Wandeletal.1999).. where the expoucut is 0zz0.5 (c.g...Vestereaard&Peterson2006).," One can then exploit this relationship, and use the broad line $FWHM$ as an estimate for $v$, obtaining virial mass estimates $\hat{M}_{BH}
 \propto L^{\theta} v^2$ \citep[e.g.,][]{wand99}, where the exponent is $\theta \approx 0.5$ \citep[e.g.,][]{vest06}."
". Uuortulaelv. the uncertainty ou the broad line estimates of Apry can be considerable. having a standard deviati1i of o,~0.1 dex (e.g...Mebure&Jarvis2002:&Peterson2006:Kellyetal."," Unfortunately, the uncertainty on the broad line estimates of $M_{BH}$ can be considerable, having a standard deviation of $\sigma_m \sim 0.4$ dex \citep[e.g.,][]{bhmmgii,vest06,kelly06a}."
 2007).. For οase 0: coniparison with previous work. 1 cstimate Afpy using oulv the IL cussion liue.," For ease of comparison with previous work, I estimate $M_{BH}$ using only the $\beta$ emission line."
 The logarithua of t1ο viria lnass estimates were calculated using tje ID) huuiuositv axd FT7M according to the relatioushi) even by Vestereaard&Peterson (2006).., The logarithm of the virial mass estimates were calculated using the $\beta$ luminosity and $FWHM$ according to the relationship given by \citet{vest06}. .
 Ay snple consists ofa subset of the sample o‘Ivelly&Bechtold(2007)., My sample consists of a subset of the sample of\citet{kelly06b}.
. These sources lave measurements Q “the N-rav photon index. Py=ay|Ll. obtained from observations. aud measurements of the QticalJUV BIuuinositv at 2500-4. denoted as Losyy. obtained from SDSS spectra.," These sources have measurements of the X-ray photon index, $\Gamma_X = \alpha_X + 1$, obtained from observations, and measurements of the optical/UV luminosity at $2500\AA$, denoted as $L_{2500}$, obtained from SDSS spectra."
 The IL} profile was modeled as a stn of (απουσίας and extracted from the SDSS xs]vectra according to the procedure cescribed in (2007).., The $\beta$ profile was modeled as a sum of Gaussians and extracted from the SDSS spectra according to the procedure described in \citet{kelly06a}.
 I estimated the I+ £TPAL aud luminosity from the line profile fits., I estimated the $\beta$ $FWHM$ and luminosity from the line profile fits.
" Tostimate the bolometric Inuinositv. L5. from the Iuininositv at 2500.4. assuming a constant bolometric correction £5,=5.6Lo5yy (Elviseta.199D.."," I estimate the bolometric luminosity, $L_{bol}$, from the luminosity at $2500\AA$, assuming a constant bolometric correction $L_{bol} = 5.6 L_{2500}$ \citep{elvis94}."
". The standard deviatici iu this bolometric correction reported by Elvisetal.(1991). is 3.l. nuplviugo an uncertainty in logL5, of σι0.25 dex."," The standard deviation in this bolometric correction reported by \citet{elvis94} is 3.1, implying an uncertainty in $\log L_{bol}$ of $\sigma_{bol} \sim
 0.25$ dex."
" Combining this with the ~0.1 dex uucertaimtv on loBMpy. tie total ueasureiment error o1i logLawπιω bocomes o,~0.17 dex."," Combining this with the $\sim 0.4$ dex uncertainty on $\log M_{BH}$, the total `measurement error' on $\log L_{bol} /
 L_{Edd}$ becomes $\sigma_x \sim 0.47$ dex."
 The distribution of Dy as a fiction o flosLi/Lg; is shown in Fiewe 10.., The distribution of $\Gamma_X$ as a function of $\log L_{bol} / L_{Edd}$ is shown in Figure \ref{f-gamx_eddrat}.
" As can be seen. the iicasureimeut errors on both Dy aud logZí4/Lr are large and make a considerable contribution to t1e observed scatter in both variables. where A,~ ü.laid Ry~OS."," As can be seen, the measurement errors on both $\Gamma_X$ and $\log L_{bol} / L_{Edd}$ are large and make a considerable contribution to the observed scatter in both variables, where $R_y \sim 0.1$ and $R_x \sim 0.8$."
 Therefore. we expect the moeasurenieit errors to have a sjeuificaut effect on the correlation aid regression analysis.," Therefore, we expect the measurement errors to have a significant effect on the correlation and regression analysis."
 I performed the regression assuwing the linear form Dya|Sloeο aud modellehue the intrinsic distribution of logLayLi; using A— 2€Guissans.," I performed the regression assuming the linear form $\Gamma_X =
 \alpha + \beta \log L_{bol} / L_{Edd}$, and modelleling the intrinsic distribution of $\log L_{bol} / L_{edd}$ using $K = 2$ Gaussians."
 Draws from the posterior were obtained using the Cabbs sampler., Draws from the posterior were obtained using the Gibbs sampler.
 The mareinal posterior distzilnitions for 2.0. and the correlation between Pay aud logLawfLead. p. are shown in Figure 11.. and the posteric imediau and 95% (26) poiutwise intervals on the regression line are show rin Fieure 10.," The marginal posterior distributions for $\beta, \sigma$ , and the correlation between $\Gamma_X$ and $\log L_{bol} / L_{edd}$, $\rho$, are shown in Figure \ref{f-posthb}, and the posterior median and $95\%$ $2\sigma$ ) pointwise intervals on the regression line are shown in Figure \ref{f-gamx_eddrat}."
 The posterior mclan estimate of the paracters are a=3.1240.11 for the coustaut. Jj=1.35+£0.51 for the slope. ¢=0.2640.11 for the intrinsic scatter about the regressiou line. µεOrrEOL| for the mean of logLy.7/Leag. aud σε0.3240.12 dex for the dispersion iu logLi/L4.," The posterior median estimate of the parameters are $\hat{\alpha} = 3.12 \pm 0.41$ for the constant, $\hat{\beta} = 1.35 \pm 0.54$ for the slope, $\hat{\sigma} = 0.26 \pm 0.11$ for the intrinsic scatter about the regression line, $\hat{\mu}_{\xi} = -0.77 \pm 0.10$ for the mean of $\log L_{bol} / L_{Edd}$, and $\hat{\sigma}_{\xi} = 0.32 \pm 0.12$ dex for the dispersion in $\log L_{bol} / L_{edd}$."
 Hore. I have used a robust estimate of the posterior standard deviation as au error bar on the parameters.," Here, I have used a robust estimate of the posterior standard deviation as an `error bar' on the parameters."
" These resuts naply that the observed scatter in logοLg; is dominated by measurement error. GfTo~1.5. as oxected. from the large value of Π,."," These results imply that the observed scatter in $\log L_{bol} / L_{Edd}$ is dominated by measurement error, $\sigma_x / \tau \sim 1.5$, as expected from the large value of $R_x$."
 For comparison. the DCES(Y[VX) estimate of the sope is Opers=329£3.31. the FITENY estimate is Jexyy=1.τος.LO. are the OLS estimate is Jops=0.5640.11: the standard error on Jpgvyy was estimated using bootstrapping.," For comparison, the $Y|X$ ) estimate of the slope is $\hat{\beta}_{BCES} = 3.29 \pm 3.34$, the FITEXY estimate is $\hat{\beta}_{EXY} = 1.76 \pm 0.49$, and the OLS estimate is $\hat{\beta}_{OLS} = 0.56 \pm 0.14$; the standard error on $\hat{\beta}_{EXY}$ was estimated using bootstrapping."
 Fieure 10 also compares the OLS. BCES. aud FITENY best-fit lines with the posterior median estimate.," Figure \ref{f-gamx_eddrat} also compares the OLS, BCES, and FITEXY best-fit lines with the posterior median estimate."
" The 95% coufideuce region ou the slo2ο implied by the posterior draws is 0.16<9<< 11. whereas the approximate 9554 confidence region implied by the BCES. FITENY. aud OLS standard errors are o2.26«P<OSL 050<>«2,12. aud 0.12«30,10. --espectivelv."," The $95\%$ confidence region on the slope implied by the posterior draws is $0.46 < \beta
 < 3.44$ , whereas the approximate $95\%$ confidence region implied by the BCES, FITEXY, and OLS standard errors are $-3.26 < \beta <
 9.84$, $0.80 < \beta < 2.72$, and $0.42 < \beta < 0.70$, respectively."
 The OLS aud FITENY cstimates aud the Davesiaiapproach eive “statistically significait” evience for a correlation between logLi;/Lead and Py: however the BCES estimate is too variable to rule out the null livpothesis of uo correlation., The OLS and FITEXY estimates and the Bayesianapproach give `statistically significant' evidence for a correlation between $\log L_{bol} / L_{Edd}$ and $\Gamma_X$; however the BCES estimate is too variable to rule out the null hypothesis of no correlation.
" As noted before. the large measurement errors on logL5,[πιω bias the OLS estimate of. towardshalloswer values aud the FITENY estiuate of JJ toward steeper values."," As noted before, the large measurement errors on $\log
 L_{bol} / L_{Edd}$ bias the OLS estimate of $\beta$ towardshallower values and the FITEXY estimate of $\beta$ toward steeper values."
 Because of this bias. coufideuce regious based O1 ons and Jpyyare not valid because they are not centered ou the true value of 2. aud thus do not," Because of this bias, confidence regions based on $\hat{\beta}_{OLS}$ and $\hat{\beta}_{EXY}$are not valid because they are not centered on the true value of $\beta$ , and thus do not"
2230.. a niuis. pulsar. was discovered in a radio survev of wnicentified EGRET eamuna rax sources using the Parkes Radio Telescope (Ilesselsal. 2005).,", a ms pulsar, was discovered in a radio survey of unidentified EGRET gamma ray sources using the Parkes Radio Telescope \citep{discovery}."
. Subsequeutlv. N-vay cussion from XAJM (Robertsetal.2007). aud 5-av enmuüssion frouTelescope (Abdo2010) was detected.," Subsequently, X-ray emission from \citep{xraypwn}
 and $\gamma$ -ray emission from \citep{fermi} was detected."
 Like most nüllisecond pulsars (MISPs}. is mn a binary svsteni.," Like most millisecond pulsars (MSPs), is in a binary system."
 The circular orbit is consistent with the pulsar having uudergoue mass trausfer aud spun up., The circular orbit is consistent with the pulsar having undergone mass transfer and spun up.
 The mass function derived from pulsar timine indicated a companion with mass Mao>O.LAL. (Iessels 2005)..," The mass function derived from pulsar timing indicated a companion with mass $M_2>0.4\,M_\odot$ \citep{discovery}."
 The svsteni recently came into prominence when Demorestetal.(2010). reported the mass of the pulsar to be LOFAOOLAL....," The system recently came into prominence when \citet{heavy} reported the mass of the pulsar to be $1.97\pm0.04\,$."
 The detection of such a massive ueutron star (NS) places very strong constraints ou the equation of state of matter at extreme uuclear deusities (sec.forexample.Lattimer&Prakash2001.2005).," The detection of such a massive neutron star (NS) places very strong constraints on the equation of state of matter at extreme nuclear densities \citep[see, for example,][]{lattimer,lattimertable}."
.. The rather exquisite precision of this mass measurement was xossible due to the orbit beiug almost perpeudicular to he plane ofthe sky., The rather exquisite precision of this mass measurement was possible due to the orbit being almost perpendicular to the plane of the sky.
 As a result. the Shapiro delay caused w the companion is very larvec. resulting iu a precise estimate of the mass of the companion. A»=0.5004006 ," As a result, the Shapiro delay caused by the companion is very large, resulting in a precise estimate of the mass of the companion, $M_2=
0.500\pm 0.006\,M_\odot$ ."
The s.7dday orbital period is siguificauth shorter AZ...than ~120 davs expected for a low-ass Nav ünary with such a massive secondary (Rappaportotal. sugecsting a peculiar evolutionary historvw for his binary.," The day orbital period is significantly shorter than $\sim120\,$ days expected for a low-mass X-ray binary with such a massive secondary \citep{massperiod} — suggesting a peculiar evolutionary history for this binary."
 Given the importance of the result of Demorestctal.(2010). additional verification or consistency checks of physicalparameters of οσα be expected to be of some value., Given the importance of the result of \cite{heavy} additional verification or consistency checks of physicalparameters of can be expected to be of some value.
 A wwhite dwarf (WD) at the interred distance of (GT~1.2 kkpe). even if a few Cyr old. is within the reach of preseut-day optical telescopes.," A white dwarf (WD) at the inferred distance of $d\sim 1.2$ kpc), even if a few Gyr old, is within the reach of present-day optical telescopes."
 It is this search for the WD that coustitutes the principal focus of this Letter., It is this search for the WD that constitutes the principal focus of this Letter.
 We observed 1)) ing and R bands using the imaging mode of +16 Low Resolutionlnaging Spectrograph (ERIS) ou the yumm I&eck-I telescope (Okeetal.1995).. with uperaded aud blue cameras (MeCarthvetal.1998:Steidelal 2001).," We observed ) in $g$ and $R$ bands using the imaging mode of the Low Resolution Imaging Spectrograph (LRIS) on the m Keck-I telescope \citep{lris}, with upgraded and blue cameras \citep{lrisb1,lrisb2}."
 Several nuages were acquired at cach target location. dithering the telescope by small amouuts between cach exposure.," Several images were acquired at each target location, dithering the telescope by small amounts between each exposure."
" The observing conditions ou UT 2010x May 15 were poor (secius 1"".1). so onlv data acquired. on UT 2010 July 5S (Ro band xccing 07.85 FWIIM) were used in this analysis."," The observing conditions on UT 2010 May 15 were poor (seeing .4), so only data acquired on UT 2010 July 8 $R$ band seeing .85 FWHM) were used in this analysis."
 The total exposure ou this night was 9608s in the R baud and LOLOss in the g baud., The total exposure on this night was s in the $R$ band and s in the $g$ band.
 The plate scale is | for both Calieras The images were processed usingIRAF., The plate scale is $^{-1}$ for both cameras The images were processed using.
. After bias correction and flat fielding. cosmic ravs wererejected using (vanDokkun 2001)..," After bias correction and flat fielding, cosmic rays wererejected using \citep{lacosmic}. ."
 The tages were then aligned with aud averaged to produce the final image for each baud (see 1))., The images were then aligned with and averaged to produce the final image for each band (see ).
 A World Coordinate Svstem was calculated using, A World Coordinate System was calculated using
blackbody contributions (hirley2000).. while the AXPs had softer spectra. with P~4 and ~0.5 keV blackbocdies contributing up to ol the X-ray flix (Mereghetti2000).,"blackbody contributions \citep{h99}, while the AXPs had softer spectra, with $\Gamma \sim 4$ and $\sim 0.5$ keV blackbodies contributing up to of the X-ray flux \citep{m99}."
. But this situation has been changing., But this situation has been changing.
 Observations ol the quiescent [found a photon index of 3.2.) closer to the nominal ANP index than to the that of the other SGRs and possible evidence for a 0.5 keV blackbody (hulkarnietal.2001a)., Observations of the quiescent found a photon index of 3.2 – closer to the nominal AXP index than to the that of the other SGRs – and possible evidence for a 0.5 keV blackbody \citep{k+00}.
. This may put the value of the photon index on a continuum related (o burst activity. ancl magnetic Ποια geometry for both groups., This may put the value of the photon index on a continuum related to burst activity and magnetic field geometry for both groups.
 Observations of iin quiescent and active states demonstrated (he presence of an underlying ~0.5 keV blackbody (Woodsοἱal.1999a.2001).," Observations of in quiescent and active states demonstrated the presence of an underlying $\sim$ 0.5 keV blackbody \citep{wkp+99a,wkg+01}."
. Updated spectral fits of archival ddata of SGRs and AXPs have shown that both groups seem to possess blackbody components whose fraction of the overall X-ray emission may constitute another unilving continuum (Pernaοἱal.2001)., Updated spectral fits of archival data of SGRs and AXPs have shown that both groups seem to possess blackbody components whose fraction of the overall X-ray emission may constitute another unifying continuum \citep{phh+01}.
. Spectral fits to the ANP LE 1043.1—5937 show that it has a hard power-law component reminiscent of the SGRs (lNaspietal.2001).., Spectral fits to the AXP 1E $-$ 5937 show that it has a hard power-law component reminiscent of the SGRs \citep{kgc+01a}.
. And finally. optical and infrared observations of SGRs (kaplanetal.2001a:πα900110) and ANPs (IIullemanetal.2000:ILulleman2000) have shown that the groups have similar rav-to-optical flux ratios. so that this ratio may be a distinguishing characteristic of the two. as à group (Ilullemanetal.2000).," And finally, optical and infrared observations of SGRs \citep{k+00c,k+01} and AXPs \citep{hvkk00,hvkvk00} have shown that the groups have similar X-ray-to-optical flux ratios, so that this ratio may be a distinguishing characteristic of the two, as a group \citep{hvkk00}."
.. All of these findings have strengthened. arguments for association between the AXNPs and SGRs., All of these findings have strengthened arguments for association between the AXPs and SGRs.
 The blackbody component of the sspeclatun. wilh Api80.5 keV and 2j215d; kim. has parameters (hat are similar (o those of other isolated NS candidates (c.f.Verbuntetal.1994).," The blackbody component of the spectrum, with $kT_{\rm
BB}\approx 0.5$ keV and $R_{\rm BB}\approx 1.5 d_5$ km, has parameters that are similar to those of other isolated NS candidates \citep[c.f.][]{vbj+94}."
. The relatively small emitting radius (hat we find. significantly smaller Caan (he nominal 210 km-radius NS. is twpically interpreted as either due to restricted emission Irom. e.g.. the NS polar caps. or as the result of temperature-dependent opacity effects in (he NS atmosphere (Rutledge 2001).," The relatively small emitting radius that we find, significantly smaller than the nominal $\approx$ 10 km-radius NS, is typically interpreted as either due to restricted emission from, e.g., the NS polar caps, or as the result of temperature-dependent opacity effects in the NS atmosphere \citep{rbb+99,phh+01}."
". The latter scenario would allow for closer distances. lower surface temperatures. and, potentially. emission [rom (he entire NS surface (Pernaetal.2001)."," The latter scenario would allow for closer distances, lower surface temperatures, and, potentially, emission from the entire NS surface \citep{phh+01}."
. The absence of anv narrow spectral features. to equivalent widths of less Chan 150 eV. is somewhat surprising given the detection bx Strohmaver&Ibrahim(2000) of à strong. 400-eV equivalent width. 226.4-keV. emission line in the sspectrum of a 1998 August 29 burst of1900--L4.," The absence of any narrow spectral features, to equivalent widths of less than 150 eV, is somewhat surprising given the detection by \citet{si00} of a strong, 400-eV equivalent width, $\approx$ 6.4-keV emission line in the spectrum of a 1998 August 29 burst of."
. Strohmaver&Ibrahim(2000) discuss two possible interpretations for the feature they observe: first. (hat it may result from fluorescence of relatively cool iron in the near vicinity of the NS: and second. that it may result [rom proton or alpha particle (He!) evclotron transitions in the SGR magnetosphere: {hese ions would have been liberated by the closely-prececling giant flare of 1998 August 27.," \citet{si00} discuss two possible interpretations for the feature they observe: first, that it may result from fluorescence of relatively cool iron in the near vicinity of the NS; and second, that it may result from proton or alpha particle $^4$ ) cyclotron transitions in the SGR magnetosphere; these ions would have been liberated by the closely-preceding giant flare of 1998 August 27."
 llowever. if the line resulted. [rom iron (norescence (hen we would expect. with," However, if the line resulted from iron fluorescence then we would expect, with"
out that most of the curve described by L passes through reeions of few electrous.,out that most of the curve described by \ref{cosphi_s} passes through regions of few electrons.
 The result is that the negative absorption coefficient are very sinall in absolute value. or that there are no negative absorption coefficients at all for these PEUT ratios.," The result is that the negative absorption coefficient are very small in absolute value, or that there are no negative absorption coefficients at all for these $\nu_p/\nu_B$ ratios."
 We performed nuucerical caleulatious of the absorption cocfiicicnt 5 using the power law distribution 7 with index à=3. and the loss cone distribution & with ουσία).=(LAL.costada)O83.," We performed numerical calculations of the absorption coefficient \ref{gyroe} using the power law distribution \ref{powerlaw}
 with index $\delta=3$, and the loss cone distribution \ref{losscone} with $\cos(\alpha)=0.81,
\cos(\alpha-\delta\alpha)=0.83$."
" The calculation were performed with a standard magnetic Ποια B=360 gauss, and an ambient munber density which was chauged to give differeut ratios of 1/rp."," The calculation were performed with a standard magnetic field $B=360\ gauss$ , and an ambient number density which was changed to give different ratios of $\nu_p/\nu_B$."
 For every cosine of enission angle cos(0) to the magnetic field the largest in absolute magnitude negative absorption coefficient was und by scamming in frequeucy from sp to (s|1)rj/p., For every cosine of emission angle $\cos(\theta)$ to the magnetic field the largest in absolute magnitude negative absorption coefficient was found by scanning in frequency from $s\nu_B$ to $(s+1)\nu_B$.
" The requency of this largest in absolute magnitude negative absorption cocficient is compared with the approximation 1l aud with our new approximation 19 in the figures 6 ο δν,"," The frequency of this largest in absolute magnitude negative absorption coefficient is compared with the approximation \ref{DM} and with our new approximation \ref{multiD} in the figures \ref{100-81-1}
 to \ref{140-81-2}."
 Iu figure 6 the comparison is mace for frequencies jetween vp sand 2»pp. and he feure shows that for sanall cos(0) the Dulk-Moelrose approximation 11 Is vorv siuilar to our new approximation.," In figure \ref{100-81-1} the comparison is made for frequencies between $\nu_B$ and $2\nu_B$, and the figure shows that for small $\cos(\theta)$ the Dulk-Melrose approximation \ref{DM} is very similar to our new approximation."
 However. for lareer cos(0). the now approximation is mach better.," However, for larger $\cos(\theta)$, the new approximation is much better."
 The new approximation is identical with the result of the full Munerical computation for most of the cos(0) range where here is negative absorption., The new approximation is identical with the result of the full numerical computation for most of the $\cos(\theta)$ range where there is negative absorption.
 In figure 7 the comparison is made for frequencies vetween 2vp and 237p., In figure \ref{100-81-2} the comparison is made for frequencies between $2\nu_B$ and $3\nu_B$.
 Here the relevant cos(0) rauge is huger. aud again for small cosines the Dulk-Moelrose approximation ll ids similar to the ummerical results aud ο our new approximation.," Here the relevant $\cos(\theta)$ range is larger, and again for small cosines the Dulk-Melrose approximation \ref{DM} is similar to the numerical results and to our new approximation."
 However. for angles smaller han about 60 degrees. the Dulk-Melrose approximation vceins to diverge. while the new approximation reais virtually identical with the numerical results.," However, for angles smaller than about $60$ degrees, the Dulk-Melrose approximation begins to diverge, while the new approximation remains virtually identical with the numerical results."
" For the ratio v,=L.I1rp our conclusions in section 5 ead us to expect that negative absorption exists ouly for requencies 7227g. aud the uuuercal computations bear lis out."," For the ratio $\nu_p=1.4\nu_B$ our conclusions in section \ref{general properties} lead us to expect that negative absorption exists only for frequencies $\nu > 2\nu_B$, and the numerical computations bear this out."
 Figure 5 shows again that for large cos(0) the approximation of equation JL begins to diverge from the miuerical results. while our new approximation reais ideutical to it.," Figure \ref{140-81-2} shows again that for large $\cos(\theta)$ the approximation of equation \ref{DM} begins to diverge from the numerical results, while our new approximation remains identical to it."
 We do not show results for the Z-1mode. since we are interested iu presenting estimates for the frequencies of occurrence of observable cuiission.," We do not show results for the Z-mode, since we are interested in presenting estimates for the frequencies of occurrence of observable emission."
 The Zauode is nuportaut in quenching the maser. but it is necessary fo compute the absorption coefficieuts for all the 11odes aud conrpare them to determine whether it docs.," The Z-mode is important in quenching the maser, but it is necessary to compute the absorption coefficients for all the modes and compare them to determine whether it does."
 We develop a new approximation for the frequencies at which the absorption cocficient is negative. aud therefore the Electron Cyclotron Maser mechanisin operates.," We develop a new approximation for the frequencies at which the absorption coefficient is negative, and therefore the Electron Cyclotron Maser mechanism operates."
 This new approxiuation is eiven by equation 13.. and is casy to compute.," This new approximation is given by equation \ref{multiD}, and is easy to compute."
 Our approximation gives results which are much more accurate than previous approximations. id are practicallythe same as the results of a full munerical calculation.," Our approximation gives results which are much more accurate than previous approximations, and are practicallythe same as the results of a full numerical calculation."
 The frequencies derived with the approxination are within 0.010.02vp of the numerically calculated frequencies., The frequencies derived with the approximation are within $0.01-0.02\ \nu_B$ of the numerically calculated frequencies.
" The paramcters cutering the approximation are the anele of emission to the magnetic field 0. the loss-cone opening angle o. the ratio r,/75. aud the ratio 1/1]."," The parameters entering the approximation are the angle of emission to the magnetic field $\theta$, the loss-cone opening angle $\alpha$, the ratio $\nu_p/\nu_B$, and the ratio $\nu/\nu_B$."
" The new approximation can be used to define a range ofpossible frequencies of millisecond spike cussion. given the ratio r,/rg."," The new approximation can be used to define a range of possible frequencies of millisecond spike emission, given the ratio $\nu_p/\nu_B$ ."
 Or. when spike cuuission is detected. the i»proxination can be used to eive the range of plivsical paraiaeters iu the emission region.," Or, when spike emission is detected, the approximation can be used to give the range of physical parameters in the emission region."
"and, accordingly, the probability distribution function of yf, is The precise functional form (11)) is not crucial, only that cannot be taken to be uniform.","and, accordingly, the probability distribution function of $\psi_0'$ is The precise functional form \ref{eq:psi1-distrib}) ) is not crucial, only that $p(\psi_0')$ cannot be taken to be uniform."
" After the sudden baryonic blowout,p(w) collisionless particles enter their new orbit in a special phase — preferentially near pericentre — so that they subsequently migrate outwards in unison."," After the sudden baryonic blowout, collisionless particles enter their new orbit in a special phase – preferentially near pericentre – so that they subsequently migrate outwards in unison."
 It is this difference in knowledge of phases before and after sudden changes that allows irreversibility in the real universe to appear in the model., It is this difference in knowledge of phases before and after sudden changes that allows irreversibility in the real universe to appear in the model.
 Only if all collisionless particles were near their pericentre just before the baryons returned would the statistical properties of the reversed picture match those of the actual model., Only if all collisionless particles were near their pericentre just before the baryons returned would the statistical properties of the reversed picture match those of the actual model.
" While this is dynamically possible, it is statistically unlikely."," While this is dynamically possible, it is statistically unlikely."
" Finally note that if the changes in potential are introduced gradually, the process should become adiabatic and hence reversible."," Finally note that if the changes in potential are introduced gradually, the process should become adiabatic and hence reversible."
" The dashed line in Figure 3 shows a numerical solution for which € changes smoothly over several orbital times from (9 to €, then back to @p."," The dashed line in Figure \ref{fig:harmonic-oscillator}
 shows a numerical solution for which $\omega$ changes smoothly over several orbital times from $\omega_0$ to $\omega_1$, then back to $\omega_0$."
" As expected from equation (4)), the final orbital amplitude is the same as its initial value, confirming the qualitatively different results to be expected from gradual variation as opposed to sudden jumps."," As expected from equation \ref{eq:adiabatic-Efinal}) ), the final orbital amplitude is the same as its initial value, confirming the qualitatively different results to be expected from gradual variation as opposed to sudden jumps."
 To test the picture expounded above we start by generating a time-dependent effective toy potential from the simulations (Section ??))., To test the picture expounded above we start by generating a time-dependent effective toy potential from the simulations (Section \ref{sec:first-sims}) ).
" This is given by equation (1)), with V(r;t) calculated from the spherically averaged density profile."," This is given by equation \ref{eq:veff}) ), with $V(r;t)$ calculated from the spherically averaged density profile."
 The starting energy Ep and the value of j can be determined by specifying initial orbital parameters., The starting energy $E_0$ and the value of $j$ can be determined by specifying initial orbital parameters.
 The angular momentum is necessarily conserved because of the spherical symmetry of the modelling (restriction 1 of Section ??))., The angular momentum is necessarily conserved because of the spherical symmetry of the modelling (restriction 1 of Section \ref{sec:virial-eqs}) ).
" In the simulations the changes in potential are not exactly symmetric (e.g. lower panel of Figure 2)); however we will see below that, for the purposes of calculating real-space density profiles, the symmetric approximation which enforces constant j actually works extremely well."," In the simulations the changes in potential are not exactly symmetric (e.g. lower panel of Figure \ref{fig:HT-fluctuations}) ); however we will see below that, for the purposes of calculating real-space density profiles, the symmetric approximation which enforces constant $j$ actually works extremely well."
" As before, the energy shift for one jump is given by averaging over possible orbital phases."," As before, the energy shift for one jump is given by averaging over possible orbital phases."
" However the potential Vphere is no longer an exact power law, so the calculation required is where the time integrals are evaluated over an orbital period; after changing variables to r this corresponds to integrating over the region where the integrand is real."," However the potential $V_{\mathrm{sphere}}$ is no longer an exact power law, so the calculation required is where the time integrals are evaluated over an orbital period; after changing variables to $r$ this corresponds to integrating over the region where the integrand is real."
 Equation (12)) agrees with equation (3)) for the special case of power-law potentials., Equation \ref{eq:delta-E1}) ) agrees with equation \ref{eq:deltaE-virial}) ) for the special case of power-law potentials.
" The remainder of this Section applies expression (12)) recursively to a time-series of potentials from the HT flattening) simulation, at each step updating AV, E and Verepriately?.."," The remainder of this Section applies expression \ref{eq:delta-E1}) ) recursively to a time-series of potentials from the HT (cusp-flattening) simulation, at each step updating $\Delta V$ , $E$ and $V_{\mathrm{eff}}$."
 The energy gain is evaluated at every stored simulation timestep; the relevant outputs are written every δί~27Myr.," The energy gain is evaluated at every stored simulation timestep; the relevant outputs are written every $\delta t\simeq 27\,\Myr$."
 Thus changes occurring on timescales <δί will implicitly be classified as “rapid” (composed of one jump) whereas those occurring on timescales >>dt will automatically be treated as “adiabatic” (composed of many small steps)., Thus changes occurring on timescales $\le \delta t$ will implicitly be classified as “rapid” (composed of one jump) whereas those occurring on timescales $\gg \delta t$ will automatically be treated as “adiabatic” (composed of many small steps).
" While the boundary between these limits cannot be uniquely defined, the change in behaviour must occur at around the orbital period for a particle, which is indeed ~25Myr."," While the boundary between these limits cannot be uniquely defined, the change in behaviour must occur at around the orbital period for a particle, which is indeed $\sim 25\, \Myr$."
 We verified by running checks with only every second timestep (dt~ 54Myr) that the results presented are insensitive to the precise time-slicing.," We verified by running checks with only every second timestep $\delta t \simeq 54\,\Myr$ ) that the results presented are insensitive to the precise time-slicing."
" The solid lines in Figure 4 show the resulting mean radius (ή) of orbits as a function of time, where The values of j and Eo for each orbit are chosen by requiring the initial motion to be circular at a range of different radii."," The solid lines in Figure \ref{fig:radial-migration} show the resulting mean radius $\langle r \rangle$ of orbits as a function of time, where The values of $j$ and $E_0$ for each orbit are chosen by requiring the initial motion to be circular at a range of different radii."
" As time progresses, the orbits starting interior to 1kpc migrate outwards, reflecting a net gain in energy."," As time progresses, the orbits starting interior to $1\,\kpc$ migrate outwards, reflecting a net gain in energy."
 Orbits outside this radius are largely unaffected., Orbits outside this radius are largely unaffected.
" In the LT run, by contrast, no tracer particles gain energy;those that start on circular orbits, for instance, are predicted to remain at the same radius for the entire run."," In the LT run, by contrast, no tracer particles gain energy;those that start on circular orbits, for instance, are predicted to remain at the same radius for the entire run."
[or Carina not to develop winds before each star formation episode is completed. again. unrealistically high numbers.,"for Carina not to develop winds before each star formation episode is completed, again, unrealistically high numbers."
 A carina.wind mocdel would give as=0.37. only mareinally acceptable from the point of view of more complete models of gas heating in small svstems. although it would: vield present cay average metallicities. matching the observed central values.," A carina.wind model would give a $\gamma=0.37$, only marginally acceptable from the point of view of more complete models of gas heating in small systems, although it would yield present day average metallicities matching the observed central values."
 The temporal evolution of model carina.dm is shown in figure 4. with the different panels. ancl curves. being totally analogous to those of figure 1.," The temporal evolution of model carina.dm is shown in figure 4, with the different panels and curves being totally analogous to those of figure 1."
 In. panel (a) we see the S£Iges. and the one used by the model. having very similar time structures. except for the total cessation in star forming activity. between the two main bursts assumed by the model. and not seen in the directly inferred star formation history.," In panel (a) we see the $SFR_{HGV}$, and the one used by the model, having very similar time structures, except for the total cessation in star forming activity between the two main bursts assumed by the model, and not seen in the directly inferred star formation history."
 This probably reflects second. order ellects. not contemplated by the simple modeling attempted here. which however. does include the broad. behaviour of the galaxy. and hence we expect to give valuable constraints on the physics of the evolution. if only at an approximate level.," This probably reflects second order effects, not contemplated by the simple modeling attempted here, which however, does include the broad behaviour of the galaxy, and hence we expect to give valuable constraints on the physics of the evolution, if only at an approximate level."
 In panel (b) we see that the average metallicity of the stars falls a little during the second. accretion. phase. which we take as being composed of primordial material.," In panel (b) we see that the average metallicity of the stars falls a little during the second accretion phase, which we take as being composed of primordial material."
 Again. the observed. metallicity range falls well within the eas metallicity during the two periods of star formation. so that our model comfortably accounts for the presence of a significant number of stars within the observed range.," Again, the observed metallicity range falls well within the gas metallicity during the two periods of star formation, so that our model comfortably accounts for the presence of a significant number of stars within the observed range."
 Panel (c) shows clearly how the star formation process is limited. by the appearance of galactic winds. once the thermal energv of the gas surpasses the gravitational potential energv.," Panel (c) shows clearly how the star formation process is limited by the appearance of galactic winds, once the thermal energy of the gas surpasses the gravitational potential energy."
 In Panel (d) we see the rates of the cdillerent types of SNea. with the memory of the first star formation episode alfecting the dynamics of the second. through the extended SNla rates.," In Panel (d) we see the rates of the different types of SNea, with the memory of the first star formation episode affecting the dynamics of the second, through the extended SNIa rates."
 Panel (ο) gives the cumulative amounts of stars and gas present. with the gas content being totally cleared. olf after the second wind.," Panel (e) gives the cumulative amounts of stars and gas present, with the gas content being totally cleared off after the second wind."
 This is seen in the final panel. where the two accretion phases and the two galactic winds are shown.," This is seen in the final panel, where the two accretion phases and the two galactic winds are shown."
 Leo Ll is very similar to Carina in its star formation history. this is rellectecd in the second set of models shown in table 2. where all numbers very closely follow what was obtained for Carina.," Leo I is very similar to Carina in its star formation history, this is reflected in the second set of models shown in table 2, where all numbers very closely follow what was obtained for Carina."
 In figure 5 we show the tempora evolution of model. ουσ., In figure 5 we show the temporal evolution of model leoi.dm.
 This figure again closely resembles the results of Carina. with the only cdillercnce being accordance in the predicted average metallicities an observed: values.," This figure again closely resembles the results of Carina, with the only difference being accordance in the predicted average metallicities and observed values."
 Our physical assumptions are shown to be consistent with observational data in predicting the tota clearing of the σας from this galaxy., Our physical assumptions are shown to be consistent with observational data in predicting the total clearing of the gas from this galaxy.
 Bowen et al. (, Bowen et al. (
1997) finc no gas around Leo | using three lines of sight towards distan QSOs. and searching for absorption features in the spectra.,"1997) find no gas around Leo I using three lines of sight towards distant QSOs, and searching for absorption features in the spectra."
 The values of fpi slightly. above unity. Le. the requirement of à core radius in the dark matter halo slightly above current observational estimates for the tidal radii of these systems. is totally consistent with the very recent dynamical studies of Ixlevna et. al. (," The values of $f_{DM}$ slightly above unity, i.e. the requirement of a core radius in the dark matter halo slightly above current observational estimates for the tidal radii of these systems, is totally consistent with the very recent dynamical studies of Kleyna et al. ("
2001) and the photomertic surveys of Odenkichen ct al. (,2001) and the photomertic surveys of Odenkichen et al. (
2001) viclcling precisely this conclusion.,2001) yielding precisely this conclusion.
 Llere the values of extra metal loss required. by. the models to agree with the ranges of observed. metallicities. are given by Asy.," Here the values of extra metal loss required by the models to agree with the ranges of observed metallicities, are given by $\Delta\gamma_Z$."
 We see that for Carina. in all cases. we require between and of metal expulsion for our models to agree with the data.," We see that for Carina, in all cases, we require between and of metal expulsion for our models to agree with the data."
 In the case of Leo 1. very little of this elfect is needed to agree with the upper limits of the measurements. although high. values of are again needed to reach the lower bound in the observed. metallicities.," In the case of Leo I, very little of this effect is needed to agree with the upper limits of the measurements, although high values of are again needed to reach the lower bound in the observed metallicities."
 In our treatment of Ursa Minor in section 3.1 we fixed our mocels in order to comply with S£Cge. however. as mentioned in the description of our method. the inference of LIGY loses all temporal resolution for ages greater than 10 Gyr.," In our treatment of Ursa Minor in section 3.1 we fixed our models in order to comply with $SFG_{HGV}$, however, as mentioned in the description of our method, the inference of HGV loses all temporal resolution for ages greater than 10 Gyr."
 In this wav. the time structure inferred for Ursa Minor. could well be only an artifact of the method.," In this way, the time structure inferred for Ursa Minor could well be only an artifact of the method."
 As remarked in the description of our present results for this ealaxy. the clement ratios we obtain are all similar to wha was obtained for Leo LL with values of O/Fe] always above about 0.2.," As remarked in the description of our present results for this galaxy, the element ratios we obtain are all similar to what was obtained for Leo II with values of [O/Fe] always above about 0.2."
 In comparing with the detailed observationa determinations for Ursa Minor. we find that for values of Η]ς- 1.6 our corresponding predictions for O/Fo] closely match observations for this galaxy.," In comparing with the detailed observational determinations for Ursa Minor, we find that for values of $<$ -1.6 our corresponding predictions for [O/Fe] closely match observations for this galaxy."
 On the other hand. the most metal-rich data point for Ursa Minor of Fe/L1]e-1.4 corresponds to an upper limit of O/Fe] 0.0.," On the other hand, the most metal-rich data point for Ursa Minor of $\approx$ -1.4 corresponds to an upper limit of [O/Fe] $\approx$ 0.0."
 Our single, Our single
There are two classes of persistent sources at cosniüc distances: galaxies and quasars/Active Calactic Nuclei (AGN).,There are two classes of persistent sources at cosmic distances: galaxies and quasars/Active Galactic Nuclei (AGN).
 Stars power galaxies while accretion outo and/or spin down of supermassive black holes power quasars., Stars power galaxies while accretion onto and/or spin down of supermassive black holes power quasars.
 τι. receutlv. ealactic and quasar phenomena were thought to be separate on both observational and theoretical grounds.," Until recently, galactic and quasar phenomena were thought to be separate on both observational and theoretical grounds."
" However. the discovery of the black hole mass — bulee stellar mass relation (Af,— AS.) iu nearby elliptical galaxies and the black hole mass stellar velocity dispersion relationteeinaineg2M. indicates that ealactic aud black hole activity are closely connected to one another."," However, the discovery of the black hole mass – bulge stellar mass relation $M_{\bullet}-M_{\star}$ ) in nearby elliptical galaxies and the black hole mass – stellar velocity dispersion relation indicates that galactic and black hole activity are closely connected to one another."
 The natural iuplication is that the energy release resulting from the »xiüld-up of the black hole mass limits auy further erowth of both the stellar bulge aud the black hole., The natural implication is that the energy release resulting from the build-up of the black hole mass limits any further growth of both the stellar bulge and the black hole.
" The fact hat the AJ,AZ, relation holds for nearlv four decades in black hole mass seecnis to suggest that aq""ieersal.sclf-sinilur oy process is at work. which acts ο self-regulate the ratio between black hole and bulge nass. irrespective of their combined mass."," The fact that the $M_\bullet-M_{\star}$ relation holds for nearly four decades in black hole mass seems to suggest that a, or process is at work, which acts to self-regulate the ratio between black hole and bulge mass, irrespective of their combined mass."
 Apparcutly. he only question that remains is with regard to the exact physical miechamisin responsible for black hole feedback and seltxegulation.," Apparently, the only question that remains is with regard to the exact physical mechanism responsible for black hole feedback and self-regulation."
 The aremuent. aloug with the work of?.. indicates hat the mass of supermassive black holes is mostly accrued ο an optically-huninous radiativelv-eficieut “quasar phase.”," The argument, along with the work of, indicates that the mass of supermassive black holes is mostly accrued during an optically-luminous radiatively-efficient “quasar phase.”"
" The energy released during the accretion oocess, Which is carried away primarily by photous and/or a “quasar wind.” may couple to the interstellaro uediu of the host galaxy aud eject it from the ealactic eravitational potential?)."," The energy released during the accretion process, which is carried away primarily by photons and/or a “quasar wind,” may couple to the interstellar medium of the host galaxy and eject it from the galactic gravitational potential."
. Tn doing so. fucl or any further galactfie aud 2227???quasar activity is removed. and the mass of the black hole. as well as of the stellar »ilee. is selt-Iuuited.," In doing so, fuel for any further galactic and quasar activity is removed and the mass of the black hole, as well as of the stellar bulge, is self-limited."
 The energy released during the accretion process lay ο carried away uot solely in the forme of photons., The energy released during the accretion process may be carried away not solely in the form of photons.
" Iu he so called ""radio-loud (as opposed to “racdio-quict™} 6jects. relativistic collumated outflows. or “jets.” put out a significant amount of enerev m mechanical form."," In the so called “radio-loud” (as opposed to “radio-quiet”) objects, relativistic collimated outflows, or “jets,” put out a significant amount of energy in mechanical form."
 Although radio-loud phenomena are also observed iu objects that are not actively accreting??7j.. the radio jet is more likely to affect the evolution of the system when a significant amount of mass is boiug built up.," Although radio-loud phenomena are also observed in objects that are not actively accreting, the radio jet is more likely to affect the evolution of the system when a significant amount of mass is being built up."
 As this work focuses on the selfreeulation of black hole growth. which occurs at high accretion rates. our itention rests on objects that are both siguificautlv accreting and racio-lIoud.," As this work focuses on the self-regulation of black hole growth, which occurs at high accretion rates, our attention rests on objects that are both significantly accreting and radio-loud."
" The kinetic power of radio jets is dissipated in sub-pc scale ""radio cores” aud kpe to AIpe scale “radio lobes.” with comparable amounts of energy dissipated at cach ste."," The kinetic power of radio jets is dissipated in sub-pc scale “radio cores” and kpc to Mpc scale “radio lobes,” with comparable amounts of energy dissipated at each site."
 The radio-loud quasar phase could be responsible for black hole selt-regulatiou if the cucrey release from the radio core. unlike that from the distant radio lobes. has the opportunity to couple to the interstellar medi of the host galaxy.," The radio-loud quasar phase could be responsible for black hole self-regulation if the energy release from the radio core, unlike that from the distant radio lobes, has the opportunity to couple to the interstellar medium of the host galaxy."
 , 
spectral tvpes were derived [rom visual classification (visual pattern matching of our smoothed program star spectra wilh standard star spectra) supported by quantitative analvsis of some spectral indices.,Spectral types were derived from visual classification (visual pattern matching of our smoothed program star spectra with standard star spectra) supported by quantitative analysis of some spectral indices.
 In (his section. we provide descriptions of (he absorption lines used to classifiv three broad groups. starting with the earliest spectral (wpe stars (D-À). moving to the F-IxX stars. ancl finally the AI stars.," In this section, we provide descriptions of the absorption lines used to classifiy three broad groups, starting with the earliest spectral type stars (B-A), moving to the F-K stars, and finally the M stars."
 We conclude the section. with a diseussion of eravitv-sensitive absorption features in the 5820-8700 sspectral range., We conclude the section with a discussion of gravity-sensitive absorption features in the 5820-8700 spectral range.
 For the purposes of matching spectral features with those of standard stars. our Ivdra spectra were smoothed using a gaussian filler to the resolution of the standaxd stars for direct comparison.," For the purposes of matching spectral features with those of standard stars, our Hydra spectra were smoothed using a gaussian filter to the resolution of the standard stars for direct comparison."
 All spectra have been normalized to 1 by dividing out a fit to the continuae. carefully. excluding regions with emission lines or broad absorption due to TiO and VO.," All spectra have been normalized to 1 by dividing out a fit to the continuae, carefully excluding regions with emission lines or broad absorption due to TiO and VO."
 Normalized spectra smoothled to a resolution of 5.7 aare shown in Fig., Normalized spectra smoothed to a resolution of 5.7 are shown in Fig.
 1 for a representative sample of program objects. wilh earlv-tvpe stars in Fig.," 1 for a representative sample of program objects, with early-type stars in Fig."
 la (B3-G9). IN stars in Fig.," 1a (B3-G9), K stars in Fig."
 Ib. early-to-micl M stars in Fig.," 1b, early-to-mid M stars in Fig."
 le. and mid-to-late M stars in Fig.," 1c, and mid-to-late M stars in Fig."
 ld., 1d.
 Both photospheric and telluric spectral features are labeled., Both photospheric and telluric spectral features are labeled.
 Two main sets of standards were used for classification (both qualitative and cuantitative)., Two main sets of standards were used for classification (both qualitative and quantitative).
 First. optical spectra from the WIYN/IIvcdra study of the Praesepe by Allen Strom (1995) were used to derive spectral types from Οδ) - ALLY. The effective resolution of these spectra was 5.7A.," First, optical spectra from the WIYN/Hydra study of the Praesepe by Allen Strom (1995) were used to derive spectral types from B8V - M4V. The effective resolution of these spectra was 5.7."
. For giants and later tvpe dwarls (AISV - M9V). optical spectra from the study of Ixirkpatrick. Ilenrv. AMeCarthy (1991) were used with an effective resolution of either 5 or 18A.," For giants and later type dwarfs (M5V - M9V), optical spectra from the study of Kirkpatrick, Henry, McCarthy (1991) were used with an effective resolution of either 8 or 18."
. In addition to these. optical spectra of very late (ype subgiants in IC 348 (Luhman 1999) and p Oph (Luhman. Liebert. Rieke 1997) were used for comparison with the coolest stars in our sample.," In addition to these, optical spectra of very late type subgiants in IC 348 (Luhman 1999) and $\rho$ Oph (Luhman, Liebert, Rieke 1997) were used for comparison with the coolest stars in our sample."
 For stars ealier (han DSV. we referred to the spectral atlas of and Weaver (1993).," For stars ealier than B8V, we referred to the spectral atlas of Torres-Dodgen and Weaver (1993)."
 Absorption lines [rom the Balmer (n22) and Paschen (1=3) series of hvdrogen are prominent in (he spectra of early (vpe stus., Absorption lines from the Balmer (n=2) and Paschen (n=3) series of hydrogen are prominent in the spectra of early type stars.
 In our spectra. we see tnblenced absorption lines from Ha (6563 À)) and Paschen 14 (8598 A)) (see Fig.," In our spectra, we see unblended absorption lines from $\alpha$ (6563 ) and Paschen 14 (8598 ) (see Fig."
 la) that reach a maximum around AO ancl weaken in warmer stars., 1a) that reach a maximum around A0 and weaken in warmer stars.
 The Ca II triplet (S498À.. 8542Α.. 8662 À)) is also observed aud decreases in strength toward early-(wpe stars until overtaken by Pa 16. 15. and 13(8502AÀ..8545A.. and 8665À)).," The Ca II triplet (8498, 8542, 8662 ) is also observed and decreases in strength toward early-type stars until overtaken by Pa 16, 15, and 13, and )."
 lence an F2 star may have a similar EWí(l1la) as a D5 star. but will be distinct by displaving stronger absorption from the Ca II triplet.," Hence an F2 star may have a similar $\alpha$ ) as a B5 star, but will be distinct by displaying stronger absorption from the Ca II triplet."
 We note that all of the alorementioned lines can appear in emission. aud (hat the observed absorption line strengths could be lower limits to the true strengths.," We note that all of the aforementioned lines can appear in emission, and that the observed absorption line strengths could be lower limits to the true strengths."
 We estimate that the uncertainties, We estimate that the uncertainties
 (Zwicky (c.g...Fabricautetal.1980)..," \citep[e.g.,][and references
therein]{cirsmf,rines08,vikhlinin09b,henry09,mantz08,rozo08}."
 1972).. (e.g.Majumdar&Moblr2001).. (Ixravtsov2008).," \citep[][]{zwicky1937} \citep[e.g.,][]{flg80}, \citep[e.g.,][]{smith05,richard10}, \citep[SZE][]{sz72}. \citep[e.g.,][]{majumdar04}. \citep[][]{kravtsov06,rozo08}."
. (Nagaictal.2007) (Motletal.2005) 2008)..," \citep{nagai07} \citep{motl05} \citep[e.g.,][]{henry09,lopes09b,mantz09b,locutushuang09}. ,"
 aud recent results from hydrodynamical simulations indicate that virial masses may have scatter as snall as ~5% (Lauctal.2010)., and recent results from hydrodynamical simulations indicate that virial masses may have scatter as small as $\sim$ \citep{lau10}.
.. Previous studies lave compared SZE signals to lydrostatic X-ray masses (Bonamenteetal.2008:Plageeetal.2010) and gravitational lensing masses (ALarroueetal.2009.hereafter AIO9).," Previous studies have compared SZE signals to hydrostatic X-ray masses \citep{bonamente08,plagge10} and gravitational lensing masses \citep[][hereafter M09]{marrone09}."
. Here. we make the first conrparison between virial masses of galaxy clusters and their SZE signals.," Here, we make the first comparison between virial masses of galaxy clusters and their SZE signals."
 We use SZE measurements from the literature and newly-measured virial masses of 15 clusters from extensive MIAIT/Tectospec spectroscopy., We use SZE measurements from the literature and newly-measured virial masses of 15 clusters from extensive MMT/Hectospec spectroscopy.
 This comparison tests the robustuess of the SZE as a proxy for cluster mass and the physical relatiouship between the SZE sienal and cluster mass., This comparison tests the robustness of the SZE as a proxy for cluster mass and the physical relationship between the SZE signal and cluster mass.
 Large SZ cluster surveys are underway aud are begining to vield cosmological constraints (Carlstrometal.2010:IBucksetal.2010:Staniszewskiot 2009).," Large SZ cluster surveys are underway and are beginning to yield cosmological constraints \citep{carlstrom10,hincks10,staniszewski09}."
". We assune a cosinoloey of 0,,20.3. O4-—0.7. aud Πιτ kins + + for all cealeulatious."," We assume a cosmology of $\Omega_m$ =0.3, $\Omega_\Lambda$ =0.7, and $H_0$ =70 km $^{-1}$ $^{-1}$ for all calculations."
 We are completing the Hectospec Cluster Survey (IleCS). a study of an N-vayv flux-limited sample of 53 ealaxy clusters at moderate redshift with extensive spectroscopy from ADIT/Iectospec.," We are completing the Hectospec Cluster Survey (HeCS), a study of an X-ray flux-limited sample of 53 galaxy clusters at moderate redshift with extensive spectroscopy from MMT/Hectospec."
 HeCS includes all clusters with ROSAT X-ray fluxes of fy>5«10 Pere tat [0.5-2.0]keV from the Bright Cluster Survey Ebelingetal.L998) or REPLEN siuvex (Bohringerctal.2001) with optical imaging in the Sixth. Data Release (DRG) of SDSS (Adchluan-AIcCarthyetal.2008)., HeCS includes all clusters with ROSAT X-ray fluxes of $f_X>5\times10^{-12}$ erg $^{-1}$ at [0.5-2.0]keV from the Bright Cluster Survey \citep[BCS][]{bcs} or REFLEX survey \citep{reflex} with optical imaging in the Sixth Data Release (DR6) of SDSS \citep{dr6}.
. We use DR6 photometry to select IHectospec targets., We use DR6 photometry to select Hectospec targets.
 The ITeCS targets are allbrighter than r=20.8 (SDSS catalogs are coniplete for point sources to 722.2)., The HeCS targets are allbrighter than $r$ =20.8 (SDSS catalogs are complete for point sources to $r$$\approx$ 22.2).
 Out of the Πος5 siuuple. 15 clusters have published SZ measurements.," Out of the HeCS sample, 15 clusters have published SZ measurements."
(oxvgen is a typical a-clement) for several models: as one can see. modela2e.. with the vields by HNV02.. predicts exactly the same behaviour of the O/Ec] ratio as model (standard case withot pair-creation SNe) except for the very carly phases.,"(oxygen is a typical $\alpha$ -element) for several models: as one can see, model, with the yields by HW02, predicts exactly the same behaviour of the [O/Fe] ratio as model (standard case withot pair-creation SNe) except for the very early phases."
 The case with poplll pair- creation SNe and LIWO2 vields. in fact. starts with a quite lower O/Fe] ratio. relative to the standard case. due to the fact that pair creation SNe favor the xoduction of Fe (at variance with the results of OFES3).," The case with popIII pair- creation SNe and HW02 yields, in fact, starts with a quite lower [O/Fe] ratio, relative to the standard case, due to the fact that pair creation SNe favor the production of Fe (at variance with the results of OFE83)."
 This ratio stavs constant while the οΗ] decreases and then increases to reach the value of the stancard case when the pair-creation supernovae disappear., This ratio stays constant while the [Fe/H] decreases and then increases to reach the value of the standard case when the pair-creation supernovae disappear.
 This inversion in case is due to the presence of infall of material of primordial chemical composition., This inversion in case is due to the presence of infall of material of primordial chemical composition.
 In fact. when the first very massive stars die. the ΠΟΠΗ] in the ISM jumps immediately at the value of be/LJ=-1.0 but soon this value decreases due to the infalling gas.," In fact, when the first very massive stars die, the [Fe/H] in the ISM jumps immediately at the value of [Fe/H]=-1.0 but soon this value decreases due to the infalling gas."
 Model is the equivalent of the standard model without infall (CB)., Model is the equivalent of the standard model without infall (CB).
 In this case. the model predicts a lower O/Fe] ratio at the beginning which increases later on. but no inversion in the Le/Ll].," In this case, the model predicts a lower [O/Fe] ratio at the beginning which increases later on, but no inversion in the [Fe/H]."
 In figure 2 we show {1e same plot of O/Fe] vs. ο] for the models ancl bl (shown again for comparison)., In figure 2 we show the same plot of [O/Fe] vs. [Fe/H] for the models and b1 (shown again for comparison).
 Model cilfers from mocel only for the nucleosvnthesis in pair-creation SNe. which is [rom OFES3.," Model differs from model only for the nucleosynthesis in pair-creation SNe, which is from OFE83."
 One can inimediately. notice the large dilference in the predictions of the two models: model predicts a very high oxveen overabundance relative to Fe in the very carly phases. due to the lack of Fe-peak elements in the ΟΙΤΗ vields.," One can immediately notice the large difference in the predictions of the two models: model predicts a very high oxygen overabundance relative to Fe in the very early phases, due to the lack of Fe-peak elements in the OFE83 yields."
 1n Figure 3 we show the O/Fe] vs. Fe/l for the strongly bimodal star formation cases (only very massive stars in the early phases) lasting for 0.1. Gar., In Figure 3 we show the [O/Fe] vs. [Fe/H] for the strongly bimodal star formation cases (only very massive stars in the early phases) lasting for 0.1 Gyr.
 Models. (yields IIWO2) and. (sields UN) show a rather constant O/Fc] ratio over the whole Fel] range., Models (yields HW02) and (yields UN) show a rather constant [O/Fe] ratio over the whole [Fe/H] range.
 This is due to the fact that the very massive pop LLL stars. in the most. recent nucleosvnthesis calculations. produce an almost solar O/Lc ratio and this will predominate also in the subsequent evolution.," This is due to the fact that the very massive pop III stars, in the most recent nucleosynthesis calculations, produce an almost solar O/Fe ratio and this will predominate also in the subsequent evolution."
 On the other hand. model shows a very high oxvgen overabundance relative to Fe. again due to the lack of Fe-peak elements in the vields of OEESS.," On the other hand, model shows a very high oxygen overabundance relative to Fe, again due to the lack of Fe-peak elements in the yields of OFE83."
 The other models not included in the Figure. where the pop LL stars form only for a very short time interval (0.01 Gyr) do not produce noticeable dillerences in the results relative to the standard el moclel.," The other models not included in the Figure, where the pop III stars form only for a very short time interval (0.01 Gyr) do not produce noticeable differences in the results relative to the standard $a1$ model."
 While the abundances in Figures |. 2 and 3 refer to the eas. in Figures 4 and 5 we show the predicted distributions of stars as functions of Fe/1] for models -age.52s. ese adc.," While the abundances in Figures 1, 2 and 3 refer to the gas, in Figures 4 and 5 we show the predicted distributions of stars as functions of [Fe/H] for models -, - ."
 In models afe-- the predicted stellar distributions are almost indistinguishable except for the absence of stars with Fe/L]«3.0 in the case with pop ILE stars., In models - the predicted stellar distributions are almost indistinguishable except for the absence of stars with $<-3.0$ in the case with pop III stars.
 Ehe reason [or this resides in the fact that the pop LL phase is very short and at the same time the star formation rate is small at carly stages when there is little gas., The reason for this resides in the fact that the pop III phase is very short and at the same time the star formation rate is small at early stages when there is little gas.
 In the closed-box cases (models and 62)) the dillerence is more noticeable since at the beginning the star formation is quite high., In the closed-box cases (models and ) the difference is more noticeable since at the beginning the star formation is quite high.
 In both models. in fact. no stars with metallicity lower than -3.0 ancl -2.0. respectively. are predicted (sce Figure 4).," In both models, in fact, no stars with metallicity lower than -3.0 and -2.0, respectively, are predicted (see Figure 4)."
 The, The
relativistic Cengine-«driven) supernovae (e.g. Chakrabortietal.301001.,relativistic (`engine-driven') supernovae (e.g. \citealt{Chakraborti+10}) ).
 The arrival directions of UHECRs provide a potentially important probe of their origin., The arrival directions of UHECRs provide a potentially important probe of their origin.
 Measurements by the Pierre Auger Observatory (PAO) rule out isotropy for the highest energy cosmic rays at ~98% confidence (Armengaud20050). and PAO has furthermore discovered a correlation between the arrival directions of UHECRs with energies E>57 EeV and nearby εἰς75 Mpc) AGN (Abrahametal. 2008)).," Measurements by the Pierre Auger Observatory (PAO) rule out isotropy for the highest energy cosmic rays at $\sim98\%$ confidence \citealt{Armengaud+08}) ), and PAO has furthermore discovered a correlation between the arrival directions of UHECRs with energies $E > 57$ EeV and nearby $\lesssim 75$ Mpc) AGN \citealt{Abraham+08}) )."
 This result does not. however. imply that UHECRSs necessarily originate from AGN. because AGN trace local Galactic structure. such that the correlation is consistent with a variety of other sources (Kashti&Waxman2008:: Ghisellinietal. 2008:: Takamietal. 2009:: Takami&Sato2009)).' At present the sources of UHECR cannot therefore be deduced from their arrival directions alone.," This result does not, however, imply that UHECRs necessarily originate from AGN, because AGN trace local Galactic structure, such that the correlation is consistent with a variety of other sources \citealt{Kashti&Waxman08}; \citealt{Ghisellini+08}; \citealt{Takami+09}; \citealt{Takami&Sato09}) At present the sources of UHECR cannot therefore be deduced from their arrival directions alone."
 The composition of UHECRs also provides important clues to their origin., The composition of UHECRs also provides important clues to their origin.
" Although the composition is measured directly at low energies €x10"" eV). at ultra-high energies it must be inferred indirectly by measuring the shower depth at maximum elongation Xy."," Although the composition is measured directly at low energies $\lesssim 10^{14}$ eV), at ultra-high energies it must be inferred indirectly by measuring the shower depth at maximum elongation $X_{\rm max}$."
" Recent measurements by PAO show that the average shower depth (X4,) and its RMS variation decrease systematically moving to the highest energies (Abrahametal.2010).", Recent measurements by PAO show that the average shower depth $\langle X_{\rm max} \rangle$ and its RMS variation decrease systematically moving to the highest energies \citep{Abraham+10}.
. This suggests that the UHECR composition transitions from being dominated by protons below the ankle to being dominated by heavier nuclei with average masses similar to Si or Fe at ~5x10! eV. We caution. however. that HiRes has not verified this finding (Abbasietal.2008).," This suggests that the UHECR composition transitions from being dominated by protons below the ankle to being dominated by heavier nuclei with average masses similar to Si or Fe at $\sim 5\times 10^{19}$ eV. We caution, however, that HiRes has not verified this finding \citep{Abbasi+05}."
 The UHECR composition measured by Auger is puzzling., The UHECR composition measured by Auger is puzzling.
" One possible explanation is that the accelerated material has an intrinsically ""mixed"" composition (with e.g. solar abundances). such that protons are accelerated to a maximum energy E=Esma~10? eV. beyond which only heavier nuclei are accelerated."," One possible explanation is that the accelerated material has an intrinsically `mixed' composition (with e.g. solar abundances), such that protons are accelerated to a maximum energy $E = E_{\rm p,max} \sim 10^{18.5}$ eV, beyond which only heavier nuclei are accelerated."
 This seems plausiblepriori because accelerator size considerations show that the maximum achievable energy increases linearly with the nuclear charge Z (Hillas1984)., This seems plausible because accelerator size considerations show that the maximum achievable energy increases linearly with the nuclear charge $Z$ \citep{Hillas84}.
. On the other hand. this explanation appears to require fine tuning because the maximum energy to which. for instance. Fe nuclei are accelerated Eia~ZXEpnay8XLo (Z/26) eV must (by coincidence) be close to the cut-off observed at ~6x10! eV and expected to occur independently from the GZK ettect.," On the other hand, this explanation appears to require fine tuning because the maximum energy to which, for instance, Fe nuclei are accelerated $E_{\rm Fe,max} \sim Z\times E_{\rm p,max} \sim 8\times 10^{19}$ (Z/26) eV must (by coincidence) be close to the cut-off observed at $\sim 6\times 10^{19}$ eV and expected to occur independently from the GZK effect."
 A “mixed” composition with metal abundance ratios similar to the Sun or Galactic cosmic rays also appears inconsistent with modeling of the propagation of UHECRs through the EBL (Allardetal. 2008). which suggest that the injected composition has a fairly narrow distribution in charge (e.g. Hooper&Taylor 2010).," A `mixed' composition with metal abundance ratios similar to the Sun or Galactic cosmic rays also appears inconsistent with modeling of the propagation of UHECRs through the EBL \citealt{Allard+08}) ), which suggest that the injected composition has a fairly narrow distribution in charge (e.g. \citealt{Hooper&Taylor10}) )."
 A second possibility is that the accelerated material is dominated by heavy nuclei., A second possibility is that the accelerated material is dominated by heavy nuclei.
 In this case the proton-dominated composition measured near the ankle may be explained as secondary particles produced by the interaction of the nuclei with the EBL (e.g. Hooper&Taylor 20101)., In this case the proton-dominated composition measured near the ankle may be explained as secondary particles produced by the interaction of the nuclei with the EBL (e.g. \citealt{Hooper&Taylor10}) ).
 A heavy-rich composition is unlikely in the case of AGN. galaxy clusters. and supernova shocks because the accelerated material originates from the interstellar medium.," A heavy-rich composition is unlikely in the case of AGN, galaxy clusters, and supernova shocks because the accelerated material originates from the interstellar medium."
 For a solar composition. the fraction of the total mass in Fe nuclei and heavier is just Xi;~LOὃς such that only for extremely super-solar metallicity (~10?Z;) could heavy nuclei dominate the total UHECR mass.," For a solar composition, the fraction of the total mass in Fe nuclei and heavier is just $X_{\rm Fe} \sim 10^{-3}$, such that only for extremely super-solar metallicity $\sim 10^{3} Z_{\odot}$ ) could heavy nuclei dominate the total UHECR mass."
 In this paper we show that UHECRs from GRBs. unlike AGN. may indeed be composed of almost entirely very heavy nuclei.," In this paper we show that UHECRs from GRBs, unlike AGN, may indeed be composed of almost entirely very heavy nuclei."
 In particular. if the outflow from the central engine is strongly magnetized we find that Fe-group nuclei and possibly heavier elements (A > 90) are synthesized during its expansion.," In particular, if the outflow from the central engine is strongly magnetized we find that Fe-group nuclei and possibly heavier elements (A $\gtrsim$ 90) are synthesized during its expansion."
 Although it is well established that long duration GRBs originate from the core collapse of massive stars Woosley&Bloom2006).. it remains debated whether the central engine is a hyper-aeereting black hole (Woosley1993) ora rapidly spinning. strongly magnetized neutron star (à proto-magnetar: e.g. Usov1992)).," Although it is well established that long duration GRBs originate from the core collapse of massive stars \citep{Woosley&Bloom06}, it remains debated whether the central engine is a hyper-accreting black hole \citep{Woosley93} or a rapidly spinning, strongly magnetized neutron star (a `proto-magnetar'; e.g. \citealt{Usov92}) )."
 We focus here on the proto-magnetar model. which recent work has shown can explain many of the observed properties of GRBs (Thompsonetal.2004:: Metzgeretal.2007:: Bucciantinietal.2007:: Metzgeretal.2010).," We focus here on the proto-magnetar model, which recent work has shown can explain many of the observed properties of GRBs \citealt{Thompson+04}; \citealt{Metzger+07}; \citealt{Bucciantini+07}; \citealt{Metzger+10}) )."
 However. similar considerations may apply to acecretion-powered models. provided that the jet is magnetically-dominated rather than a thermally-driven fireball ($2.2).," However, similar considerations may apply to accretion-powered models, provided that the jet is magnetically-dominated rather than a thermally-driven fireball $\S\ref{sec:BH}$ )."
 The high temperatures 7> | MeV near the central engine imply that all nuclei are dissociated into free neutrons and protons., The high temperatures $T >$ 1 MeV near the central engine imply that all nuclei are dissociated into free neutrons and protons.
 Heavier elements form only once lower temperatures and densities are reached at larger radii in the outflow., Heavier elements form only once lower temperatures and densities are reached at larger radii in the outflow.
" If the outflow forms as a fireball dominated by thermal energy (as would occur if the jet is powered by neutrino annihilation along the rotational axis: e.g. Eichleretal. 1989)). its entropy is necessarily high S.>10"" &, '."," If the outflow forms as a fireball dominated by thermal energy (as would occur if the jet is powered by neutrino annihilation along the rotational axis; e.g. \citealt{Eichler+89}) ), its entropy is necessarily high $S \gtrsim 10^{5}$ $k_{\rm b}$ $^{-1}$."
 Free nuclei recombine into Helium only once the deuterium bottleneck is broken., Free nuclei recombine into Helium only once the deuterium bottleneck is broken.
 Since this occurs at low densities when the entropy is high. few elements heavier than He are formed. similar to Big Bang nucleosynthesis (Lemoine2002:: Beloborodov 2003).," Since this occurs at low densities when the entropy is high, few elements heavier than He are formed, similar to Big Bang nucleosynthesis \citealt{Lemoine02}; \citealt{Beloborodov03}) )."
 Pure fireballs are therefore unlikely to produce jets enriched in heavy elements., Pure fireballs are therefore unlikely to produce jets enriched in heavy elements.
 The situation is different if the jet is accelerated magnetically. as occurs from proto-magnetars or magnetized accretion disk winds.," The situation is different if the jet is accelerated magnetically, as occurs from proto-magnetars or magnetized accretion disk winds."
" In this case most of the energy is stored in the magnetic field (Poynting flux) at small radii and the flow has a much lower entropy S.-LO—300K, ! (see Fig."," In this case most of the energy is stored in the magnetic field (Poynting flux) at small radii and the flow has a much lower entropy $S \sim 10-300\,k_{\rm b}$ $^{-1}$ (see Fig."
 |. and eg. [IT] , \ref{fig:wind} and eq. \ref{eq:SNRNM}] ]
below)., below).
 Under these conditions Helium recombination occurs at higher densities. such that heavier nuclei can be formed efficiently via e.g. the triple-a reaction and subsequent a captures.," Under these conditions Helium recombination occurs at higher densities, such that heavier nuclei can be formed efficiently via e.g. the $\alpha$ reaction and subsequent $\alpha$ captures."
 Below we focus on the nucleosynthesis in proto-magnetar winds because the outflow properties can be calculated with relative contidence (Metzgeretal.2010):: however. in $2.2 we briefly discuss the composition of accretion-powered outflows.," Below we focus on the nucleosynthesis in proto-magnetar winds because the outflow properties can be calculated with relative confidence \citep{Metzger+10}; however, in $\S\ref{sec:BH}$ we briefly discuss the composition of accretion-powered outflows."
 When a massive star runs out of nuclear fuel. its core undergoes gravitational collapse.," When a massive star runs out of nuclear fuel, its core undergoes gravitational collapse."
" This results in a hot ""proto-neutron' star (proto-NS). which radiates the energy released during the collapse in neutrinos (e.g. Burrows&Lattimer |986))."," This results in a hot `proto-neutron' star (proto-NS), which radiates the energy released during the collapse in neutrinos (e.g. \citealt{Burrows&Lattimer86}) )."
 As neutrinos escape. they heat the material above the proto-NS surface. potentially powering a supernova (SN) explosion during the first few hundred milliseconds after core bounce (e.g. Bethe&Wilson 19855).," As neutrinos escape, they heat the material above the proto-NS surface, potentially powering a supernova (SN) explosion during the first few hundred milliseconds after core bounce (e.g. \citealt{Bethe&Wilson85}) )."
 However. regardless of how the star explodes. if the core does not collapse into a black hole. neutrinos continue to heat the proto-NS atmosphere on longer timescales t ~|—100 s. This drives mass from the proto-NS into the expanding cavity behind the outgoing," However, regardless of how the star explodes, if the core does not collapse into a black hole, neutrinos continue to heat the proto-NS atmosphere on longer timescales t $\sim 1-100$ s. This drives mass from the proto-NS into the expanding cavity behind the outgoing"
some of these sources. a change in the column density of the absorber. rather than a switchingolf of the source. cannot be completely ruled. out. ancl indeed. Hisaliti ct al. (,"some of these sources, a change in the column density of the absorber, rather than a switching–off of the source, cannot be completely ruled out, and indeed Risaliti et al. ("
2002) claimed that variations in the absorbing column densitv are common in Sevfert. 2 galaxies.,2002) claimed that variations in the absorbing column density are common in Seyfert 2 galaxies.
 However. the changing absorbers in the Risaliti et al. (," However, the changing absorbers in the Risaliti et al. ("
2002) sample are all Comptonthin. and in most cases the variations are too small to rule out the possibility that they are an artifact due to comparing spectra obtained: with dilleren instruments.,"2002) sample are all Compton–thin, and in most cases the variations are too small to rule out the possibility that they are an artifact due to comparing spectra obtained with different instruments."
 Moreover. this solution is clearly untenable for GC 2992 (Cail et al.," Moreover, this solution is clearly untenable for NGC 2992 (Gilli et al."
 2000: see Sec., 2000; see Sec.
 2.4). which has been well monitored over the vears. showing a gradual change of 10 nuclear [lux and a costant absorber.," 2.4), which has been well monitored over the years, showing a gradual change of the nuclear flux and a costant absorber."
 We find. cillieul o imagine a situation in which a Comptonthick absorber on the pescale (as suggeste by the lack of variability. of e reflection components) and. with a large covering factor (to allow for the rather large reflection components) can Pvary so dramatically on timescales of vears., We find difficult to imagine a situation in which a Compton–thick absorber on the pc–scale (as suggested by the lack of variability of the reflection components) and with a large covering factor (to allow for the rather large reflection components) can vary so dramatically on time–scales of years.
 Pherefore. in 1e following we will assume that the observed. variations re due to the switchingolf of the nucleus.," Therefore, in the following we will assume that the observed variations are due to the switching–off of the nucleus."
 After reviewing jo current observation status of this field (Sect., After reviewing the current observation status of this field (Sect.
 2). we will discuss some possible implications (Sect.," 2), we will discuss some possible implications (Sect."
 3), 3).
 UGC 4203 (a.k.a., UGC 4203 (a.k.a.
 Mkn 1210) has been recently observed by NAIALNewton (Guainazzi ct al., Mkn 1210) has been recently observed by XMM–Newton (Guainazzi et al.
 2002). unveiling a XNταν bright nucleus. absorbed by Ny2«1077 em7.," 2002), unveiling a X–ray bright nucleus, absorbed by $N_H \simeq 2 \times 10^{23}$ $^{-2}$."
 Llowever. in an ASCA observation performed about five and half vears earlier (Awaki et al.," However, in an ASCA observation performed about five and half years earlier (Awaki et al."
 2000). the prominent iron line 1 keV) and the factor of 5 lower 210 keV [ux indicated a rellection.dominated spectrum. (Fig. 5)).," 2000), the prominent iron line $EW \simeq 1$ keV) and the factor of 5 lower 2–10 keV flux indicated a reflection–dominated spectrum (Fig. \ref{polittico}) ),"
 with the nuclear emission too faint to be visible., with the nuclear emission too faint to be visible.
 The limited. bandpass of ASCA. along with the low Ilux of the source. does not permit to distinguish bewteen different column densities of the rellecting matter. provided that it exceeds about 107 7 (see Fie. 4)).," The limited bandpass of ASCA, along with the low flux of the source, does not permit to distinguish bewteen different column densities of the reflecting matter, provided that it exceeds about $10^{23}$ $^{-2}$ (see Fig. \ref{refl}) )."
 Ht is therefore possible that in this case the absorbing and rellecting materials are one and the same., It is therefore possible that in this case the absorbing and reflecting materials are one and the same.
 NGC 6300. discovered: serendipitouslv by (Awaki et al.," NGC 6300, discovered serendipitously by (Awaki et al."
 1991). was observed in a Comptonthick rellection-dominated state by RAPE on February 1997 (Leighly et al.," 1991), was observed in a Compton–thick reflection-dominated state by RXTE on February 1997 (Leighly et al."
 1999)., 1999).
 Two and half vears later a remarkably strong Seviert nucleus (210 keV Hux ~1.3.10.+ erg em7s +) seen through a column density with Nyc21075 en7 (Fig. 5)).," Two and half years later a remarkably strong Seyfert nucleus (2–10 keV flux $\sim 1.3 \times 10^{-11}$ erg $^{-2}$ $^{-1}$ ) seen through a column density with $N_H \simeq 2 \times 10^{23}$ $^{-2}$ (Fig. \ref{polittico}) ),"
 was discovered. in a BeppoSAN observation., was discovered in a BeppoSAX observation.
 An XMMNewton observation performed carly in. 2001 caught the source still in the high lux. Comptonthin state (Maclelox et al.," An XMM–Newton observation performed early in 2001 caught the source still in the high flux, Compton–thin state (Maddox et al."
 2002)., 2002).
 As the RAPE bandpass extends up to 20 keV. for this source jt is possible to distinguish between Comptonthin and Comptonthick reflection (Fig. 4)).," As the RXTE bandpass extends up to 20 keV, for this source it is possible to distinguish between Compton–thin and Compton–thick reflection (Fig. \ref{refl}) )."
 The detection in the PCA highest energy. band is too strong to be explained as pure reflection by matter with Ny&2107 7., The detection in the PCA highest energy band is too strong to be explained as pure reflection by matter with $N_H\simeq 2 \times 10^{23}$ $^{-2}$.
 In this case. therefore. the (thick) reflector must be dillerent from the (thin) absorber.," In this case, therefore, the (thick) reflector must be different from the (thin) absorber."
 A BeppoSAX observation on August 1997 detected in this source a bright Sevfert nucleus. seen through a Comptonthin (Ngc41077: Risaliti et al.," A BeppoSAX observation on August 1997 detected in this source a bright Seyfert nucleus, seen through a Compton--thin $N_H \simeq 4 \times 10^{23}$$^{-2}$; Risaliti et al."
 2000) absorber., 2000) absorber.
 On the contrary. an ASCA observation. performed three vears earlier. detected a very [lat X-ray continuum (E2 0.8) and a 2.1 keV We iron line. both indicating a rellection.dominated state (Fie. 5)).," On the contrary, an ASCA observation, performed three years earlier, detected a very flat X-ray continuum $\Gamma \simeq 0.8$ ) and a 2.1 keV $\alpha$ iron line, both indicating a reflection–dominated state (Fig. \ref{polittico}) )."
 Due to the limited bandwidth of ASCA. and similarly to UGC 4203. the possibility that the reflector. is simply the inner wall of the absorber cannot be ruled out.," Due to the limited bandwidth of ASCA, and similarly to UGC 4203, the possibility that the reflector is simply the inner wall of the absorber cannot be ruled out."
 The brightest anc best studied source in our little sample is NGC 2992. a Sevlert 1.9 galaxy with an X.ταν absorbing column clensity Ng9107 ?.," The brightest and best studied source in our little sample is NGC 2992, a Seyfert 1.9 galaxy with an X–ray absorbing column density $_{H}\sim9\times10^{21}$ $^{-2}$ ."
 The Xrav Dux of NGC 2992 steadily declined since LOTS. when it was observed," The X–ray flux of NGC 2992 steadily declined since 1978, when it was observed"
on identified clusters from ??????.. ,"on identified clusters from \cite{2001AJ....122.1796M, 2002ChJAA...2..197M, 2002AJ....123.3141M,
 2002AcA....52..453M, 2004ChJAA...4..125M, 2004A&A...413..563M}."
The SM catalogue contains 595 objects of which 428 are classified as high-confidence clusters (based on7/9T and high-resolution ground-based imaging)., The SM catalogue contains 595 objects of which 428 are classified as high-confidence clusters (based on and high-resolution ground-based imaging).
 The most recent work. currently not within the SM catalogue. includes work based on CFIUT/MeeaCam imagingby ? and? and JST imagine by ?.. thal contain 3554. 599 and 91 new star cluster candidates. respectively. (," The most recent work, currently not within the SM catalogue, includes work based on CFHT/MegaCam imagingby \cite{2008AcA....58...23Z} and \cite{2010arXiv1007.1042S} and $HST$ imaging by \cite{2009AcA....59...47Z}, that contain 3554, 599 and 91 new star cluster candidates, respectively. ("
All of these M33 studies cover only the inner one square degree.),All of these M33 studies cover only the inner one square degree.)
 ? claim that. 222054. of the 3554 cluster candidates identified in ? are likely to be genuine clusters.," \cite{2009AcA....59...47Z}
 claim that $\approx$ $\%$ of the 3554 cluster candidates identified in \cite{2008AcA....58...23Z} are likely to be genuine clusters."
 Unlike the GCSs of the MW and M31. AI33 is host to intermediate-age clusters (?2).. suggesting that the evolution of M33 was ditferent rom that of both the MW or M3I.," Unlike the GCSs of the MW and M31, M33 is host to intermediate-age clusters \citep{1998ApJ...508L..37S, 2002ApJ...564..712C}, suggesting that the evolution of M33 was different from that of both the MW or M31."
 Studying the Local Group gives us the best chance to observe the remnants of galaxy formation in detail. but. M33 remains to be scrutinized in as much detail as either of its larger neighboring galaxies. or the Magellanic Clouds.," Studying the Local Group gives us the best chance to observe the remnants of galaxy formation in detail, but M33 remains to be scrutinized in as much detail as either of its larger neighboring galaxies, or the Magellanic Clouds."
 The work on the M33 GCS has so [ar been constrained to the classical disk regions. with the exception of the four outer halo clusters found by ? between projected radii of 9.6 and 28.5 kpe and one cluster by ? al a projected radius of 12.5 kpc.," The work on the M33 GCS has so far been constrained to the classical disk regions, with the exception of the four outer halo clusters found by \cite{2009ApJ...698L..77H} between projected radii of 9.6 and 28.5 kpc and one cluster by \cite{2008AJ....135.1482S} at a projected radius of 12.5 kpc."
 The outer halo clusters ave important. not least because the most distant clusters may be the last that were accreted (e.g.. 2)).," The outer halo clusters are important, not least because the most distant clusters may be the last that were accreted (e.g., \citealt{2005MNRAS.360..631M}) )."
 ? have shown that AI3I's outer halo is rich with clusters., \cite{2010ApJ...717L..11M} have shown that M31's outer halo is rich with clusters.
 ? undertook a search for M33 outer halo clusters through 12 sq., \cite{2009ApJ...698L..77H} undertook a search for M33 outer halo clusters through 12 sq.
 degrees of the Isaac Newton Telescope Wide-Field Camera data reaching to Ve-24.5 and 123.5., degrees of the Isaac Newton Telescope Wide-Field Camera data reaching to $\sim$ 24.5 and $\sim$ 23.5.
 The PÁndAS data allow this search to be extended to larger τας. deeper depths aud better image quality.," The PAndAS data allow this search to be extended to larger radii, deeper depths and better image quality."
 This is the project that we undertake in (his paper., This is the project that we undertake in this paper.
 We define outer halo clusters to be those which are projected bevond the isophotal radius of M33 (ο kpc. 2)).," We define outer halo clusters to be those which are projected beyond the isophotal radius of M33 $\sim$ 9 kpc, \citealt{2011Cockcroftinprep}) )."
 Such objects are sufficiently. remote Chat they are unlikely to be associated with the main disk component of the ealaxv: ? find little evidence Irom clirect stellar photometry that the disk extends bevond (hat point., Such objects are sufficiently remote that they are unlikely to be associated with the main disk component of the galaxy: \cite{2010ApJ...723.1038M} find little evidence from direct stellar photometry that the disk extends beyond that point.
 For comparison. the isophotal radius of NGC 253. an Sc-tvpe galaxy of similar size. is re9.8 kpe (?).. ," For comparison, the isophotal radius of NGC 253, an Sc-type galaxy of similar size, is $r \sim 9.8$ kpc \citep{2003AJ....125..525J}. ."
Ultimately however. we will require metallicity ancl velocity measurements to determine more definitely," Ultimately however, we will require metallicity and velocity measurements to determine more definitely"
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\tikzmark{mainBodyCitationStart622}\citet[][; hereafter referred to as Paper I]{kato08} "
Extra bodies can be added in circular orbits about the center of mass.,Extra bodies can be added in circular orbits about the center of mass.
 Llaves(2003) demonstrates that as the number of moving bodies in a Ηχος potential increases the shadow clurations decrease., \cite{Hayes} demonstrates that as the number of moving bodies in a fixed potential increases the shadow durations decrease.
 Le would be of interest to determine if a similar relationship holds for the Sitnikoy N-body problem., It would be of interest to determine if a similar relationship holds for the Sitnikov $N$ -body problem.
 It should. be stressed. again that the failure ol the refinement procedure does not necessarily mean that a shaclow does not exist for a given pseudo-orbit., It should be stressed again that the failure of the refinement procedure does not necessarily mean that a shadow does not exist for a given pseudo-orbit.
 It may very well be that shadows do exist for orbits in regions where the refinement procedure fails., It may very well be that shadows do exist for orbits in regions where the refinement procedure fails.
 We are encouraged that this may Ȣ the case., We are encouraged that this may be the case.
 Both the Sitnikovy problem and the approximate 'oincaré map possess à hyperbolic invariant set. A. near he escape boundaries (see Moser(1973). and Urminsky respectively).," Both the Sitnikov problem and the approximate Poincaré map possess a hyperbolic invariant set, $\Lambda$, near the escape boundaries (see \cite{Moser} and \cite{UrminskyThesis}
 respectively)."
 Despite the fact that iX is near the x»undary. ODy. the shadowing theorems by Anosov(1967) and Bowen(1972). guarantee that any pseudo-orbit on A has an associated shacdow-orbit.," Despite the fact that $\Lambda$ is near the boundary $\partial \mathcal{D}_0$, the shadowing theorems by \cite{Anosov} and \cite{Bowen} guarantee that any pseudo-orbit on $\Lambda$ has an associated shadow-orbit."
 Phis demonstrates that being in he vicinity on the escape boundary does not necessarily rule out the existence of shadow-orbits., This demonstrates that being in the vicinity on the escape boundary does not necessarily rule out the existence of shadow-orbits.
 DU was supported by the National Acronautics and Space Administration through grant. NNX-07ALIII5CG. The author would like to thank D. Hegegie. D. Merritt and D. Dicken for their helpful suggestions.," DU was supported by the National Aeronautics and Space Administration through grant NNX-07AH15G. The author would like to thank D. Heggie, D. Merritt and D. Dicken for their helpful suggestions."
 In addition. the author would like to thank the anonymous referee for his/her careful reading of the manuscript and useful suggestions.," In addition, the author would like to thank the anonymous referee for his/her careful reading of the manuscript and useful suggestions."
"The main point of the following analysis is the determination of the Bollzmann entropy SCM.Q.V) and of the total οποιον £7,V.Q.V) along Che sequence of models. ie. as a [function of the concentration parameter V. defined above.","The main point of the following analysis is the determination of the Boltzmann entropy $S(M, Q, \Psi)$ and of the total energy $E_{tot}(M, Q, \Psi)$ along the sequence of models, i.e. as a function of the concentration parameter $\Psi$ defined above."
 These functions. at constant AL and Q. ave illustrated in Fig.," These functions, at constant $M$ and $Q$, are illustrated in Fig."
"1. They have been obtained by noting that. from (he definitions of S and f"". From the definitions Q=Adabd!YrOb) and M=Aa2dV""ALOV) andthe definition of . we can express (he variables (21.0.d) in terms of the variables (V/.Q.Y) anclthus find that (he entropy per unit mass can be written as $/M=$40M.Q)+ o(V). where Sy is constant when the values of M. and Q are fixed. with Here E=E(w) is the dimensionless (otal energy defined from Ej,=AaFd?""E."," They have been obtained by noting that, from the definitions of $S$ and $f^{(\nu)}$, From the definitions $Q = A a^{-9/4}d^{-1 - 3/\nu} \hat{Q}(\Psi)$ and $M = A a^{-9/4}d^{- 3/\nu} \hat{M}(\Psi)$ andthe definition of $\gamma$, we can express the variables $(A, a, d)$ in terms of the variables $(M, Q, \Psi)$ andthus find that the entropy per unit mass can be written as $S/M=S_0(M,Q)+\sigma(\Psi)$ , where $S_0$ is constant when the values of $M$ and $Q$ are fixed, with Here $\hat{E} = \hat{E}(\Psi)$ is the dimensionless total energy defined from $E_{tot} = A a^{-13/4} d^{-3/\nu}\hat{E}$."
" From the identity «f),,/Al=E/M and the expression of a=a(M.Q.V) obtained previously. we lind £,,/M=H(M.Q)e(V). with: The factor Z£(M.Q) is à constant when A aud (Q are taken to be constant."," From the identity $a E_{tot}/M = \hat{E}/\hat{M}$ and the expression of $a = a(M, Q, \Psi)$ obtained previously, we find $E_{tot}/M = H(M,Q)\epsilon(\Psi)$, with: The factor $H(M,Q)$ is a constant when $M$ and $Q$ are taken to be constant."
" The quantities ο], VOY). QQU). and LOY) that enter the expression of σ ancl € depend only on V and are evaluated numerically on the equilibrium sequence."," The quantities $\gamma (\Psi)$, $\hat{M}(\Psi)$, $\hat{Q}(\Psi)$, and $\hat{E}(\Psi)$ that enter the expression of $\sigma$ and $\epsilon$ depend only on $\Psi$ and are evaluated numerically on the equilibrium sequence."
 This completes the derivation that allows us to craw the analogy with the classical paper of Lynden-BellandWood(1963)., This completes the derivation that allows us to draw the analogy with the classical paper of \citet{lyn68}.
". This step. straightforward [or the /"" models. is by itself interesting ancl new."," This step, straightforward for the $f^{(\nu)}$ models, is by itself interesting and new."
 In fact. other attempts at applving theparadigm of the gravothermal cabastrophe to stellar ανασα equilibrium sequences were either based on an unjustified ansalz [οι the identification of the relevant temperature (e.g.. see Appendix V in the article bv Lynelen-BellanclWood 1968: Katz 1980: Magliocchettiοἱal. 1993)) or onthe use of non-standard entropies (for less realistic models: Chavanis 2002)).," In fact, other attempts at applying theparadigm of the gravothermal catastrophe to stellar dynamical equilibrium sequences were either based on an unjustified for the identification of the relevant temperature (e.g., see Appendix V in the article by \citealt{lyn68}; ; \citealt{kat80}; ; \citealt{mag98}) ) or onthe use of non-standard entropies (for less realistic models; \citealt{cha02}) )."
 When the /' models wereconstructed (StiavelliandBertin 1987). it was immediatelv," When the $f^{(\nu)}$ models wereconstructed \citep{sti87}, , it was immediately"
metal ionization in the convection zone and the mean molecular weight of the solar core.,metal ionization in the convection zone and the mean molecular weight of the solar core.
" In à comprehensive review, Basu&Antia(2008) found that increases in neon did not reduce the need for a high oxygen in the convection zone (c.f."," In a comprehensive review, \citet{BA08} found that increases in neon did not reduce the need for a high oxygen in the convection zone (c.f."
 their Figure 21)., their Figure 21).
" One would therefore need to invoke multiple errors in distinct interiors and atmospheric input physics, with the same sign and magnitude, to produce an acceptable high Ne - low O mixture."," One would therefore need to invoke multiple errors in distinct interiors and atmospheric input physics, with the same sign and magnitude, to produce an acceptable high Ne - low O mixture."
" Our best interiors estimate for composition is A(O)=8.86+0.04, A(Ve)=8.15£0.17 and A(Fe)=7.50£0.05."," Our best interiors estimate for composition is $A(O)=8.86\pm0.04 $ , $A(Ne)=8.15\pm0.17 $ and $A(Fe)=7.50\pm0.05$."
 Loddersetal.(2009) combined recent atmospheres measurements for estimates of A(O)=8.732:0.07. A(Ne)=8.05£0.10 and A(Fe)=7.45+£0.08: these two scales are consistent within the errors.," \citet{Lodders2009} combined recent atmospheres measurements for estimates of $A(O)=8.73\pm0.07 $, $A(Ne)=8.05\pm0.10 $ and $A(Fe)=7.45\pm0.08$; these two scales are consistent within the errors."
 A weighted mean of the two leads to A(O)=8.83+0.04. A(Ne)=8.08£0.09 and A(Fe)=7.49--0.05 which we contend is the most precise current estimate for abundances.," A weighted mean of the two leads to $A(O)=8.83\pm0.04 $, $A(Ne)=8.08\pm0.09 $ and $A(Fe)=7.49\pm0.05$ which we contend is the most precise current estimate for abundances."
" In DPO6 we derived relationships quantifying combinations of C. Ν. O, and Ne with the same Rey and Y,,,."," In DP06 we derived relationships quantifying combinations of C, N, O, and Ne with the same $R_{CZ}$ and $Y_{surf}$."
" In the case of Ne, nonlinear effects can be induced when the change in the Ne/O ratio is too large and a better fitting relationship 1s A(Q)2—71.52x(Ne/OY49.03(Ne/Oy—3.10(Ne/O)—0.32O)-8.98 Our abundance scale is consistent with some photospheric abundance measurements (Caffauetal.2010;Pinsonneault&Delahaye2009) but not the lower AGSSO9 values."," In the case of Ne, nonlinear effects can be induced when the change in the Ne/O ratio is too large and a better fitting relationship is $A(O)= -7.52\times (Ne/O)^4 +9.03\times (Ne/O)^3 - 3.10\times (Ne/O)^2 -0.32\times (Ne/O) + 8.98$ Our abundance scale is consistent with some photospheric abundance measurements \citep{caffau10,PD09} but not the lower AGSS09 values."
" This is not a conflict between modern and primitive atmospheres treatments, or between one- and three-dimensional studies, but rather reflects differences between competing atmospheres models and judgment calls on the choice of indicators, continuum levels, and the proper treatment of blending features."," This is not a conflict between modern and primitive atmospheres treatments, or between one- and three-dimensional studies, but rather reflects differences between competing atmospheres models and judgment calls on the choice of indicators, continuum levels, and the proper treatment of blending features."
 We are therefore hopeful that helioseismology may provide an absolute abundance standard which can be used to discriminate between competing models and which can be used to calibrate the appropriate composition diagnostics for the next generation of stellar models., We are therefore hopeful that helioseismology may provide an absolute abundance standard which can be used to discriminate between competing models and which can be used to calibrate the appropriate composition diagnostics for the next generation of stellar models.
 FD would like to thank C. Stehlé for financial support and A. Formicola for information onupdated nuclear reaction cross-sections., FD would like to thank C. Stehlé for financial support and A. Formicola for information onupdated nuclear reaction cross-sections.
 MP would like to acknowledge support from DOE, MP would like to acknowledge support from DOE
radio eniission.,radio emission.
 Because of the claims in the literature that SMCis are radio bright. we do not favor these models.," Because of the claims in the literature that SMGs are radio bright, we do not favor these models."
" The sugeested relative radio brightuess of high-: SAIGs therefore provides some evidence that. in fact. Box rather than Boxpl?95, "," The suggested relative radio brightness of $z$ SMGs therefore provides some evidence that, in fact, $B \propto \Sigma_g^{0.7 - 0.8}$ rather than $B \propto \rho^{0.5 - 0.6}$."
Towever. the matter of whether high-: SALGs are in fact radio bright is not vet settled.," However, the matter of whether $z$ SMGs are in fact radio bright is not yet settled."
 Although LTQ couclided that D iust increase dramatically from normal galaxies to dense starbursts 2)). they were unable to distinguish between these two possibilities with the :z0 FRC aloue.," Although LTQ concluded that $B$ must increase dramatically from normal galaxies to dense starbursts \ref{sec:Theory}) ), they were unable to distinguish between these two possibilities with the $z \approx 0$ FRC alone."
 For this reason. higeh-: starbursts and their qualitatively ciffercut morphologies compared to those at z&0 can distiuguisli theories of the FRC.," For this reason, $z$ starbursts and their qualitatively different morphologies compared to those at $z \approx 0$ can distinguish theories of the FRC."
 A prediction of all of our variants is that puffv starbursts like submillimeter galaxies should have steep nou-thermal radio spectra. with azOA1 (Seo Table À2))," A prediction of all of our variants is that puffy starbursts like submillimeter galaxies should have steep non-thermal radio spectra, with $\alpha \approx 0.8 - 1.0$ (see Table \ref{table:Models}) )."
 The steep spectra are caused by strong synchrotron cooling in the BoxLVFYS case aud the relatively stronger IC cooling off starlight in the Dxpe?0 case., The steep spectra are caused by strong synchrotron cooling in the $B \propto \Sigma_g^{0.7-0.8}$ case and the relatively stronger IC cooling off starlight in the $B \propto \rho^{0.5 - 0.6}$ case.
" Tu. general. putty starbursts should lave roughly the same a as normal salaxies in the local universe, which tends to be somewhat higher (aστ0.7. 1.0) than in compact starbursts (aτς 0.7)."," In general, puffy starbursts should have roughly the same $\alpha$ as normal galaxies in the local universe, which tends to be somewhat higher $\alpha \approx 0.7 - 1.0$ ) than in compact starbursts $\alpha \la 0.7$ )."
 The slope should hold even out to extremely ligh X. as loug as starbursts are putty.," The slope should hold even out to extremely high $\Sigma_g$, as long as starbursts are puffy."
 In coutrast. we fud that a=0.5 In compact starbursts because of efficieut ionization aud emisstraliluug losses. which flatten the equilibiiunii CR spectruni because of their euergev dependence.," In contrast, we find that $\alpha \approx 0.5$ in compact starbursts, because of efficient ionization and bremsstrahlung losses, which flatten the equilibrium CR spectrum because of their energy dependence."
 As we rote in LTQ. our predicted spectral iudex for normal ealaxies is somewhat too high. and this difference in a nay carry over to the puffv starbursts.," As we note in LTQ, our predicted spectral index for normal galaxies is somewhat too high, and this difference in $\alpha$ may carry over to the puffy starbursts."
 Towever. the sienificaut difference iu a between compact and putty starbursts should remain as a general prediction of our uodel: compact starbursts should have flatter spectra han puffy starbursts.," However, the significant difference in $\alpha$ between compact and puffy starbursts should remain as a general prediction of our model: compact starbursts should have flatter spectra than puffy starbursts."
 The high spectral slopes can be observed either with direct ieasurements of iultifrequency data of individual subimüilluneter galaxies. or with single frequency observations at a variety of redshifts.," The high spectral slopes can be observed either with direct measurements of multifrequency data of individual submillimeter galaxies, or with single frequency observations at a variety of redshifts."
" There are relatively few measurements of à for submillimeter ealaxies specifically: faint radio sources have a220.5.0.7 (IIuvuh.etal.2007:Dondict2007).. though. that sample iucludes both compact starbursts and Αν,"," There are relatively few measurements of $\alpha$ for submillimeter galaxies specifically; faint radio sources have $\alpha \approx 0.5 - 0.7$ \citep{Huynh07,Bondi07}, though that sample includes both compact starbursts and AGNs."
 Sajinaunmaetet;al.(2008)SOS do find that; ozLi1CGIlzNile~0.?ge for SMCs. comparable to our predictions.," \citet{Sajina08} do find that $\alpha_{610~\MHz}^{1.4~\GHz} \approx 0.8$ for SMGs, comparable to our predictions."
 They also find that subiuüllineter ealaxies have a radio-excess. m agereenmient with Figure 1..," They also find that submillimeter galaxies have a radio-excess, in agreement with Figure \ref{fig:LFIRRadioRest}."
 More τοσα]. Ibaroetal.(2010) found oean averageο aLIGIEZMIIZLzcοτεo6. which: is). somewji flatter than610 our imo0dels.," More recently, \citet{Ibar10} found an average $\alpha_{610~\MHz}^{1.4~\GHz} \approx 0.75 \pm 0.06$, which is somewhat flatter than our models."
 These spectral slopes arc not different from normal star-forming ealaxies. but are noticeably steeper than local ULIRGs (Clemensetal.2Kos)...," These spectral slopes are not different from normal star-forming galaxies, but are noticeably steeper than local ULIRGs \citep{Clemens08}."
 ILowever. we do not account for frec-free absorpti3. Which probably flattens the spectra of local ULIRGs ike Arp 220 at low frequency (Condoetal. 1991).. aud is not well uuderstood in SAICs.," However, we do not account for free-free absorption, which probably flattens the spectra of local ULIRGs like Arp 220 at low frequency \citep{Condon91}, and is not well understood in SMGs."
" Since puffv sfarbursts lave steeper spectra than compact starburs swe expect their inferred Loup/L7, will increase witi redshift: if the true radio spectral slopes of SAICs are greater than the assumed à by Ao. they will appear to become radio dinuuer by a factor (11:)59"", or up to ~IU at 2=2."," Since puffy starbursts have steeper spectra than compact starbursts, we expect their inferred $L_{\rm TIR}^{\prime}/L_{\rm radio}^{\prime}$ will increase with redshift: if the true radio spectral slopes of SMGs are greater than the assumed $\alpha$ by $\Delta\alpha$, they will appear to become radio dimmer by a factor $(1 + z)^{\Delta \alpha}$, or up to $\sim 40\%$ at $z = 2$."
 Iu Figure 3.. we show the expected radio svuchrotrou spectra of starburst galaxies. without correcting for thermal absorption or thermal cmission.," In Figure \ref{fig:StarburstSpectra}, , we show the expected radio synchrotron spectra of starburst galaxies, without correcting for thermal absorption or thermal emission."
 At a rest-frame frequency of 1 CIIz. putty starbursts (dashed) have steeper radio spectra than compact starbursts (solid).," At a rest-frame frequency of 1 GHz, puffy starbursts (dashed) have steeper radio spectra than compact starbursts (solid)."
 Note that at high frequencies 0”210 GIIz). the ratio of the radio huuinosities per unit star formation of the conrpact and putty starbursts asvinptotes to a value sot by the ratio of Up aud Uy iu these starbursts.," Note that at high frequencies $\nu^{\prime} \ga 10\ \GHz$ ), the ratio of the radio luminosities per unit star formation of the compact and puffy starbursts asymptotes to a value set by the ratio of $U_B$ and $U_{\rm ph}$ in these starbursts."
 At these lieh frequencies. ouly svuchrotron aud IC cooling are effective. aud IC cooling would be the same for puffy aud conipact starbursts because of the Schinidt Law 2)).," At these high frequencies, only synchrotron and IC cooling are effective, and IC cooling would be the same for puffy and compact starbursts because of the Schmidt Law \ref{sec:Theory}) )."
 For BxSy. Up is the same for putty aud. compact starbursts. but for Boxp. puffy starbursts have iumch snaller τω.," For $B \propto \Sigma_g^a$, $U_B$ is the same for puffy and compact starbursts, but for $B \propto \rho^a$, puffy starbursts have much smaller $U_B$."
 Thus. measuremcuts of the svuchrotron radio cuuission of SAIGs at hieh ν΄ could determine the magnetic feld streneth of SAIGs and determine which scenario applies.," Thus, measurements of the synchrotron radio emission of SMGs at high $\nu^{\prime}$ could determine the magnetic field strength of SMGs and determine which scenario applies."
 Of course. there are unlikely to be two perfectly distinct. populations of compact starbursts and putty starhbursts.," Of course, there are unlikely to be two perfectly distinct populations of compact starbursts and puffy starbursts."
 DIusteacd. there may be a continu variatio1 iji scale. heights from teis fo thousands of parsecs.," Instead, there may be a continuum variation in scale heights from tens to thousands of parsecs."
" We would then expectto see a larger scatter. both i Spf, aud o. ina ""ll sample of both the mos conrpact and the most pifiv starbursts."," We would then expectto see a larger scatter, both in $L_{\rm TIR}^{\prime}/L_{\rm radio}^{\prime}$ and $\alpha$, in a full sample of both the most compact and the most puffy starbursts."
 πηρανctal. fud that ποτιueter galaxies do have a largeLad scatter d ὅτε than other galaxies., \citet{Murphy09} find that submillimeter galaxies do have a larger scatter in $q_{\rm TIR}^{\prime}$ than other galaxies.
 However. Tharο fud a relatively simall scatterof ~0.3 1- SMG radio spectral index.," However, \citet{Ibar10} find a relatively small scatterof $\sim 0.3$ in SMG radio spectral index."
" huportautlv. for lareer Ph both Linn/radio aud a’ asviuptote as CR electron ai positron losses are entirely determined by svuchrotre- and IC: fx1.kpe starbursts are already: near this Πατ,"," Importantly, for larger $h$ , both $L_{\rm TIR}^{\prime}/L_{\rm radio}^{\prime}$ and $\alpha^{\prime}$ asymptote as CR electron and positron losses are entirely determined by synchrotron and IC; $h \approx 1~\kpc$ starbursts are already near this limit."
thick. ecometrically thin state and. viclels both an observable radio jet anc narrow-line emission.,"thick, geometrically thin state and yields both an observable radio jet and narrow-line emission."
 ‘Tables 2. and 3 suggest this is not the case for the later (control early types) stages of evolution., Tables \ref{tab:fracOpticalAGN_dustlanes} and \ref{tab:fracAllAGN} suggest this is not the case for the later (control early types) stages of evolution.
 This is likely to be because the merger-driven GN. activity. has mostly expelled the available gas. and the galaxy migrated to a quiescent state.," This is likely to be because the merger-driven AGN activity has mostly expelled the available gas, and the galaxy migrated to a quiescent state."
 Phe mean starburst age for the control sample is TOO Myr. which is much larger than both the transition time of ~200 Myr between the starburst and transition (i.c. starburst|ACGN) phase argued by and the maximum AGN lifetime of up to a few hundred. Myr2008).," The mean starburst age for the control sample is $700$ Myr, which is much larger than both the transition time of $\sim 200$ Myr between the starburst and transition (i.e. starburst+AGN) phase argued by and the maximum AGN lifetime of up to a few hundred Myr."
. Therefore. most of these objects are not. expected to host. AGN that would have a common origin with the starburst.," Therefore, most of these objects are not expected to host AGN that would have a common origin with the starburst."
 Conversely. with their mean starburst ages of 300 Myr. dust. lane early. types fall in the “sweet spot of being old enough to host an AGN. but not so old that AGN activity has been terminated.," Conversely, with their mean starburst ages of $\sim 300$ Myr, dust lane early types fall in the `sweet spot' of being old enough to host an AGN, but not so old that AGN activity has been terminated."
 We test this paradigm by deriving radio AGN ages for those galaxies that were classified as having excess 1.4-Cillz luminosity., We test this paradigm by deriving radio AGN ages for those galaxies that were classified as having excess 1.4-GHz luminosity.
 This was done by generating a library of tracks in racio buminositv- linear size space., This was done by generating a library of tracks in radio luminosity- linear size space.
 As discussed at length in(2008).. given a clensity profile for the atmosphere into which the radio source is expanding. anc aspect ratio of the source. its Luminosity and size allow je power and age of the radio source to be determined uniquely.," As discussed at length in, given a density profile for the atmosphere into which the radio source is expanding, and aspect ratio of the source, its luminosity and size allow jet power and age of the radio source to be determined uniquely."
 We use observations of local early-type galaxies for the X-ray gas density. profile. and adopt an axia ratio of Ap=2.," We use observations of local early-type galaxies for the X-ray gas density profile, and adopt an axial ratio of $R_{\rm T}=2$."
 X number of caveats are associated: with this analysis., A number of caveats are associated with this analysis.
 Our assumption for the gas density profile may not be applicable for dust. lane earlv-tvpe galaxies. most of which show disturbed morphologies (Paper D).," Our assumption for the gas density profile may not be applicable for dust lane early-type galaxies, most of which show disturbed morphologies (Paper I)."
 Furthermore. hotspot contribution to radio Luminosity can be comparable with that of the lobes at 1.4 1.," Furthermore, hotspot contribution to radio luminosity can be comparable with that of the lobes at 1.4 GHz."
 However. estimate that the derived time-scales are accurate to better than a factor of two.," However, estimate that the derived time-scales are accurate to better than a factor of two."
 Importantly. the radio source models of and are only applicable for edge-brightened: Fanaroll- tvpe EL (ETC LD] sources.," Importantly, the radio source models of and are only applicable for edge-brightened [Fanaroff-Riley type II (FR II)] sources."
 On the other hand. most radio AGN in the local universe are core-dominated ETC 1 objects.," On the other hand, most radio AGN in the local universe are core-dominated FR I objects."
 Llowever. most of the resolved objects in our sample are FR IL objects. for which the modelling is applicable.," However, most of the resolved objects in our sample are FR II objects, for which the modelling is applicable."
 Only upper limits on ages can be placed for unresolved radio sources., Only upper limits on ages can be placed for unresolved radio sources.
 ‘Table 4 eives the size distribution of radio AGN [or the dust lanes ancl the matched. control sample., Table \ref{tab:radioSizesAges} gives the size distribution of radio AGN for the dust lanes and the matched control sample.
 The distributions are statistically indistinguishable at the LO per cent level., The distributions are statistically indistinguishable at the 10 per cent level.
 However. the number of dust lane racio AGN in our sample (46) is small. and most of these are unresolved.," However, the number of dust lane radio AGN in our sample (46) is small, and most of these are unresolved."
 Lighresolution imaging will be required to address this issue properly., High resolution imaging will be required to address this issue properly.
 lig., Fig.
 S compares the derived radio source ages with the ages of stellar populations for dust lane early-type galaxies., \ref{fig:SFtimescales} compares the derived radio source ages with the ages of stellar populations for dust lane early-type galaxies.
 Sevlerts ancl LINERS have similar age distributions. with median values around 300 Myr.," Seyferts and LINERs have similar age distributions, with median values around 300 Myr."
 There is no cillerence between radio-Ioud and. radio-quiet emission-line ACN., There is no difference between radio-loud and radio-quiet emission-line AGN.
 We interpret this as evidence for the simultaneous trigeering of the radio jet and line emission., We interpret this as evidence for the simultaneous triggering of the radio jet and line emission.
 Sevfert-lHike line emission will come from the raciativelv ellicient AGN clise. while LINER-like emission can either come from the disce or be distributed throughout the radio cocoon2011).," Seyfert-like line emission will come from the radiatively efficient AGN disc, while LINER-like emission can either come from the disc or be distributed throughout the radio cocoon."
. This picture is further supported. by the 100150 Mr olfset between radio and starburst ages in Fig. N((, This picture is further supported by the $100-150$ Myr offset between radio and starburst ages in Fig. \ref{fig:SFtimescales}( (
b) being consistent with the age dillerence. between starburst) and ransition objects in Fig S((a).,b) being consistent with the age difference between starburst and transition objects in Fig \ref{fig:SFtimescales}( (a).
 Our interpretation of the origin and evolution of dust ane earlv-tv galaxies is then as follows., Our interpretation of the origin and evolution of dust lane early-type galaxies is then as follows.
 (X. quiescent earlv-tvpe galaxy.pe. undergoes a gaserich. minor merger., A quiescent early-type galaxy undergoes a gas-rich minor merger.
 The resulting gas inllow triggers a starburst., The resulting gas inflow triggers a starburst.
 Some time later (tvpically <200 Myr) the AGN switches on., Some time later (typically $<200$ Myr) the AGN switches on.
 Cireumnuclear starbursts could be the reason Lor this delav. with powerful winds driven by OB stars and supernovae suppressing ACN uclling temporarily2007).," Circumnuclear starbursts could be the reason for this delay, with powerful winds driven by OB stars and supernovae suppressing AGN fuelling temporarily."
.. Alternatively. mass loss from newly formed stars could. help [απο AGN activity. providing a possible mechanism of [Tunneling gas owards the central engine.," Alternatively, mass loss from newly formed stars could help fuel AGN activity, providing a possible mechanism of funneling gas towards the central engine."
 Once the ACGN switches on. the accretion rate onto the central black hole is high. enough (tvpically at least a few per cent of the Eddington value) for he accretion cise to be in a classical. raclatively ellicient hin disc state. and the AGN is observed in both emission ines and at radio wavelengths.," Once the AGN switches on, the accretion rate onto the central black hole is high enough (typically at least a few per cent of the Eddington value) for the accretion disc to be in a classical, radiatively efficient thin disc state, and the AGN is observed in both emission lines and at radio wavelengths."
 This phase is accompanied w the presence of significant amounts of dust. ancl the earlv-tvpe galaxy. is observed as having a dust lane feature.," This phase is accompanied by the presence of significant amounts of dust, and the early-type galaxy is observed as having a dust lane feature."
" Eventually, AGN feedback heats up and/or expels the gas rom the galaxy2009)."," Eventually, AGN feedback heats up and/or expels the gas from the galaxy."
 oth the star formation and AGN activity are truncated. with the galaxy returning to à quiescent state.," Both the star formation and AGN activity are truncated, with the galaxy returning to a quiescent state."
 In Paper Lowe showed that the mass of dust. in the observed. features is too great to come from stellar mass oss. sugeesting stronely an external origin for this dust.," In Paper I we showed that the mass of dust in the observed features is too great to come from stellar mass loss, suggesting strongly an external origin for this dust."
 Since dust ane gas are typically coupled. one might expect dust lane earlv-type galaxies to be gas rich. consistent with he gas-rich merger origin of these objects.," Since dust and gas are typically coupled, one might expect dust lane early-type galaxies to be gas rich, consistent with the gas-rich merger origin of these objects."
" As shown in ""aper L dust lane carly types do show enhanced. levels of star formation compared to a control saniple of all early-vpe galaxies."," As shown in Paper I, dust lane early types do show enhanced levels of star formation compared to a control sample of all early-type galaxies."
 However. Fig.," However, Fig."
 7 shows that galaxies with dust eatures also have vounger starburst ages. and it is therefore possible that the higher star formation rates in these objects," \ref{fig:t2control} shows that galaxies with dust features also have younger starburst ages, and it is therefore possible that the higher star formation rates in these objects"
momentum of the accretor. i.e. we mimick in this way helium nova explosions) yielded virtually identical abundances as function of orbital period.,"momentum of the accretor, i.e. we mimick in this way helium nova explosions) yielded virtually identical abundances as function of orbital period."
 We also ran two sets of calculations. completely conservative and completely non-conservative. for the case of a neutron star accretor with initial helium donor masses of 0.35. 0.65 and 1M. and again found no significant difference in the abundances as function of orbital period.," We also ran two sets of calculations, completely conservative and completely non-conservative, for the case of a neutron star accretor with initial helium donor masses of 0.35, 0.65 and $1\,\msun$ and again found no significant difference in the abundances as function of orbital period."
 Because the onset of mass transfer from the helium star the helium burning3imainBodyCitationEnd1576]skh86. the chemical composition of the core of the donor depends sensitively on the at which the helium star fills its Roche lobe.," Because the onset of mass transfer from the helium star the helium burning, the chemical composition of the core of the donor depends sensitively on the at which the helium star fills its Roche lobe."
 We use evolutionary calculations for binaries that start mass transfer almost immediately after the common-envelope phase in which the helium star is formed. as well as systems that fill their Roche lobe just before core helium exhaustion.," We use evolutionary calculations for binaries that start mass transfer almost immediately after the common-envelope phase in which the helium star is formed, as well as systems that fill their Roche lobe just before core helium exhaustion."
 In this way the complete range of expected abundances can be probed. although we would like to note that the extremes of this range will likely be rare in practice. because the periods need to be fine tuned.," In this way the complete range of expected abundances can be probed, although we would like to note that the extremes of this range will likely be rare in practice, because the periods need to be fine tuned."
 We discuss this in more detail in Section 3.., We discuss this in more detail in Section \ref{population}.
 In Fig., In Fig.
 + we show a representative set of evolutionary sequences for binaries with helium donors and white dwarf/neutron star accretors., \ref{fig:heabund} we show a representative set of evolutionary sequences for binaries with helium donors and white dwarf/neutron star accretors.
 The full set of sequences is published on-line., The full set of sequences is published on-line.
 For each sequence we plot the abundances of the transferred material. the mass transfer rate. donor mass and orbital period as function of the time since the onset of RLOF.," For each sequence we plot the abundances of the transferred material, the mass transfer rate, donor mass and orbital period as function of the time since the onset of RLOF."
 This allows full assessment of the evolution. both in the initial phase when the binary evolves to shorter periods at almost constant mass transfer rate of a few times LO7M.vr.J| as. well as later. when the system has passed its period minimum.," This allows full assessment of the evolution, both in the initial phase when the binary evolves to shorter periods at almost constant mass transfer rate of a few times $10^{-8}\,\myr$, as well as later, when the system has passed its period minimum."
 Before the period minimum stars lose matter that was outside the convective core in the helium-burning stage so the material is heltum rich. with CNO abundances corresponding to the CNO cycle equilibrium for relatively massive stars that are typical progenitors of heliumstars!.," Before the period minimum stars lose matter that was outside the convective core in the helium-burning stage so the material is helium rich, with CNO abundances corresponding to the CNO cycle equilibrium for relatively massive stars that are typical progenitors of helium."
. Shortly after the »eriod. minimum tor for the most evolved donors already before oriod. minimum) helium burning products. most notably C and ater O. come to the surface.," Shortly after the period minimum (or for the most evolved donors already before period minimum) helium burning products, most notably C and later O, come to the surface."
 Depending on the initial period of he binary the enrichment by C and O can be very mild. or C and O can dominate even He.," Depending on the initial period of the binary the enrichment by C and O can be very mild, or C and O can dominate even He."
 In the top row of Fig., In the top row of Fig.
 4. we plot sequences for an initially VAAL. helium star donor transferring material to an initially 6M. white dwarf aecretor.," \ref{fig:heabund} we plot sequences for an initially $0.4\,\msun$ helium star donor transferring material to an initially $0.6\,\msun$ white dwarf accretor."
 The leftmost plot is for the case of a post-common-envelope period of 20 min. in which RLOF starts almost immediately after formation of the helium star and very ittle helium burning occurs. so He. N. O and Νο abundances are virtually unchanged.," The leftmost plot is for the case of a post-common-envelope period of 20 min, in which RLOF starts almost immediately after formation of the helium star and very little helium burning occurs, so He, N, O and $^{22}$ Ne abundances are virtually unchanged."
 However. after the period minimum the carbon abundance increases substantially.," However, after the period minimum the carbon abundance increases substantially."
 For an intermediate post-common-envelope period of mmin. there is a dramatic change of abundances around period minimum because the layers which were in the core and experienced some He-burning are exposed and for most of the AM CVn evolution C and © (and even 77 Ne) dominate over N. The He abundance is noticeably reduced.," For an intermediate post-common-envelope period of min, there is a dramatic change of abundances around period minimum because the layers which were in the core and experienced some He-burning are exposed and for most of the AM CVn evolution C and O (and even $^{22}$ Ne) dominate over N. The He abundance is noticeably reduced."
 Finally for the most evolved donor. with initial period of 130 min (helium abundance in the core at RLOF Y;z 0.07). the changes are even more dramatic and O and C dominate even over He.," Finally for the most evolved donor, with initial period of 130 min (helium abundance in the core at RLOF $Y_{\rm c} \approx 0.07$ ), the changes are even more dramatic and O and C dominate even over He."
" Nitrogen becomes extinguished even before /2,,.4,4."," Nitrogen becomes extinguished even before $ P_{\rm orb, min} $."
 Note that in the most extreme cases He-burning continues for some time after RLOF but He is still not completely burnt so significant amounts of He should still be detectable. contrary to hybrid white dwarf donors.," Note that in the most extreme cases He-burning continues for some time after RLOF but He is still not completely burnt so significant amounts of He should still be detectable, contrary to hybrid white dwarf donors."
 In the bottom row of Fig., In the bottom row of Fig.
 + a number of evolutionary sequences for systems with neutron star accretors are shown., \ref{fig:heabund} a number of evolutionary sequences for systems with neutron star accretors are shown.
 The leftmost panel is again for fairly short initial period of mmin and shows quite a change in C but not much in the other elements., The leftmost panel is again for fairly short initial period of min and shows quite a change in C but not much in the other elements.
 The middle plot is for the sequence with the longest initial period for which RLOF starts when the star almost totally burnt helium in its core (3.zz0.06) and looks very similar to the most evolved donor shown for the sequence with a white dwarf accretor., The middle plot is for the sequence with the longest initial period for which RLOF starts when the star almost totally burnt helium in its core $ Y_{\rm c} \approx 0.06 $ ) and looks very similar to the most evolved donor shown for the sequence with a white dwarf accretor.
 The bottom right plot is for a He star plus neutron star system with initial donor mass of 0.65M. and initial orbital period close to the minimum possible for such a system. mmin It shows that at a given orbital period there may be a scatter of a factor of several in abundance ratios. depending on the initial mass of the donor.," The bottom right plot is for a He star plus neutron star system with initial donor mass of $0.65\,\msun$ and initial orbital period close to the minimum possible for such a system, min It shows that at a given orbital period there may be a scatter of a factor of several in abundance ratios, depending on the initial mass of the donor."
 A peculiar evolutionary path is followed by initially relatively massive helium stars (more than about 0.65ΔΕ. ) with neutron star companions.," A peculiar evolutionary path is followed by initially relatively massive helium stars (more than about $0.65\,\msun$ ) with neutron star companions."
 These overfill their Roche lobes after burning of a substantial fraction of helium in the core and continue He-burning during the semidetached stage of evolution., These overfill their Roche lobes after burning of a substantial fraction of helium in the core and continue He-burning during the semidetached stage of evolution.
" For instance. the system with Mg.=O.SOAL. and £4=TOmmin starts mass loss when Y,=0.643 and proceeds along a conventional evolutionary track."," For instance, the system with $M_{\rm He}=0.80\,\msun$ and $
 P_0=70$ min starts mass loss when $Y_c \approx 0.643$ and proceeds along a conventional evolutionary track."
" In a slightly wider initially system with /3, 75mmin. RLOF occurs when 3;=0.56 (see Fig. 59)."," In a slightly wider initially system with $ P_0=75$ min, RLOF occurs when $Y_c \approx 0.56$ (see Fig. \ref{fig:evolHe}) )."
" In this system the donor detaches from its Roche lobe when its mass ws decreased to 0.52M. and Y,z0.006."," In this system the donor detaches from its Roche lobe when its mass has decreased to $0.52\,\msun$ and $Y_c \approx 0.006$."
 The orbit continues o shrink and mass exehange resumes in the helium-shell burning stage., The orbit continues to shrink and mass exchange resumes in the helium-shell burning stage.
 However. when the donor mass is more than about 0.45M.. mass loss cannot be stabilised by mass and angular-momentum oss from the system and the ensuing mass loss proceeds on a dynamical time scale (see Fig. 59).," However, when the donor mass is more than about $0.45\,\msun$, mass loss cannot be stabilised by mass and angular-momentum loss from the system and the ensuing mass loss proceeds on a dynamical time scale (see Fig. \ref{fig:evolHe}) )."
 Thus. such a system does not contribute to the heltum star channel or UCXBs.," Thus, such a system does not contribute to the helium star channel for UCXBs."
 Evolutionary sequences for 1M... donor stars follow a similar path irrespective of the amount of He burnt prior to RLOF.," Evolutionary sequences for $1\,\msun$ donor stars follow a similar path irrespective of the amount of He burnt prior to RLOF."
 The details of this type of evolution will be discussed in a forthcoming paper (Yungelson et al., The details of this type of evolution will be discussed in a forthcoming paper (Yungelson et al.
 in Thus. there are two factors limiting the helium star channel for the formation of ultra-compact binaries. the maximum post-common-envelope period for which mass transfer still starts during core helium burning and a limiting mass above which the system detaches as described above (see also Fig. 105.," in Thus, there are two factors limiting the helium star channel for the formation of ultra-compact binaries, the maximum post-common-envelope period for which mass transfer still starts during core helium burning and a limiting mass above which the system detaches as described above (see also Fig. \ref{fig:hechannel_MP}) )."
 The formation of ultra-compact binaries from main-sequence donors requires two conditions., The formation of ultra-compact binaries from main-sequence donors requires two conditions.
 First. the initial periods are such that the progenitors fill their Roche lobe close to the end of the main sequence.," First, the initial periods are such that the progenitors fill their Roche lobe close to the end of the main sequence."
 Second. angular momentum loss drives the components together at a sufficient rate that the ultra-short periods can. be reached within the Hubble time (222?)..," Second, angular momentum loss drives the components together at a sufficient rate that the ultra-short periods can be reached within the Hubble time ."
 Radiative transfer is an important phenomena in star formation., Radiative transfer is an important phenomena in star formation.
 Radiation sets the temperature of the gas during the collapse of a molecular cloud core., Radiation sets the temperature of the gas during the collapse of a molecular cloud core.
 This both influences the degree of fragmentation of the eloud. and sets the minimum mass of brown dwarfs (theopacitylimitforfragmentation:2)..," This both influences the degree of fragmentation of the cloud, and sets the minimum mass of brown dwarfs \citep[the opacity limit for fragmentation;][]{LL1976}."
 Once protostars have formed in a cloud. radiative and mechanical feedback from them can affect subsequent star formation.," Once protostars have formed in a cloud, radiative and mechanical feedback from them can affect subsequent star formation."
 Such feedback mechanisms include protostellar jets and outflows from low-mass stars. and ionisation from massive stars which creates HII regions and destroys a cloud.," Such feedback mechanisms include protostellar jets and outflows from low-mass stars, and ionisation from massive stars which creates HII regions and destroys a cloud."
 Computer simulations are vital in our efforts to understand the complex problem of star formation., Computer simulations are vital in our efforts to understand the complex problem of star formation.
 Many previous simulations have used the smoothed particle hydrodynamics method (SPH) (e.g.222222).," Many previous simulations have used the smoothed particle hydrodynamics method (SPH) \citep[e.g.][]{GM1981,PCDNDW1991MmSAI,BMBAB1991,NP1993,BBP1995,KBB1998}."
 Other methods used are typically based around grid-based codes (e.g.222222).," Other methods used are typically based around grid-based codes \citep[e.g.][]{L1969b,BB1979,BM1992,BB1993,TKMHHG1997,BFKM2000}."
 SPH is a Lagrangian method firs developed by ? and ? (see?.forareview)..., SPH is a Lagrangian method first developed by \citet{L1977} and \citet{GM1977} \citep[see][for a review]{M1992}.
 It approximates the fluid as a series of discrete fluid elements denoted by individual SPH particles and uses interpolation to obtain the fluic variables at any point in the simulation., It approximates the fluid as a series of discrete fluid elements denoted by individual SPH particles and uses interpolation to obtain the fluid variables at any point in the simulation.
 SPH is conceptually simple to understand. and can naturally adapt its resolution to the local density distribution. unlike grid-based codes which require complex adaptive-mesh refinement algorithms to perform the same task.," SPH is conceptually simple to understand, and can naturally adapt its resolution to the local density distribution, unlike grid-based codes which require complex adaptive-mesh refinement algorithms to perform the same task."
 This property makes it ideal for use in star formation. where densities may range over many orders of magnitude in a single simulation.," This property makes it ideal for use in star formation, where densities may range over many orders of magnitude in a single simulation."
 Despite these advantages. few attempts have been made to include radiative transfer into SPH (22227). and until recently SPH with radiative transfer has not been applied to star formation.," Despite these advantages, few attempts have been made to include radiative transfer into SPH \citep{L1977,B1985,B1986,OW2003,WB2004}, and until recently \citep*{BCV2004,BCV2005} SPH with radiative transfer has not been applied to star formation."
 Instead. many past simulations have simply used isothermal or barotropic equations of state to model the collapse of a molecular cloud.," Instead, many past simulations have simply used isothermal or barotropic equations of state to model the collapse of a molecular cloud."
 The former is only valid up to densities of ~10U & ? at which point the cloud traps radiation etficiently enough for the cloud to begin to heat up., The former is only valid up to densities of $\sim 10^{-13}$ g $^{-3}$ at which point the cloud traps radiation efficiently enough for the cloud to begin to heat up.
 The latter is usually based on the evolution of the temperature at the highest density. during the collapse of spherically symmetric clouds as calculated using radiative transfer (e.g.222)..," The latter is usually based on the evolution of the temperature at the highest density during the collapse of spherically symmetric clouds as calculated using radiative transfer \citep[e.g.][]{L1969b,WN1980,MI2000}."
 However. a barotropic equation of state can at best only hope to provide an adequate description of temperature at the density maximum: it is unlikely to give an accurate temperature distribution during a three-dimensional calculation with complex density and velocity structure.," However, a barotropic equation of state can at best only hope to provide an adequate description of temperature at the density maximum; it is unlikely to give an accurate temperature distribution during a three-dimensional calculation with complex density and velocity structure."
 Indeed. ? performed grid-based calculations of the collapse of a molecular cloud core both with a barotropic equation of state and with radiative transfer in the Eddington approximation and found they dittered somewhat.," Indeed, \citet{BFKM2000} performed grid-based calculations of the collapse of a molecular cloud core both with a barotropic equation of state and with radiative transfer in the Eddington approximation and found they differed somewhat."
 However. they did not examine in detail how the relation between temperature and density dittered from that of the barotropic equation of state spatially and temporally. or its dependence on initial conditions.," However, they did not examine in detail how the relation between temperature and density differed from that of the barotropic equation of state spatially and temporally, or its dependence on initial conditions."
 ? recently presented an implicit algorithm for calculating radiative transfer using the flux-limited diffusion approximation within the SPH formalism., \citet*{WBM2005} recently presented an implicit algorithm for calculating radiative transfer using the flux-limited diffusion approximation within the SPH formalism.
 This paper describes a three-dimensional implementation of this algorithm) and uses it to examine the thermodynamics during the collapse of molecular cloud cores., This paper describes a three-dimensional implementation of this algorithm and uses it to examine the thermodynamics during the collapse of molecular cloud cores.
 Section 2. describes the changes necessary to the radiative transfer algorithm of 2? for use in three dimensions and the initial conditions for our star formation calculations.," Section \ref{sec:method}
 describes the changes necessary to the radiative transfer algorithm of \citet{WBM2005} for use in three dimensions and the initial conditions for our star formation calculations."
 Section 3 presents the results of simulations of the collapse of molecular cloud cores with different initial conditions and examines the evolution of their temperature structure., Section \ref{sec:results} presents the results of simulations of the collapse of molecular cloud cores with different initial conditions and examines the evolution of their temperature structure.
 Finally. section 5 summarises the main conclusions of this paper.," Finally, section \ref{sec:conclusions} summarises the main conclusions of this paper."
Each person has their own favorite list of future observational needs.,Each person has their own favorite list of future observational needs.
 Here is mine: ο We need a high rate (>100 GRBs 1) of bursts with good locations. in order to change the sociology of ground-based optical and radio observations.," Here is mine: $\bullet$ We need a high rate $> 100$ GRBs $^{-1}$ ) of bursts with good locations, in order to change the sociology of ground-based optical and radio observations."
 This many good GRB positions to follow-up each year would make it possible to propose and carry out GRB afterglow monitoring programs at many medium-to-large aperture telescopes., This many good GRB positions to follow-up each year would make it possible to propose and carry out GRB afterglow monitoring programs at many medium-to-large aperture telescopes.
 e The diversity of GRBs. GRB afterglows. and host galaxies means that we need a large number (71000) of good GRB positions in order to be able to study the correlations between these properties.," $\bullet$ The diversity of GRBs, GRB afterglows, and host galaxies means that we need a large number $> 1000$ ) of good GRB positions in order to be able to study the correlations between these properties."
 This is important for determining whether or not there are distinct subclasses of bursts. and more than one burst mechanism.," This is important for determining whether or not there are distinct subclasses of bursts, and more than one burst mechanism."
 Any correlations found will also impose important constraints on burst mechanisms and models., Any correlations found will also impose important constraints on burst mechanisms and models.
 e We need many rapid (near real time) one areminute GRB positions in order to determine whether or not significant optical emission accompanies the bursts (Park 1999). and to make it possible to take spectra of the burst afterglows while the afterglows are still bright — and thereby obtain redshifts of the bursts themselves from absorption line systems. and if there are bursts at high redshifts. from the Ενα break.," $\bullet$ We need many rapid (near real time) one arcminute GRB positions in order to determine whether or not significant optical emission accompanies the bursts (Park 1999), and to make it possible to take spectra of the burst afterglows while the afterglows are still bright – and thereby obtain redshifts of the bursts themselves from absorption line systems, and if there are bursts at high redshifts, from the $\alpha$ break."
" e All of the GRBs that BeppoSAX has detected are ""long"" bursts.", $\bullet$ All of the GRBs that BeppoSAX has detected are “long” bursts.
" Currently we know nothing about the afterglow properties. the distance scale. and the hosts (1f any) of ""short"" bursts."," Currently we know nothing about the afterglow properties, the distance scale, and the hosts (if any) of “short” bursts."
 Therefore we need good/quick positions for short bursts. in order to determine these properties for short bursts in the same way that BeppoSAX has enabled us to determine these properties for long bursts.," Therefore we need good/quick positions for short bursts, in order to determine these properties for short bursts in the same way that BeppoSAX has enabled us to determine these properties for long bursts."
 e Currently. there is a largely unexplored gap in our knowledge of the X-ray and optical behavior of burst afterglows of ~10!Lo’ seconds immediately following the bursts. corresponding to the time needed to bring the BeppoSAX NFlIs to bear on a burst.," $\bullet$ Currently, there is a largely unexplored gap in our knowledge of the X-ray and optical behavior of burst afterglows of $\sim 10^4 - 10^5$ seconds immediately following the bursts, corresponding to the time needed to bring the BeppoSAX NFIs to bear on a burst."
 We need to fill in this unexplored gap. in order to see if bursts always. often. or rarely join smoothly onto their X-ray and optical afterelows. and to explore the geometry and kinematics of GRB afterglows (Sari 1999).," We need to fill in this unexplored gap, in order to see if bursts always, often, or rarely join smoothly onto their X-ray and optical afterglows, and to explore the geometry and kinematics of GRB afterglows (Sari 1999)."
 e We also need to search for variability in the X-ray and optical afterglows., $\bullet$ We also need to search for variability in the X-ray and optical afterglows.
" Observations of such variability would impose severe constraints on. models. including the widely-discussed relativistic fireball model of burst afterelows (see. e.g.. Fenimore 1999),"," Observations of such variability would impose severe constraints on models, including the widely-discussed relativistic fireball model of burst afterglows (see, e.g., Fenimore 1999)."
 The Rome Workshop provided a feast of observational and theoretical results. and the opportunity to discuss them.," The Rome Workshop provided a feast of observational and theoretical results, and the opportunity to discuss them."
 On behalf of all of the Workshop participants. | would like to thank Enrico Costa. Luigi Piro. Filippo Fontana. and everyone else who helped to organize this meeting for bringing all of us together and for providing us with such “fine dining.”," On behalf of all of the Workshop participants, I would like to thank Enrico Costa, Luigi Piro, Filippo Fontana, and everyone else who helped to organize this meeting for bringing all of us together and for providing us with such “fine dining.”"
Physical properties of cometary dust can be obtained from the solar radiation scattered by the cometary dust. which. in the process. gets. polarised.,"Physical properties of cometary dust can be obtained from the solar radiation scattered by the cometary dust which, in the process, gets polarised."
 The degree of polarisation and its direction mainlv depend. on the size clistribution. composition of the particles. phase angle and the wavelength of the incident. solar radiation.," The degree of polarisation and its direction mainly depend on the size distribution, composition of the particles, phase angle and the wavelength of the incident solar radiation."
 However. the real situation is not that straight forward.," However, the real situation is not that straight forward."
 In an attempt to study the detailed behaviour of polarisation with phase. Dolllusetal.CI988). synthesized the polarisation observation data on l1P/LHallev by various researchers ancl derived: curves of polarisation as a function of the phase angle.," In an attempt to study the detailed behaviour of polarisation with phase, \citet{Dollfus1988} synthesized the polarisation observation data on 1P/Halley by various researchers and derived curves of polarisation as a function of the phase angle."
 Phase curve below about 20° shows negative polarisation while it is positive at. higher. phase angles., Phase curve below about $ 20^{\circ}$ shows negative polarisation while it is positive at higher phase angles.
 They find slight modification in the polarisation phase curve as one moves away [rom the nucleus indicating the change in the nature of the dust. particles., They find slight modification in the polarisation phase curve as one moves away from the nucleus indicating the change in the nature of the dust particles.
 Also an anomalously high transient polarisation was noted between October 17 and 80. 1985. at phase angle 25°. attributed to sudden. release of large number of smaller particles.," Also an anomalously high transient polarisation was noted between October 17 and 30, 1985, at phase angle $25^{\circ}$, attributed to sudden release of large number of smaller particles."
 On the basis of this work on 1P/Hallev. Dollfus(1989). derived the physical properties of the dust. erains indicating the presence of large particles - aggregates comprising of submicron sized grains. very rough and dark.," On the basis of this work on 1P/Halley, \citet{Dollfus1989} derived the physical properties of the dust grains indicating the presence of large particles - aggregates comprising of submicron sized grains, very rough and dark."
 These rather large grains are mixed with the clouds of small particles ancl they are usually responsible for almost all the polarisation elfects in. visible light. except. during temporary specific dust. release events (as seen in case of 1P/Halley curing October 17-30. 1985) bv the nucleus. (Dollfus1989).," These rather large grains are mixed with the clouds of small particles and they are usually responsible for almost all the polarisation effects in visible light, except during temporary specific dust release events (as seen in case of 1P/Halley during October 17-30, 1985) by the nucleus \citep{Dollfus1989}."
. The complex. behaviour of the dust is also seen in comet ς105 (Llale-Bopp). especially the region. around. the nucleus shows complex structure (Llacdameik&Levasseur-Itegourd.2003).," The complex behaviour of the dust is also seen in comet C/1995 (Hale-Bopp), especially the region around the nucleus shows complex structure \citep{hadamcik2003}."
. Though the comets. in general. show similar polarisation behaviour with phase angle. they are divided into three classes based on the maximum in polarisation(Levasseur-Itegourd.1999).," Though the comets, in general, show similar polarisation behaviour with phase angle, they are divided into three classes based on the maximum in \citep{levasseur1999}."
. Varving polarisation observed in the coma or in the features (jets. shells. etc.)," Varying polarisation observed in the coma or in the features (jets, shells, etc.)"
 indicates a diversity of dust. particles., indicates a diversity of dust particles.
 Issues related to the dust. characteristics are adequately reviewed by Ixolokolovactal.(2005)., Issues related to the dust characteristics are adequately reviewed by \citet{kolokolova2005}.
. One of the main objective behind the study of comets is to uncerstand the origin of the solar svstem., One of the main objective behind the study of comets is to understand the origin of the solar system.
 Since comets spend. substantial part of their life away from the sun. their sub-surface material is considered. pristine.," Since comets spend substantial part of their life away from the sun, their sub-surface material is considered pristine."
" Space mission Deep Lupact was launehed on January 12. 2005 to study the composition of the interior of the comet 9P""Tempel 1 by colliding a part of the spacecraft with the comet CXLearnetal.2005b:Aleech 2000)."," Space mission Deep Impact was launched on January 12, 2005 to study the composition of the interior of the comet 9P/Tempel 1 by colliding a part of the spacecraft with the comet \citep{hearn2005b, meech2000}."
. At 5:52 UPC on July 4. 2005. the impactor of the Deep Impact. probe successfully collicled with the comet's nucleus. excavating huge amount," At 5:52 UTC on July 4, 2005, the impactor of the Deep Impact probe successfully collided with the comet's nucleus, excavating huge amount"
The orientation of the white loops as drawn is based upon comparison between the HMI magnetic field contours and the loops actually observed by for at least 2.5 days following the initial eruption.,The orientation of the white loops as drawn is based upon comparison between the HMI magnetic field contours and the loops actually observed by for at least 2.5 days following the initial eruption.
" We attribute the slight difference in loop appearance between those in panels c) and d) in this core region as being due to using a potential extrapolation near the limb, 2 days of additional rotation, and possible shear (a typical characteristic of pre-eruptive regions along polarity inversion lines (Su,Golub,&VanBallegooijen||2007))))."," We attribute the slight difference in loop appearance between those in panels c) and d) in this core region as being due to using a potential extrapolation near the limb, 2 days of additional rotation, and possible shear (a typical characteristic of pre-eruptive regions along polarity inversion lines \citep{su-golub-vanball_2007}) )."
 The green lines correspond to the development of the field in the green region of interest., The green lines correspond to the development of the field in the green region of interest.
 These loops appear kinked and stretched in the SECCHI images and are therefore highly non-potential., These loops appear kinked and stretched in the SECCHI images and are therefore highly non-potential.
" The western footpoints of these loops migrate after several hours possibly due to reconnection higher in the corona where the field lines have different footpoints, as is noted above in the PFSS extrapolations."," The western footpoints of these loops migrate after several hours possibly due to reconnection higher in the corona where the field lines have different footpoints, as is noted above in the PFSS extrapolations."
 The actual flaring site is located within the red region of interest where the first post-eruption arcade develops., The actual flaring site is located within the red region of interest where the first post-eruption arcade develops.
 This region also corresponds to the strongest polarity inversion., This region also corresponds to the strongest polarity inversion.
" Due to limb obscuration, the original erupting flux rope is not imaged by AIA nor is it readily apparent in the SECCHI observations. ("," Due to limb obscuration, the original erupting flux rope is not imaged by AIA nor is it readily apparent in the SECCHI observations. ("
A long filamentary structure is seen in the SECCHI images in the far south of the region but is not noticeably involved with the eruption.),A long filamentary structure is seen in the SECCHI images in the far south of the region but is not noticeably involved with the eruption.)
" Fortuitously, two eruptions originate from this region over the following several days wherein a flux rope is observed to erupt from the red core region. ("," Fortuitously, two eruptions originate from this region over the following several days wherein a flux rope is observed to erupt from the red core region. ("
Refer to the available online movie.),Refer to the available online movie.)
" By analyzing the subsequent eruptions and combining this information with the PFSS extrapolations, we infer that a relatively small flux rope erupts from within the core red region."," By analyzing the subsequent eruptions and combining this information with the PFSS extrapolations, we infer that a relatively small flux rope erupts from within the core red region."
" As it travels into the corona, it encounters the overlying field traversing the entire region in the general east-west direction (Figure 6| b))."," As it travels into the corona, it encounters the overlying field traversing the entire region in the general east-west direction (Figure \ref{mag_pfss} b))."
 This overlying field combines with the original erupting flux rope (possibly via reconnection or entanglement) to create a larger erupting flux rope which disrupts the field above the entire region., This overlying field combines with the original erupting flux rope (possibly via reconnection or entanglement) to create a larger erupting flux rope which disrupts the field above the entire region.
" This process may be likened to the flare breakout model (MacNeiceetal.| 20047; DeVore,&Klimchuk| [1999)) although not a necessary scenario due to the uncertainty in the original flux rope and overlying field structure; however, a truly 3-D model addressing the various overlying field line orientations, bend in the core polarity inversion line, and complex photospheric polarity organization is necessary in order understand the evolution of the flux rope."," This process may be likened to the flare breakout model \citeauthor{macneiceEA_2004} \citeyear{macneiceEA_2004}; \citeauthor{antiochos-devore-klimchuk_1999} \citeyear{antiochos-devore-klimchuk_1999}) ) although not a necessary scenario due to the uncertainty in the original flux rope and overlying field structure; however, a truly 3-D model addressing the various overlying field line orientations, bend in the core polarity inversion line, and complex photospheric polarity organization is necessary in order understand the evolution of the flux rope."
" 'The focus of this paper is to discuss the flows occurring in the wake of the flux rope; therefore, we will refer to the erupting flux rope for this region as the large, combined one traversing the region since it is required to create the post-eruption arcade in the green domain — where the flows of interest are occurring. ("," The focus of this paper is to discuss the flows occurring in the wake of the flux rope; therefore, we will refer to the erupting flux rope for this region as the large, combined one traversing the region since it is required to create the post-eruption arcade in the green domain – where the flows of interest are occurring. ("
The post-eruption arcade within the core red region could develop independently due to the original smaller flux rope eruption.),The post-eruption arcade within the core red region could develop independently due to the original smaller flux rope eruption.)
" Using these image sets and the magnetic field topological information as aguide combined with the basic reconnection scenario depicted in Figure D], a simplified description of the orientations of the most"," Using these image sets and the magnetic field topological information as aguide combined with the basic reconnection scenario depicted in Figure \ref{cartoon_all}, , a simplified description of the orientations of the most"
overestimated by up to 0.5 mag 7.,overestimated by up to 0.5 mag $^{-2}$.
 Late type spirals ane irregulars will only have a negligible effect., Late type spirals and irregulars will only have a negligible effect.
 However late types will be more elfected by inclination., However late types will be more effected by inclination.
 To model this. we have assumed. that a disk galaxy is optically thin and has no internal extinction.," To model this, we have assumed that a disk galaxy is optically thin and has no internal extinction."
 A galaxy of area Aj... was assumed to have an isophotal radius ri;;. given bv: The isophotal magnitude ancl radius were used to calculate the total magnitude and the cllective surface brightness yr as described in Cross et al.," A galaxy of area $A_{iso}$, was assumed to have an isophotal radius $r_{iso}$, given by: The isophotal magnitude and radius were used to calculate the total magnitude and the effective surface brightness $\mu_e$ as described in Cross et al."
 2001., 2001.
 However a disk galaxy. inclined at an angle. / with major axis e and minor axis > has an area where The surface brightness has increased at cach point by a [actor i zs inereasing the semi-major axis until pla) Mtm ," However a disk galaxy, inclined at an angle, $i$ with major axis $a$ and minor axis $b$ has an area where The surface brightness has increased at each point by a factor $\frac{1}{\cos(i)}$ , increasing the semi-major axis until $\mu(a)=\mu_{lim}$ ."
where à is the disk scale length., where $\alpha$ is the disk scale length.
 This increases the isophotal [lux of the galaxy. as well as the isophotal radius.," This increases the isophotal flux of the galaxy, as well as the isophotal radius."
 An exponential profile is fitted to these parameters as in Cross et al. (, An exponential profile is fitted to these parameters as in Cross et al. (
2001).,2001).
 The central surface brightness and total magnitude are calculated for this galaxy assuming that the ealaxy is [ace on., The central surface brightness and total magnitude are calculated for this galaxy assuming that the galaxy is face on.
 The error in the central surface surface brightness is the difference between the true central surface brightness and the caleulated: central surface. brightness., The error in the central surface surface brightness is the difference between the true central surface brightness and the calculated central surface brightness.
 The error in the elective surface brightness is exactly the same. as the dillerence between central and effective surface brightness is a constant for an exponential profile.," The error in the effective surface brightness is exactly the same, as the difference between central and effective surface brightness is a constant for an exponential profile."
Phe total,The total
fast. reliable. and sufficiently accurate compared to line-by-line calculations.,"fast, reliable, and sufficiently accurate compared to line-by-line calculations."
 We compared spectra computed by line-by-line and correlated-A models. and found that the mean difference between the two techniques is less than 5 per cent. which is considerably smaller than the range of uncertainties on available exoplanet measurements.," We compared spectra computed by line-by-line and $k$ models, and found that the mean difference between the two techniques is less than 5 per cent, which is considerably smaller than the range of uncertainties on available exoplanet measurements."
 In the &-distribution method. the absorption spectrum over an interval is sorted in order of increasing absorption and the fraction of the interval with absorption less than a certain value A(g) is represented in terms of the fraction of the interval g.," In the $k$ -distribution method, the absorption spectrum over an interval is sorted in order of increasing absorption and the fraction of the interval with absorption less than a certain value $k(g)$ is represented in terms of the fraction of the interval $g$."
 Since A(g) is a smoothly varying function of g. it may be integrated with relatively few quadrature points /V to determine the mean transmission over the interval as: where m is the amount of a molecule: and A; and Ag; are the k-coefficients and quadrature weights at each quadrature point (Irwinetal.2008).," Since $k(g)$ is a smoothly varying function of $g$, it may be integrated with relatively few quadrature points $N$ to determine the mean transmission over the interval as: where $m$ is the amount of a molecule; and $k_{i}$ and $\Delta g_{i}$ are the $k$ -coefficients and quadrature weights at each quadrature point \citep{irw08}."
. The A-coefficients are calculated from line data for a range of temperatures and pressures expected in the atmosphere in advance of the retrieval to enable rapid calculation of the transmission during each iteration., The $k$ -coefficients are calculated from line data for a range of temperatures and pressures expected in the atmosphere in advance of the retrieval to enable rapid calculation of the transmission during each iteration.
 The molecular line lists used in this study were taken from HITEMP2010 (Rothmanetal.2010) for H»O. the Carbon Dioxide Spectroscopic Databank (CDSD-(Tashkunetal.2003) (also used for HITEMP2010) for CO». HITEMP1995 (Rothmanetal.1995) for CO. and the Spherical Top Data System (STDS) (Wenger&Champion1998). for CH.," The molecular line lists used in this study were taken from HITEMP2010 \citep{rot10} for $_{2}$ O, the Carbon Dioxide Spectroscopic Databank (CDSD-1000)\citep{tas03} (also used for HITEMP2010) for $_{2}$, HITEMP1995 \citep{rot95} for CO, and the Spherical Top Data System (STDS) \citep{wen98} for $_{4}$."
 The absorption coefficients for the extreme atmospheric temperatures found within exoplanetary atmospheres are continuously being improved. so the conclusions of this study are considerably limitec by the quality of the available spectroscopic data.," The absorption coefficients for the extreme atmospheric temperatures found within exoplanetary atmospheres are continuously being improved, so the conclusions of this study are considerably limited by the quality of the available spectroscopic data."
 We have made use of the best available spectroscopic parameters at the time of writing. and our line database will be updated as new sources of data become available.," We have made use of the best available spectroscopic parameters at the time of writing, and our line database will be updated as new sources of data become available."
 The combination of the optimal estimation retrieval scheme and the correlated-# method enhances the efficiency of the retrieva process in terms of time and computational resources., The combination of the optimal estimation retrieval scheme and the $k$ method enhances the efficiency of the retrieval process in terms of time and computational resources.
 Furthermore. NEMESIS calculates the matrix. of the partial derivatives of radiances at each wavelength with respect to each retrieved variable. which are called the frnedional derivatives (or Jacobrans). in order that the contribution of different atmospheric yurameters at each wavelength can be easily interpreted by comparing the elements of this matrix.," Furthermore, NEMESIS calculates the matrix of the partial derivatives of radiances at each wavelength with respect to each retrieved variable, which are called the $functional$ $derivatives$ (or $Jacobians$ ), in order that the contribution of different atmospheric parameters at each wavelength can be easily interpreted by comparing the elements of this matrix."
" Our e priori dayside atmosphere of HD 19597330 extends Tom "" to IO bar to capture the full atmospheric range of »otential spectral contributions.", Our $a$ $priori$ dayside atmosphere of HD 189733b extends from $^{-9}$ to 10 bar to capture the full atmospheric range of potential spectral contributions.
 For initial modelling we assumed hat all species were well-mixed throughout the atmosphere and he molecular abundanee was defined in terms of a single sealing yarameter., For initial modelling we assumed that all species were well-mixed throughout the atmosphere and the molecular abundance was defined in terms of a single scaling parameter.
 This is because the retrieval of a continuous profile of composition would be under-constrained considering the small number of data points and the low spectral resolutions available. eading to non-physical oscillations in the retrieved profile.," This is because the retrieval of a continuous profile of composition would be under-constrained considering the small number of data points and the low spectral resolutions available, leading to non-physical oscillations in the retrieved profile."
 The « priori estimate for the abundance scaling parameter is assumed to dave a large uncertainty so that retrieved values are not weighted by initial guesses since a simple scaling parameter already includes vertical smoothing., The $a$ $priori$ estimate for the abundance scaling parameter is assumed to have a large uncertainty so that retrieved values are not weighted by initial guesses since a simple scaling parameter already includes vertical smoothing.
 The @ priori for temperature. however. is assumed to have a continuous profile and the assumed error on he « priori profile was adjusted to achieve the optimal balance between the quality of the fit to the measured data and the vertical smoothing.," The $a$ $priori$ for temperature, however, is assumed to have a continuous profile and the assumed error on the $a$ $priori$ profile was adjusted to achieve the optimal balance between the quality of the fit to the measured data and the vertical smoothing."
 This will be further discussed in Section 5.3.1., This will be further discussed in Section 5.3.1.
 As an important source of absorption in exoplanetary atmospheres. collisional-induced absorption (CIA) between principle gases are included in the model atmosphere.," As an important source of absorption in exoplanetary atmospheres, collisional-induced absorption (CIA) between principle gases are included in the model atmosphere."
 We consider the interactions between H»—-H» and He—-He and the coefficients are taken from Borysowetal.(1989). Borysow&Frommhold(1989). Borysow&Frommhold(1990).. Zheng&Borysow(1995)... Borysowetal.(1997).. and Borysow(2002).," We consider the interactions between $_{2}$ $_{2}$ and $_{2}$ –He and the coefficients are taken from \citet{bor89}, \citet{bor892}, \citet{bor90}, \citet{zhe95}, \citet{bor97}, and \citet{bor02}."
. The mole fractions of H» and He are assumed to be related to the fractions of atomic H and He. which are close to the typical solar value of 0.91 and 0.0887 each (Burrows&Sharp1999).," The mole fractions of $_{2}$ and He are assumed to be related to the fractions of atomic H and He, which are close to the typical solar value of 0.91 and 0.0887 each \citep{bur99}."
. The implications of this assumption will be tested in Section 5.5., The implications of this assumption will be tested in Section 5.5.
 Seattering by clouds and hazes has been reported from fiS7 observations at visible wavelengths. yielding a featureless ransmission spectrum of HD 189733b (Pontetal.2008:Singetal.2011).," Scattering by clouds and hazes has been reported from $HST$ observations at visible wavelengths, yielding a featureless transmission spectrum of HD 189733b \citep{pon08,sin11}."
 Recently. Hengetal.(2011) showed that the effect of scattering by clouds and hazes modifies the inferred emperature profile. and that an isothermal temperature structure above an adiabatic troposphere could also be caused by a cloud-op or haze layer.," Recently, \citet{hen11} showed that the effect of scattering by clouds and hazes modifies the inferred temperature profile, and that an isothermal temperature structure above an adiabatic troposphere could also be caused by a cloud-top or haze layer."
 We consider. however. that the inclusion of scattering. clouds and hazes in our retrieval model is beyond he scope of the present study for a number of reasons: (1) the scattering properties of possible clouds and hazes is insufficiently known: (1) adding such particles would greatly enlarge an already arge and poorly constrained parameter space and: (i) modelling scattering. processes will significantly increase the computation ime.," We consider, however, that the inclusion of scattering clouds and hazes in our retrieval model is beyond the scope of the present study for a number of reasons: (i) the scattering properties of possible clouds and hazes is insufficiently known; (ii) adding such particles would greatly enlarge an already large and poorly constrained parameter space and; (iii) modelling scattering processes will significantly increase the computation time."
 In summary. the addition of scattering would increase the complexity of the model and does not appear to be warranted by he dayside emission spectra studied here.," In summary, the addition of scattering would increase the complexity of the model and does not appear to be warranted by the dayside emission spectra studied here."
 We intend to assess the ikely effects of scattering on our modelling of folJipilers in a ollow-up study., We intend to assess the likely effects of scattering on our modelling of $hot Jupiters$ in a follow-up study.
 To retrieve the 7-7 profile and compositional abundances. we used three measurement. sets of the secondary. eclipse for HD 189733b. representing all of the available infrared dayside eclipse measurements at the present time. ranging over a wide wavelength range from [.45 jim to 24 jim: Go eighteen 7/5 77NICMOS (Swain channels covering the range 1.45—2.5 jim: (1) forty seven Spitzer IRS (Grillmairetal.2008) channels covering the range 5—14.5 jim and one IRS photometry (Deming channel at 16 sam: and (i) five Spi/zer TRAC (Infrared. Array Camera) and MIPS (Multiband Imaging Photometer for pil zer) (Charbonneauetal.2008). photometry channels at 3.6. 4.5. 5.8. 8.0. and 24 jim. The measurement errors on the three observational datasets are directly obtained from the studies referenced above.," To retrieve the $P$ $T$ profile and compositional abundances, we used three measurement sets of the secondary eclipse for HD 189733b, representing all of the available infrared dayside eclipse measurements at the present time, ranging over a wide wavelength range from 1.45 $\mu$ m to 24 $\mu$ m: (i) eighteen $HST$ /NICMOS \citep{swa09} channels covering the range 1.45–2.5 $\mu$ m; (ii) forty seven $Spitzer$ IRS \citep{gri08} channels covering the range 5–14.5 $\mu$ m and one IRS photometry \citep{dem06} channel at 16 $\mu$ m; and (iii) five $Spitzer$ IRAC (Infrared Array Camera) and MIPS (Multiband Imaging Photometer for $Spitzer$ ) \citep{cha08} photometry channels at 3.6, 4.5, 5.8, 8.0, and 24 $\mu$ m. The measurement errors on the three observational datasets are directly obtained from the studies referenced above."
 To integrate the binned flux in each channel. we use (i) filter widths taken from the literature for Sp/fzer broadband channels (Fazioetal.2004:RiekeDeming2006): (i) ILSTINICMOS widths of ~10 nm at 2 jim: and (ii) Spit zer/IRS widths of ~ 100 nm at 8 jam. The reference stellar spectrum for HD 189733 is taken from the Kurucz grid model’.," To integrate the binned flux in each channel, we use (i) filter widths taken from the literature for $Spitzer$ broadband channels \citep{faz04,rie04,dem06}; (ii) $HST$ /NICMOS widths of $\sim$ 10 nm at 2 $\mu$ m; and (iii) $Spitzer$ /IRS widths of $\sim$ 100 nm at 8 $\mu$ m. The reference stellar spectrum for HD 189733 is taken from the Kurucz grid ."
 Despite the substantial efforts towards finding and reducing the errors from the data. uncertainties on the given datasets still remain and are widely distributed over the wavelengths due to various error sources.," Despite the substantial efforts towards finding and reducing the errors from the data, uncertainties on the given datasets still remain and are widely distributed over the wavelengths due to various error sources."
 First of all. techniques to decorrelate a transit light curve from the combined light of a planetary system are not consistent each other fe.g. Swainetal.(2008). vs. Gibsonetal. D]].," First of all, techniques to decorrelate a transit light curve from the combined light of a planetary system are not consistent each other [e.g. \citet{swa08} vs. \citet{gib11}] ]."
 We will discuss the effects of underestimated errors on, We will discuss the effects of underestimated errors on
We simulate the scenario described. setting the background cleusity to be 5%Ut of the value in the dense chimps. aud the initial velocity of this gas to be zero.,"We simulate the scenario described, setting the background density to be $5\%$ of the value in the dense clumps, and the initial velocity of this gas to be zero."
 We focus just ou the collision region. so that the right and left sides ο ‘the domain are set to “clump” conditions. as in Equation (19)).," We focus just on the collision region, so that the right and left sides of the domain are set to “clump"" conditions, as in Equation \ref{cflIC}) )."
 When the two dense «lumps meet each other. a stroug shock forms (Fig. 5)).," When the two dense clumps meet each other, a strong shock forms (Fig. \ref{CC3}) )."
 Since all [uid variables (à. By. Up) are s and continuous prior to shock formation. the features produced are not a consequence of discontilous initial coucitious.," Since all fluid variables $n$, $B_y$, $v_x$ ) are smooth and continuous prior to shock formation, the features produced are not a consequence of discontinuous initial conditions."
 This test case eventually evolves to profile similar to that in the simpe convergent [low test (Fie. 6))., This test case eventually evolves to profile similar to that in the simple convergent flow test (Fig. \ref{CC4}) ).
 The central peaks in deusity auc ratio show up as well., The central peaks in density and mass-to-flux ratio show up as well.
 Susequent evolution leads to a decline in the central peak iu 5 aud niBy (Fig. 6))., Subsequent evolution leads to a decline in the central peak in $n$ and $n/B_y$ (Fig. \ref{CC4}) ).
 Peaks in density above the “steady” shock solution have also been observed in other ambipolar diffusion simulations using different MHD codes (e.g..Choietal.2009)..," Peaks in density above the “steady"" shock solution have also been observed in other ambipolar diffusion simulations using different MHD codes \citep[e.g.,][]{2009ApJS..181..413C}."
 In. addition. similar transient behavior of C shocks has been noted i1 models with more complex chemistry impleineuted (e.g.Chiezeοἱa.1998:vanLooetal.2009:Ashinore2010)..," In addition, similar transient behavior of C shocks has been noted in models with more complex chemistry implemented \citep[e.g.][]{1998MNRAS.295..672C, 2009MNRAS.395..319V, 2010A&A...511A..41A}."
" Physically. we believe these peaks arise becaise the neutrals are effectivey Cunmaegnetized"" when the shock first forms."," Physically, we believe these peaks arise because the neutrals are effectively “unmagnetized"" when the shock first forms."
 As a COllsecquence. 1ie ueutrals cau be very stroiely compressed. forming what is seen as a central density peak in Figs.," As a consequence, the neutrals can be very strongly compressed, forming what is seen as a central density peak in Figs."
 1-6.., \ref{cvg1}$ $-$ \ref{CC4}.
 The magnetic field. however. does not follow the inital strong compression of the ueutrals.," The magnetic field, however, does not follow the initial strong compression of the neutrals."
 Iustead. the overv-conupressed ueutrals generate higher pressure in the central regions. inhibiting the magnetic [lux from gettiug in.," Instead, the overly-compressed neutrals generate higher pressure in the central regions, inhibiting the magnetic flux from getting in."
 Fig., Fig.
 7 slows the total p 5 ⋅ ↠∖⊽∖⊽⊳∖↕↩⋯⋅∖∖↽∐≺↵↕⋅↩∫⋟≜≕⊳⋮↳∖∶∣↗⊔≬⋡⊽∖−⋜↕∐≺⇂∫⋟∠↗∶≵≩−↙∕∕↖∖∣⊤⋅∖∖↽∐∐∣⋅∣∣∐↕≺↵⋜↕⊳∖⋯⋅≺↵≺⇂∐⊔∐≺↵↥⋜↕∣⋉∐⋅⋜↕↕∩↥⋅⊽∖⊽⊔⋅⋜⋃∐≺↵⋅⊺∐≺↵⊳∖≺↵ tlire," \ref{pressure} shows the total pressure $P_\mathrm{tot} = \rho_n {v_n}^2 + P_\mathrm{gas} + P_B$ of the system, where $P_\mathrm{gas} = \rho_n {c_s}^2$ and $P_B = B^2/8\pi$, with $v_n$ measured in the “laboratory"" frame."
"e terms correspoxd to ACUgrj. ry. and rg7/?Jg respective""y. ln. Equation (11))."," These three terms correspond to ${\cal M}^2/r_n$, $r_n$, and ${r_B}^2/\beta$ respectively, in Equation \ref{up_down}) )."
 Since uses the conservative [ori1 of the momentum equation (OUA(pe)/0f+OP/Or= 0). P iust become coustant in the posttock region at late times.," Since uses the conservative form of the momentum equation $ \partial \left(\rho v\right) / \partial t + \partial P_\mathrm{tot} / \partial x = 0$ ), $P_\mathrm{tot}$ must become constant in the post-shock region at late times."
 For stroiο shocks. the maguetic pressure term dominates at late tiides in the post-shock region.," For strong shocks, the magnetic pressure term dominates at late times in the post-shock region."
 At early times. there is a slight depression of the magnetic field streng at the center of the shock. in order to balauce the extremely Ligh neutral £as pressure in the «sity peak.," At early times, there is a slight depression of the magnetic field strength at the center of the shock, in order to balance the extremely high neutral gas pressure in the density peak."
 Combining he strong ueutral compression aL slight magnetic exclusion. the —jass-Lo-Iux ratio is elevated in tle ce “when a shock forms.," Combining the strong neutral compression and slight magnetic exclusion, the mass-to-flux ratio is elevated in the center when a shock forms."
 The αςlisious between neutrals aud lous will gradually slow down the iucoο neutrals auc compress iois and maguetic field to the center., The collisions between neutrals and ions will gradually slow down the incoming neutrals and compress ions and magnetic field to the center.
 Meauwhile. the neutrals in the 'eutral peak diffuse outward in order to balance tle increasing maguetic pressure auc keep the total )ressture constalr.," Meanwhile, the neutrals in the central peak diffuse outward in order to balance the increasing magnetic pressure and keep the total pressure constant."
 Eventually. the ious aud weutrals interact sufficiently that a steady-state C shiocJ structure deveops.," Eventually, the ions and neutrals interact sufficiently that a steady-state C shock structure develops."
" The post-shock n/B, is the same as the upstream value.", The post-shock $n/B_y$ is the same as the upstream value.
 However. the ambivolar diffusion process takes tiue. aud during the trausieut stage. a region of very strongly compressed ueutrals will be present.," However, the ambipolar diffusion process takes time, and during the transient stage, a region of very strongly compressed neutrals will be present."
 Our finding that there is a transieut stage of very strong deusity compression. with au enliauced," Our finding that there is a transient stage of very strong density compression, with an enhanced"
Aapni~A20.,$\lambda_{MRI}\sim\Delta/20$.
 Ou general erounds. however. we note hat we would expect that the transition poiut would lic at Aarg&SA.," On general grounds, however, we note that we would expect that the transition point would lie at $\lambda_{MRI}\ll 8\Delta$."
 A eiven vertical field is unstable uot just o the fastest erowing mode. but also to a whole spectrum of slower-growing modes that have longer wavelengths hat are more casily resolvable ποσααν.," A given vertical field is unstable not just to the fastest growing mode, but also to a whole spectrum of slower-growing modes that have longer wavelengths that are more easily resolvable numerically."
 Plausibe he transition point will then correspond to the condition hat we resolve the slowest erowing mode that grows appreciably before it is truncated by non-linear coupling o other MRI modes or some other aspect of the plivsics (ee. a dynamo evele).," Plausibly, the transition point will then correspond to the condition that we resolve the slowest growing mode that grows appreciably before it is truncated by non-linear coupling to other MRI modes or some other aspect of the physics (e.g. a dynamo cycle)."
 Tf this is the case. then it is uusirprising that the transition poiut varies between ocal and elobal simulations. since the time scale available or a mode to grow may well depeud on the presence or absence of a low density disk corona within which the AIRI is not active.," If this is the case, then it is unsurprising that the transition point varies between local and global simulations, since the time scale available for a mode to grow may well depend on the presence or absence of a low density disk corona within which the MRI is not active."
" To consider this more quautitativelv. consider purely vertical AIRT modes (4,=&, 0) in a thin accretiou disk (so that radial eracieuts of pressure and eutropy can be neglected)."," To consider this more quantitatively, consider purely vertical MRI modes $k_r=k_\phi=0$ ) in a thin accretion disk (so that radial gradients of pressure and entropy can be neglected)."
 Let the spacetime dependence of the modes he (€{οςobz), Let the spacetime dependence of the modes be $e^{i(\omega t-kz)}$.
 The dispersion relation for these modes (Balbus&Hawley1991). reads where 2?=4heer and # is the radial epicyclic (angular) frequency., The dispersion relation for these modes \citep{mri1} reads where $\tilde{\omega}^2=\omega^2-k^2v_A^2$ and $\kappa$ is the radial epicyclic (angular) frequency.
 We wish to examine modes with wavelengths much longer than the fastest growing mode. ic. with |keV/O9|< 1.," We wish to examine modes with wavelengths much longer than the fastest growing mode, i.e., with $|kv_A/\Omega_0|\ll 1$ ."
 Rewriting in terms of the growth rate. d=fw and expanding the dispersion relation to lowest order in k?267/08 gives Suppose that a eiven mode cau erow exponcutially for a time 7 before it is truncated by mode coupling or some other unspecified physical process.," Rewriting in terms of the growth rate, $\sigma=-i\omega$ and expanding the dispersion relation to lowest order in $k^2v_A^2/\Omega_0^2$ gives Suppose that a given mode can grow exponentially for a time $\tau$ before it is truncated by mode coupling or some other unspecified physical process."
 Then. the slowest erowing inode that actually experiences siguificaut erowth (hereafter. the slowest appreciably eroxwing mode [SAGAL)) has ση=20/7 and a saveumuber given by Our lypothesis is that the transition point iu the fiux-stress relation correspouds to the point where the slowest appreciably erowiug mode is just resolvable. ie.. where Asayin=28hugs~SA.," Then, the slowest growing mode that actually experiences significant growth (hereafter, the slowest appreciably growing mode [SAGM]) has $\sigma_{sagm}=2\pi/\tau$ and a wavenumber given by Our hypothesis is that the transition point in the flux-stress relation corresponds to the point where the slowest appreciably growing mode is just resolvable, i.e., where $\lambda_{sagm}\equiv
2\pi/k_{sagm}\sim 8\Delta$."
 This predicts a transition point at where r'—Titan. topp bouge the orbital period at that radius.," This predicts a transition point at where $\tau^\prime\equiv\tau/t_{orb}$, $t_{orb}$ being the orbital period at that radius."
 Equation 13. offers some insight iuto the transition point found in the flux-stress relations that we have been considering., Equation \ref{eq:lmri} offers some insight into the transition point found in the flux-stress relations that we have been considering.
" In the local uustratified simulations of Pessaletal.(2007).. the implicit potential is Newtonian (ne= OF) and we find that Agra, A. imiplving 7~ f,5,."," In the local unstratified simulations of \citet{boxscaling}, the implicit potential is Newtonian $\kappa^2=\Omega_0^2$ ) and we find that $\lambda_{MRI}\sim \Delta$ , implying $\tau\sim t_{orb}$ ."
" lu the global sinuuations presented here. we find the transition point at Aaya,cA/20 which (accounting for the fact that #7?<o2 in the pseudo-Newtoniau poteutial) eijvoes το—θέ "," In the global simulations presented here, we find the transition point at $\lambda_{MRI}\sim \Delta/20$ which (accounting for the fact that $\kappa^2<\Omega_0^2$ in the pseudo-Newtonian potential) gives $\tau\sim 5-10t_{orb}$."
Thus. within the framework of this argument. the difference iu the location of the transition point between the local wustratified and the global siuulations is due to a difference iu the robustness of the long waveleneth aud slowly erowiug modes: slowing erowius nodes appear to be able to erow for longer withiu the elobal simulation before being truncated.," Thus, within the framework of this argument, the difference in the location of the transition point between the local unstratified and the global simulations is due to a difference in the robustness of the long wavelength and slowly growing modes; slowing growing modes appear to be able to grow for longer within the global simulation before being truncated."
 The nature of this differeuce. which must be closely related to the saturation of the turbulent state. is bevoud the scope of this paper aud will be explored in future work.," The nature of this difference, which must be closely related to the saturation of the turbulent state, is beyond the scope of this paper and will be explored in future work."
 The results of refsec:fluxstress sugecsts that the Guctuating) magnetic Hux threading a local patch of the disk determines the ro0 component of the magnetic stress generated by he turbulence in that patch., The results of \\ref{sec:fluxstress} suggests that the (fluctuating) magnetic flux threading a local patch of the disk determines the $r-\phi$ component of the magnetic stress generated by the turbulence in that patch.
 If the vertical maguetic Hux is indeed the causal ageut i determine the stress. we expect a temporal lag between fluctuations in the maeuectic flux and the resulting variations in he stress," If the vertical magnetic flux is indeed the causal agent in determining the stress, we expect a temporal lag between fluctuations in the magnetic flux and the resulting variations in the stress."
 On the basis of experiments with local sinulations (ILuwvlevetal.1996).. we expect this lag to © approximately two (local) orbital periods.," On the basis of experiments with local simulations \citep{hawley96}, we expect this lag to be approximately two (local) orbital periods."
 Thus. we expect the temporal lag to imerease with radius in the disk due to the increasing orbital period.," Thus, we expect the temporal lag to increase with radius in the disk due to the increasing orbital period."
 To search for this lag. we use and output the 3-d structure of the disk once every 0.1 ISCO orbits during the interval between 5090 ISC'O orbits (this is 10 times the nominaldata output rate).," To search for this lag, we use and output the 3-d structure of the disk once every 0.1 ISCO orbits during the interval between 50–90 ISCO orbits (this is 10 times the nominaldata output rate)."
" Using these 100 snapshots of the disk structure. we then colputed the instantaneous vertical magnetic faxes aud maguctic ro stresses in families of co-moving wedges at three radi rC[SryLey.12r,}."," Using these 400 snapshots of the disk structure, we then computed the instantaneous vertical magnetic fluxes and magnetic $r-\phi$ stresses in families of co-moving wedges at three radii $r\in
\{8r_g, 10r_g, 12r_g\}$."
" The azunutballv averaged value of c, at cach radius was used to track a given comoving wedee between tinesteps.", The azimuthally averaged value of $v_\phi$ at each radius was used to track a given comoving wedge between timesteps.
 This procedure is uot fully lagrangian. because it docs not account for the radial movement or fluctuating azimuthal velocity of a comoving patch. but we expect these effects to be uceleible for the short timeframe uudoer consideration.," This procedure is not fully lagrangian, because it does not account for the radial movement or fluctuating azimuthal velocity of a comoving patch, but we expect these effects to be negligible for the short timeframe under consideration."
 The time-series of maguetic fux and stress for cach wedge were then cross-correlated aud. finally. the cross-correlations for all wedges at a given radius were averaged.," The time-series of magnetic flux and stress for each wedge were then cross-correlated and, finally, the cross-correlations for all wedges at a given radius were averaged."
 The resulting averaged temporal cross-correlations are shown in Fie. L., The resulting averaged temporal cross-correlations are shown in Fig. \ref{crosscorr}.
 At each radius we see a strong lustautaneous correlation. Likely due to the nmuuediate sheariug of perturbed vertical fields.," At each radius we see a strong instantaneous correlation, likely due to the immediate shearing of perturbed vertical fields."
 However. in general. 16 cross-correlation is biased toward positive lag.," However, in general, the cross-correlation is biased toward positive lag."
 Tis is consistent with what we would expect. namely that the xeseuce of vertical flux will feed the MBI aud result in rosthanced transport.," This is consistent with what we would expect, namely that the presence of vertical flux will feed the MRI and result in enhanced transport."
 Of note is the fact that the imner-lost radius considered. R=sry exhibits a double peak structure whereas this is unresolved at higher radii.," Of note is the fact that the inner-most radius considered, $R=8r_g$ exhibits a double peak structure whereas this is unresolved at higher radii."
" Also seculiar is the fact that the outermost radius. R=I2r, is senificautlv less biased towards positive lag than the other radii under consideration."," Also peculiar is the fact that the outer-most radius, $R=12r_g$ is significantly less biased towards positive lag than the other radii under consideration."
 A further exploration of hese issues is bevoud the scope of this paper. aud will ο explored iu future work cmplovine orbital advection algoritlins aud test-particle tracers in order to correctly ollow the evolution of a local patc[u," A further exploration of these issues is beyond the scope of this paper, and will be explored in future work employing orbital advection algorithms and test-particle tracers in order to correctly follow the evolution of a local patch."
"m Whenconsidering the correlation between vertical fux and stress we workdirectly with AM, but when studving possible trends in average stress with domain size we instead define au effective o-paraneter"," Whenconsidering the correlation between vertical flux and stress we workdirectly with $M_{r\phi}$ , but when studying possible trends in average stress with domain size we instead define an effective $\alpha$ -parameter"
shear signal may suffer [rom such a bias (Whiteetal.2002)..),shear signal may suffer from such a bias \citep{white:clusters-completeness}. .)
 On intermediate scales. the orientation of nearby filamentary structure correlated with the cluster lens (as well as the angular distribution of nearby correlated halos} must also be random for a large statistical sample.," On intermediate scales, the orientation of nearby filamentary structure correlated with the cluster lens (as well as the angular distribution of nearby correlated halos) must also be random for a large statistical sample."
 Since such structure is correlated wilh the lens. itdoes contribute to (he tangential shear on large scales: indeed. as we have argued above. it dominates the signal from clusters on scales above afew ft\Mpe and simply represents the large-scale chister-mass correlation funelion.," Since such structure is correlated with the lens, it contribute to the tangential shear on large scales; indeed, as we have argued above, it dominates the signal from clusters on scales above a few $h^{-1}$ Mpc and simply represents the large-scale cluster-mass correlation function."
 Again. as shown above. the effects of correlated structure along the line of sight appear to be negligible out to the virial radius of massive clusters.," Again, as shown above, the effects of correlated structure along the line of sight appear to be negligible out to the virial radius of massive clusters."
 The only residual problem for the statistical inversion method is anisotropic sample variance ab large scales (discussed above in Section 6)): Chis finite-volume ellect creates a small scatter in the inferred mass but does not create a bias., The only residual problem for the statistical inversion method is anisotropic sample variance at large scales (discussed above in Section \ref{section:tests}) ); this finite-volume effect creates a small scatter in the inferred mass but does not create a bias.
 We have developed. a non-parametric method lor inverüing cross-correlation lensing measurements to obtain average densitv and mass profiles of halos., We have developed a non-parametric method for inverting cross-correlation lensing measurements to obtain average density and mass profiles of halos.
 We have demonstrated the method on an N-body simulation and shown that it successfully recovers the 3D proliles., We have demonstrated the method on an N-body simulation and shown that it successfully recovers the 3D profiles.
 We argued that asphericitv of halos and Iarge-scale structure along the line of sight do not introduce bias or substantial uicertaintv in the inferred mass estimates., We argued that asphericity of halos and large-scale structure along the line of sight do not introduce bias or substantial uncertainty in the inferred mass estimates.
 This method should find several useful applications in survevs., This method should find several useful applications in surveys.
 Applied to galaxv-galaxy lensing measurements. this method can be used to measure «Αρη=Ὁροῃ£y(r) Lor samples of lens galaxies (e.g.. Sheldon. et al.," Applied to galaxy-galaxy lensing measurements, this method can be used to measure $\Delta\rho(r) = \Omega_m\rho_{crit}~\xi_{gm}(r)$ for samples of lens galaxies (e.g., Sheldon, et al."
 2004)., 2004).
 On small scales. (his provides information about galaxy dark matter halos and should thereby constrain models of galaxy formation ancl evolution. including the details of hierarchical structure formation and merging.," On small scales, this provides information about galaxy dark matter halos and should thereby constrain models of galaxy formation and evolution, including the details of hierarchical structure formation and merging."
 Statistical nass profiles around galaxies can provide virial mass measurements and. combined with the average lieht profile around these galaxies. determine mass-to-lieht ratios as a Iunction of scale.," Statistical mass profiles around galaxies can provide virial mass measurements and, combined with the average light profile around these galaxies, determine mass-to-light ratios as a function of scale."
 On larger scales. galaxv-galaxy lensing inversions should tell us more about cosmology.," On larger scales, galaxy-galaxy lensing inversions should tell us more about cosmology."
" The autocorrelation function of galaxies allows us to measure £4,(7)=b(r)&,,,(r). where b(r) is the scale-dependent bias and £,,;00) is the mass autocorrelation function."," The autocorrelation function of galaxies allows us to measure $\xi_{gg}(r) = b^2(r) \xi_{mm}(r)$, where $b(r)$ is the scale-dependent bias and $\xi_{mm}(r)$ is the mass autocorrelation function."
" Lensine allows us to measure O,,£,,(r)=OQ,bir)r«(r)Sianeli). where the eross-bias. κ). is sometimes referred to as stochastic bias (a term we do not recommend using in this context. since it is not bounded by 41)."," Lensing allows us to measure $\Omega_m~\xi_{gm}(r) =
\Omega_m~b(r)~\mbox{r}_{\times}(r)~\xi_{mm}(r)$, where the cross-bias, $\mbox{r}_{\times}(r)$, is sometimes referred to as stochastic bias (a term we do not recommend using in this context, since it is not bounded by $\pm1$ )."
 The presence of these (vo bias functions. b(r) aad ry(7). allows for the most general model of the relative clustering of galaxies and mass. (," The presence of these two bias functions, $b(r)$ and $\mbox{r}_{\times}(r)$, allows for the most general model of the relative clustering of galaxies and mass. ("
See also (2005)..),See also \cite{neyrinck:halo-model}. .)
 Simulations indicate that rz(7) is consistent with unity on scales larger than, Simulations indicate that $\mbox{r}_{\times}(r)$ is consistent with unity on scales larger than
by svuchrotvon photons plavs an important role in the internal shock model (Guetta. Spada Waxman 2001: Asano Kobavashi 2002).,"by synchrotron photons plays an important role in the internal shock model (Guetta, Spada Waxman 2001; Asano Kobayashi 2002)."
 In order to obtain hieh radiative efficiency. ancl the characteristic clustering of spectral break energies of GRBs in the range 0.1-1 MeV. the collision radii are required to be similar to the photosphere radius.," In order to obtain high radiative efficiency and the characteristic clustering of spectral break energies of GRBs in the range 0.1-1 MeV, the collision radii are required to be similar to the photosphere radius."
 Since collisions producing narrow pulses occur at small radii xD. the photosphere might obscure these. leaving only the wider pulses visible.," Since collisions producing narrow pulses occur at small radii $\propto D$, the photosphere might obscure these, leaving only the wider pulses visible."
 This will make (he temporal profile smooth., This will make the temporal profile smooth.
 In a wider jet. (he tvpical Lorentz lactor Dis smaller. a larger fraction of the collisions occur at small radii xE? below the photosphere. therefore. the smoothing effect is expected to be stronger.," In a wider jet, the typical Lorentz factor $\Gamma$ is smaller, a larger fraction of the collisions occur at small radii $\propto \Gamma^2$ below the photosphere, therefore, the smoothing effect is expected to be stronger."
 This might explain the correlation between luminosity ancl variability in GRBs., This might explain the correlation between luminosity and variability in GRBs.
 We discuss the smoothing effect. by using a multiple-shell model., We discuss the smoothing effect by using a multiple-shell model.
 We represent the irregular wind bv relativistic shells in a manner similar to that in INPS., We represent the irregular wind by relativistic shells in a manner similar to that in KPS.
" Because of the relativistic beaming effect. we can study (he emission [rom a jet by using spherical shells with an isotropic explosion energv. Ej,(0)=46/6? where E is the geometrically corrected explosion energv and 8 is the opening angle."," Because of the relativistic beaming effect, we can study the emission from a jet by using spherical shells with an isotropic explosion energy $E_{iso}(\theta)=4 E/\theta^2$ where $E$ is the geometrically corrected explosion energy and $\theta$ is the opening angle."
" The eanuna-ray energy released troma GRD is narrowly clustered around 10"" ere (Frail et al.", The gamma-ray energy released froma GRB is narrowly clustered around $10^{51}$ erg (Frail et al.
 2001). and the conversion elliciency [rom (he explosion energv into the gamma-rays is about 1056(Guetta. Spada Waxman 2001: IXobavashi Sari 2001).," 2001), and the conversion efficiency from the explosion energy into the gamma-rays is about $10\%$(Guetta, Spada Waxman 2001; Kobayashi Sari 2001)."
 Then. E~LO” erg.," Then, $E \sim 10^{52}$ erg."
 Cousider a two shell collision which is the elementary process in the multiple shell evolution., Consider a two shell collision which is the elementary process in the multiple shell evolution.
" A rapid shell with Lorentz [actor D, and mass m, catches up to a slower one with D; and m. and the two merge to temporarily form a single shell."," A rapid shell with Lorentz factor $\Gamma_r$ and mass $m_r$ catches up to a slower one with $\Gamma_s$ and $m_s$, and the two merge to temporarily form a single shell."
" Using conservation ol energv and momentum we calculate (he Lorentz factor of the merged shell to be The internal energy. of the merged shell is the difference of kinetic energy before ancl alter the collision. Ej;=m,c(T,—Ty)-mic(D. D,)."," Using conservation of energy and momentum we calculate the Lorentz factor of the merged shell to be The internal energy of the merged shell is the difference of kinetic energy before and after the collision, $E_{int}=m_rc^2(\Gamma_r-\Gamma_m)+m_sc^2(\Gamma_s-\Gamma_m)$ ."
" We assume that electrons are accelerated to a power-law distribution of Lorentz factor ορ. with a minimum. Lorentz factor mi Nelda,ox8,Uds,2 am."," We assume that electrons are accelerated to a power-law distribution of Lorentz factor $\gamma_e$, with a minimum Lorentz factor $\gamma_m$ : $N(\gamma_e)d\gamma_e \propto \gamma_e^{-p}d\gamma_e, \gamma_e \ge
\gamma_m$ ."
 Throughout this paper. we use thestandard choice," Throughout this paper, we use thestandard choice"
comparison with observational data.,comparison with observational data.
 We estimate the cllipcity coniponeuts of cach halo by nicasmiug the second moments of the projected mass distribution m a vay analogous to that used for the surface brightuess distribution of galaxies (Valdes 1983. Alivalda-Escudé 1991): where and Tere. i6 and g ave the coordinates of the ceuter of mass of the halo iu question. which contains VY particles of equal mass.," We estimate the ellipcity components of each halo by measuring the second moments of the projected mass distribution in a way analogous to that used for the surface brightness distribution of galaxies (Valdes 1983, Miralda-Escudé 1991): where and Here, $\overline{x}$ and $\overline{y}$ are the coordinates of the center of mass of the halo in question, which contains $N$ particles of equal mass."
 In the observational case. these moments are often calculated using the surface brightuess weielted by a function chosen to maximize the S/N of the measurement (see c.g. WW).," In the observational case, these moments are often calculated using the surface brightness weighted by a function chosen to maximize the S/N of the measurement (see e.g., VW)."
 In our case. as we have mentioned above. we will probe the seusitivitv of our iueasuremeuts to varatious of this sort by comparing halos chosen with a laree FOF linking leueth to those using a small value which cfectively give a non-zero weight oulv to the dense central reelous.," In our case, as we have mentioned above, we will probe the sensitivity of our measurements to variations of this sort by comparing halos chosen with a large FOF linking length to those using a small value which effectively give a non-zero weight only to the dense central regions."
 When working with ραπνο statistics. such as the ellipticity correlation fiction (see 833). it will be advantageous to redefine ey aud e» for cach pair of halos. with the w-axis defined to be line joinine the halo ceutres aud the y-axis a line perpendicular to it.," When working with pairwise statistics, such as the ellipticity correlation function (see 3), it will be advantageous to redefine $e_{1}$ and $e_{2}$ for each pair of halos, with the $x$ -axis defined to be line joining the halo centres and the $y$ -axis a line perpendicular to it."
 The quantity (61.609) defines a pseudovector. of leneth C=VG| 3.," The quantity $(e_{1},e_{2})$ defines a pseudovector, of length $e=\sqrt{e_{1}^{2}+e_{2}^{2}}$ ."
 A positive c4 component indicates a stretching along the. axis. aud a negative compoucut a stretching along the yo axis.," A positive $e_{1}$ component indicates a stretching along the $x-$ axis, and a negative component a stretching along the $y-$ axis."
 The e» component is likewise a nieasure of the stretching along axes at Lodeg to the .c aud y axes., The $e_{2}$ component is likewise a measure of the stretching along axes at $45 \deg$ to the $x$ and $y$ axes.
 If we have au cllipse of ellipticity e (6—(1.«λα147). where q is the ratio of winmor and major axis leneths). then where s is the augle between the major axis aud the CO naxbs," If we have an ellipse of ellipticity $e$ $e \equiv (1-q^{2})/(1+q^{2})$, where $q$ is the ratio of minor and major axis lengths), then where $\beta$ is the angle between the major axis and the $x-$ axis."
 Iu the bottom three panels of Fig. l..," In the bottom three panels of Fig. \ref{halo},"
 we plot the elliptidties of the halos chosen by the three eroupfinders., we plot the ellipticities of the halos chosen by the three groupfinders.
 For each halo. we show a bar with leneth proportional to c. oriented along the direction of the major axis.," For each halo, we show a bar with length proportional to $e$, oriented along the direction of the major axis."
 It is dathicult to pick out bv eve auv tendency for halos to be aligned., It is difficult to pick out by eye any tendency for halos to be aligned.
 One worry which we mieht have had concerus the tendency of close halos to be joined together bv thin vuecks” of particles and for the eroupfinder to count such pairs as oue halo., One worry which we might have had concerns the tendency of close halos to be joined together by thin “necks” of particles and for the groupfinder to count such pairs as one halo.
 If these halos are distributed along fiunenuts. then we might expect au artificial alieuneut to arise.," If these halos are distributed along filaments, then we might expect an artificial alignment to arise."
 This does not seen to be au obvious problem here., This does not seem to be an obvious problem here.
" For example. if we look at halos picked out bv the FOE,4 aud FOF)» eroupfiuders. the ellipticities and. directions appear to be faitly similar even though differeut parts of the halo are being used to caleulate them."," For example, if we look at halos picked out by the $_{0.1}$ and $_{0.2}$ groupfinders, the ellipticities and directions appear to be fairly similar even though different parts of the halo are being used to calculate them."
 We will check the statistical endencies for aliguiments iu the next section., We will check the statistical tendencies for alignments in the next section.
 The TOP alos do in several cases appear to be made frou several sanall halos joined together., The HOP halos do in several cases appear to be made from several small halos joined together.
 Presuinablv this could be alleviated by tuning the free parameters which govern the uethod. but we have not attempted to do this.," Presumably this could be alleviated by tuning the free parameters which govern the method, but we have not attempted to do this."
 Tn Fig. 2..," In Fig. \ref{ehist},"
 we show the probabilitv distribution of « values for the halos (again with >20 member particles) Xcked out bv the three groupfuders., we show the probability distribution of $e$ values for the halos (again with $>20$ member particles) picked out by the three groupfinders.
" They are all fairly simular. with a mean e around 0.Ll. slightly higher than the 43 typical of real galaxies (σοι, Wittinan 2000)."," They are all fairly similar, with a mean $e$ around $0.4$, slightly higher than the $0.3$ typical of real galaxies (e.g., Wittman 2000)."
 Observationally. ellipticity correlations are measured as a function of augular separation. ou the plane of the sky (at least for weak lensing survevs).," Observationally, ellipticity correlations are measured as a function of angular separation, on the plane of the sky (at least for weak lensing surveys)."
 With distance information (for example. using redshifts}. it would be possible to measure them as a function of separation iu three dimensions. and this is what it is most natural to do using our Nbody simulations.," With distance information (for example, using redshifts), it would be possible to measure them as a function of separation in three dimensions, and this is what it is most natural to do using our Nbody simulations."
 Iu this section. we will do this. and later convert the measurements to angular correlations which cau be compared to weak lensing survey results.," In this section, we will do this, and later convert the measurements to angular correlations which can be compared to weak lensing survey results."
 This conversion will be done iu two different wavs., This conversion will be done in two different ways.
 The first is an analytical projection of our three dimensional results using Linwber’s equation (Limber. 1959].," The first is an analytical projection of our three dimensional results using Limber's equation (Limber, 1959)."
 The second is a direct measurement from sinmlated surveys made bv projecting the halo distributions im the box. and applying a radial selection fuuctiou.," The second is a direct measurement from simulated surveys made by projecting the halo distributions in the box, and applying a radial selection function."
 In the prescut section. although we will be dealing with halo separations in three dimensions. if is worth bearing im mind that we restrict ourselves to quantities which can be measured directly observationally (albeit with redshifts). so that the cllipticitics wewill be correlating are projected ellipticities (defined in Equation 1)).," In the present section, although we will be dealing with halo separations in three dimensions, it is worth bearing in mind that we restrict ourselves to quantities which can be measured directly observationally (albeit with redshifts), so that the ellipticities wewill be correlating are projected ellipticities (defined in Equation \ref{e1}) )."
 Two cllipticity correlations cau be defined (following Aliralda-Escudé 1991) as, Two ellipticity correlations can be defined (following Miralda-Escudé 1991) as
during the evolution of structure at early times. anc visible as an isotropic eamma-rav backeround at the present cay.,"during the evolution of structure at early times, and visible as an isotropic gamma-ray background at the present day."
 In this paper. we calculate. the Dux produced. by annihilations in simple CDAL halos. relative to the tux ooduced in a uniform background.," In this paper, we calculate the flux produced by annihilations in simple CDM halos, relative to the flux produced in a uniform background."
 We correct this result or halo substructure. and. integrate over a large. volume o determine the cosmological background. from WIAIP annihilation as a function of redshift.," We correct this result for halo substructure, and integrate over a large volume to determine the cosmological background from WIMP annihilation as a function of redshift."
 The outline of this yaper is as follows., The outline of this paper is as follows.
 In section 2. we define a climensionless lux multiplier f. that accounts for the enhanced. rate of two-body interactions produced. by inhomogeneities in he dark matter distribution. and. determine its value for simple virialisecd halos.," In section 2, we define a dimensionless flux multiplier $f$ that accounts for the enhanced rate of two-body interactions produced by inhomogeneities in the dark matter distribution, and determine its value for simple virialised halos."
 In section 3 we caleulate { for cosmological volumes. using analytic estimates of the halo mass function and halo concentrations. ancl determine its redshift’ dependence.," In section 3 we calculate $f$ for cosmological volumes, using analytic estimates of the halo mass function and halo concentrations, and determine its redshift dependence."
 Finally. in section. 4 we study. the contribution to f. from substructure within virialisecl halos. and calculate f for à set of realistic halos generated: using a semi-analvtie model of halo substructure.," Finally, in section 4 we study the contribution to $f$ from substructure within virialised halos, and calculate $f$ for a set of realistic halos generated using a semi-analytic model of halo substructure."
 Throughout this paper we assume a Lambda-CDM (LODM) cosmology with a cosmological constant Ag=0.7. a matter density Oo=0.3 and a Hubble parameter 40=fhlO00knmis with 5—0.65.," Throughout this paper we assume a Lambda-CDM (LCDM) cosmology with a cosmological constant $\Lambda_0 = 0.7$, a matter density $\Omega_{{\rm m},0} = 0.3$ and a Hubble parameter $H_0 = h \times\ 100 {\rm km\,s^{-1}}$, with $h = 0.65$."
 In current hierarchical models. cold dark matter is expected to. form centrally concentrated: halos with a characteristic density profile.," In current hierarchical models, cold dark matter is expected to form centrally concentrated halos with a characteristic density profile."
 Dense substructure is abundant within these halos. as a relic from earlier stages of the hierarchical merging process.," Dense substructure is abundant within these halos, as a relic from earlier stages of the hierarchical merging process."
 Since the annihilation Hux is quadratic in the density. these inhomogeneities will increase the πας from a halo of a given mean density.," Since the annihilation flux is quadratic in the density, these inhomogeneities will increase the flux from a halo of a given mean density."
 We begin hy computing a cimensionless quantity that describes this enhancement., We begin by computing a dimensionless quantity that describes this enhancement.
 In the next section we will then stuck the evolution ofthis quantity with epoch., In the next section we will then study the evolution of this quantity with epoch.
 If clark matter consists of neutralinos with a mass my ancl a velocity-averaged cross-section for annihilation (ec then the annihilation Uux produced within a volume V. will be where p is the local density ofCDM.," If dark matter consists of neutralinos with a mass $m_\chi$ and a velocity-averaged cross-section for annihilation ${\langle \sigma v \rangle}$, then the annihilation flux produced within a volume $V$ will be where $\rho$ is the local density of CDM."
 For non-relativistic xuticles. £00? is approximately independent of e so the Dux will just be proportional to p.," For non-relativistic particles, ${\langle \sigma v \rangle}$ is approximately independent of $v$ so the flux will just be proportional to $\rho^2$."
 Since the rate depends equadratically on the density. the otal rate from a given mass within a given. volume will be ueher ifthe dark matter is distributed inhomogeneously.," Since the rate depends quadratically on the density, the total rate from a given mass within a given volume will be higher if the dark matter is distributed inhomogeneously."
 We can study thisenhancement. by defining the dimensionless lux multiplier for a distribution within a volume V. where p is the average density within this volume.," We can study thisenhancement by defining the dimensionless flux multiplier for a distribution within a volume $V$, where ${\bar \rho}$ is the average density within this volume."
 This function can also be written as à mass-weighted density average: where A is the total mass within V., This function can also be written as a mass-weighted density average: where $M$ is the total mass within $V$ .
 Clearly. f=1 for a homogeneous distribution. while for a power-law density profile pxrmm provided à«1.5.," Clearly $f = 1$ for a homogeneous distribution, while for a power-law density profile $\rho \propto r^{-\alpha}$, provided $\alpha < 1.5$."
" For à—1.5. integrating equation (2)) from luin lO Mux ΟΙΝΟΝ: so f diverges logarithmically as rii, goes to zero."," For $\alpha = 1.5$, integrating equation \ref{fluxm1}) ) from $r_{\rm min}$ to $r_{\rm max}$ gives: so $f$ diverges logarithmically as $r_{\rm min}$ goes to zero."
 Two analytic density profiles are commonly used to fit the spherically averaged: properties of dark matter halos. the NEW profile (Navarro. Frenk White 1996. 1997). and the Moore profile (Moore et 11998).," Two analytic density profiles are commonly used to fit the spherically averaged properties of dark matter halos, the NFW profile (Navarro, Frenk White 1996, 1997), and the Moore profile (Moore et 1998)."
 We can specify. these generically as with à=31.5—2 for the NEW profile and a= for the Moore profile.," We can specify these generically as with $\alpha = \beta = 1, \gamma = 2$ for the NFW profile and $\alpha = \beta = 1.5, \gamma = 1$ for the Moore profile."
 Ehe total mass within radius r is with for the NEW profile and for the Moore. profile. while the mean density. within this radius is simply where.—rfr.," The total mass within radius $r$ is with for the NFW profile and for the Moore profile, while the mean density within this radius is simply where $x \equiv r/r_s$."
" 1n a cosmological setting. halos are virialised out to a radius r« corresponding to an overdensity A, of roughly 200 relative to the background."," In a cosmological setting, halos are virialised out to a radius $r_v$ corresponding to an overdensity $\Delta_c$ of roughly 200 relative to the background."
" The concentration ο—refr, ofa halo describes the size of this radius relative tor...", The concentration $c \equiv r_v/r_s$ of a halo describes the size of this radius relative to $r_s$ .
 Calculating the (ux multiplier over the virialised region of a halo. we get a function which depends only on e," Calculating the flux multiplier over the virialised region of a halo, we get a function which depends only on$c$ :"
"To properly assess the contribution of the various mono-abundance sub-populations to the mass or surface-mass budget in the Solar neighborhood, we must convert the number of spectroscopically-observed stars in each bin to a surface-mass density at the Solar neighborhood.","To properly assess the contribution of the various mono-abundance sub-populations to the mass or surface-mass budget in the Solar neighborhood, we must convert the number of spectroscopically-observed stars in each bin to a surface-mass density at the Solar neighborhood."
" On the one hand, this requires incorporation of a model for the sselection function (see A of B11)) and our exponential-disk fits (B11))."," On the one hand, this requires incorporation of a model for the selection function (see A of ) and our exponential-disk fits )."
" On the other hand, this requires the use of stellar-population models that can relate the observed number of G-type dwarfs to the total mass of the stellar population, given the metallicity of the sub-population and an assumed star formation history for it."," On the other hand, this requires the use of stellar-population models that can relate the observed number of G-type dwarfs to the total mass of the stellar population, given the metallicity of the sub-population and an assumed star formation history for it."
 This is described in detail in ??.., This is described in detail in \ref{sec:surfmass}.
" Briefly, for G-type dwarfs in a given ([Fe/H],,[a/Fe])) bin we calculate their total number per square pc over vertical height) by adjusting the normalization(integrated of the number-density profile such that after running it through our model for the sselection function it predicts the observed number of stars in each bin."," Briefly, for G-type dwarfs in a given ) bin we calculate their total number per square pc (integrated over vertical height) by adjusting the normalization of the number-density profile such that after running it through our model for the selection function it predicts the observed number of stars in each bin."
 Then we relate the number density of G-type dwarfs to the total stellar surface-mass density by multiplying the number density by the average mass of a G-type dwarf—calculated using Padova isochrones (Marigoetal. dividing it by the fraction of the mass in a stellar 2008)-—andpopulation in G-type dwarfs (calculated using the same isochrones and assuming a lognormal Chabrier(2001) initial mass function)., Then we relate the number density of G-type dwarfs to the total stellar surface-mass density by multiplying the number density by the average mass of a G-type dwarf—calculated using Padova isochrones \citep{Marigo08a}- —and dividing it by the fraction of the mass in a stellar population in G-type dwarfs (calculated using the same isochrones and assuming a lognormal \citet{Chabrier01a} initial mass function).
" At a given abundance, this fraction of course depends on the age of the population, and here is calculated by marginalizing over a flat age distribution between 0.5 and 10 Gyr for each bin."," At a given abundance, this fraction of course depends on the age of the population, and here is calculated by marginalizing over a flat age distribution between 0.5 and 10 Gyr for each bin."
" However, averaging only over older ages for a-enhanced stars would be appropriate, as a-enhanced stars likely represent the oldest part of the disk."," However, averaging only over older ages for $\alpha$ -enhanced stars would be appropriate, as $\alpha$ -enhanced stars likely represent the oldest part of the disk."
" As we show in ??,, this gives similar results, with a slightly steeper decline in X(HRo) with h,."," As we show in \ref{sec:surfmass}, this gives similar results, with a slightly steeper decline in $\Sigma(R_0)$ with $h_z$."
 We calculate uncertainties on the surface-mass densities by varying the density-parameters according to the posterior probability distribution for the parameters inB11.., We calculate uncertainties on the surface-mass densities by varying the density-parameters according to the posterior probability distribution for the parameters in.
 These uncertainties do not include systematic uncertainties due to the use of the stellar isochrones., These uncertainties do not include systematic uncertainties due to the use of the stellar isochrones.
" This procedure results in an estimate of the stellar surface-mass density contribution at the Solar radius for any abundance-selected sub-population, which in turn has a vertical scale height associated with it."," This procedure results in an estimate of the stellar surface-mass density contribution at the Solar radius for any abundance-selected sub-population, which in turn has a vertical scale height associated with it."
" The relative total stellar surface-mass densities of different mono-abundance bins are not affected by assuming a different initial mass function although assuming a different IMF can systematically (IMF),shift all surface-mass densities by a few percent (see below)."," The relative total stellar surface-mass densities of different mono-abundance bins are not affected by assuming a different initial mass function (IMF), although assuming a different IMF can systematically shift all surface-mass densities by a few percent (see below)."
 'The results from this mass estimation are shown in s1 and 2.., The results from this mass estimation are shown in s \ref{fig:mass_afe_feh} and \ref{fig:mass_hz_afe}.
" The left panel of 1 is simply a more coarsely binned version of the unweighted GG-dwarf sample abundance distribution, which shows two distinct maxima, one considerably more metal poor and o-enhanced than the other, seemingly reflecting a chemically-distinct thick-disk component."," The left panel of \ref{fig:mass_afe_feh} is simply a more coarsely binned version of the unweighted G-dwarf sample abundance distribution, which shows two distinct maxima, one considerably more metal poor and $\alpha$ -enhanced than the other, seemingly reflecting a chemically-distinct thick-disk component."
" It is important to note that the marginalized mmetallicity distribution (left [Fe/H]panel, top) shows no hint of any bi-modality."," It is important to note that the marginalized metallicity distribution (left panel, top) shows no hint of any bi-modality."
" There is distinct bi-modality in the marginalized ddistribution, but [a/Fe]Schónrich&Binney already showed that even for a smooth age distribution(2009a) such bi-modality arises, simply separating stars that formed before and after enrichment by SN Ia became important."," There is distinct bi-modality in the marginalized distribution, but \citet{Schoenrich08a} already showed that even for a smooth age distribution such bi-modality arises, simply separating stars that formed before and after enrichment by SN Ia became important."
" The right panel of 1 shows the stellar mass density at the solar radius Ro in each elemental-abundance bin, corrected for selection effects due to the spectroscopic sselection as described above."," The right panel of \ref{fig:mass_afe_feh} shows the stellar surface-mass density at the solar radius $R_0$ in each elemental-abundance bin, corrected for selection effects due to the spectroscopic selection as described above."
" It represents the properly mass-weighted, underlying distribution of disk stars in the elemental-abundance space spanned by aand[a/Fe],, and dramatically differs from the raw sample distribution in the left panel: it does not have the strong bi-modality apparent in the raw nnumber distribution."," It represents the properly mass-weighted, underlying distribution of disk stars in the elemental-abundance space spanned by and, and dramatically differs from the raw sample distribution in the left panel: it does not have the strong bi-modality apparent in the raw number distribution."
 It also shows no hint of a bi-modal mmetallicity distribution and the remaining hint of bbi-modality is explained as in the left panel., It also shows no hint of a bi-modal metallicity distribution and the remaining hint of bi-modality is explained as in the left panel.
 The right panel of 1 now provides the relevant weights of each mono-abundance bin to a surface-mass- distribution of disk scale heights., The right panel of \ref{fig:mass_afe_feh} now provides the relevant weights of each mono-abundance bin to a surface-mass-weighted distribution of disk scale heights.
 The colored symbols in 2 show exactly these surface-mass, The colored symbols in \ref{fig:mass_hz_afe} show exactly these surface-mass
lass concentration iu the local universe.,mass concentration in the local universe.
" Throughout. the following cosmological parameters are adopted: Q,,=0.3. O4=O7. and IZ,=100h=T0 5. which Πάρος a scale of 1.6 Alpe + (77 kpe t) at the ~20.000 nuuean redshift of the IRS."," Throughout, the following cosmological parameters are adopted: $\Omega _m = 0.3$ , $\Omega _\Lambda = 0.7$ , and $H_o = 100h
= 70 $ $^{-1}$, which implies a scale of 4.6 Mpc $^{-1}$ (77 kpc $^{-1}$ ) at the $\sim$ 20,000 mean redshift of the HRS."
 The UK Schinidt Telescope (UIST) six-deerce field (dE). multi-fiber system is uniquely suited to survey large supercluster regious im the nearby universe.," The UK Schmidt Telescope (UKST) six-degree field (6dF), multi-fiber system is uniquely suited to survey large supercluster regions in the nearby universe."
 6dE deploys 150 fibers over a circular field of caameter wwith a nuninnun required spacing between fibers of 677. set bv the magnetic prism buttons.," 6dF deploys 150 fibers over a circular field of diameter with a minimum required spacing between fibers of 7, set by the magnetic prism buttons."
 Light is fed from the fibers iuto a fast f/0.9 CCD spectrograph (Parkeretal.19908)., Light is fed from the fibers into a fast f/0.9 CCD spectrograph \citep{par98}.
. Two interchangeable field plate units allow for the simultaneous observation of the current feld aud configuration of the next., Two interchangeable field plate units allow for the simultaneous observation of the current field and configuration of the next.
 A practical limiting magnitude for the svstem is b; = 17.5., A practical limiting magnitude for the system is $_J$ = 17.5.
 All of these attributes taken together iuply that the 6dF is most effectively used to probe the large-scale iuter-cluster euvironnieuts of local superclusters. while avoiding the more densely crowded cluster members.," All of these attributes taken together imply that the 6dF is most effectively used to probe the large-scale inter-cluster environments of local superclusters, while avoiding the more densely crowded cluster members."
 Consequently. in studviug the TRS our goal was to produce a catalog of galaxies iu the iuter-cluster region.," Consequently, in studying the HRS our goal was to produce a catalog of galaxies in the inter-cluster region."
" Galaxy solectiou took place in the following manner,", Galaxy selection took place in the following manner.
" A . aarea of the sky centered upon a=3""19.5 wwas chosen for the region of observation based upon previously published literature (Zuccaetal. 1993)."," A $\times$ area of the sky centered upon $\alpha = 3^h19^m, \delta =
-$ was chosen for the region of observation based upon previously published literature \citep{zuc93}."
. A complete catalog of all galaxies dow ο à by imaegnitude of 17.5 was extracted iu four equivalent « 67) regions from the UKST survey ates previously scanned by. the SuperCOSMOS nachine (ILuublvetal.2001)., A complete catalog of all galaxies down to a $_J$ magnitude of 17.5 was extracted in four equivalent $\times$ ) regions from the UKST survey plates previously scanned by the SuperCOSMOS machine \citep{ham01}.
". There was also he addition of a fifth rectaugular region « 6"")) in the far Southeru portion to incorporate he feld surrounding ACO clusters 3106 αμα 3161.", There was also the addition of a fifth rectangular region $\times$ ) in the far Southern portion to incorporate the field surrounding ACO clusters 3106 and 3164.
 The ealaxy classification flag assigned bv SuperCOSAIOS was used for the initial suuple selection., The galaxy classification flag assigned by SuperCOSMOS was used for the initial sample selection.
 The b;=17.5 maguitude lait was adopted as a practical limiting magnitude for the L2an aperture UNST., The $_J= 17.5$ magnitude limit was adopted as a practical limiting magnitude for the 1.2-m aperture UKST.
 To avoid expeuding fibers on galaxies within clusters. our original intention was to excise from the catalog all ealaxies within a 17 radius circle of sixteen ACO clusters listed by Zuccaetal.(1993). απο νους of the URS aud intersecting our observing region.," To avoid expending fibers on galaxies within clusters, our original intention was to excise from the catalog all galaxies within a $^{\circ}$ radius circle of sixteen ACO clusters listed by \cite{zuc93} as members of the HRS and intersecting our observing region."
" The 1"" radius exclusion corresponds to ~2 Abell radi (where 1 Ry»= 2 Alpe) at the mean redshift of the TRS.", The $^{\circ}$ radius exclusion corresponds to $\sim$ 2 Abell radii (where 1 $_A=$ 2 Mpc) at the mean redshift of the HRS.
 This would cusure that new spectroscopic information relates onlv to the iuter-cluster regions of the IIRS., This would ensure that new spectroscopic information relates only to the inter-cluster regions of the HRS.
 Iowever. a coding error was discovered in the programs. that excises galaxies from the cluster regions oulv after the observations were made.," However, a coding error was discovered in the program that excises galaxies from the cluster regions only after the observations were made."
 The cos(ó) conversion factor in the Right Ascension (RA) coordinate. when expressed in degrees. was not included. in the calculation of angular distances of galaxies from cluster ceuters.," The $\cos$ $\delta$ ) conversion factor in the Right Ascension (RA) coordinate, when expressed in degrees, was not included in the calculation of angular distances of galaxies from cluster centers."
 As a result. the actual excision regious are fhtened in the RA coordinate and correspondingly more flattened at higher Declination.," As a result, the actual excision regions are flattened in the RA coordinate and correspondingly more flattened at higher Declination."
 The typical fattening is a factor of 1.6., The typical flattening is a factor of 1.6.
 Nevertheless. the result remains that we have generated a sample that is alinost entirely conrprised. of iuter-cluster galaxies.," Nevertheless, the result remains that we have generated a sample that is almost entirely comprised of inter-cluster galaxies."
 After the above coustraiuts were applied. there relmained 28LS ealaxies (Figure |).," After the above constraints were applied, there remained 2848 galaxies (Figure \ref{f1}) )."
 The πιακα uunber of optical ealaxy redshifts that could be obtained uncer optimal observing conditions was estimated at 1500., The maximum number of optical galaxy redshifts that could be obtained under optimal observing conditions was estimated at 1500.
 Consequeuth. we produced a subcatalog of 1500 targets from the original list of 2818.," Consequently, we produced a subcatalog of 1500 targets from the original list of 2848."
 This was accomplished as follows., This was accomplished as follows.
 Calaxies in cach ος region were assigued a random) Πο aud then arranged in ascending order., Galaxies in each $\times$ region were assigned a random number and then arranged in ascending order.
 This ordering provides a basis for selecting an unbiased subsample from the larger complete sample., This ordering provides a basis for selecting an unbiased subsample from the larger complete sample.
" The uuuberius schemes from the iadividual SO"" iregious were merged into a final catalog of 1500 objects with cach region weighted according to the fraction of galaxies found in that region.", The numbering schemes from the individual $\times$ regions were merged into a final catalog of 1500 objects with each region weighted according to the fraction of galaxies found in that region.
 That is. if of the ealaxies in the original catalog came from a particular region. the subcatalog of 1500 ealaxics also contaiuc from that region.," That is, if of the galaxies in the original catalog came from a particular region, the subcatalog of 1500 galaxies also contained from that region."
 Hence the method preserves natural galaxy overdeusities while randomlysuupling theeutire extracted region., Hence the method preserves natural galaxy overdensities while randomlysampling theentire extracted region.
 Finally. a Digitized Sky Survey (DSS) nuage of cach target was cxamuned to further reduce the uuuber of misclassified galaxies iun the," Finally, a Digitized Sky Survey (DSS) image of each target was examined to further reduce the number of misclassified galaxies in the"
but not exactly: it is not possible to derive C! when the two dips are perfectly superposed).,but not exactly: it is not possible to derive $C$ when the two dips are perfectly superposed).
 lt appears [rom Fig., It appears from Fig.
 1 that C could be represented with a polynomial with terms less and less significant when the order increases: the most important is Cy. aid a slope may possibly be added. since C may be as large as Co|O.1 when [V1E is large.," \ref{fig:C_V2} that $C$ could be represented with a polynomial with terms less and less significant when the order increases: the most important is $C_0$, and a slope may possibly be added, since $C$ may be as large as $C_0 + 0.1$ when $|V_1-V_2|$ is large."
 Then. €'CoΟτοvs].," Then, $C=C_0 + C_1|V_1-V_2|$."
 The (6η.ον) coellicients. were calculated: as terms of he orbital solution for 7 of the S svstems with more than 10 blended measurements that are presented in. Fig.," The $(C_0, C_1)$ coefficients were calculated as terms of the orbital solution for 7 of the 8 systems with more than 10 blended measurements that are presented in Fig."
 2 (the triple system 2:54B was set aside)., \ref{fig:orblend} (the triple system 2:54B was set aside).
 It appeared. t Cy was never significant. since its maximum value. t was obtained for system. 2:74B. was only 2.35 times its uncertainty.," It appeared that $C_1$ was never significant, since its maximum value, that was obtained for system 2:74B, was only 2.35 times its uncertainty."
 Moreover. negative values were found for 3 systems (1:19D. 2:58D. and 2:79B). although this should no happen in theory.," Moreover, negative values were found for 3 systems (1:19B, 2:58B, and 2:79B), although this should not happen in theory."
 We conclude then that the Ct term. is sullicient for deriving the blended. velocities., We conclude then that the $C_0$ term is sufficient for deriving the blended velocities.
 This is easilv confirmed by a visual inspection of Eig. 2..," This is easily confirmed by a visual inspection of Fig. \ref{fig:orblend},"
 where the mocde velocity curves derived by assuming only the €t term in the expression of Care represented., where the model velocity curves derived by assuming only the $C_0$ term in the expression of $C$ are represented.
 The blended measurements are equally distributed: around the theoretical curves. anc no deviation related with |Viio is visible in practice.," The blended measurements are equally distributed around the theoretical curves, and no deviation related with $|V_1-V_2|$ is visible in practice."
 Lt was possible to derive the orbital elements for 51 stars. inclucling a triple-system solution which consists of 2 orbits.," It was possible to derive the orbital elements for 51 stars, including a triple-system solution which consists of 2 orbits."
 The spectroscopic orbits are presented. in Tables 40 to 6.., The spectroscopic orbits are presented in Tables \ref{tab:orb1} to \ref{tab:orb3}.
 We count 40 first orbits: 27 91 and 13 SB2., We count 40 first orbits: 27 SB1 and 13 SB2.
 The 12 other orbits are. distributed. as follows: Six are new orbits that were computed taking into account other measurements in addition to ours: these caleulations were done including the olfset oetween the 2 RW sources as a free. parameter of the model., The 12 other orbits are distributed as follows: Six are new orbits that were computed taking into account other measurements in addition to ours; these calculations were done including the offset between the 2 RV sources as a free parameter of the model.
 Two other orbits are. published orbits which were partly based on our measurements: they are just expressed with the same conventions as the others: RY in the system. and epochs in Julian davs: the last 4 orbits are new orbits derived [rom our measurements alone. since including the others did not ameliorate the solution.," Two other orbits are published orbits which were partly based on our measurements; they are just expressed with the same conventions as the others: RV in the system, and epochs in Julian days; the last 4 orbits are new orbits derived from our measurements alone, since including the others did not ameliorate the solution."
 The phase plots of 49 orbits ave available in electronic form., The phase plots of 49 orbits are available in electronic form.
 “Phe two orbits that were already. published are not drawn again., The two orbits that were already published are not drawn again.
 In Table 3.. we count 15 stars for which it was not possible to derive a spectroscopic orbit.," In Table \ref{tab:Vmoy}, we count 15 stars for which it was not possible to derive a spectroscopic orbit."
 In addition. we still found two triple svstems (2:16 ancl 2:981). with a short-period SB and a clit in the residual RV.," In addition, we still found two triple systems (2:16A and 2:98B), with a short-period SB and a drift in the residual RV."
 Ehe figures showing the racial velocities of these 17 stars as functions of the epochs are given in electronic form., The figures showing the radial velocities of these 17 stars as functions of the epochs are given in electronic form.
 In addition to the two triple svstems alreacly mentioned. we still count 7 long period SBI (2:sL.," In addition to the two triple systems already mentioned, we still count 7 long period SB1 (2:8B,"
can be described by a Kings law (Lada&Lacla2003).. although the present objects are not virialised.,"can be described by a King's law \citep{Lada03}, although the present objects are not virialised."
 Llowever. the differential dust absorption produces conspicuous variations in the RDPs of the present objects.," However, the differential dust absorption produces conspicuous variations in the RDPs of the present objects."
 Wing profile describes the structure of clusters close to spherical symmetry ancl centrally concentrated., King profile describes the structure of clusters close to spherical symmetry and centrally concentrated.
 However. many voung clusters are substructurecl or asymmetric. deviating significantly. fron this shape. and therefore. cannot be fitted hy Ixing's. law (Cartwright&Whitworth2004:Cutermuthctal," However, many young clusters are substructured or asymmetric, deviating significantly from this shape, and therefore cannot be fitted by King's law \citep{Cartwright04, Gutermuth05}."
.2005).. 1n Fig. 14..," In Fig. \ref{fig:14},"
 we show the spatial distribution of stars in the decontaminated photometry of the three EC's that follow a Ixine-like wolile (FSR 784. Sh2-235 Cluster ancl Sh2-235E2)} and two representative cases of objects that do not (AWC 11 and BDSB 73).," we show the spatial distribution of stars in the decontaminated photometry of the three ECs that follow a King-like profile (FSR 784, Sh2-235 Cluster and Sh2-235E2) and two representative cases of objects that do not (KKC 11 and BDSB 73)."
 The former are centrally concentrated and nearly circularly svimmetric., The former are centrally concentrated and nearly circularly symmetric.
 Phe cavities anc overdensities in the stellar distribution (Eig. 14)), The cavities and overdensities in the stellar distribution (Fig. \ref{fig:14}) )
 can be seen as bumps and dips in the RDP (Fig. 11))., can be seen as bumps and dips in the RDP (Fig. \ref{fig:11}) ).
 On the other hand. objects like BDSB τὸ that are not centrally condensed and WAC 11 with more elongated shape do not follow a Wing profile.," On the other hand, objects like BDSB 73 that are not centrally condensed and KKC 11 with more elongated shape do not follow a King profile."
 Phe multiple peaks in the RDPs of these ECs may be a fractal ellect., The multiple peaks in the RDPs of these ECs may be a fractal effect.
 these objects survive the primordial gas expulsion. they may undergo merging evolving into a relatively smooth structure.," If these objects survive the primordial gas expulsion, they may undergo merging evolving into a relatively smooth structure."
 The angular distribution of decontaminated stars (Fig. 14)).," The angular distribution of decontaminated stars (Fig. \ref{fig:14}) ),"
 used. in the CALD construction. reprocuces the distribution of stars in the RDPs built with filtered photometry (Figs.," used in the CMD construction, reproduces the distribution of stars in the RDPs built with filtered photometry (Figs."
 12. and 13)). supporting the consistency οἱ our results.," \ref{fig:12} and \ref{fig:13}) ), supporting the consistency of our results."
 Given the poorly-populated nature of the MS. we simply counted stars in the CAIDs (within the region P<Rapp). and summed their masses as estimated. from the mass-luminosity relation implied by the respective isochrone solution (Sect. 3.2)).," Given the poorly-populated nature of the MS, we simply counted stars in the CMDs (within the region $R<R_{RDP}$ ), and summed their masses as estimated from the mass-luminosity relation implied by the respective isochrone solution (Sect. \ref{sec:3.2}) )."
 The results are given in Table 4.., The results are given in Table \ref{tab8}.
 All the ECs clearly present. clistinet populations of MIS and PAIS stars (Figs., All the ECs clearly present distinct populations of MS and PMS stars (Figs.
 4 o S))., \ref{fig:04} to \ref{fig:08}) ).
 However. given the cillcrential recldening. it is not pessible to attribute a prec“Isc mass value for cach PAm star.," However, given the differential reddening, it is not possible to attribute a precise mass value for each PMS star."
 Thus. we simply count the number of PAIS stars anc adop an average mass value [ου he PMS stars to estin enpars and mpars.," Thus, we simply count the number of PMS stars and adopt an average mass value for the PMS stars to estimate $n_{PMS}$ and $m_{PMS}$."
 Assuming that he mass distribution of the PMS stars also follows Ixroupa's (2001) ME. the zsc. average PAS mass - for masses wilin he range 0.08m(A.)i is <ΠΕΙ=~06A.," Assuming that the mass distribution of the PMS stars also follows Kroupa's (2001) MF, the zsc, average PMS mass - for masses within the range $0.08\la m(\ms)\la7$ - is $<m_{PMS}>\approx0.6\ms$."
 ‘Thus. we simply multiply the number of PAIS stars (Pablc S) w this value to estimate the PAS mass.," Thus, we simply multiply the number of PMS stars (Table 8) by this value to estimate the PMS mass."
 Finally. we add the alter value to the AIS mass to obtain an estimate of the otal stellar mass.," Finally, we add the latter value to the MS mass to obtain an estimate of the total stellar mass."
 Vhese values should. be taken as lower inits., These values should be taken as lower limits.
 N-body simulations of massive star clusters that include the οσοι of eas removal (e.g.Goodwin&Bastian2006) show that the phase of dramatic core radii increase may Last about 10-30 Myr., N-body simulations of massive star clusters that include the effect of gas removal \citep[e.g.][]{Goodwin06} show that the phase of dramatic core radii increase may last about 10-30 Myr.
 Mass segregation may also lead to a phase of core contraction. with high mass stars more concentrated in the core while low mass stars are transferred to outer parts of the cluster.," Mass segregation may also lead to a phase of core contraction, with high mass stars more concentrated in the core while low mass stars are transferred to outer parts of the cluster."
 In this context. we suggest that most objects of our," In this context, we suggest that most objects of our"
The satellite (2)). launched in October 2002. ts an ESA space mission specifically designed to study the gamma-ray sky.,"The satellite \citealt{winkler}) ), launched in October 2002, is an ESA space mission specifically designed to study the gamma-ray sky."
 In particular the IBIS (Imager on Board of the Satellite)telescope (2)) 1s the main hard X-ray/soft gamma-ray coded aperture Imaging mstrument (?)). and is responsible for surveying and cataloguing the sky above 17 keV. Here we discuss the point source location accuracy (PSLA) of the ISGRI low energy detector (15-1000 keV) of IBIS (?)).," In particular the IBIS (Imager on Board of the Satellite)telescope \citealt{ubertini03}) ) is the main hard X-ray/soft gamma-ray coded aperture imaging instrument \citealt{goldwurm03}) ), and is responsible for surveying and cataloguing the sky above 17 keV. Here we discuss the point source location accuracy (PSLA) of the ISGRI low energy detector (15-1000 keV) of IBIS \citealt{lebrun03}) )."
 Due to the continuing necessity to follow-up the growing unidentified source population in other wavebands (particularly optical and infrared). it is of great interest to assure the correctness of the IBIS/ISGRI PSLA.," Due to the continuing necessity to follow-up the growing unidentified source population in other wavebands (particularly optical and infrared), it is of great interest to assure the correctness of the IBIS/ISGRI PSLA."
 In particular. it is hoped that with the release of new and updated Offline Science Analysis (OSA 7.0) software the PSLA could have improved substantially compared to the already published estimates based on early mission data and software releases (?.con-structedthePSLAbaseduponOSA 3.0)..," In particular, it is hoped that with the release of new and updated Off-line Science Analysis (OSA 7.0) software the PSLA could have improved substantially compared to the already published estimates based on early mission data and software releases \citep[][constructed the PSLA based upon OSA 3.0]{gros}."
 Any improvement in the PSLA is important in reducing the chance of random or multiple source associations in other wavebands., Any improvement in the PSLA is important in reducing the chance of random or multiple source associations in other wavebands.
 In order to empirically determine the PSLA of the IBIS/ISGRI telescope. we can extract the positions of objects from the IBIS/ISGRI Science (SeWs) and compare these with their best known positions.," In order to empirically determine the PSLA of the IBIS/ISGRI telescope, we can extract the positions of objects from the IBIS/ISGRI Science (ScWs) and compare these with their best known positions."
 This can then be used to estimate the 90% error offset as a function of detected significance. allowing us to define the 90% PSLA.," This can then be used to estimate the $90\%$ error offset as a function of detected significance, allowing us to define the $90\%$ PSLA."
 In particular. the IBIS/ISGRI telescope (and coded mask telescopes in general) PSLA depends strongly on detection significance. but also on the position of the source within the field of view.," In particular, the IBIS/ISGRI telescope (and coded mask telescopes in general) PSLA depends strongly on detection significance, but also on the position of the source within the field of view."
 We therefore require for our analysis a set of sources spanning a wide range of significances and off-axis positions in order to allow our analysis to be useful for all detected IBIS/ISGRI sources., We therefore require for our analysis a set of sources spanning a wide range of significances and off-axis positions in order to allow our analysis to be useful for all detected IBIS/ISGRI sources.
 We begin by compiling a list of sources with good positions., We begin by compiling a list of sources with good positions.
 It is best at this stage to be conservative and only select sources where accurate nominal positions are known rather than to bias our sample by also including sources with large nominal error radi. like many newly discovered sources.," It is best at this stage to be conservative and only select sources where accurate nominal positions are known rather than to bias our sample by also including sources with large nominal error radii, like many newly discovered sources."
" To do this we take the latest General Reference Catalog (Version 30. ?.. http://isdc.unige.ch/Data/cat/latest/)) and make a selection on the error radius. selecting those sources with an error less than 30"": at this level. the error on the true position should give a negligible contribution to the measured offsets to the IBIS position."," To do this we take the latest General Reference Catalog (Version 30, \citealt{ebisawa}, ) and make a selection on the error radius, selecting those sources with an error less than $''$; at this level, the error on the true position should give a negligible contribution to the measured offsets to the IBIS position."
 Thus the total number of objects used in our sample is 332. spanning a wide range of detection significances and off-axis angles.," Thus the total number of objects used in our sample is 332, spanning a wide range of detection significances and off-axis angles."
 This number might seem small when compared to the 72] sources detected in ?.. however we note that many objects in that catalog are newly discovered sources for which X-ray follow-up is not yet available. and that therefore have relatively large nominal error radi.," This number might seem small when compared to the 721 sources detected in \cite{cat4}, however we note that many objects in that catalog are newly discovered sources for which X-ray follow-up is not yet available, and that therefore have relatively large nominal error radii."
 After performing an imaging analysis with the OSA 7.0 pipeline. we inspected all available ScW images and extracted the fitted positions which. resulted from the image deconvolution: these are the columns FFIN and .FFIN from the file for all pointings where any of the 332 objects was present.," After performing an imaging analysis with the OSA 7.0 pipeline, we inspected all available ScW images and extracted the fitted positions which resulted from the image deconvolution; these are the columns FIN and FIN from the file for all pointings where any of the 332 objects was present."
 The dataset was divided into fully coded field of view (FCFOV) and partially coded field of view (PCFOV)., The dataset was divided into fully coded field of view (FCFOV) and partially coded field of view (PCFOV).
 The IBIS coded aperture mask has a field of view of 30°.. the FCFOV is the central 9*x9* region:," The IBIS coded aperture mask has a field of view of , the FCFOV is the central $\degr\times$ $\degr$ region;"
After separating variables. and projecting the forcing terms appropriately. as done above for the interactions in à warpec disc. the equations describing the coupling between the trapped r mode. eccentric dise and à=1 intermediate mode read,"After separating variables, and projecting the forcing terms appropriately, as done above for the interactions in a warped disc, the equations describing the coupling between the trapped r mode, eccentric disc and $n=1$ intermediate mode read"
XX-ray Center (CXC. operated for NASA by SAO). and the ESA'S SScience Archive (XSA).,"X-ray Center (CXC, operated for NASA by SAO), and the ESA's Science Archive (XSA)."
" PE acknowledges financial support from the Autonomous Region of Sardinia through a research grant under the program PO Sardegna FSE 2007-2013. Τ.Κ. 7/2007 “Promoting scientific research and innovation technology in Sardinia""."," PE acknowledges financial support from the Autonomous Region of Sardinia through a research grant under the program PO Sardegna FSE 2007–2013, L.R. 7/2007 “Promoting scientific research and innovation technology in Sardinia”."
 This work was partially supported by the ASI/INAF contract 1009/10/0., This work was partially supported by the ASI/INAF contract I/009/10/0.
one would see multiple scintles. possibly in different stages of development.,"one would see multiple scintles, possibly in different stages of development."
 Our observing bandwidth of 192 MHz is at least four times more narrow than the expected Avpiss (see Table 3)., Our observing bandwidth of 192 MHz is at least four times more narrow than the expected $\Delta\nu_{\rm DISS}$ (see Table 3).
 It means that at any given time we are usually able to see only a fraction of a single seintle., It means that at any given time we are usually able to see only a fraction of a single scintle.
 Following Cordes Lazio (1991)) we estimated our number of seintles in both frequency and time. and calculations yielded ΔΝ to be very close to unity (as expected) and ΑΔ54.6.," Following Cordes Lazio \cite{cordes91}) ), we estimated our number of scintles in both frequency and time, and calculations yielded $N_f$ to be very close to unity (as expected) and $N_t \simeq 4.6$."
 The latter leads to the expected DISS modulation index of 0.466. much lower than the observed values.," The latter leads to the expected DISS modulation index of 0.466, much lower than the observed values."
 However. after including the contribution from RISS (εις=0.56: the value found via structure. function analysis. see next sections of the paper). and using the total intensity variance formula (Rickett 1990)). we obtained the final expected value of mu=0.889.," However, after including the contribution from RISS $m_{\rm RISS}=0.56$; the value found via structure function analysis, see next sections of the paper), and using the total intensity variance formula (Rickett \cite{rick}) ), we obtained the final expected value of $m_{\rm tot}=0.889$."
 This value is still somewhat lower than the observed modulation indices. but not by a huge margin.," This value is still somewhat lower than the observed modulation indices, but not by a huge margin."
 Because every individual session was at least several hours long. we believe that diffractive scintillations (which happen at the timescale of ca.," Because every individual session was at least several hours long, we believe that diffractive scintillations (which happen at the timescale of ca."
 40 minutes. see subsection 2.4)) should not affect our average flux measurements.," 40 minutes, see subsection \ref{sect_sf}) ) should not affect our average flux measurements."
 On the other hand. refractive timescales. which are significantly longer. may have affected our results.," On the other hand, refractive timescales, which are significantly longer, may have affected our results."
 Figure 2. shows the results of the average flux density measurement versus the observing epoch., Figure \ref{flux_long} shows the results of the average flux density measurement versus the observing epoch.
 Clearly (conf., Clearly (conf.
 Table 1) there is significant variation in the average values., Table 1) there is significant variation in the average values.
 Using these data. we calculated the modulation index of the average flux density measurements. which yielded the value of m=0.57.," Using these data, we calculated the modulation index of the average flux density measurements, which yielded the value of $m=0.57$."
 Assuming that the pulsar itself is not varying in intensity. this modulation could be caused by only the refractive scintillations. which happen at significantly long timescales. close to the length of a single observing session.," Assuming that the pulsar itself is not varying in intensity, this modulation could be caused by only the refractive scintillations, which happen at significantly long timescales, close to the length of a single observing session."
" To take that into account. we calculated the errors of average flux measurements following Κάπρι Stinebring (1992.. KS92) as where Top, 1s [fiasthe length of a given observing session. and fjiss Is the RISS timescale."," To take that into account, we calculated the errors of average flux measurements following Kaspi Stinebring \cite{kasp}, KS92) as where $T_{\rm obs}$ is the length of a given observing session, and $t_{\rm RISS}$ is the RISS timescale."
 To calculate the error values shown on Figure 2 we used the value of rss obtained via the structure function analysis (305 minutes. see subsection 2.4))," To calculate the error values shown on Figure \ref{flux_long} we used the value of $t_{\rm RISS}$ obtained via the structure function analysis (305 minutes, see subsection \ref{sect_sf}) )."
 Using those error estimates. we were able to calculate the weighted average flux density (wj=o ) for our entire observing session. which ts (Fi)=11.69 mJy.," Using those error estimates, we were able to calculate the weighted average flux density $w_i = \sigma_i^{-2}$ ) for our entire observing session, which is $\left<F_{\rm tot}\right> = 11.69$ mJy."
 This value ts shown in Figure 2 as a dashed-dotted horizontal line. along with the simple arithmetic average from Table | (dashed line).," This value is shown in Figure \ref{flux_long} as a dashed-dotted horizontal line, along with the simple arithmetic average from Table \ref{t1} (dashed line)."
 As we mentioned above. the average flux values and their respective uncertainties for the last two sessions in our project differ significantly from the remaining observations.," As we mentioned above, the average flux values and their respective uncertainties for the last two sessions in our project differ significantly from the remaining observations."
 However. a proper calculation of the error estimates improves the picture.," However, a proper calculation of the error estimates improves the picture."
 Por relatively short sessions. which we had towards the end of the project. RISS can play a huge role. but the duration of the session is included in the uncertainty estimates. which makes them more reliable than the formal errors cited in Table 1.," For relatively short sessions, which we had towards the end of the project, RISS can play a huge role, but the duration of the session is included in the uncertainty estimates, which makes them more reliable than the formal errors cited in Table 1."
 The, The
All the necessary X-ray data were obtained from theChandra (NGC 6543:630: DD--303639: 587: and NGC 7027: 588)) and (NGC 7009). science archives.,All the necessary X-ray data were obtained from the (NGC 6543:; $+$ 303639:; and NGC 7027: ) and (NGC 7009) science archives.
" The Chandra observations were carried out with the ACIS-S cleteetor and the corresponding Science Threads for Image Spectroscopy in CLAO were used for extracting the X-ray spectra,", The Chandra observations were carried out with the ACIS-S detector and the corresponding Science Threads for Image Spectroscopy in CIAO were used for extracting the X-ray spectra.
 For (ae NMM-Newton observation of NGC 7009. the most recent version of SAS was used for the spectral extraction.," For the XMM-Newton observation of NGC 7009, the most recent version of SAS was used for the spectral extraction."
 Given the goal of our study. only the EPIC PN spectrum was used in the following analysis due to its better photon statistics compared to those of the EPIC MOSI and MOS2.," Given the goal of our study, only the EPIC PN spectrum was used in the following analysis due to its better photon statistics compared to those of the EPIC MOS1 and MOS2."
 The periods of high background count rates were excluded [rom the EPIC PN data., The periods of high background count rates were excluded from the EPIC PN data.
 Finally. for each object. X-ray spectra were extracted corresponding to the entire object and that part within the optical slit.," Finally, for each object, X-ray spectra were extracted corresponding to the entire object and that part within the optical slit."
 The X-ray spectra were fit within XSPEC 11.3 (Arnancd1996) using theΠάρος model for the emission from opticallv-thin plasma., The X-ray spectra were fit within XSPEC 11.3 \citep{arnaud1996} using the model for the emission from optically-thin plasma.
 We note that the derived parameters of the A-ray emitting plasma do not depend on this choice and using a different model (e.g. mekal) leads to the same results within the expected., We note that the derived parameters of the X-ray emitting plasma do not depend on this choice and using a different model (e.g. ) leads to the same results within the expected.
 Thus. the model was chosen to make caleulations of the aand eenissions technically more straightforward.," Thus, the model was chosen to make calculations of the and emissions technically more straightforward."
 The hot plasma abundances were taken [rom Alaness&Vrtilek(2003) for NCC: 65423. from Manessetal.(2003). lor NGC 1027. and from Guerreroetal...(2002). [or NGC 1009.," The hot plasma abundances were taken from \citet{manessvrtilek2003} for NGC 6543, from \citet{manessetal2003} for NGC 7027, and from \citet{guerreroetal2002} for NGC 7009."
 ForBD+30°3639.. two abundance sets were used. the first being a variant of the Manessetal...(2003). abundance set (fourth column in Table 3:: henceforth. we denote this set as ‘wind’ abundances) and the second a new set derived here (last column in Table 3: nebular’ abuucdances) based partly upon the nebular abundances of Aller&Ivune(1995) and which have values similar to the ‘nebular’ abundances derived bv Arnaudοἱal.(1996) from the analvsis of the ASC'A data of this object.," For, two abundance sets were used, the first being a variant of the \citet{manessetal2003} abundance set (fourth column in Table \ref{table_xrayfit3}; henceforth, we denote this set as `wind' abundances) and the second a new set derived here (last column in Table \ref{table_xrayfit3}; `nebular' abundances) based partly upon the nebular abundances of \citet{allerhyung1995} and which have values similar to the `nebular' abundances derived by \citet{arnaud1996a} from the analysis of the ASCA data of this object."
 The two fits are incdistinguishable in fitting the data., The two fits are indistinguishable in fitting the data.
 Figure 3. presents our fits to the X-ray spectra., Figure \ref{figure_xrayspec} presents our fits to the X-ray spectra.
 The parameters used in these fits. including elemental abundances. are listed in Tables 3. and 4..," The parameters used in these fits, including elemental abundances, are listed in Tables \ref{table_xrayfit3} and \ref{table_7027par}."
 Our fits to the total X-ray spectra olf these PNe are consistent with those obtained by other eroups using the same data sets., Our fits to the total X-ray spectra of these PNe are consistent with those obtained by other groups using the same data sets.
 Finally. we note that. if the Gorenstein(1975). conversion," Finally, we note that, if the \citet{gorenstein1975} conversion"
At the end of the 20th ceutury. observations of tvpe In supernovae (δα) revealed that the Universe expansion is accelerating (27)..,"At the end of the 20th century, observations of type Ia supernovae (SNIa) revealed that the Universe expansion is accelerating \citep{riess98,perlmutter99}."
 Siuce these publications. several efforts wave been mace to explain these observatious (?????? aud references therem).," Since these publications, several efforts have been made to explain these observations \citet{cunha09, frieman08, linder08, linder05, samsing10, freaza02, ishida05,ishida08} and references therein)."
 Iu a standard analvsis. dark-euergy models are characterized by a small set of paramcters.," In a standard analysis, dark-energy models are characterized by a small set of parameters."
 These are placed iuto the cosmic expansion rate by means oftje Friedman equations. iu substitution for the conveutional cosinological-constant terim.," These are placed into the cosmic expansion rate by means of the Friedman equations, in substitution for the conventional cosmological-constant term."
 This approach assmmes a specific dependence of the dark-euergy equation of state (0) on redshift :xd provides some insight iuto the probable valies of the parameters involved., This approach assumes a specific dependence of the dark-energy equation of state $w$ ) on redshift and provides some insight into the probable values of the parameters involved.
 IHowever. the results remain restricted to that particular parametrization.," However, the results remain restricted to that particular parametrization."
 An interesting question to attempt to answer is what can be inferred about the cosmic expansion rate from observations withmit any reference to a specific model for the energy. coutent of the Universe?, An interesting question to attempt to answer is what can be inferred about the cosmic expansion rate from observations without any reference to a specific model for the energy content of the Universe?
 To perforui an iudepeudenu analvsis we used principal component analysis (PCA).," To perform an independent analysis, we used principal component analysis (PCA)."
 Iu simple terms. PCA idcutifies the cirectious of daa poluts c‘lustering in the phase space defined by the parameters of a given model.," In simple terms, PCA identifies the directions of data points clustering in the phase space defined by the parameters of a given model."
 Cousequeutly. it allows a cimenusioialitv redction with as nininimn an information loss as possible (2)..," Consequently, it allows a dimensionality reduction with as minimum an information loss as possible \citep{tegmark97}."
 The importance of a reconsruction o the cosmüc expansion rate las already been investigated iu the literature (72???7)..," The importance of a model-independent reconstruction of the cosmic expansion rate has already been investigated in the literature \citep{Huterer99,Huterer00,Tegmark02,Wang05,mignone08}."
" Tn this coutest. PCA has be4 used to reconstruct the dark-cnerey equation of state (777) and the deceleration paraueter (7) asa ""unction o “redshift."," In this context, PCA has been used to reconstruct the dark-energy equation of state \citep{huterer03,Crittenden09,simpson06} and the deceleration parameter \citep{shapiro06} as a function of redshift."
 The use of PCA was also proposed in the interpretation of future experiments results by ?.., The use of PCA was also proposed in the interpretation of future experiments results by \citet{albretch09}.
 In the face of growing interest in the application of PCA to cosinology. ?. recall that some care must be taken in choosing the basic expansion functiois and the interpretation assigned to the components.," In the face of growing interest in the application of PCA to cosmology, \citet{kitching09} recall that some care must be taken in choosing the basic expansion functions and the interpretation assigned to the components."
 The main goal of this work is to apply PCA to reconstruct directly the piriuneter redshift dependence without auv reference to a specific cosmological model., The main goal of this work is to apply PCA to reconstruct directly the parameter redshift dependence without any reference to a specific cosmological model.
 In this contest. the eieeuvectors aud cigenvalucs of the Fisher matrix form a new basis in which the parameter is expanded.," In this context, the eigenvectors and eigenvalues of the Fisher matrix form a new basis in which the parameter is expanded."
 For the first time. we show that it is possible," For the first time, we show that it is possible"
Preprint The hwdrogeu between galaxies was reiouized more than 12.5 Cer ago (7)., The hydrogen between galaxies was reionized more than $12.5~$ Gyr ago \citep{fan06}.
 After this event. a larecly nuiforma t-ionizing radiatiou backeround pervaded the intergalactic πουπα (IGM) and kept the hydrogen lughly ionized everywhere except within rare. dense pockets (2222??2)..," After this event, a largely uniform -ionizing radiation background pervaded the intergalactic medium (IGM) and kept the hydrogen highly ionized everywhere except within rare, dense pockets \citep{gunn65, cen94, miralda96, hernquist96, haardt96, katz96}."
 The amplitude of this backerouud appears to have declined quickly above a redshift of :=6 (o.@.. 2)) and to havestaved relatively coustaut over Doc Lee. ?)).," The amplitude of this background appears to have declined quickly above a redshift of $z=6$ (e.g., \citealt{fan06}) ) and to havestayed relatively constant over $2<z<4$ (e.g., \citealt{faucher08b}) )."
 It is of some debate whether this decline owed to the overlap stage of cosinological reionization or to something more nundaue (e.g... ?77]).," It is of some debate whether this decline owed to the overlap stage of cosmological reionization or to something more mundane (e.g., \citealt{gnedin00, fan06, becker07}) )."
 Iu addition. the deuse selfshiclding svstenis that remained in the TGAL called Exauau-Init svstems. provide a window iuto structure formation in a different density regime than explored by other large-scale structure probes.," In addition, the dense self-shielding systems that remained in the IGM, called Lyman-limit systems, provide a window into structure formation in a different density regime than explored by other large-scale structure probes."
 This paper aims to study interealactic radiative transfer aud the properties of Lsaiuan-liuüt| svstenis (systems with ccoluuus of 100152.Nyy<1019 2)., This paper aims to study intergalactic radiative transfer and the properties of Lyman-limit systems (systems with columns of $10^{17.2} < N_{\rm HI} < 10^{19}~$ $^{-2}$ ).
 To do so. we post-process cosmological simulations (both with and without ealactic feedback muiplemeutations) with lonizing radiative transfer.," To do so, we post-process cosmological simulations (both with and without galactic feedback implementations) with ionizing radiative transfer."
 The abuudance of Lyman-iuit svstenis in the post-processed simulations is then compared with measurements of their abundance frou quasar absorption line studies., The abundance of Lyman-limit systems in the post-processed simulations is then compared with measurements of their abundance from quasar absorption line studies.
 This comparison tests row cosmological simulations fare at the outskirts of ealactic halos., This comparison tests how cosmological simulations fare at the outskirts of galactic halos.
 Previous attempts to model Exaiuau-Inuit systelus in simulations had reported various levels of agreement with observations aud had used much smaller ON Sizes and particle uuuboers than are preseutly feasible QUT). ," Previous attempts to model Lyman-limit systems in simulations had reported various levels of agreement with observations and had used much smaller box sizes and particle numbers than are presently feasible \citep{katz96, gardner01, kohler07}. ."
Tu addition. recent observations have significantly nuproved the coustraints on the abundance of svstems (2?7)..," In addition, recent observations have significantly improved the constraints on the abundance of systems \citep{prochaska09, prochaska10, songaila10}."
 Our caleulatious also provide au estimate for the -ionizius cuussivity of star-forming galaxies and quasars since. after cosmological reionization. the production rate of ionizing photons was in balance with the number of absorptions.," Our calculations also provide an estimate for the -ionizing emissivity of star-forming galaxies and quasars since, after cosmological reionization, the production rate of ionizing photons was in balance with the number of absorptions."
 These absorptious occured primarily witlin ΓιατιΠΠ systems., These absorptions occurred primarily within Lyman-limit systems.
 Our determinations of the cluissivity using the simulations aeree with observations at redshifts where it can be measured observationallv (2<2« 1l) and they also provide a meaus to estimate the enüssivitv at higher redshifts.," Our determinations of the emissivity using the simulations agree with observations at redshifts where it can be measured observationally $2 < z<4$ ), and they also provide a means to estimate the emissivity at higher redshifts."
 We also study the relationship between the cuissivity of ionizing photons and the intensity of the ioniziug background., We also study the relationship between the emissivity of ionizing photons and the intensity of the ionizing background.
 luterestiuglv. we find that small changes in the cuissivity at hieh redshifts could have resulted iu much larger changes in the ioniziug backeround.," Interestingly, we find that small changes in the emissivity at high redshifts could have resulted in much larger changes in the ionizing background."
 This finding has interesting nuplicatious for the observed trends in this backerounud., This finding has interesting implications for the observed trends in this background.
 This paper is organized as follows., This paper is organized as follows.
 Section 2. describes the observations aud our nuucerical techniques., Section \ref{sec:thecase} describes the observations and our numerical techniques.
 Section 3 compares our ininerical calculatious of 02A/0:ONY the umuber of absorption svstcms per uuit redshift per ccoluun (Nyy) with observations.," Section \ref{sec:comparison} compares our numerical calculations of $\partial^2{\cal N}/\partial z\,\partial N_{\rm HI}$ – the number of absorption systems per unit redshift per column $N_{\rm HI}$ ) – with observations."
 Section | studies the relationship between dense absorption systelus. the jouizine cnussivity. and the ionizing vackerouncd.," Section \ref{sec:relationship} studies the relationship between dense absorption systems, the ionizing emissivity, and the ionizing background."
 Section 5 considers how plivsical effects that may not be captured in the simulations could alter our conclusions., Section \ref{sec:other} considers how physical effects that may not be captured in the simulations could alter our conclusions.
 Finally. Appendix A xoviles supplementary information reearcding the xoperties of the simulated Lyian-limit systems.," Finally, Appendix \ref{ap:properties} provides supplementary information regarding the properties of the simulated Lyman-limit systems."
 Our calculations assume a flat ACDAL cosmolosv with (QuQn.ηνσε.ο).=(0.28.0.016.0.70.0.82.0.96) ancl a lvdrogen nass fraction of 0.75. consistent with the atest Wilkinson Microwave Anisotropy Probe analysis (?)..," Our calculations assume a flat $\Lambda$ CDM cosmology with $(\Omega_m,\Omega_b,h,\sigma_8,n_s) = (0.28,0.046,0.70,0.82,0.96)$ and a hydrogen mass fraction of $0.75$ , consistent with the latest Wilkinson Microwave Anisotropy Probe analysis \citep{komatsu10}."
derived roni the cross-correlations (see section 6) cau be fou in the Áppeudix. A (table A.l. A.2. A.B. ALL X ALND).,"derived from the cross-correlations (see section 6) can be found in the Appendix A (table A.1, A.2, A.3, A.4 and A.5)."
" The cataloge lists all the sources detected iu all the 303. fields analvzed. ie. we did not associate ""OIECON «etected iu dittvent observations of the same area he sky."," The catalogue lists all the sources detected in all the 4,303 fields analyzed, i.e. we did not associate sources detected in different observations of the same area of the sky."
" Am estimae of the nuuber of indepenudoeut rees m the catalogue can be obtained compressing the uber of sources using an error on the source position 10"" (a conservative init dictated bv the uncertainties conneced wih the boxsight correction).", An estimate of the number of independent sources in the catalogue can be obtained compressing the number of sources using an error on the source position of $10''$ (a conservative limit dictated by the uncertainties connected with the boresight correction).
 We fouud 20.153 epeideut objects.," We found 20,453 independent objects."
 Iu Fig., In Fig.
 1 we plot the distribution of source offaxis aneles showiis the typical trend due to the increase of the coIectingo area together with the decreasing «ft the sensitivity with offaxis., 4 we plot the distribution of source off–axis angles showing the typical trend due to the increase of the collecting area together with the decreasing of the sensitivity with off–axis.
" Clearly. the peak at zero o[axis is due to poiited sources,"," Clearly, the peak at zero off–axis is due to pointed sources."
 Oue of the most importait characteristics of wavelet algoritlas is the ability to determine the source exteusion. Le. the scale of the wavelet ransforiu where the S/N is maximized after the application of the fitting refinement procedure (see Lazzati ct al. 1909).," One of the most important characteristics of wavelet algorithms is the ability to determine the source extension, i.e. the scale of the wavelet transform where the S/N is maximized after the application of the fitting refinement procedure (see Lazzati et al. \cite{lazzati99}) ),"
 aud. if a criterion is eiven. also to discutangle poiit and extended sources (see Campana et al. 19993).," and, if a criterion is given, also to disentangle point and extended sources (see Campana et al. \cite{campana99}) )."
 To assess the source extension criterion we considered al SOTIEECS detecte iu the observatious that have a star(5) as a target (ROR ποτ beeiuning with 2) axd that were available im the public archives iu a preliminary phase of our catalogue: 6.013 sources iu 756 IIRI fields.," To assess the source extension criterion we considered all sources detected in the observations that have a star(s) as a target (ROR number beginning with 2) and that were available in the public archives in a preliminary phase of our catalogue: 6,013 sources in 756 HRI fields."
 Tιο clistribution of the source extension as a fiction of fie source offaxis angele has ceu divided iuto bius of 1: wemln each., The distribution of the source extension as a function of the source off–axis angle has been divided into bins of 1 arcmin each.
 To each bin we jen applied a c cli»piug aleorithiu ο ciscarcl iteratively rulv exteuded sources aud ο derive the mean value of je source extension in the iu for pointed sources., To each bin we then applied a $\sigma-$ clipping algorithm to discard iteratively truly extended sources and to derive the mean value of the source extension in the bin for pointed sources.
 We ien determined the 39 dis]version on the mean for cach iu.," We then determined the $3\,\sigma$ dispersion on the mean for each bin."
 The mean value plus he 36 dispersion provides je threshold for the source extension.," The mean value plus the $3\,\sigma$ dispersion provides the threshold for the source extension."
 We conservativev lassifv a source as extended if it ies 1uore than 20 fro us limit (1.6. if the source csxtension error bar lies more an twice from the 20 limit described above).," We conservatively classify a source as extended if it lies more than $2\,\sigma$ from this limit (i.e. if the source extension error bar lies more than twice from the $3\,\sigma$ limit described above)."
 Combining us threshold with the 36 ou the iutrinsic dispersion. we obtain a Lo coufkence level for the exteusion classification (sce also Rosaict al. 1995)).," Combining this threshold with the $3\,\sigma$ on the intrinsic dispersion, we obtain a $\sim\,4.5\,\sigma$ confidence level for the extension classification (see also Rosati et al. \cite{rosati95}) )."
 In Fie., In Fig.
 5 we alot the distribution of the source extension versus offaxis angle for the 29.089 DÀW-IRI sources (small dots).," 5 we plot the distribution of the source extension versus off--axis angle for the 29,089 BMW-HRI sources (small dots)."
 The solid line in Fig., The solid line in Fig.
 5 represeuts the mean value of source extension for poiute sources as described above. while the dashed line is the 3c dispersion on this nean.," 5 represents the mean value of source extension for pointed sources as described above, while the dashed line is the $3\,\sigma$ dispersion on this mean."
 Open squares in Fie., Open squares in Fig.
 5 repreent sources we classified as extended. ic. with a confence level for the extensiou classification of ~1ωσ.," 5 represent sources we classified as extended, i.e. with a confidence level for the extension classification of $\sim\,4.5\,\sigma$."
 We eud up with 2.717 exended sources (open squares in Fig.," We end up with 2,717 extended sources (open squares in Fig."
 5) containing supernova remnants. galaxies. cluster of galaxies ete; (," 5) containing supernova remnants, galaxies, cluster of galaxies etc. ("
as well as blering of nearby sources).,as well as blending of nearby sources).
" The distributions of source extensious for the truv exteided SOTCCS, he full sai.je aud only the high-ealactic latitude MU| ) one (2.139 sources). are reported in Figο"," The distributions of source extensions for the truly extended sources, the full sample and only the high-galactic latitude $|b| > 20^{\degr}$ ) one (2,139 sources), are reported in Fig."
", ϐ,", 6.
 We assuued that a] the ligrh-ealactic latitude exteided sources are oxtragalactic in natire., We assumed that all the high-galactic latitude extended sources are extragalactic in nature.
 In the plot we also report the onaxis angular resolution of t1ο ROSAT PSPC., In the plot we also report the on–axis angular resolution of the ROSAT PSPC.
 Xrav exte1s1olis are calculated subtractiig du quacrattre the relaive PSF extension at a given offaxis anele (solid. line in Fie., X–ray extensions are calculated subtracting in quadrature the relative PSF extension at a given off–axis angle (solid line in Fig.
 5)., 5).
 The extended sources were used to select a is of candidate Xrav selected cluster of galaxies that we then studie wih optical folkW-Up (Moretti et al., The extended sources were used to select a list of candidate X–ray selected cluster of galaxies that we then studied with optical follow-up (Moretti et al.
 2002 in prepuation: €uzzo et al., \cite{moretti02} in preparation; Guzzo et al.
 209 iu preparation)., \cite{guzzo02} in preparation).
 As the scusitivity of the WRI Πισίναment is not uniforui over the entire field of view. for a eiven Πιο flux the surveved area does not coiucide with he detector oue but it is ecnerally simaller.," As the sensitivity of the HRI instrument is not uniform over the entire field of view, for a given limiting flux the surveyed area does not coincide with the detector one but it is generally smaller."
 We calculated the sky coverage of the cutive survey as a fuuctiou of the flux. (calculated with the full cohuun deusity) by mncaus of simulations., We calculated the sky coverage of the entire survey as a function of the flux (calculated with the full column density) by means of simulations.
 To, To
shifts and asymmetric profiles is an obvious indication that their winds are much more transparent in N-ravs than originally believed.,shifts and asymmetric profiles is an obvious indication that their winds are much more transparent in X-rays than originally believed.
 This result. along with the small filling [actors of the X-ray emitting plasma in OB stars as found in previous AT) studies (Ixudritzkietal.1996).. lends support to a physical picture in which the stellar winds are clumpy and/or porous.," This result, along with the small filling factors of the X-ray emitting plasma in OB stars as found in previous ) studies \citep{ku_96}, lends support to a physical picture in which the stellar winds are clumpy and/or porous."
 Alternatively. if the winds are smooth and homogeneous. the mass-oss rates should have values appreciably smaller (a factor ~5 or more) than those presently accepted (Ixramer.Cohen&Owocki2003: Cohenetal. 2006)).," Alternatively, if the winds are smooth and homogeneous, the mass-loss rates should have values appreciably smaller (a factor $\sim$ 5 or more) than those presently accepted \citealt{kr_03}; \citealt{co_06}) )."
 Wind clumping ancl porosity elfects. on the X-ray emission from OB stars. anc specifically on the shape of the line profiles. have been explored in recent analytical and numerical models (e.g. Feldmeier.Oskinova& 2003:: Oskinova. Feldmoeier Lamann 2004. 2006: Owocki&Cohen 2006)).," Wind clumping and porosity effects on the X-ray emission from OB stars, and specifically on the shape of the line profiles, have been explored in recent analytical and numerical models (e.g., \citealt{feld_03}; Oskinova, Feldmeier Hamann 2004, 2006; \citealt{ow_06}) )."
 It has been shown that under given conditions (such as reduced mass loss and. assumed tvpical! distance between clumps) their inclusion may lead. to results much more Consistent with the observations., It has been shown that under given conditions (such as reduced mass loss and assumed `typical' distance between clumps) their inclusion may lead to results much more consistent with the observations.
 Given the complexity of these models and the fact that they are not vet fully selt-consistent. it is important to gather empirical information about the physical conditions in the regions responsible for the X-ray. emission to put further constraints on numerical nmoclels.," Given the complexity of these models and the fact that they are not yet fully self-consistent, it is important to gather empirical information about the physical conditions in the regions responsible for the X-ray emission to put further constraints on numerical models."
 Various diagnostics have been used in this respect: analvsis of helium-like triplet ratios. global [its with ciscrete-tompcrature models. constructing a distribution of cmussion measures as function. of temperature. based.on fits to individual-lineHuxes and to the total X-ray spectra (Cassinellietal. 2001: Wahnctal.2001: Walelron&Cassinelli 2000: Miller.ctal. 2002: Cohenetal. 2003:: Schulzetal. 2003:: Sanz-Forcada.Franciosini&Pallavicini 2004: Wojdowski Schulz 2004. 2005: Gagnectal.2005:: Leuteneggeretal. 2006)).," Various diagnostics have been used in this respect: analysis of helium-like triplet ratios, global fits with discrete-temperature models, constructing a distribution of emission measures as function of temperature, basedon fits to individual-linefluxes and to the total X-ray spectra \citealt{cass_01}; ; \citealt{kahn_01}; \citealt{wa_00}; \citealt{mi_02}; \citealt{co_03}; \citealt{schu_03}; \citealt{sa_04}; Wojdowski Schulz 2004, 2005; \citealt{ga_05}; \citealt{leu_06}) )."
 These studies indicate that X-rays are produced close to the stellar surface (likely. in the wind acceleration zone). and that the corresponding hot. plasma has a temperature stratification.," These studies indicate that X-rays are produced close to the stellar surface (likely, in the wind acceleration zone), and that the corresponding hot plasma has a temperature stratification."
 The latter point reinforces the idea that the X-ray emission originates in an ensemble of shocks., The latter point reinforces the idea that the X-ray emission originates in an ensemble of shocks.
 Guided by this background. we have developed a simple model which bears all the basic charateristics of the X-ray production in wind shocks.," Guided by this background, we have developed a simple model which bears all the basic charateristics of the X-ray production in wind shocks."
 Ehe model is described in 2: the data sample is given in 3: the results are presented in 4 and diseussed in 5.., The model is described in \ref{sec:mod}; the data sample is given in \ref{sec:obs}; the results are presented in \ref{sec:res} and discussed in \ref{sec:dis}.
 Phe conclusions close the paper., The conclusions close the paper.
 As in the case of the RDL and ALICWS models. our basic assumption is that the X-ray. emission of hot massive stars originates in shocks.," As in the case of the RDI and MCWS models, our basic assumption is that the X-ray emission of hot massive stars originates in shocks."
 Given the relatively high densities in he wind. the energy losses by the shock-heated plasma are considerable: thus. shocks should. beradialive.," Given the relatively high densities in the wind, the energy losses by the shock-heated plasma are considerable: thus, shocks should be."
" ""This conclusion follows from simple estimates which show that he tvpical cooling time of a parcel of σας at the postshock emperature and density issmaller than the typical dynamic ime of the Dow.", This conclusion follows from simple estimates which show that the typical cooling time of a parcel of gas at the postshock temperature and density issmaller than the typical dynamic time of the flow.
" Namely. for the shock position at a distance r [rom the star. the ratio of the cooling time (£. πριν.) o the dvnamic time of the [low (5,= rfe) is: ⋜⋯∠⇂↿↓↕⋖⋅↓⋅⋜∐⊲⊓⋯⇂⋅↿↓∐⊾↿↓↥⊀⊔⇍↓∡⊔⋖⊾⊳∖⊳∖∪⇂⋅↿↓∐⋅↓⋅⋯⇂⊲↓⋜⊔⊀↓∖⇁⋖⋅⊳∖↓↕⋯⇍↳↿∖⊀⊔⊾↦ the cooling length of a parcel of hot eas at the postshock temperature: f=Logd, ) and shock ""radius! is: where 72, is the postshock temperature given in keV. CQouo is the terminal stellar wind. velocity. (in units of 9). and Aly is the stellar niass loss (in units of 1O""MA. P "," Namely, for the shock position at a distance $r$ from the star, the ratio of the cooling time $t_c = \frac{5}{2}p_{sh}/Q_c$ ) to the dynamic time of the flow $t_d = r/v$ ) is: and the ratio of the thickness of the radiative shock (i.e., the cooling length of a parcel of hot gas at the postshock temperature: $l_c = \frac{1}{4} v_{sh} t_c$ ) and shock `radius' is: where $T_{sh}$ is the postshock temperature given in keV, $v_{1000}$ is the terminal stellar wind velocity (in units of ), and $\dot{M}_6$ is the stellar mass loss (in units of $10^{-6}$ $_\odot$ $^{-1}$ )."
"Fora strong shock. the relation between the postshock temperature (in keV) and the shock velocity (in units of )) is given by Zi,=1.9565/02,,,4. and pois the mean atomic weight."," For a strong shock, the relation between the postshock temperature (in keV) and the shock velocity (in units of ) is given by $T_{sh} = 1.956\mu v_{shock}^2$, and $\mu$ is the mean atomic weight."
 The relative number density of hydrogen is assumed rg=0.9: thus. the relative electron number density is or.=1.1 for a fully ionized. plasma.," The relative number density of hydrogen is assumed $x_H = 0.9$; thus, the relative electron number density is $x_e = 1.1$ for a fully ionized plasma."
 The gascvnamical quantities and the cooling function. are described in 2.1.., The gasdynamical quantities and the cooling function are described in \ref{subsec:shock}.
 For typical mass-loss rates (AlyxsOL1) anc wind velocities (P100071.5. 2.5). the cooling time ancl cooling ength of a shock developed in a wind. are appreciably smaller than the corresponding twpical characteristics of he flow (see eqs. 1]. 2]].," For typical mass-loss rates $\dot{M}_6 \approx 0.1 - 1$ ) and wind velocities $v_{1000} \approx 1.5 - 2.5$ ), the cooling time and cooling length of a shock developed in a wind are appreciably smaller than the corresponding typical characteristics of the flow (see eqs. \ref{eq:tcool}] ], \ref{eq:lcool}] ])."
" This is true for postshock emperatures below 1 keV. ναι, shock velocities smaller than for solar abundances) ancl shock locations rom a lew to several tens of stellar radii."," This is true for postshock temperatures below 1 keV (i.e., shock velocities smaller than for solar abundances) and shock locations from a few to several tens of stellar radii."
 VPherelore. the asstunption of a steady-state. plane-parallel raciative shock is a good approximation for our analysis.," Therefore, the assumption of a steady-state, plane-parallel, radiative shock is a good approximation for our analysis."
 This conclusion inds support in numerical simulations of both the RDI and ALOWS models., This conclusion finds support in numerical simulations of both the RDI and MCWS models.
 In fact. as à result of ellicient cooling. the shocks ‘collapse’ in geometrically thin shells and. clisk-like structures. respectively (e.g. Feldimeicr.Puls&Pauldrach 1997:: ud-Doula&Owocki 2002: Gagneetal. 2005)).," In fact, as a result of efficient cooling, the shocks `collapse' in geometrically thin shells and disk-like structures, respectively (e.g. \citealt{feld_97}; \citealt{ud_02}; \citealt{ga_05}) )."
 The description of our shock mocdel and the ensemble of shocks is given below., The description of our shock model and the ensemble of shocks is given below.
 The models were then used. in he recent version (11.3.2). of the software. package for analysis of X-ray spectra. (Arnaucl1996).," The models were then used in the recent version (11.3.2) of the software package for analysis of X-ray spectra, \citep{a_96}."
.. A elobal-it approach was adopted in our analysis for the following reasons: (i) the fit can automatically take into account he quasi-continuum due to numerous weak lines: (ii) by itting the shape of the underlying continuum. the mocel laces: additional constraints on the plasma temperature: (ii) the model can constrain the column clensitw of the X-rav absorbing eas: and (iv) it can vield estimates of relative clement abuncances.," A global-fit approach was adopted in our analysis for the following reasons: (i) the fit can automatically take into account the quasi-continuum due to numerous weak lines; (ii) by fitting the shape of the underlying continuum, the model places additional constraints on the plasma temperature; (iii) the model can constrain the column density of the X-ray absorbing gas; and (iv) it can yield estimates of relative element abundances."
" Finally. the X-ray emission is assumed o originate [rom a hot optically-thin plasma in collisional ionization equilibrium (CLE). as is the case for various types of astrophysical objects. including hot massive stars (e.g. ""uerels&Ixahn 2006))."," Finally, the X-ray emission is assumed to originate from a hot optically-thin plasma in collisional ionization equilibrium (CIE), as is the case for various types of astrophysical objects, including hot massive stars (e.g., \citealt{pa_03}) )."
 We consider a steady-state. plane-parallel. shock moving in a gas with adiabatic index ~= 5/3.," We consider a steady-state, plane-parallel, shock moving in a gas with adiabatic index $\gamma = 5/3$ ."
" The basic phwsical quantities of the Dow (density or nucleon number density. velocity. ancl pressure) have their standard postshock values for a strong shock: pi,=Apo (n.i,= dno). Pu,= ca/4. and po,= 3/4por5. where the subscript7Q denotes the preshock values (the gas velocity is given in the rest [rame of the shock front)."," The basic physical quantities of the flow (density or nucleon number density, velocity, and pressure) have their standard postshock values for a strong shock: $\rho_{sh} = 4\rho_0$ $n_{sh} = 4n_0$ ), $v_{sh} = v_0/4$ , and $p_{sh} = 3/4\rho_0v_0^2$ , where the subscript`0' denotes the preshock values (the gas velocity is given in the rest frame of the shock front)."
 The mass and. momentum, The mass and momentum
For a given surlace brightness profile twpe. there exists a relation between. absolute magnitude and apparent axis ratio.,"For a given surface brightness profile type, there exists a relation between absolute magnitude and apparent axis ratio."
 Contour plots of mean apparent axis ratio as a function ol both and AZ. are given in Figure 2.. (, Contour plots of mean apparent axis ratio as a function of both and $M_r$ are given in Figure \ref{fig:fdevmag}. (
"To give a feel for the absolute magnitude scale. fitting a Schechter function to the luminosity function of SDSS galaxies vields 3,,sz—21.4 (Nakamuraetal. 2003)..)","To give a feel for the absolute magnitude scale, fitting a Schechter function to the luminosity function of SDSS galaxies yields $M_{*,r} \approx -21.4$ \citep{na03}. .)"
" The upper panel of the figure shows that the trend in (q,,,) runs from the flattest galaxies at fracDeV©0 and M,z—18. where (qu)20.52. to the roundest galaxies al fracDeVzz| and M,ez—23. where (qu)7&0.83."," The upper panel of the figure shows that the trend in $\langle q_{\rm am} \rangle$ runs from the flattest galaxies at $\texttt{fracDeV} \approx 0$ and $M_r \approx -18$, where $\langle q_{\rm am} \rangle \approx 0.52$, to the roundest galaxies at $\texttt{fracDeV} \approx 1$ and $M_r \approx -23$, where $\langle q_{\rm am} \rangle \approx 0.83$."
 This result is not surprising. since moderatelv bright galaxies with exponential profiles are intrinsically fattened clisk ealaxies. while extremely bright galaxies with de Vaucouleurs profiles are intrinsically nearly spherical eiant elliptical galaxies (Tremblay&Merritt1996)..More surprising are (he results shown in the bottom panel of Figure 2.. which shows (02). the mean isophotal axis ratio.," This result is not surprising, since moderately bright galaxies with exponential profiles are intrinsically flattened disk galaxies, while extremely bright galaxies with de Vaucouleurs profiles are intrinsically nearly spherical giant elliptical galaxies \citep{tm96}.More surprising are the results shown in the bottom panel of Figure \ref{fig:fdevmag}, which shows $\langle q_{25} \rangle$, the mean isophotal axis ratio."
" ]lere. we find that the apparently [attest galaxies. as nieasured by qos. are nol exponential ealaxies. bul galaxies wilh fracDeVzz0.7 (corresponding (o Sérrsic index nx2.5) and M,£g—20.5: these galaxies have (qo3)220.53."," Here, we find that the apparently flattest galaxies, as measured by $q_{25}$, are not exponential galaxies, but galaxies with $\texttt{fracDeV} \approx 0.7$ (corresponding to Sérrsic index $n \approx 2.5$ ) and $M_r \approx -20.5$; these galaxies have $\langle q_{25} \rangle \approx 0.53$."
" The roundest galaxies. measured by (055). are nol bright de Vaucouleurs galaxies. but bright exponential galaxies: the masxinnun value of (qos) is 8Ο.Τ. at fracDeVzz0. M,zz—22.5."," The roundest galaxies, measured by $\langle q_{25} \rangle$, are not bright de Vaucouleurs galaxies, but bright exponential galaxies; the maximum value of $\langle q_{25} \rangle$ is $\approx 0.74$, at $\texttt{fracDeV} \approx 0$, $M_r \approx -22.5$."
" Note also that [or bright galaxies (M,S —21). the contours of constant (001) in Figure 2) are nearly horizontal."," Note also that for bright galaxies $M_r \la -21$ ), the contours of constant $\langle q_{25} \rangle$ in Figure \ref{fig:fdevmag}
 are nearly horizontal."
 That is. among bright galaxies. the flattening of the outer isophotes doesnt depend strongly on the surface briehtuess profile (wpe.," That is, among bright galaxies, the flattening of the outer isophotes doesn't depend strongly on the surface brightness profile type."
 A view of the dependence of (q) on absolute magnitude for each of our four profile types. ον. ‘ex/de’. de/ex'. and ‘ce’. is given in Figure 3..," A view of the dependence of $\langle q \rangle$ on absolute magnitude for each of our four profile types, `ex', `ex/de', 'de/ex', and `de', is given in Figure \ref{fig:mag}."
 In the upper panel. which shows (qa). note that for each prolile type. there is a critical absolute magnitude Mag at which(qu) is al a minimum.," In the upper panel, which shows $\langle q_{\rm am} \rangle$ , note that for each profile type, there is a critical absolute magnitude $M_{\rm crit}$ at which$\langle q_{\rm am} \rangle$ is at a minimum."
" This οσα] absolute magnitude ranges from M4;~—20.6 for the ‘de’ galaxies to Magc—19.4 for the ""ex! galaxies.", This critical absolute magnitude ranges from $M_{\rm crit} \sim -20.6$ for the `de' galaxies to $M_{\rm crit} \sim -19.4$ for the `ex' galaxies.
" At AL.<Mag. the value of (gun) increases relatively rapidly with increasing luminosity: at M,>Mig. the value of (gy) increases less rapidly with decreasing Iuminositv."," At $M_r < M_{\rm crit}$, the value of $\langle q_{\rm am} \rangle$ increases relatively rapidly with increasing luminosity; at $M_r > M_{\rm crit}$, the value of $\langle q_{\rm am} \rangle$ increases less rapidly with decreasing luminosity."
" At a fixed absolute magnitude. the average axis ratio of ‘cle’ galaxies is always greater than (hat of ex’ galaxies: however. for M,<<—20. the flattest ealaxies. on average. al a given absolute magnitude are not (he ‘ex galaxies. butthose with the mixed ‘de/ex’ and ‘ex/de’ prolile types."," At a fixed absolute magnitude, the average axis ratio of `de' galaxies is always greater than that of `ex' galaxies; however, for $M_r \la -20$, the flattest galaxies, on average, at a given absolute magnitude are not the `ex' galaxies, butthose with the mixed `de/ex' and `ex/de' profile types."
" The bottom panel of Figure 3. shows (q@5) versus AM, for the different profile types.", The bottom panel of Figure \ref{fig:mag} shows $\langle q_{25} \rangle$ versus $M_r$ for the different profile types.
" In the interval —20<V,—22. we see that ‘ex’ galaxies have values of (q55) than galaxies of other profile types."," In the interval $-20 \la M_r \la -22$, we see that `ex' galaxies have values of $\langle
q_{25} \rangle$ than galaxies of other profile types."
 Tremblay&Merritt(1996). divided elliptical galaxies into (wo classes: galaxies brighter than Aj8—20 are rounder on average than [anter ellipticals., \citet{tm96} divided elliptical galaxies into two classes; galaxies brighter than $M_B \approx -20$ are rounder on average than fainter ellipticals.
" Using a tvpical color for elliplical galaxiesof 5—V.zz 0.9. this corresponds (to an absolute magnitude in the r band ol M,zz —21.2. using the transformation r=2—1.4408V)20.12 (Smithetal. 2002).."," Using a typical color for elliptical galaxiesof $\bv \approx 0.9$ , this corresponds to an absolute magnitude in the $r$ band of $M_r \approx -21.2$ , using the transformation $r = B - 1.44 (B-V) + 0.12$ \citep{sm02}. ."
 In, In
luterestinglw. both x-ray sources are transient and have been much brighter iu the past 6)).,"Interestingly, both x-ray sources are transient and have been much brighter in the past \ref{s:temporal}) )."
" No known backeround or foreground objects were found coincident with the $3 ssources,", No known background or foreground objects were found coincident with the S3 sources.
 Up to ~25 backerouncd objects are expected im he $3 field above S/N =3.5 based upon our analysis of he calibration Ποια CRSS J0030.5|2618 (see also Draudt 22000)., Up to $\sim$ 25 background objects are expected in the S3 field above S/N $=3.5$ based upon our analysis of the calibration field CRSS J0030.5+2618 (see also Brandt 2000).
 This estimate ignores anv obscurtion bv MSI., This estimate ignores any obscurtion by M81.
 Based on results of the stellar survey (Topka 11982). we estimate roughly 0.1 F- and C-stars could be detectable in our field but these should all be extremely soft (Alageio 11987. Todekin Pye 1991. Scliunitt 1997).," Based on results of the stellar survey (Topka 1982), we estimate roughly 0.4 F- and G-stars could be detectable in our field but these should all be extremely soft (Maggio 1987, Hodgkin Pye 1994, Schmitt 1997)."
 Α soft spectrum was observed οι 3 catalogued stars detected ou CCD S2., A soft spectrum was observed from 3 catalogued stars detected on CCD S2.
 The total source counts in the hard xay baud (2.0. 8.0 keV) is shown in figure 2 against the soft (0.2 2.0 keV) baud for the subset of sources with S/N223.5 excluding the uucleus., The total source counts in the hard x-ray band (2.0 – 8.0 keV) is shown in figure 2 against the soft (0.2 – 2.0 keV) band for the subset of sources with $\ge$ 3.5 excluding the nucleus.
 There is no adjustineut for pileup which tends to harden spectra of the brightest sources., There is no adjustment for pileup which tends to harden spectra of the brightest sources.
 Also shown are the hardness ratios of the stun of all the bulge sources. and of all the disk sources. the harduess ratio ofthe excess bulge emission (backeround aud unclear PSF subtracted). aud of the typical background (scaled to a 2« area).," Also shown are the hardness ratios of the sum of all the bulge sources, and of all the disk sources, the hardness ratio of the excess bulge emission (background and nuclear PSF subtracted), and of the typical background (scaled to a $2\arcmin \times 2\arcmin$ area)."
 The nucleus aud the bright soft source are olutted from the summed bulee poiut. aud N-6 is omitted from the sununaed disk poiut as these sources are clearly spectroscopically atypical.," The nucleus and the bright soft source are omitted from the summed bulge point, and X-6 is omitted from the summed disk point as these sources are clearly spectroscopically atypical."
 The line shown corresponds to an absorbed power law of spectral iudex D=1.6. typical of x-ray binaries. with the Galactic absorbing cola deusity of Ny=bx102? cm? (Stark 11992).," The line shown corresponds to an absorbed power law of spectral index $\Gamma = 1.6$, typical of x-ray binaries, with the Galactic absorbing column density of $N_H = 4 \times 10^{20}$ $^{-2}$ (Stark 1992)."
 Though iudicative aud useful to euide the eve. this canonical spectral shape is not a model fit result.," Though indicative and useful to guide the eye, this canonical spectral shape is not a model fit result."
 Most sources fall near this curve., Most sources fall near this curve.
 Notable exceptions are the soft sources at the lower right of figure 2., Notable exceptions are the soft sources at the lower right of figure 2.
 The softest bulee sources have no known associations but the 3 softest disk sources all lie on spiral axis., The softest bulge sources have no known associations but the 3 softest disk sources all lie on spiral arms.
 The softest source is also tinme-variable 63) aud is the third brightest source in the entire sample., The softest source is also time-variable \ref{s:temporal}) ) and is the third brightest source in the entire sample.
 The harduess ratios for the sunmuued bulee and sumuued disk spectra are also typical of x-rav binarics., The hardness ratios for the summed bulge and summed disk spectra are also typical of x-ray binaries.
 The excess bulge cinissiou is softer indicating some of the excess chussion is from truly diffuse. hot. iuterstellar gas and not cutirely from unresolved point sources.," The excess bulge emission is softer indicating some of the excess emission is from truly diffuse, hot, interstellar gas and not entirely from unresolved point sources."
 The ummber of sources above a lanitine brieltucss. NSS). is shown against the 0.2. δ keV. count rate. S. iu figure 3.," The number of sources above a limiting brightness, $N(>$$S)$, is shown against the 0.2 – 8.0 keV count rate, $S$, in figure 3."
 The nucleus is excluded as are sources with S/N<3.5., The nucleus is excluded as are sources with $<$ 3.5.
 Backerouncl sources were taken from the calibration field CRSS J0030.5|2618 3)). by applying our sale analysis methods. aud have been subtracted after scaling to the fractional areas of the bulee and disk.," Background sources were taken from the calibration field CRSS J0030.5+2618 \ref{s:spatial}) ), by applying our same analysis methods, and have been subtracted after scaling to the fractional areas of the bulge and disk."
" The curve shown in figure 3 represcuts the best power law fit to the disk distribution, Vo=0.135 """", over the entire range of S,"," The curve shown in figure 3 represents the best power law fit to the disk distribution, $N =0.43S^{-0.50}$ , over the entire range of $S$."
" Similarly, N=0.119"""" for the bulee source distribution for S-«0.001 1."," Similarly, $N = 0.44 S^{-0.57}$ for the bulge source distribution for $S$$<$ 0.004 $^{-1}$."
 There is a break iu the slope of the distribution of bulge sources above this point but no such break iu the disk source distribution., There is a break in the slope of the distribution of bulge sources above this point but no such break in the disk source distribution.
 Of the LO disk sources with 570.001 s1. 7 ave coincident with spiral arius.," Of the 10 disk sources with $S$$>$ 0.004 $^{-1}$, 7 are coincident with spiral arms."
 The huuinositv of cach source can be estimated by assunune the P—1.6 power law spectral model of and a distance of 3.6 AIpe to M81 (Freecanan 11991)., The luminosity of each source can be estimated by assuming the $\Gamma = 1.6$ power law spectral model of \ref{s:spectral} and a distance of 3.6 Mpc to M81 (Freedman 1994).
" For our 50-ks observation. therefore. S=0.001 3 in the 0.2 — 8.0 keV baud corresponds to a huuinosity Ly=3.04107 cress | ane the faintest source iu the field corresponds to a luminosity of ~3«1076 eres st,"," For our 50-ks observation, therefore, $S=0.004$ $^{-1}$ in the 0.2 – 8.0 keV band corresponds to a luminosity $L_X = 3.7 \times 10^{37}$ ergs $^{-1}$ and the faintest source in the field corresponds to a luminosity of $\sim 3 \times 10^{36}$ ergs $^{-1}$."
 Excluding the nucleus. the total bulee luuinosity is Ly~2.1.10 eres bof which is excess (unresolved) cluission.," Excluding the nucleus, the total bulge luminosity is $L_X \sim 2.4 \times 10^{39}$ erg $^{-1}$ of which is excess (unresolved) emission."
 The total hmuuimositv of the disk sources is Ly~3.9&1079 cre 1., The total luminosity of the disk sources is $L_X \sim 3.9 \times 10^{39}$ erg $^{-1}$.
 The bluniuositv of the nucleus is Ly~Ls10! cre ft based ou spectral fits to the trailed image., The luminosity of the nucleus is $L_X \sim 4 \times 10^{40}$ erg $^{-1}$ based on spectral fits to the trailed image.
 This is within the ASCA-obscrved range of huninositics (Ishisali 11996) and comparable to the BeppoSAX observed huninosity (Pellegrini 22000)., This is within the ASCA-observed range of luminosities (Ishisaki 1996) and comparable to the BeppoSAX observed luminosity (Pellegrini 2000).
 The uucleus coutributes approximately of the total luminosity in the 87.3.8:3 S3 field of view., The nucleus contributes approximately of the total luminosity in the $8\arcmin.3 \times 8\arcmin.3$ S3 field of view.
 A simple test for time viuiabilitv was made by binning the light curves of all sources on 1000. 2000. and 000 second intervals.," A simple test for time variability was made by binning the light curves of all sources on 1000, 2000, and 4000 second intervals."
 Applying a 4? test for the lypothesis of consistency with a mean value found only oue. clearly siguificant deviation ou all three timescales., Applying a $\chi^2$ test for the hypothesis of consistency with a mean value found only one clearly significant deviation on all three timescales.
 This source is also the softest source oi $3 aud the third brightest., This source is also the softest source on S3 and the third brightest.
 It is present iu at least 3 of 6 IIIRI observations spauuiug 1993 - 1998 but is too close to the nucleus to be identified in amy other previous x-ray observation., It is present in at least 3 of 6 HRI observations spanning 1993 - 1998 but is too close to the nucleus to be identified in any other previous x-ray observation.
 ΑΙΣ1 has been observed in x-ravs at moderate spatial resolution by in 1979 and bv oover the period 19911998., M81 has been observed in x-rays at moderate spatial resolution by in 1979 and by over the period 1991–1998.
 There are 6 ssources (E88) within the $3 field of view and another 1 ROSAT—detected sources according to our analysis., There are 6 sources (F88) within the S3 field of view and another 4 -detected sources according to our analysis.
 rresolves six ssource regions into two or more sources inchiding ssources X-7 and N-1O. and. of course. the nucleus.," resolves six source regions into two or more sources including sources X-7 and X-10, and, of course, the nucleus."
 Long-term variability has been detected in the vvariahle source. the nucleus. N-2 and the ROSAT—detected source comcideut with au uudocumoeuted star-like object 3)).," Long-term variability has been detected in the variable source, the nucleus, X-2 and the -detected source coincident with an undocumented star-like object \ref{s:spatial}) )."
 Both of the latter sources are moderately weak in the preseut observations (~3 and ~6<107 sto respectively) but have been much brighter in the past.," Both of the latter sources are moderately weak in the present observations $\sim$ 3 and $\sim 6 \times 10^{37}$ $^{-1}$, respectively) but have been much brighter in the past."
" The location of another ssource, X-12. places it on the eastern edge of 53."," The location of another source, X-12, places it on the eastern edge of S3."
" There are no ssources Within the portion of the 15"" error circle of this IPC source that falls on $3 andit is not detectable iu the oobservations.", There are no sources within the portion of the $\arcsec$ error circle of this IPC source that falls on S3 andit is not detectable in the observations.
 The global properties of the ssources ideutified in the ceutral 8/23:&8:3 region of M8SI have been examined., The global properties of the sources identified in the central $8\arcmin.3 \times 8\arcmin.3$ region of M81 have been examined.
 There is a high density of sources, There is a high density of sources
lis veal. it is probably caused by the lower starlight dilution at that wavelength. aud so does not represcut an actual increase m polarized fiux.,"is real, it is probably caused by the lower starlight dilution at that wavelength, and so does not represent an actual increase in polarized flux."
 The RIAF candidate shows siguificaut continuum volarization. while the diluted candidate does not.," The RIAF candidate shows significant continuum polarization, while the diluted candidate does not."
 The uean continua polarization of 0958[9|013220 is 0.784LOT at oan anele of LOLκ degrees., The mean continuum polarization of 095849+013220 is $0.78 \pm 0.07\%$ at an angle of $104 \pm 4\%$ degrees.
 The mean continui polarization m the blue. measured at L100.< ΡΟΡΟΔ.. is even stronger: 1.37c0.16%.," The mean continuum polarization in the blue, measured at $4400<\lambda<5050$ , is even stronger: $1.37 \pm 0.16\%$."
 Subtracting he host galaxy coniponent. which contributes ~h55% of the total enission at Aya~SOOOA.. sugecsts that he mean poluization of the ACN component is —1.74 overall aud ~3.0% in the blue.," Subtracting the host galaxy component, which contributes $\sim$ of the total emission at $\lambda_{\rm rest} \sim 5000$, suggests that the mean polarization of the AGN component is $\sim$ overall and $\sim$ in the blue."
 The polarized coutimmn roughly doubles fom ~6500A to ~I500À.. consistent with a typical uunobseured quasar coutinmun of fy~A19 (VandenBerketal.2001).," The polarized continuum roughly doubles from $\sim$ to $\sim$, consistent with a typical unobscured quasar continuum of $f_{\lambda} \sim \lambda^{-1.6}$ \citep{van01}."
. Iu addition to the two optically dull AGNs from COSMOS. we observed the cluster BL Lac candidate MS 1155-N2 from Hartetal.(2009) and. discuss. its spectropolarimetry results in the Appendix.," In addition to the two optically dull AGNs from COSMOS, we observed the cluster BL Lac candidate MS 1455-X2 from \citet{har09} and discuss its spectropolarimetry results in the Appendix."
 Neither of our targets show evidence for polarized broad or narrow cussion lines. which sugeests that the lues are not missus due to anisotropic absorption as in a standard AGN “torus” model (Autonucci1993).," Neither of our targets show evidence for polarized broad or narrow emission lines, which suggests that the lines are not missing due to anisotropic absorption as in a standard AGN “torus” model \citep{ant93}."
. Iustead. our RIAF candidate has a blue polarized contimmun.," Instead, our RIAF candidate has a blue polarized continuum."
 Here we place linits on polarized ciission liue flux and investigate the two possible causes of the coutiuumun polarization: svuchrotron cussion or scatterie., Here we place limits on polarized emission line flux and investigate the two possible causes of the continuum polarization: synchrotron emission or scattering.
 For both objects the mean polarization in the waveleugth regious of oor bbroad lines are below or consistent with the measured contimmuu polarization iu the surrounding regions., For both objects the mean polarization in the wavelength regions of or broad lines are below or consistent with the measured continuum polarization in the surrounding regions.
 However a nou-detectioun does not necessarily mean that there are no polarized enüssion lines. since very weak lines might be undetected within our errors.," However a non-detection does not necessarily mean that there are no polarized emission lines, since very weak lines might be undetected within our errors."
 We approximate each line as a 2000 Em/s wide top hat. aud estimate upper limits ou the degree of polarization for ad bbroad lines above the polarized coutimuun using the Jc errors across the line region.," We approximate each line as a 2000 km/s wide top hat, and estimate upper limits on the degree of polarization for and broad lines above the polarized continuum using the $3\sigma$ errors across the line region."
 For the RIAF cauclicdate 095819|013220 the 30 upper limits are 2.0% for aand 2.1% forΠο and for the diluted AGN 100036|021929 the30 upper lanits are 1.7% for aand 3. for ον.," For the RIAF candidate 095849+013220 the $3\sigma$ upper limits are $2.0\%$ for and $2.1\%$ for, and for the diluted AGN 100036+024929 the$3\sigma$ upper limits are $1.7\%$ for and $3.3\%$ for ."
" These upper limits decrease for broad lines wider than 2000 αιέν, Ta contrast. “lidden-BLR” ÁACNs typically have reflected broad emission Hines detected at P25% (Antonucci1993:Barthetal.1999:Aloyanetal.2000)."," These upper limits decrease for broad lines wider than 2000 km/s. In contrast, ``hidden-BLR'' AGNs typically have reflected broad emission lines detected at $P \simeq
5\%$ \citep{ant93,bar99,mor00}."
. The iutrinsic fuses from both broad dines and confinmun are presumably represented in polarized cussion., The intrinsic fluxes from both broad lines and continuum are presumably represented in polarized emission.
 For the RIAF candidate 095819|013220 we detect a polarized contiuuuu. aud can determine upper limits in equivalent width (EW) for broad lines in this intrinsic ACN eiission.," For the RIAF candidate 095849+013220 we detect a polarized continuum, and can determine upper limits in equivalent width (EW) for broad lines in this intrinsic AGN emission."
" The 36 upper limits on EW in the intrinsic cussion are EWqp,<SOÀ aand EWqp;ΓΑ.", The $3\sigma$ upper limits on EW in the intrinsic emission are ${\rm EW}_{\Ha}<80$ and ${\rm EW}_{\Hb}<47$.
" These Tanits are less than those of typical quasars. which have EWqp;~50A aand EWp,~200A citepvautl.."," These limits are less than those of typical quasars, which have ${\rm EW}_{\Hb}
\sim 50$ and ${\rm EW}_{\Ha} \sim 200$ \\citep{van01}."
 Au anisotropic scatterer (e.g... the canonical AGN “torus) is not the reason our two sources lack broad cussion lines.," An anisotropic scatterer (e.g., the canonical AGN “torus”) is not the reason our two sources lack broad emission lines."
 Tustead both ACNs have weaker broad lines than typical Type 1 or hidden-BLR ACNs. consistent with a RIAF or lost galaxy starlight dilution.," Instead both AGNs have weaker broad lines than typical Type 1 or hidden-BLR AGNs, consistent with a RIAF or host galaxy starlight dilution."
 Svuchrotron emission from the anisotropic maguctic ficld associated with a radio jet has a hieh deeree of linear olarization., Synchrotron emission from the anisotropic magnetic field associated with a radio jet has a high degree of linear polarization.
 The RIAF candidate 095819|013220 could ο analogous to DL Lacertac ACGNs. which similarly ack both direct aud reflected emission lines.," The RIAF candidate 095849+013220 could be analogous to BL Lacertae AGNs, which similarly lack both direct and reflected emission lines."
 BL Lacs vpicallv exhibit linear polarization at ~2.114. witha uediau of στ (CJannuuzictal.1994:Capettiet2007).," BL Lacs typically exhibit linear polarization at $\sim2-11\%$, with a median of $\sim$ \citep{jan94,cap07}."
.. At first elance the ~L.7% polarization for the ACN-only component of 195819|013220. as in the low ranee of measured volarization for DL Lacs., At first glance the $\sim$ polarization for the AGN-only component of 095849+013220 is in the low range of measured polarization for BL Lacs.
 Ποπονοα the starlight dilution in 095819)013220 is not well coustrained. and the polarization of the AGN compoucut iuight o0 xienificautlv hieher.," However the starlight dilution in 095849+013220 is not well constrained, and the polarization of the AGN component might be significantly higher."
 Sxuchrotron cussion should extend from radio wavelengths to the optical/UV as a yower-law. and so the resultant polarized fix should ollow the same power-law (e.g.Cohenetal.1997).," Synchrotron emission should extend from radio wavelengths to the optical/UV as a power-law, and so the resultant polarized flux should follow the same power-law \citep[e.g.,][]{coh97}."
. The observed increase of polarized flux. towards blue wavelengths for 095819]013220 seems tfo coutradict lis. but could again be explained by dilution from an old. red stellar population.," The observed increase of polarized flux towards blue wavelengths for 095849+013220 seems to contradict this, but could again be explained by dilution from an old, red stellar population."
 The svuchrotron polarization of DL Lacs is also strongly variable in amplitude aud anele (.Januuzictal.1991:Cohenet 1997).. but we do not have multi-cpochpolarization measurements aud," The synchrotron polarization of BL Lacs is also strongly variable in amplitude and angle \citep{jan94,coh97}, , but we do not have multi-epochpolarization measurements and"
Significant in our models with higher mass-loss rates.,significant in our models with higher mass-loss rates.
 Reliable predictions for these phenomena in non-spherical outflows requires multi-dimensional radiative transfer and thus our methods are particularly well-suited to describing them., Reliable predictions for these phenomena in non-spherical outflows requires multi-dimensional radiative transfer and thus our methods are particularly well-suited to describing them.
 To illustrate the types of line profile. Fig.," To illustrate the types of line profile, Fig."
 12. shows a montage of spectra from models with relatively high mass-loss rates., \ref{fig_emiss} shows a montage of spectra from models with relatively high mass-loss rates.
" There is a wide range of Ka profile types in these models and their emission EWs, measured relative to a power-law continuum fit. can be in excess of 200 eV: At low inclination angles. our model Fe Ka profiles have a near-P-Cygni character: they shows some blueshifted absorption and broad. redshifted emission. the strength of which grows with AL."," There is a wide range of $\alpha$ profile types in these models and their emission EWs, measured relative to a power-law continuum fit, can be in excess of 200 eV; At low inclination angles, our model Fe $\alpha$ profiles have a near-P-Cygni character: they shows some blueshifted absorption and broad, redshifted emission, the strength of which grows with $\dot{M}$."
 Such protiles are qualitatively similar to those observed in some Seyfert | galaxies (2).., Such profiles are qualitatively similar to those observed in some Seyfert 1 galaxies \citep{done07}.
 At intermediate angles. when the line-of-sight is close to looking down the wind cone. the narrow absorption lines are strongest but are accompanied by moderately strong. emission lines with fairly extended red wings.," At intermediate angles, when the line-of-sight is close to looking down the wind cone, the narrow absorption lines are strongest but are accompanied by moderately strong, emission lines with fairly extended red wings."
 We note that the narrow absorption lines become less prominent at very high M.a consequence of the increased contribution of scattered radiation in the spectra.," We note that the narrow absorption lines become less prominent at very high $\dot{M}$, a consequence of the increased contribution of scattered radiation in the spectra."
 For the highest inclinations and AL -values. the Fe Ka line is almost pure emission. peaking at the rest energy of the line and having a long red tail.," For the highest inclinations and $\dot{M}$ -values, the Fe $\alpha$ line is almost pure emission, peaking at the rest energy of the line and having a long red tail."
 The redward extension of the line emission is not a consequence of gravitational redshift but arises from electron scattering in the outflow., The redward extension of the line emission is not a consequence of gravitational redshift but arises from electron scattering in the outflow.
 These emission line properties show that highly ionized outflow may affect the Fe K region beyond imprinting narrow. blueshifted absorption lines.," These emission line properties show that highly ionized outflow may affect the Fe K region beyond imprinting narrow, blueshifted absorption lines."
 This may have consequences for study of the Fe K fluorescent emission which originates in AGN accretion disks (e.g. 2222?) since it may contaminate the disk emission and/or lead to an apparently multi-component emission line (see e.g. 2)).," This may have consequences for study of the Fe K fluorescent emission which originates in AGN accretion disks (e.g. \citealt{fabian89,nandra97,fabian00,jmiller07,nandra07}) ), since it may contaminate the disk emission and/or lead to an apparently multi-component emission line (see e.g. \citealt{oneill07}) )."
 In this context. the red wings predicted for the emission lines may be of particular relevance in view of the potential for confusion with the effects of gravitational redshift — however. more sophisticated 3D models going beyond the smooth-flow assumption of our parameterised wind models must be examined before such a possibility can be considered in much greater detail.," In this context, the red wings predicted for the emission lines may be of particular relevance in view of the potential for confusion with the effects of gravitational redshift – however, more sophisticated 3D models going beyond the smooth-flow assumption of our parameterised wind models must be examined before such a possibility can be considered in much greater detail."
 Absorption by outflowing material has been discussed as a possible explanation for the so-called soft-excess in X-ray spectra (22222).," Absorption by outflowing material has been discussed as a possible explanation for the so-called soft-excess in X-ray spectra \citep{gierlinski04,gierlinski06,middleton07,schurch07,schurch08}."
 In this picture. the decrease in flux above ~1 keV is attributed to absorption by light or intermediate mass elements.," In this picture, the decrease in flux above $\sim 1$ keV is attributed to absorption by light or intermediate mass elements."
 Such absorption does occur in our models. particularly. for vigh Al-values (see Fig. 125 ," Such absorption does occur in our models, particularly for high $\dot{M}$ -values (see Fig. \ref{fig_emiss}) )"
however the scale of the effect is oo small in the models presented here: the typical observed soft excess requires a drop in flux of nearly a factor of two between about | and 2 keV (?)., however the scale of the effect is too small in the models presented here: the typical observed soft excess requires a drop in flux of nearly a factor of two between about 1 and 2 keV \citep{middleton07}.
 Furthermore. ? have argued that very arge. relativistic velocities are required for the absorption model to work since the observed soft excess spectra are very smooth.," Furthermore, \citet{schurch08} have argued that very large, relativistic velocities are required for the absorption model to work since the observed soft excess spectra are very smooth."
 Our current models support their conclusions since all cases in which significant spectral curvature arises are accompanied by discreet features., Our current models support their conclusions since all cases in which significant spectral curvature arises are accompanied by discreet features.
 Further consideration must await extension of our calculations to low-ionization states since more significant. absorption at relatively soft energies can then be expected., Further consideration must await extension of our calculations to low-ionization states since more significant absorption at relatively soft energies can then be expected.
 One of the AGN whose X-ray spectrum provided some of the initial motivation for this work is 7766 (22)..," One of the AGN whose X-ray spectrum provided some of the initial motivation for this work is 766 \citep{miller07, turner07}."
 As well as the possible outflow absorption lines already mentioned in refsect os.28nggesledthallheertremcelghard2 10 keVspectruminthelowflieestate:coupledwiththestrongP'ecdgqeandweak ragspeclraof ACN.," As well as the possible outflow absorption lines already mentioned in \\ref{sect_obs}, \citet{miller07} suggested that the extremely hard 2 -- 10 keV spectrum in the low flux state, coupled with the strong Fe edge and weak Fe emission line, might be a signature of either ionised absorption or ionised reflection, possibly from an obscuring wind."
 ineludingM T766(e.g. 7).," Such a picture provides an alternative explanation to the hypothesis of “relativistic blurring” that has been invoked to explain the apparent “red wing” at energies below 7 keV in the mean X-ray spectra of AGN, including 766 \citep[e.g.][]{nandra07}."
. We now make a direct comparison between the wind model in this paper and the mean spectrum of 7766. obtained by summing the data from the 2000. 2001 and 2006 observations.," We now make a direct comparison between the wind model in this paper and the mean spectrum of 766 obtained by summing the data from the 2000, 2001 and 2006 observations."
 Observation IDs were 0096020101. 0096020101 and 0304030[1-7]01. and data were processed as described by (2)..," Observation IDs were 0096020101, 0096020101 and 0304030[1-7]01 and data were processed as described by \citep{miller07}."
 As in that work. energy bins of width equal to the energy-dependent HWHM were adopted.," As in that work, energy bins of width equal to the energy-dependent HWHM were adopted."
 The model was fitted to the data by first creating a multiplicative table (*mtable”) covering a range of parameters that could be used with (2)..., The model was fitted to the data by first creating a multiplicative table (“mtable”) covering a range of parameters that could be used with \citep{arnaud96}.
 The table was used to multiply an assumed intrinsic yower-law whose photon index was a free parameter in the fit: in winciple the fitting is valid only for the choice of power-law index used to create the model. but in practice we tind that there is little variation between models whose power-law indices differ. and that he variation that is seen is to some extent degenerate with the effects caused by varying other wind parameters (see Section 5).," The table was used to multiply an assumed intrinsic power-law whose photon index was a free parameter in the fit: in principle the fitting is valid only for the choice of power-law index used to create the model, but in practice we find that there is little variation between models whose power-law indices differ, and that the variation that is seen is to some extent degenerate with the effects caused by varying other wind parameters (see Section \ref{sect_grid}) )."
 We expect also some amount of spectral curvature caused by the ionised absorption previously identified in 7766. so we add a low column of ionised absorption calculated by (2) and we model the narrow 6.4 FFe emission line as arising in mildly ionised reflection calculated by the model (2). with no relativistic blurring.," We expect also some amount of spectral curvature caused by the ionised absorption previously identified in 766, so we add a low column of ionised absorption calculated by \citep{kallman04} and we model the narrow 6.4 Fe emission line as arising in mildly ionised reflection calculated by the model \citep{ross05} with no relativistic blurring."
 Owing to computing time limitations the model grid was necessarily rather coarse in parameter space. so to check the fit of the best-fitting model. interpolated between models. we then recalculated the model with precisely the parameter values indicated by fitting.," Owing to computing time limitations the model grid was necessarily rather coarse in parameter space, so to check the fit of the best-fitting model, interpolated between models, we then recalculated the model with precisely the parameter values indicated by fitting."
 The wind best-fit parameter were rjiB85 ru. d=1μα. Al= 0.4MM. yyr Land cos@=0.77. while the best-fit photon index was 52.08 and absorber parameters were Ny=5.3107 ? and ionization parameter £=Weergcemss .," The wind best-fit parameter were $r_{\rm min}=385$ $_{\rm g}$, $d = 1.6 r_{\rm min}$, $\dot{M} = 0.4 $ $_\odot$ $^{-1}$ and $\cos\theta=0.77$, while the best-fit photon index was $\Gamma=2.03$ and absorber parameters were $_{\rm H} = 5.3 \times 10^{21}$ $^{-2}$ and ionization parameter $\xi = 40$ $^{-1}$."
 These parameters yielded a goodness-of-fit V=145 for 108 degrees of freedom (no component of systematic error is assumed)., These parameters yielded a goodness-of-fit $\chi^2 = 145$ for 108 degrees of freedom (no component of systematic error is assumed).
 The fit to the spectrum in the 2 — 10 keV range is shown in refmrk766.. where we plot the model spectrum at the resolution together with “unfolded” data points.," The fit to the spectrum in the 2 – 10 keV range is shown in \\ref{mrk766}, where we plot the model spectrum at the resolution together with “unfolded” data points."
 The fit of the wind model to the data does indicate that this may well provide a good explanation of the shape of X-ray spectra in AGN such as 7766., The fit of the wind model to the data does indicate that this may well provide a good explanation of the shape of X-ray spectra in AGN such as 766.
 With the exception of the narrow Fe emission line. most of the spectral shape in the Fe regime is provided by the wind model. with the reflection component contributing only 3pper cent of the total 2 — 10 keV energy flux and only 4pper cent of the flux density in the continuum around Fe.," With the exception of the narrow Fe emission line, most of the spectral shape in the Fe regime is provided by the wind model, with the reflection component contributing only per cent of the total 2 – 10 keV energy flux and only per cent of the flux density in the continuum around Fe."
 It is unlikely that the current model can explain the, It is unlikely that the current model can explain the
profiles. while galaxies are niodellecl as a cise plus spheroid.,"profiles, while galaxies are modelled as a disc plus spheroid."
 The satellite galaxies. plus halo are then advanced by calculating the combined gravitation forces of the host and satellite haloes. as well as the ellects of dvnamical friction. calculated using Chandrasekhar's formula.," The satellite galaxies plus halo are then advanced by calculating the combined gravitation forces of the host and satellite haloes, as well as the effects of dynamical friction, calculated using Chandrasekhar's formula."
 The code keeps track of tidal stripping. to remove mass from the satellites.," The code keeps track of tidal stripping, to remove mass from the satellites."
 The new halo mass is then used for the next iteration of the orbit equations., The new halo mass is then used for the next iteration of the orbit equations.
 This integration of the orbital equations continues until one of three conditions are met: (1) the final redshift. (i) the host merges with a new larger halo. in which case the satellite becomes a part of the new halo and has new orbital parameters assigned. or (iii) the satellite merges with the central galaxy.," This integration of the orbital equations continues until one of three conditions are met: (i) the final redshift, (ii) the host merges with a new larger halo, in which case the satellite becomes a part of the new halo and has new orbital parameters assigned, or (iii) the satellite merges with the central galaxy."
 In this last case. merging takes place when the orbital radius falls below River. which is the sum of the half mass radius of the central and. satellite galaxies.," In this last case, merging takes place when the orbital radius falls below $R_{merge}$ which is the sum of the half mass radius of the central and satellite galaxies."
 The mass of the merged. satellite is then added to the central ealaxy., The mass of the merged satellite is then added to the central galaxy.
 Phe merging times match largely match those of ?.. however some satellites have very long merging times. as they loose a great deal of mass through tidal stripping. meaning that dynamical friction forces become very weak.," The merging times match largely match those of \citet{Cole2000}, however some satellites have very long merging times, as they loose a great deal of mass through tidal stripping, meaning that dynamical friction forces become very weak."
" The Durham mocels relate supernova reheating directLy to the circular velocity of the galaxy disk according to ?:: where Alpena is the rate of change of mass of reheated gas. ιν ἰξ the disk circular velocity. Ad, is the time derivative of stellar mass and AZ.;; is the change in the mass of ejected eas."," The Durham models relate supernova reheating directly to the circular velocity of the galaxy disk according to \citet{Cole2000}: where $\dot{M}_{reheat}$ is the rate of change of mass of reheated gas, $V_{disk}$ is the disk circular velocity, $\dot{M}_*$ is the time derivative of stellar mass and $\dot{M}_{eject}$ is the change in the mass of ejected gas."
 In haloes with a shallow potential. this has the cllect of reducing the amount of cold. gas available to form stars. bv heating the gas back into the hot gas reservoir.," In haloes with a shallow potential, this has the effect of reducing the amount of cold gas available to form stars, by heating the gas back into the hot gas reservoir."
 The hot gas is dominated. by ejection for low mass haloes. ancl by reheating without ejection for large haloes.," The hot gas is dominated by ejection for low mass haloes, and by reheating without ejection for large haloes."
 For low mass galaxies. supernova feedback is an important mechanism bv which galactic winds are driven and star formation is cquenched.," For low mass galaxies, supernova feedback is an important mechanism by which galactic winds are driven and star formation is quenched."
 The Durham models implement AGN feedback is such ἃ wav às to regulate the cooling of hot gas., The Durham models implement AGN feedback is such a way as to regulate the cooling of hot gas.
 In large haloes with laree Ecelington Luminosity. the ΑΝ. feedback. is assumed to balance heating ancl cooling. thus truncating star formation.," In large haloes with large Eddington Luminosity, the AGN feedback is assumed to balance heating and cooling, thus truncating star formation."
 This prevents the formation of over-Iuminous ealaxies., This prevents the formation of over-luminous galaxies.
 While feedback is active. the black hole is assumed ο grow proportionally to the cooling luminosity. and gas is accreted due to disk imstabilities.," While feedback is active, the black hole is assumed to grow proportionally to the cooling luminosity, and gas is accreted due to disk instabilities."
" The model assumes quasi-ivdrostatie cooling for AGN active galaxies. and has a strict ransition between AGN ""active"" and. ""inactive"" phases."," The model assumes quasi-hydrostatic cooling for AGN active galaxies, and has a strict transition between AGN “active” and “inactive” phases."
 AGN feedback becomes ellicient in galaxies of mass greater han ~2. 10h+ M., AGN feedback becomes efficient in galaxies of mass greater than $\sim$ $\times 10 ^{11}h^{-1}$ $_{\odot}$.
 The essential dillerence in the 1060 and. 05 moclels is the implementation of ram-pressure stripping of he hot eas., The essential difference in the B06 and F08 models is the implementation of ram-pressure stripping of the hot gas.
 In the D-BBOG model. along with both Munich models. the hot gas is instantaneously stripped. when it enters à halo already containing à central galaxy.," In the B06 model, along with both Munich models, the hot gas is instantaneously stripped when it enters a halo already containing a central galaxy."
 In the EFIFOS model this process happens gradually and. depends on the orbit of the galaxy., In the F08 model this process happens gradually and depends on the orbit of the galaxy.
 This has the οσοι of reducing he population of faint red ealaxies., This has the effect of reducing the population of faint red galaxies.
 The Munich merger trees (2). used in M.DDO6 and ALBBO7 follow the positions of subhaloes for as long as they can be identified., The Munich merger trees \citep{Springel2005} used in D06 and B07 follow the positions of subhaloes for as long as they can be identified.
 Identification is possible only when the number of xuticles bound to a subhalo exceeds the minimum number of particles set hy SUDEIND., Identification is possible only when the number of particles bound to a subhalo exceeds the minimum number of particles set by SUBFIND.
 The trees are constructed by ollowing the most bound halo particles and. searching for he descendant halo in the next output., The trees are constructed by following the most bound halo particles and searching for the descendant halo in the next output.
 One of the kev dillerences between the Munich mocels and those of Durham is that the Munich. models explicitlv ollow dark matter haloes even after they are accreted onto arger systems. allowing the dynamics of satellite galaxies residing in the infalling haloes to be followed: until the dark matter substructure is destroved.," One of the key differences between the Munich models and those of Durham is that the Munich models explicitly follow dark matter haloes even after they are accreted onto larger systems, allowing the dynamics of satellite galaxies residing in the infalling haloes to be followed until the dark matter substructure is destroyed."
 The galaxy position is calculated by assigning the galaxy to the most. bounc xwticle of a (sub)halo at each time step., The galaxy position is calculated by assigning the galaxy to the most bound particle of a (sub)halo at each time step.
 This is done unti he (sub)halo is no longer identifiable. whereupon the galaxy is assigned to the most bound. particle of the (sub)halo a he last time the (sub)halo could be identified.," This is done until the (sub)halo is no longer identifiable, whereupon the galaxy is assigned to the most bound particle of the (sub)halo at the last time the (sub)halo could be identified."
 An analytic countdown to galaxy merging begins when the satellite subhalo can no longer be identified. and resets if the paren milo undergoes a major merger.," An analytic countdown to galaxy merging begins when the satellite subhalo can no longer be identified, and resets if the parent halo undergoes a major merger."
 Εις. in the Munich mocels. he lifetime of galaxies in groups depends on the amount of ime the (sub)halo finder can identify the subhalo. plus the analytic countdown.," Thus, in the Munich models, the lifetime of galaxies in groups depends on the amount of time the (sub)halo finder can identify the subhalo, plus the analytic countdown."
" The analytic merging follows that of ?:: where In(1∣|a4M is the coulomb logarithm. Vg is the halo circular velocity. ro, is the distance of the subhalo from the halo centre at the time it is last identified. mi; is the mass of the satellite dark halo at the time it was last identified and AZg is the halo mass."," The analytic merging follows that of \citet{Croton2006}: where $\ln \left(1+\frac{M_H}{M_{sat}} \right)$ is the coulomb logarithm, $V_H$ is the halo circular velocity, $r_{sat}$ is the distance of the subhalo from the halo centre at the time it is last identified, $m_{sat}$ is the mass of the satellite dark halo at the time it was last identified and $M_H$ is the halo mass."
 In ALDDOG. the amount of reheated (by. supernovae) cold gas is proportional to the stellar mass. and the mass ejected from. the halo is inversely proportional to the host halo's circular velocity squared (2): The important parameters. 35; and AZ. have the same relationship in the Durham models.," In D06, the amount of reheated (by supernovae) cold gas is proportional to the stellar mass, and the mass ejected from the halo is inversely proportional to the host halo's circular velocity squared \citep{Croton2006}; The important parameters, $V_{vir}$ and $M_*$, have the same relationship in the Durham models."
 In the Munich. models. the reheatecl mass is proportional to the mass of the halo. while in the Durham model it is proportional to the clisk mass.," In the Munich models, the reheated mass is proportional to the mass of the halo, while in the Durham model it is proportional to the disk mass."
 M.DDBOT adopts a more sophisticated treatment of supernova feedback., B07 adopts a more sophisticated treatment of supernova feedback.
 Rather than simply. parametrising the ellect of supernovae feedback. the NBDBOT model follows the dynamical evolution of the wind as an adiabatic expansion ollowed: by snowploughing.," Rather than simply parametrising the effect of supernovae feedback, the B07 model follows the dynamical evolution of the wind as an adiabatic expansion followed by snowploughing."
 This implementation has the ellect of increasing the luminosity ofthe brightest. galaxies., This implementation has the effect of increasing the luminosity of the brightest galaxies.
 As also occurs with the Durham mocels. some of the gas will » ejected from low mass haloes in both Munich mocels.," As also occurs with the Durham models, some of the gas will be ejected from low mass haloes in both Munich models."
" In the Munich models. a tvo-mocdoe formalism is adopted or AGN. wherein a high-energv. or uasar mode occurs subsequent to mergers. and a constant low-energy ""radio"" mode suppresses cooling [lows cue το the interaction tween the gas and the central black hole ?.."," In the Munich models, a two-mode formalism is adopted for AGN, wherein a high-energy, or “quasar” mode occurs subsequent to mergers, and a constant low-energy “radio” mode suppresses cooling flows due to the interaction between the gas and the central black hole \citet{Croton2006}."
 In the quasar model. accretion of gas onto the black hole peaks at z 3. while the radio mode reaches a plateau at z ~2.," In the quasar model, accretion of gas onto the black hole peaks at z $\sim$ 3, while the radio mode reaches a plateau at z $\sim$ 2."
 AGN eedback is assumed to be efficient only in massive halocs. with supernova feedback being more dominant in lower-niass jaloes.," AGN feedback is assumed to be efficient only in massive haloes, with supernova feedback being more dominant in lower-mass haloes."
 The properties of groups in the Munich models. ALDDOG and M.DBDOT. are similar to one another in many Cases. as are the properties of the two Durham. mocels. D_BBOG and D.EFOs.," The properties of groups in the Munich models, D06 and B07, are similar to one another in many cases, as are the properties of the two Durham models, B06 and F08."
 Thus. in some analysis. we just discuss," Thus, in some analysis, we just discuss"
In this section we present the returned masses. time delays. and integrated yields for the chemical elements considered in this work.,"In this section we present the returned masses, time delays, and integrated yields for the chemical elements considered in this work."
 In Table 3 we show the time delays of LIMS for the returned mass and for the integrated. vields of eight chemical elements., In Table 3 we show the time delays of LIMS for the returned mass and for the integrated yields of eight chemical elements.
" These data are independent of the M,,, value adopted in the IMF. but they are dependent on the initial stellar metallicity Z;."," These data are independent of the $M_{up}$ value adopted in the IMF, but they are dependent on the initial stellar metallicity $Z_i$."
" The data shown in Tables 4 and 5 were obtainedby assuming M,=40 and M,,,=80 M... respectively."," The data shown in Tables 4 and 5 were obtainedby assuming $M_{up}=40$ and $M_{up}=80$ , respectively."
We thank Anatoly Whypin for. providing the simulation and Jeremy Tinker for providing the halo finder used in this study.,We thank Anatoly Klypin for providing the simulation and Jeremy Tinker for providing the halo finder used in this study.
 Additionally. we thank Matthew 1t. Decker and Andrey. Ixravtsoy. for providing us with the weak lensing shear and convergence maps from. the simulations.," Additionally, we thank Matthew R. Becker and Andrey Kravtsov for providing us with the weak lensing shear and convergence maps from the simulations."
 The authors are grateful. to Matthew: It... Becker. for. helpful comments and Joseph J. Mohr for discussions on the MOR., The authors are grateful to Matthew R. Becker for helpful comments and Joseph J. Mohr for discussions on the MOR.
 TYE would like to thank the hospitality of the University of PennsyIvania where this collaboration started., TYL would like to thank the hospitality of the University of Pennsylvania where this collaboration started.
 DG acknowledges: support by Stuclienstiltune des deutschen Volkes. and DAAD (Deutscher. Akaclemischer Austauschedienst)., DG acknowledges support by Studienstiftung des deutschen Volkes and DAAD (Deutscher Akademischer Austauschdienst).
 GB was supported by NSE erant. ASTT-000827 and DOL grant DE-FCO2-95ERA40893., GB was supported by NSF grant AST-090827 and DOE grant DE-FG02-95ER40893.
" ΕΥ, was supported. by the World. Premier. International. Research Center Initiative. (WPL Initiative). ALENT. Japan and a JSPS travel grant (Institutional. Program for Young Researcher Overseas. Visits)."," TYL was supported by the World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan and a JSPS travel grant (Institutional Program for Young Researcher Overseas Visits)."
" SS acknowledges support by TR33 ""Phe Dark Universe and the DEG Cluster. of Excellence on the ""Origin and Structure of the Universe""."," SS acknowledges support by TR33 ""The Dark Universe"" and the DFG Cluster of Excellence on the ""Origin and Structure of the Universe""."
 In the following we describe the halo model usec to calculate the covariance of shear in annuli due to structure uncorrelated to the central halo and outside the integration region of the simulations., In the following we describe the halo model used to calculate the covariance of shear in annuli due to structure uncorrelated to the central halo and outside the integration region of the simulations.
" For à general review of halo models of large scale structure. we refer the reader to ὃν,"," For a general review of halo models of large scale structure, we refer the reader to \citet{2002PhR...372....1C}."
 Approaches similar to ours are presented in 2.inFourierspace and. inreal space.. , Approaches similar to ours are presented in \citet[][in Fourier space]{2000ApJ...535L...9C} and \citet[][in real space]{2003MNRAS.344..857T}. .
Consider a set of halos 1.....⊔∖∖⋰↓↿↓↥∪∖∕∶∣∖∕∶∖∶ ↓⊳↾∐↥⋖⊾∣⋜↧⊔⋏∙≟∢," Consider a set of halos $1,\ldots,i,\ldots,n$ with $0<z_i<z_s=1$ ."
"⊾↓∐⊀↓⋜↧↓⊳∖↓∐⊾∥↓⋅⋎∕∣↿↓↕⋖⋅∙∖⇁↕↓↥↿↓⋅⋖⋟∠⇂⇂⇂≼↛⋖⋅↕↓↕⋜↧↓↕⋜↧↓↥↓↕⇂⇂↓⇂⇂≻∕⋅⋅↕≻ where Ay=(6,;.05;.5;Miei...) is à set of parameters describing the position and mass prolile of halo i and +}(hy) is the tangential shear such a halo introduces on annulus A."," The tangential shear $\gamma_t^k$ they introduce in an annulus $k$ is where $\bmath{h_i}=(\theta_{1,i},\theta_{2,i},z_i,M_i,c_{i},\ldots)$ is a set of parameters describing the position and mass profile of halo $i$ and $\gamma_t^k(\bmath{h_i})$ is the tangential shear such a halo introduces on annulus $k$."
 We can then straightforwardly calculate the covariance of) ands! in annuli & and / as where the average goes over random realizations of the halo populations., We can then straightforwardly calculate the covariance of $\gamma_t^k$ and $\gamma_t^l$ in annuli $k$ and $l$ as where the average goes over random realizations of the halo populations.
 Note that we have separated. the terms caused by the shear of a single halo on both annuli (one-halo term) from correlations caused by pairs of halos (tsvo-halo term)., Note that we have separated the terms caused by the shear of a single halo on both annuli (one-halo term) from correlations caused by pairs of halos (two-halo term).
 The latter will not be zero. because halo positions are correlated.," The latter will not be zero, because halo positions are correlated."
 The average in Eqn., The average in Eqn.
 A2 can be expressed as an integral over probabilities «P(R8) of finding a halo with properties h and αι.ha) of simultaneously finding two halos with properties Ay and Aa. ic. The probability density απ(ο) of finding a halo with properties A has as its cosmological ingredients the halo mass function. πα--(A4.z). and the distribution. of halo clensity profiles.," \ref{eqn:gtcov1} can be expressed as an integral over probabilities $dP_1(\bmath{h})$ of finding a halo with properties $\bmath{h}$ and $dP_2(\bmath{h_1},\bmath{h_2})$ of simultaneously finding two halos with properties $\bmath{h_1}$ and $\bmath{h_2}$, i.e. The probability density $dP_1(\bmath{h})$ of finding a halo with properties $\bmath{h}$ has as its cosmological ingredients the halo mass function, $\frac{dn}{dM}(M,z)$, and the distribution of halo density profiles."
" We can write assuming NEW halo profiles and a probability density pic]M.z) of finding concentration parameters e=eoggsrogusArs Which only depends on the halo mass Al=Moo, and recshift ο."," We can write assuming NFW halo profiles and a probability density $p(c|M,z)$ of finding concentration parameters $c\equiv c_{200m}=r_{200m}/r_s$ which only depends on the halo mass $M\equiv M_{200m}$ and redshift $z$ ."
" Note that we use definitions of rogis Conon- ancl Alsou, with respect to 200 times the mean matter density at the halo's redshift."," Note that we use definitions of $r_{200m}$, $c_{200m}$, and $M_{200m}$ with respect to 200 times the mean matter density at the halo's redshift."
 We assume as a fixed concentration relation the best-fitformula from ? (usingWALADPScosmology.cf. 2).. For the halo mass function and the mass dependent biases b(AL) used. below. we adopt the Sheth-Tormen mass Function and the associated. bias (7)..," We assume as a fixed mass-concentration relation the best-fitformula from \citet{2008MNRAS.390L..64D} \citep[using WMAP5 cosmology, cf.][]{2009ApJS..180..330K}, For the halo mass function and the mass dependent biases $b(M)$ used below, we adopt the Sheth-Tormen mass function and the associated bias \citep{1999MNRAS.308..119S}."
 The combined probability. density αι.R3) of simultaneously| finding two halos with properties Aa and Re must include the interdependence of their probabilities.," The combined probability density $dP_2(\bmath{h_1},\bmath{h_2})$ of simultaneously finding two halos with properties $\bmath{h_1}$ and $\bmath{h_2}$ must include the interdependence of their probabilities."
 Assume that halos cluster. according to α 3d. two-point correlation function £íí(r.Ra.R3) which is generally dependent on halo properties A4/5 in addition to their separation r.," Assume that halos cluster according to a 3d two-point correlation function $\xi_{hh}(r,\bmath{h_1},\bmath{h_2})$ which is generally dependent on halo properties $\bmath{h_{1/2}}$ in addition to their separation $r$."
" We can then write Because the tangential shear due to a randomly placed halo does not. correlate with the shear due to another randomly. placed halo. the only contribution to the covariance stems from the term. proportional to &),;,. and we can drop the other part for the numerical integration."," We can then write Because the tangential shear due to a randomly placed halo does not correlate with the shear due to another randomly placed halo, the only contribution to the covariance stems from the term proportional to $\xi_{hh}$, and we can drop the other part for the numerical integration."
 In our halo model. what contributes to the correlation of halo positions is the evolution of structures.," In our halo model, what contributes to the correlation of halo positions is the evolution of structures."
 The non-linear evolution. Le. the formation of halos themselves. need not be considered. to this end. as it is dominant only on scales smaller than the ones considered here.," The non-linear evolution, i.e. the formation of halos themselves, need not be considered to this end, as it is dominant only on scales smaller than the ones considered here."
" The halo-halo correlation function £5, cantherefore bewritten as (?) giving the excess probability of finding a pair of halos with masses Al; ancl Alo at comoving separation r at redshift zz2,8: ce."," The halo-halo correlation function $\xi_{hh}$ cantherefore bewritten as \citep{2010MNRAS.408..300G}
 giving the excess probability of finding a pair of halos with masses $M_1$ and $M_2$ at comoving separation $r$ at redshift $z\approx z_1\approx z_2$ ."
 We use a linear power spectrum that is consistent with WALAP 5 results., We use a linear power spectrum that is consistent with WMAP 5 results.
 We relate this to an angular correlation at redshift) z by performing one redshilt-space integral over the three-climensionaltwo-point, We relate this to an angular correlation at redshift $z$ by performing one redshift-space integral over the three-dimensionaltwo-point
" The iouiziugC» spectral enerevOo, distribution (liereafter ISED) of nearby active ealactic nuclei cauno be observed directly. due to the large Calactic absorption ονομα the Lii limit.","} The ionizing spectral energy distribution (hereafter ISED) of nearby active galactic nuclei cannot be observed directly, due to the large Galactic absorption beyond the Lyman limit."
 Owing to the redshift effect. however. we cau eget a elimpse of the ISED from the spectra of distant quasars.," Owing to the redshift effect, however, we can get a glimpse of the ISED from the spectra of distant quasars."
 The pioneering work of Zheng ct al (, The pioneering work of Zheng et al. (
1997. ZIN97). using IIST-FOS archived data. ideutifie a noticeable steepening iu quasar SEDs near 1200A.,"1997, ZK97), using HST-FOS archived data, identified a noticeable steepening in quasar SEDs near 1200."
". The power-law index a (Fy xv"") in quasarssteepeus from 1 for A>1200 Πο ~2 at shorterwaveleuethst.. down to (Avext))."," The power-law index $\alpha$ $F_{\nu}\propto
\nu^{\alpha}$ ) in quasarssteepens from $ -1$ for $\lambda >
1200\,$ to $\simeq -2$ at shorter, down to )."
 Iu a more recent study. Telfer et al.," In a more recent study, Telfer et al."
 2002 (hereafter TZ02) found similar results with a colposite spectrum characterized by a imieau nearUV index of 0.7 stecpening to ~1.7 iu the far-UV., 2002 (hereafter TZ02) found similar results with a composite spectrum characterized by a mean near-UV index of $-0.7$ steepening to $\simeq -1.7$ in the far-UV.
 The above authors favor the interpretation that the steepeniug is intrinsic to quasars and that it is the signature of a coniptonized accretion cixk., The above authors favor the interpretation that the steepening is intrinsic to quasars and that it is the signature of a comptonized accretion disk.
 Ou the other hand. certain distributious of intergalactic absorption Sas can also cause al (apparent) steepening of the SED. startiug at (Apest}}. as shown in this Paper. that is. located in the same position as the break encountered bv TZ02.," On the other hand, certain distributions of intergalactic absorption gas can also cause an (apparent) steepening of the SED, starting at ), as shown in this Paper, that is, located in the same position as the break encountered by TZ02."
" As lone as ISED studies are based ou detectors that do not exteud bevoud (Ας,)). one cannot readily distinguish. between the contribution to the observed steepeniue from an intervene absorption model and that of a purely intrinsic break im the quasar SED as proposed by ZISO7 or TZ02."," As long as ISED studies are based on detectors that do not extend beyond ), one cannot readily distinguish between the contribution to the observed steepening from an intervening absorption model and that of a purely intrinsic break in the quasar SED as proposed by ZK97 or TZ02."
 FUSE or UST-STIS spectra. however. extend unich farther iuto the UV and will be shown here to provide us with a colmpclling test to discriminate between the two interpretations.," FUSE or HST-STIS spectra, however, extend much farther into the UV and will be shown here to provide us with a compelling test to discriminate between the two interpretations."
" Large scale structure formation models predict hat a substautial fraction of barvous should reside in a wvarnrhot phase at current epoch. possibly up ο of Or, (Daveetal.2001).."," Large scale structure formation models predict that a substantial fraction of baryons should reside in a warm-hot phase at current epoch, possibly up to of $\Omega_{bar.}$ \citep{davea}."
 Depending on its temperature and distribution with redshift. lis warnrhot gas may contribute to a reduced raction of the observed steepenimgo of the ISED.," Depending on its temperature and distribution with redshift, this warm-hot gas may contribute to a reduced fraction of the observed steepening of the ISED."
 Iu this Paper. we show that this compoucut nay eive rise to a sieht flux increase (that is. a discoutimuty) iu the region (Αν) with respect to longer wavelengths.," In this Paper, we show that this component may give rise to a slight flux increase (that is, a discontinuity) in the region ) with respect to longer wavelengths."
 The calibration of the level of this discoutiuuous excess flux can be used as a technique to set useful limits for the barvonic mass coutributiou from the wari-hot intergalactic coniponeut., The calibration of the level of this discontinuous excess flux can be used as a technique to set useful limits for the baryonic mass contribution from the warm-hot intergalactic component.
 The objective (and results) of this Paper is threcfold: The Comn-Peterson (CP) effect (Com&Pe-terson1965) sets stringent limits for tle presence of neutral diffuse eas at hieh redshifts., The objective (and results) of this Paper is threefold: The Gunn-Peterson (GP) effect \citep{gunn} sets stringent limits for the presence of neutral diffuse gas at high redshifts.
 Iu the siuplest form of the GP test. the absorption gas produces a flux decrement between aand Ly.," In the simplest form of the GP test, the absorption gas produces a flux decrement between and ."
.. The decrement ix measured against a continu level. usually taken to be a power-law," The decrement is measured against a continuum level, usually taken to be a power-law"
levels (2.04+0.37)% for DP-IP and (0.86+0.24)% for SP-IP.,levels $2.04\pm0.37$ for DP–IP and $(0.86\pm0.24)\%$ for SP–IP.
" Such solutions we called ""a good solution"".", Such solutions we called ”a good solution”.
" In summary, each simulation run including 50000 detection attempts resulted in about 5500 detections (satisfying geometrical detection conditions (Section ??)))."," In summary, each simulation run including 50000 detection attempts resulted in about 5500 detections (satisfying geometrical detection conditions (Section \ref{sec.detect.cond}) ))."
 'There were 641360 simulation runs including all possible combinations of the probability distribution functions (Section ??))., There were $641\:360$ simulation runs including all possible combinations of the probability distribution functions (Section \ref{sec.prob.density.funct}) ).
 Among the resulting 4x10? geometrical detections we found 827 good solutions (as described in item 16 above)., Among the resulting $\sim4\times10^9$ geometrical detections we found 827 good solutions (as described in item 16 above).
 These solutions are listed in Tables B2 — B7 and 9 best examples with relatively high K-S probability values are shown in Table 2 (items 1a — 9a)., These solutions are listed in Tables B2 – B7 and 9 best examples with relatively high K–S probability values are shown in Table \ref{tab.2} (items 1a – 9a).
 Items 1b — 9b correspond to tests of the luminosity problem described in the Appendix ?? in the on-line materials., Items 1b – 9b correspond to tests of the luminosity problem described in the Appendix \ref{sec.appendix.luminosity} in the on–line materials.
 Our new statistical analysis is based on a number of new databases that we compiled from the recently published data., Our new statistical analysis is based on a number of new databases that we compiled from the recently published data.
" We followed methodology developed by KGM04, which relies on comparison of synthetic and real pulsar data."," We followed methodology developed by KGM04, which relies on comparison of synthetic and real pulsar data."
 Our period P database contains 1520 items (we excluded all recycled and binary pulsars) in the range of 0.02 — 8.51 seconds., Our period $P$ database contains 1520 items (we excluded all recycled and binary pulsars) in the range of 0.02 – 8.51 seconds.
 This database includes 355 more pulsars than recently analysed database compiled by KGM04., This database includes 355 more pulsars than recently analysed database compiled by KGM04.
 We also compiled a new database of pulse-widths Wio measured at intensity level., We also compiled a new database of pulse–widths $W_{10}$ measured at intensity level.
" This database contains 414 items, which is 176 more than that of KGM04."," This database contains 414 items, which is 176 more than that of KGM04."
" There are many more pulse-width measurements available these days, but our database is restricted to pulsars with DM«150 pc cm""?, which guarantees avoiding significant external pulse broadening."," There are many more pulse–width measurements available these days, but our database is restricted to pulsars with $DM<150$ pc $^{-3}$, which guarantees avoiding significant external pulse broadening."
" Most importantly, we created the largest ever database of interpulse occurrence in pulsar emission."," Most importantly, we created the largest ever database of interpulse occurrence in pulsar emission."
" Our database contains 44 pulsars (compared with 14 pulsars in Tayloretal. (1993)//KGM04 database), including 31 cases of DP- and 13 cases of SP-IP."," Our database contains 44 pulsars (compared with 14 pulsars in \citet{taylor93}/ /KGM04 database), including 31 cases of DP--IP and 13 cases of SP–IP."
" Although our IP database is more numerous than any of the previous ones (e.g. KGM04, WJO08a), it seems that the frequencies of occurrence are similar to those occurring in previous smaller databases."," Although our IP database is more numerous than any of the previous ones (e.g. KGM04, WJ08a), it seems that the frequencies of occurrence are similar to those occurring in previous smaller databases."
" In fact, we have of total number of IP cases, divided into DP-IP and SP-IP cases, respectively, in a population of 1520 pulsars."," In fact, we have of total number of IP cases, divided into DP–IP and SP–IP cases, respectively, in a population of 1520 pulsars."
" This can be compared with2.71%,, and0.78%,, respectively, found in KGM04 database."," This can be compared with, and, respectively, found in KGM04 database."
" One can therefore firmly state that there should be about of IP cases in the population of normal pulsars, including about and of DP-IP and SP-IP cases, respectively."," One can therefore firmly state that there should be about of IP cases in the population of normal pulsars, including about and of DP–IP and SP–IP cases, respectively."
the proto-cluster (Peuteriecietal. 2002)..,the proto-cluster \citep{pen02}. .
 Suuületal.(20031) also report the detection of four ταν sources coiucidenut with subnmuui sources surrounding three IIzRCis., \citet{sma03b} also report the detection of four X-ray sources coincident with submm sources surrounding three HzRGs.
" These proto-clusters also contaiu a large amount of eas, ax revealed by the hhaloes surrounding the UzRCs. which have plysical scales up to 2200 kpe (forarecentreview.seevanBreugeletal.. 2003)."," These proto-clusters also contain a large amount of gas, as revealed by the haloes surrounding the HzRGs, which have physical scales up to $>$ 200 kpc \citep[for a recent review, see][]{wvb03}."
. Together with the mereased mereer rates in higher redshift clusters (e.y.vanDokkietal. 1999).. this provides the ingrecdcieuts to induce wide-spreack starbursts and ACN. which could) be revealed by thei thermal dust (sub-)mua euission.," Together with the increased merger rates in higher redshift clusters \citep[\eg][]{dok99}, this provides the ingredients to induce wide-spread starbursts and AGN, which could be revealed by their thermal dust (sub-)mm emission."
 Statistical overdeusitics of (sub-)nunu galaxies (SACS) lave indeed been fouud fromSCUBA bolometer imaging of the fields surromidline the :—3.09 ‘redshift spike’ (Chapmanetal.2001).. the 2=2.59 radio galaxy S3WO02 (αἱetal..2003a).. the X58 radio galaxv LC 11.17 (Ivisonetabl.2000).. a τι QSO (Stevensetal.2001). and six other IIZRC: fields (Stevensetab.2003).," Statistical overdensities of (sub-)mm galaxies (SMGs) have indeed been found fromSCUBA bolometer imaging of the fields surrounding the $z$ =3.09 'redshift spike' \citep{cha01}, the $z$ =2.39 radio galaxy 53W002 \citep{sma03a}, the $z$ =3.8 radio galaxy 4C 41.17 \citep{ivi00}, a $z$ =1.8 QSO \citep{ste04} and six other HzRG fields \citep{ste03}."
. In this paper. we present a 1.2 mun map covering the central 25 arcmin? of the nost distant proto-cluster known to date. surrounding he :=L1 radio galaxy TN 1912.," In this paper, we present a 1.2 mm map covering the central 25 $^2$ of the most distant proto-cluster known to date, surrounding the $z$ =4.1 radio galaxy TN $-$ 1942."
 We find au overdensity of 1.2 nuu sources. which we identify with optically very faint galaxies usine a deep 1.1 GIIz map. ut Gud no overlap between the population of excess celitters and the 1.2 11a sources.," We find an overdensity of 1.2 mm sources, which we identify with optically very faint galaxies using a deep 1.4 GHz map, but find no overlap between the population of excess emitters and the 1.2 mm sources."
" Throughout this paper. we use a X cosmology with Uy=71 lau + | Q4,—0.27. and O420.73 (Spergeletal.2003:Toury 2003)."," Throughout this paper, we use a $\Lambda-$ cosmology with $_0$ =71 km $^{-1}$ $^{-1}$ , $\Omega_{\rm M}$ =0.27 and $\Omega_{\Lambda}$ =0.73 \citep[][]{spe03,ton03}."
" At :—L1. the huuinositv clistance Is D, 6.65 Gpc. and ccorresponuds to 7.0 kpe."," At $z$ =4.1, the luminosity distance is $D_L$ =37.65 Gpc, and corresponds to 7.0 kpc."
 To image the field of TN 1912 at nun waveleneths. we used the 37- and 117-chanunel Max Planck Bolometer arravs (ATAMBO-LaudMAMDBO-2:Isveysaetal..1998) at the IRAM. 301m telescope ou Pico Veleta. Spain.," To image the field of TN $-$ 1942 at mm wavelengths, we used the 37- and 117-channel Max Planck Bolometer arrays \citep[MAMBO-1 and MAMBO-2;][]{kre98} at the IRAM 30m telescope on Pico Veleta, Spain."
 ALAMDBO has a half-power spectral bandwidth from 210 to 290 CIIz. with an effective bandwidth ceutre for steep thermal spectra of ~250 CIIz (1.2 nuu).," MAMBO has a half-power spectral bandwidth from 210 to 290 GHz, with an effective bandwidth centre for steep thermal spectra of $\sim$ 250 GHz (1.2 mm)."
 The effective beam EWIINL is 10777 with au array size of, The effective beam FWHM is 7 with an array size of.
 The observations were done in a pooled observing mode during he winter 2011-2002 season., The observations were done in a pooled observing mode during the winter 2001-2002 season.
 Die to the low DNeclination. fιο field couk oulv be observed for 1 hours La elven nieit. with clevetions between uud3.," Due to the low declination, the field could only be observed for 4 hours in a given night, with elevations between and."
 The atinospheric zenith opacities at 1.2 nuu varied between 0.12 auk 0.25., The atmospheric zenith opacities at 1.2 mm varied between 0.12 and 0.25.
 The total on source oeitegration fine was 17.0 hours. of which 2.6 h were obtained with the I1l7-chanuel array. and 146 h with the 37-channel array.," The total on source integration time was 17.0 hours, of which 2.6 h were obtained with the 117-channel array, and 14.6 h with the 37-channel array."
 We used the standard ou-the-flv mapping technique. comprised of LL subscaus of 10 s each. while chopping the secondary mirror in azimuth at 2 Uz.," We used the standard on-the-fly mapping technique, comprised of 41 subscans of 40 s each, while chopping the secondary mirror in azimuth at 2 Hz."
" To nuininuize he residual effects of the double eam poiut spread function. we used different ehop (Gvobbler) throws(30. oor 15"") iuxfor scan directions for cach map."," To minimize the residual effects of the double beam point spread function, we used different chop (wobbler) throws, or ) and/or scan directions for each map."
 We used 7 different pointines offset bv tto cover the eutire field of the VLT/FORS2 imagine to Πο., We used 7 different pointings offset by to cover the entire field of the VLT/FORS2 imaging to uniform depth.
" We checked the poimtiug aud focus at least once per hour using the bright poiut source 127. and found the poiutiug to be stable to within κ,"," We checked the pointing and focus at least once per hour using the bright point source $-$ 127, and found the pointing to be stable to within $<$."
 The absolute flux calibration is based on observatious of several standard calibration sources. iucliding plaucts. resulting in an estinatec acetracy ofC.," The absolute flux calibration is based on observations of several standard calibration sources, including planets, resulting in an estimated accuracy of."
 We analyzed the data using the AIOPSI software oickaee (Zvlka.1998)., We analyzed the data using the MOPSI software package \citep{zyl98}.
. We subtracted the skvuoise. aud conmibined the couble-beam maps using a shift-aud-add xocedure. producing for cach map a positive nage xacketed by two negative mages of half the inteusitv ocated oue chop throw away.," We subtracted the skynoise, and combined the double-beam maps using a shift-and-add procedure, producing for each map a positive image bracketed by two negative images of half the intensity located one chop throw away."
 Because we combined our 37 maps obtained with differeut chop throws. the effect of he confusion due to negative sidelobes is niuimuized. but still preseut iu regious of high source density.," Because we combined our 37 maps obtained with different chop throws, the effect of the confusion due to negative sidelobes is minimized, but still present in regions of high source density."
 The noise evel increases outward in our co-added map of the field., The noise level increases outward in our co-added map of the field.
 Du Fie., In Fig.
 l two coutours show the region within which the riis roise level is less than 1.2 aud 0.6 1i1Jy (before sioothing). chclosing areas of 25.6 and 2.6 square arcu. respectively.," \ref{RMAMBOSN} two contours show the region within which the rms noise level is less than 1.2 and 0.6 mJy (before smoothing), enclosing areas of 25.6 and 2.6 square arcmin, respectively."
 To obtain accurate positions of the nuu sources. and to search for possible radio-loud ACN counterparts. we observed the TN 1912 field with the Very Large Arrav (VLA:Napier.Thompson&Ekers.1983) on UT 2002 April 1 to 12 for a total of 12 hours iu the A-array at 20 cin.," To obtain accurate positions of the mm sources, and to search for possible radio-loud AGN counterparts, we observed the TN $-$ 1942 field with the Very Large Array \citep[VLA; ][]{nap83} on UT 2002 April 1 to 12 for a total of 12 hours in the A-array at 20 cm."
 We observed iu a pseudo-coutinmun. spectral hueσα mode with 7«3.125 MIIz channels.," We observed in a pseudo-continuum, spectral line mode with $7 \times 3.125$ MHz channels."
 We monitored the point source L18 every 10 ain to provide amplitude. phase and baudpass calibration. aud used an observation of 3€ 286 to provide the absolute fiux calibration.," We monitored the point source $-$ 148 every 40 min to provide amplitude, phase and bandpass calibration, and used an observation of 3C 286 to provide the absolute flux calibration."
 We performed: standard spectral-line calibration and (ditius of the data usiue the NRAO ppackage. and enloved standard wide field imagine techniques (Tavlor.Carilli.&Perley.1999).," We performed standard spectral-line calibration and editing of the data using the NRAO package, and employed standard wide field imaging techniques \citep{tay99}."
. The final «1155 image has au ruis noise level of 15 jTy 1 except in the area close to the ceutral radio galaxy. which is Limited by the ability to clean the bright radio source (sce Fig. 2)).," The final $\times$ 5 image has an rms noise level of 15 $\mu$ Jy $^{-1}$, except in the area close to the central radio galaxy, which is limited by the ability to clean the bright radio source (see Fig. \ref{KMAMBOSNVLA}) )."
" The FWIIM resolution of the restoring beam is 273«173 at a position augle PA=0"".", The FWHM resolution of the restoring beam is $2\farcs3 \times 1\farcs3$ at a position angle .
. We obtained 850 yan and 150 gan photometry of 6 AMAABO sources previouslyidentified using the VLA and VLT tagging (83.1) using the Subimillimetre Conmion- Bolometer Array (SCUBA:Illudetal.1999) ou the 15m James Clerk Maxwell Telescope on UT2003 February 17 to 22. for a total iutegration tine of," We obtained 850 $\mu$ m and 450 $\mu$ m photometry of 6 MAMBO sources previouslyidentified using the VLA and VLT imaging 3.1) using the Submillimetre Common-User Bolometer Array \citep[SCUBA;][]{hol99} on the 15m James Clerk Maxwell Telescope on UT2003 February 17 to 22, for a total integration time of"
(hereafter PAID) at wavelengths 3.5<A13jan for redshifts ;=0.5 while still avoiding the coutamiuation your the Galactic IR cimus at longer wavelengths (Shupeetal. 1998).,"(hereafter PAH) at wavelengths $3.3 <
\lambda < 13~\micron$ for redshifts $z \lesssim 0.5$ while still avoiding the contamination from the Galactic IR cirrus at longer wavelengths \citep{Shupe98}."
". Measurements of IR πο counts aud luuinosity ""uctious (LPs) of salaxies demonstrate thle rapid evolution of the SER.", Measurements of IR number counts and luminosity functions (LFs) of galaxies demonstrate the rapid evolution of the SFR.
 Early IR stucies utilized the (RAS Bright Galaxy Sample (Soiferetal.1989). and theJRAS calibration field (lacking&TouckL987) aud focused on accounting for the cosmic infrared background (CIB:ELizetal.1999) with discreet MIR sources (Iackiugetal.1987:Saundersetal.1990:Soifer&Neugebauer 1991).," Early IR studies utilized the Bright Galaxy Sample \citep{Soifer89} and the calibration field \citep{HackingHouck87} and focused on accounting for the cosmic infrared background \citep[CIB;][]{Elbaz99}
 with discreet MIR sources \citep{Hacking87, Saunders90,
 SoiferNeu91}."
. Studies of galaxy number counts with750 aucSpitzer je shown that IR galaxies evolve very strongly (Elbazeauetal.2001Papovich200 1).," Studies of galaxy number counts with and have shown that IR galaxies evolve very strongly \citep{Elbaz99, Frances01, Dole04, Marleau04, Papovich04}."
.. As redshifts1: of IR siuuples became available. the LE. could be constructed in the nearby volume (Ricke&Lebofskv1986:Suiunu-al.1997:Shupoect1998). and subsequently exteuded to intermediate aud hieh redshifts (5> 3) with large siuuples of photometric redshifts (Pozzietal.2001:Leal.2007:Miuleauet2007) to study the SER evolution over cosuuic tine.," As redshifts of IR samples became available, the LF could be constructed in the nearby volume \citep{RiekeLebof86, Saunders90,
 Yahil91, RowRob97, Shupe98} and subsequently extended to intermediate and high redshifts $z > 3$ ) with large samples of photometric redshifts \citep{Pozzi04, LeFloch05, PPG05, Caputi07,
 Marleau07} to study the SFR evolution over cosmic time."
 However. most of these studies were either in the local regine (2<<0.1) or at intermediate to lüeh redshifts (2> 0.5).," However, most of these studies were either in the local regime $z<0.1$ ) or at intermediate to high redshifts $z>0.5$ )."
 The redshift rauge 0.1<20.6 includes more than | Cr of galaxy. evolution aud needs to be characterized iu detail to understand this process.," The redshift range $0.1 < z
< 0.6$ includes more than 4 Gyr of galaxy evolution and needs to be characterized in detail to understand this process."
 Observations of ealaxy evolution at low and intermediate redshifts face two challenges., Observations of galaxy evolution at low and intermediate redshifts face two challenges.
 First. a large area. deep survey is required to access a huge volume at lower redshifts. which is critical to minimize cosnic variance.," First, a large area, deep survey is required to access a large volume at lower redshifts, which is critical to minimize cosmic variance."
 For example. a deep survey at redshift +~0.5 would need five times the solid angle coverage of oue at Dod to probe a comparable comoving volune.," For example, a deep survey at redshift $z
\sim 0.5$ would need five times the solid angle coverage of one at $z
\sim 1$ to probe a comparable comoving volume."
 Second. while photo-i* are generally adequate for distance deterumination at hieh redshift. accurate spectroscopic redshifts are critical at low aud intermediate redshifts.," Second, while $z$ 's are generally adequate for distance determination at high redshift, accurate spectroscopic redshifts are critical at low and intermediate redshifts."
 To illustrate this issue. consider typical photometric redshift uncertainties of A:/(l|:)=0.05.," To illustrate this issue, consider typical photometric redshift uncertainties of $\Delta z/(1+z) = 0.05$."
 The corresponding fractional uncertainty in 2 would be a tolerable at zolL. while at 2~0.25 and below. the fractional nnecrtainty would exceed25%.," The corresponding fractional uncertainty in $z$ would be a tolerable at $z \sim 1$, while at $z \sim 0.25$ and below, the fractional uncertainty would exceed."
. Also. photometric redshifts are oulv accurate for galaxies whose spectral energv distributions (SEDs) are well matched by either SED templates or. iu the case of eumipirical photometric redshifts. ealaxies iu a spectroscopic training set.," Also, photometric redshifts are only accurate for galaxies whose spectral energy distributions (SEDs) are well matched by either SED templates or, in the case of empirical photometric redshifts, galaxies in a spectroscopic training set."
 This may not be the case for heavily obscured galaxies such as huninous IR ealaxies (LIRGs)., This may not be the case for heavily obscured galaxies such as luminous IR galaxies (LIRGs).
 Iu addition. the accuracy of some photometric redshift samples in the literature has not been verified with spectroscopy. and uukuowu systelatic errors nav be present.," In addition, the accuracy of some photometric redshift samples in the literature has not been verified with spectroscopy, and unknown systematic errors may be present."
 Iu this work we construct the rest-frame 21 pan LFs of star-forming galaxies. nieasure their evolution. and estimate the SFR for OO<ix0.6) from the LF.," In this work we construct the rest-frame 24 $\micron$ LFs of star-forming galaxies, measure their evolution, and estimate the SFR for $0.0 \leq z \leq 0.6$ from the LF."
 Iu addition to constraining the LF aud SER evolution accurately at low and iuteriuediate redshifts. we combine our results with those from the Far-IR Deep Extragalactic Legacy Survey (FIDEL.Magnellictal.2009) to extend the constraints to 5—1.2.," In addition to constraining the LF and SFR evolution accurately at low and intermediate redshifts, we combine our results with those from the Far-IR Deep Extragalactic Legacy Survey \citep[FIDEL,][]{Magnelli09} to extend the constraints to $z \sim 1.2$."
 We organize this paper as follows., We organize this paper as follows.
 Iu Section 77 we diseuss ourSpitzer 21 pau saunple. the spectroscopic— redshitts from ACN and Galaxies Evolution Survey (AGES). the optical catalog from NDWES. aud address theircompleteness.," In Section \ref{sec:data} we discuss our 24 $\micron$ sample, the spectroscopic redshifts from AGN and Galaxies Evolution Survey (AGES), the optical catalog from NDWFS, and address theircompleteness."
 Tn Section ?? we deseribe the coustruction of the LF. starting with the derivatious of galaxy Iuninosities (Section ??)). identification and exclusion of ACN from star-forming galaxies (Section ??)). estimators for the LF (Section ??7)). parameterization of the evolution of the LF with redshift (Section ??)). aud the estimation of the SFR from the LF (Section ??)).," In Section \ref{sec:IR_LF} we describe the construction of the LF, starting with the derivations of galaxy luminosities (Section \ref{sec:IR_LF_LI_L24_Estimation}) ), identification and exclusion of AGN from star-forming galaxies (Section \ref{sec:IR_LF_AGN}) ), estimators for the LF (Section \ref{sec:IR_LF_Method}) ), parameterization of the evolution of the LF with redshift (Section \ref{sec:IR_LF_evolparam}) ), and the estimation of the SFR from the LF (Section \ref{sec:IR_LF_SFR}) )."
 We then prescut the results iu Section ?? for the local LF (LLF: Section 229). the LF evolution at low aud imtcrmeciate redshifts (Section ??)). the LF evolution coustraint out to :~1.2 (Section 7?7)). as well as the SFR evolution of star-forming galaxies (Section ??)).," We then present the results in Section \ref{sec:results} for the local LF (LLF; Section \ref{sec:results_LLF}) ), the LF evolution at low and intermediate redshifts (Section \ref{sec:results_LF_evol}) ), the LF evolution constraint out to $z
\sim 1.2$ (Section \ref{sec:results_LF_evol_to_FIDEL}) ), as well as the SFR evolution of star-forming galaxies (Section \ref{sec:results_SFRD}) )."
 We stiuarize our results in Section ??7.., We summarize our results in Section \ref{sec:conclusions}.
" We asstume a ACDAL cosmology with Q,,—(3. O4=0.7. and {η=70 lan tMpe | throughout this paper."," We assume a $\Lambda$ CDM cosmology with $\Omega_m = 0.3$, $\Omega_{\Lambda} = 0.7$, and $H_0 =
70$ km $^{-1}$ $^{-1}$ throughout this paper."
 We construct the 21 ga LF of galaxies in the 9 sq., We construct the 24 $\micron$ LF of galaxies in the 9 sq.
 deeree field of the NOAO Deep Wide-Field Survey (NDWESJauuuzi&Dev 1999).," degree field of the NOAO Deep Wide-Field Survey \citep[NDWFS Data
 Release 3; $\alpha=14^h32^m$, $\delta=+34\arcdeg16\arcmin$,
 J2000.0;][]{Jannuz99}."
. The field las optical iaagiusg reaching the 5-5 fux limit of 26.6. 26.0. 26.0. and 21.1 in By. R. f. and AK (Veea uaenitudes). respectively.," The field has optical imaging reaching the $\sigma$ flux limit of 26.6, 26.0, 26.0, and 21.4 in $B_W$, $R$, $I$, and $K$ (Vega magnitudes), respectively."
 Details aboutthe NDWES can )e found iu Januuzi&Dev(1999). aud at the survevs Website.," Details aboutthe NDWFS can be found in \citet{Jannuz99}
 and at the survey's Web."
 Iu addition to the deep optical imaging. the feld has extensive panchromatic coverage from other survers. including au X-ray survey with the Space Telescope (NBoottes: comprising of 5 ks ACIS-I exposures. near-UV deep 30 ks iuaeiug fomGALEN (Alartinctal.2003).. the IRAC Shallow Survey in the 23.6. 1.5. 5.8. and 8.0 san bands reaching magnitudes of 19.1. 18.3. 15.9. aud 15.2 mag (Vega mae). respectively. (Eisenhardtetal.2001).. the deeperSpitzer Deep. Wide-Field Survey (SDWFS.Ashbyetal. 2009)... aud MIPS MIR photometry at 21421 down to 0.27 mJy (Papovichetal.2001:ITouck2005).," In addition to the deep optical imaging, the field has extensive panchromatic coverage from other surveys, including an X-ray survey with the Space Telescope \citep[XBo\""{o}t comprising of 5 ks ACIS-I exposures, near-UV deep 30 ks imaging from \citep{Martin03}, the IRAC Shallow Survey in the 3.6, 4.5, 5.8, and 8.0 $\micron$ bands reaching magnitudes of 19.1, 18.3, 15.9, and 15.2 mag (Vega mag), respectively, \citep{Eisenhardt04}, the deeper Deep, Wide-Field Survey \citep[SDWFS,][]{Ashby09}, and MIPS MIR photometry at 24 $\micron$ down to 0.27 mJy \citep{Papovich04, Houck05}."
. These nuagius data are complemented by the set of optical spectroscopic redshifts from the ACN and Calaxy Evolution Survey (ACES: Kochauek et abl.," These imaging data are complemented by the set of optical spectroscopic redshifts from the AGN and Galaxy Evolution Survey (AGES; Kochanek et al.,"
 iun prep.).," in prep.),"
 which is a statistically complete redshift survey of the field at low aud intermediate redshifts (0<2.< O8) (C. S. Ixochauck et al..," which is a statistically complete redshift survey of the field at low and intermediate redshifts $0 \leq z \leq 0.8$ ) (C. S. Kochanek et al.,"
 in preparation)., in preparation).
 The 21 jan photometry in the field was obtained by the Telescope wide- survey with the Multibaund huagiug Photometer for Spitzer (MIPS:Riekeetal.2001). and processed with the MIPS GTO Data AnalysisTool (DAT:Cordonot 2005)., The 24 $\micron$ photometry in the field was obtained by the wide-area survey with the Multiband Imaging Photometer for \citep[MIPS;][]{Rieke04} and processed with the MIPS GTO Data AnalysisTool \citep[DAT;][]{Gordon05}.
.The mean exposure tine per pixel across the ffüeld is 87 s. The pixel size of the final mosaic is 1725 with a PWIA of 577., .The mean exposure time per pixel across the field is 87 s. The pixel size of the final mosaic is $1\farcs25$ with a FWHM of $5\farcs7$ .
(Bdeefal.2010:Peeret2010).,"\citep{Ryde+10,
 Peer+10}."
. Oue clear prediction of the results preseuted here. is that if a thermal component contributed siguificantly to the observed spectrum. than the low cnereyv spectral slope (0) is expected to vary with time: at carly times. when the inuer eugine is active. à is expected to be steep (close to the Ravleigh-Jeaus tail). while at later times. FoxP.," One clear prediction of the results presented here, is that if a thermal component contributed significantly to the observed spectrum, than the low energy spectral slope $\alpha$ ) is expected to vary with time: at early times, when the inner engine is active, $\alpha$ is expected to be steep (close to the Rayleigh-Jeans tail), while at later times, $F_\nu \propto \nu^0$."
 This result is expected to be correlated with a decay of the peak energy aud the thermal flux. aud was possibly observed iu several bursts (e.e..Crider2009:etctαἱ 2009).," This result is expected to be correlated with a decay of the peak energy and the thermal flux, and was possibly observed in several bursts \citep[e.g.,][]{Crider97,
 Ghirlanda+07, RP09, Page+09}."
. We note though that in many CRBs. the data is not easily compared to the theoretical prediction. suce the outflows in CRBs are generally. not smooth. but very fluctuative.," We note though that in many GRBs, the data is not easily compared to the theoretical prediction, since the outflows in GRBs are generally not smooth, but very fluctuative."
 Therefore. photons originating from different emission episodes (cach characterized by its own mass loss ejection rate and its own Lorentz factor) are often superimposed.," Therefore, photons originating from different emission episodes (each characterized by its own mass loss ejection rate and its own Lorentz factor) are often superimposed."
 Thus. in fact. thermal photons may oulv be identified using detailed time resolved spectral analysis of separate flares.," Thus, in fact, thermal photons may only be identified using detailed time resolved spectral analysis of separate flares."
 This task was carried by Ryde&Peer. (2009).. which were indeed able to identify the thermal component in a large sample of bursts.," This task was carried by \citet{RP09}, which were indeed able to identify the thermal component in a large sample of bursts."
 Once done. this analysis method can be further used to deduce the properties of the CRB outflow (Peerc£af2001).," Once done, this analysis method can be further used to deduce the properties of the GRB outflow \citep{PRWMR07}."
. Iu addition to the prompt cuussion phase in CRBs. thermal activity may occur as part of the flaring activity observed in the carly afterglow phase of many CRBs (Burrowscfal2005:Falconeetal.2007:Chincarinictαἱ.2010:Maveuttiefal2010).," In addition to the prompt emission phase in GRBs, thermal activity may occur as part of the flaring activity observed in the early afterglow phase of many GRBs \citep{Burrows05, Falcone07, Chincarini+10,
 Margutti+10}."
. The exact nature of these flares is currently not vet clear., The exact nature of these flares is currently not yet clear.
 As it is plausible that the flares result from renewed ciission from the Inner core. a renewed thermal cussion may occur.," As it is plausible that the flares result from renewed emission from the inner core, a renewed thermal emission may occur."
 The analysis presented here may therefore apply iu the study of the late time fares as well., The analysis presented here may therefore apply in the study of the late time flares as well.
 While analysis of GRB prompt enuüssion spect serves as a niajor motivation to this work. we note that the analvsis presented here is general.," While analysis of GRB prompt emission spectrum serves as a major motivation to this work, we note that the analysis presented here is general."
 With inodifed paraueters values. the analvsis apply to emission frou anv astronomical transient. which is characterized ly relativistic outflow speeds aud high density core.," With modified parameters values, the analysis apply to emission from any astronomical transient, which is characterized by relativistic outflow speeds and high density core."
 Thus. we expect the analysis carricd here to be valid for objects such as ACGNs and mvicroquasars.," Thus, we expect the analysis carried here to be valid for objects such as AGNs and microquasars."
 This is valid because the exact nature of the radiative process that produces the thermal photous is of no importance. as long as the emission occurs deep enough iu the flow. where the optical depth is high. r1.," This is valid because the exact nature of the radiative process that produces the thermal photons is of no importance, as long as the emission occurs deep enough in the flow, where the optical depth is high, $\tau \gg 1$."
 We would also like to thauk Mario Livio. Zecljka Dosuujak. Chryvssa IKouveliotou. Charles Deriner. Ralph AALJ. Wijers. Peter Méssziuros. Martin Rees aud Bing Zhang for mauv useful discussions aud comments.," We would also like to thank Mario Livio, Z̆eeljka Bos̆nnjak, Chryssa Kouveliotou, Charles Dermer, Ralph A.M.J. Wijers, Peter Mésszárros, Martin Rees and Bing Zhang for many useful discussions and comments."
 AP is supported by the Riccardo Ciacconi Fellowship award of the Space Telescope Scicuce Institute., AP is supported by the Riccardo Giacconi Fellowship award of the Space Telescope Science Institute.
systematic errors associated with the cluster masses.,systematic errors associated with the cluster masses.
 For exiunple. estimates of cluster masses frou gravitational lensing τον often differ from those determined frou A-aav data (sometimes by up fo a factor of two: eS. Mivalda-Escudé Babul 1995: Wu Fane 1997: 0.1iu Ota et al.," For example, estimates of cluster masses from gravitational lensing very often differ from those determined from X-ray data (sometimes by up to a factor of two; e.g., Miralda-Escudé Babul 1995; Wu Fang 1997; 0.1in Ota et al."
 2002) and it is not vet clear why this happeus to be the case., 2002) and it is not yet clear why this happens to be the case.
 There could be a systematic problem with the N-rav. determined masses (ee. Allen 1998).," There could be a systematic problem with the X-ray determined masses (e.g., Allen 1998)."
 Aside from the yyM(rsoo) relation. we also explore the Sharedfo ALCrsyo) treud.," Aside from the $y_0-M(r_{500})$ relation, we also explore the $S_{\nu,arc}/f_{\nu}-M(r_{500})$ trend."
 This trend could provide a consistency clieck of the results discussed inunediatelv above., This trend could provide a consistency check of the results discussed immediately above.
 Unfortunately. the available SZ effect flux deusitv - nase data do not coustrain Ag.," Unfortunately, the available SZ effect flux density - mass data do not constrain $K_0$."
" This is not completely unexpected since the S,/f,—M(rsog) relation is mach less sensitive than the yyAMírsog) relation to the eutropv floor level of clusters (MDIIDBUS)."," This is not completely unexpected since the $S_{\nu,arc}/f_{\nu}-M(r_{500})$ relation is much less sensitive than the $y_0-M(r_{500})$ relation to the entropy floor level of clusters (MBHB03)."
" Future data from. for exiuuple. the will allow for much more precise determinations of the SZ effect. flux deusities of clusters aud we expect that the future Spiesf,συ) relation will place much tighter coustraiuts on Ay."," Future data from, for example, the will allow for much more precise determinations of the SZ effect flux densities of clusters and we expect that the future $S_{\nu,arc}/f_{\nu}-M(r_{500})$ relation will place much tighter constraints on $K_0$."
 Plotted in Figure bis the observed wyLy relation., Plotted in Figure 4 is the observed $y_0-T_X$ relation.
 Iu addition to the observational data. we have also plotted the +=0.2 (thick lines) aud +=0.5 (thin lines) theoretical relations.," In addition to the observational data, we have also plotted the $z = 0.2$ (thick lines) and $z = 0.5$ (thin lines) theoretical relations."
 The various line types (ce... solid. dotted) have the same meanings as in Figure 2.," The various line types (e.g., solid, dotted) have the same meanings as in Figure 2."
 Again. because the theoretical relations are not a strong function of redshift. a preliminary visual comparison of the data aud models is feasible.," Again, because the theoretical relations are not a strong function of redshift, a preliminary visual comparison of the data and models is feasible."
 First. it is clear that the data indicate that there is a fairly tight correlation between a clusters central Compton paraimcter and its oenmiüsson-welehted eax temperature.," First, it is clear that the data indicate that there is a fairly tight correlation between a cluster's central Compton parameter and its emission-weighted gas temperature."
 This correlation secius to hold true irrespective οἳ whether the clusters lave cooling flows or not., This correlation seems to hold true irrespective of whether the clusters have cooling flows or not.
 This may be expected since we lave used cooling flow-corrected temperatures., This may be expected since we have used cooling flow-corrected temperatures.
 However. even without this correction the data still 0.1iu exhibit a. tight trend.," However, even without this correction the data still 0.1in exhibit a tight trend."
 For the clusters in this sinple. cooling flow correction only slightly iucreascs the temperatire of a cluster (by about or L keV) aud. therefore. does not siguificautly affect the results.," For the clusters in this sample, cooling flow correction only slightly increases the temperature of a cluster (by about or 1 keV) and, therefore, does not significantly affect the results."
 There is also a lint of a slight svsteinatic difference iu the yoTx relatious for low and high redshift clusters (filled and open svnibols. respectively).," There is also a hint of a slight systematic difference in the $y_0-T_X$ relations for low and high redshift clusters (filled and open symbols, respectively)."
 In particular. the relation for the high redshift clusters has a normalization that is slightly higher than that of the relation for low redshift clusters Gguoring RNJISLF. whose temperature is quite uncertain).," In particular, the relation for the high redshift clusters has a normalization that is slightly higher than that of the relation for low redshift clusters (ignoring RXJ1347, whose temperature is quite uncertain)."
 This is discussed iu more detail below., This is discussed in more detail below.
 Comparing the various theoretical relations το the observational data. it is apparent that the standard sclsimular model predicts a yy£x relation which is à poor match to the data.," Comparing the various theoretical relations to the observational data, it is apparent that the standard self-similar model predicts a $y_0-T_X$ relation which is a poor match to the data."
 In particular. the uormalizatious of the selfsimuilar relations are roughly 2.5 times Larger thu observed (over the range 6 keV zTy=12 keV).," In particular, the normalizations of the self-similar relations are roughly 2.5 times larger than observed (over the range 6 keV $\lesssim T_X \lesssim 12$ keV)."
 The eutropy floor models with Ny~300 keV eng. on the other haud. provide a very good qualitative fit to the data.," The entropy floor models with $K_0 \sim 300$ keV $^2$, on the other hand, provide a very good qualitative fit to the data."
 Iu MDIIDBUS. we found that injecting the ICAL with nou-eravitational cutropy teuds to decrease yy aud. at the same time. increase Ty.," In MBHB03, we found that injecting the ICM with non-gravitational entropy tends to decrease $y_0$ and, at the same time, increase $T_X$."
 This reduces the normalization of the predicted trend between these two quantities aud. as is apparent frou Fieure 1. brings close agrecment between theory and observations.," This reduces the normalization of the predicted trend between these two quantities and, as is apparent from Figure 4, brings close agreement between theory and observations."
 Fitting all 21 clusters in Tables 1 aud 2 via the method outlined iu 83.1. we obtain a best-fit cutropy floor of Ky=3000 keV can? with 42=26.36/201.32.," Fitting all 21 clusters in Tables 1 and 2 via the method outlined in 3.1, we obtain a best-fit entropy floor of $K_0 = 300^{+80}_{-60}$ keV $^2$ with $\chi^2_{\nu} = 26.36/20
= 1.32$."
 The residuals of a comparison between the Ay=300 keV au? cutropy floor model and the self-similar model with the observation data are plotted in Figure 5., The residuals of a comparison between the $K_0 = 300$ keV $^2$ entropy floor model and the self-similar model with the observation data are plotted in Figure 5.
 Note the striking separation between the residuals of the two models., Note the striking separation between the residuals of the two models.
" Theoretically, the yyFy relation is the most"," Theoretically, the $y_0-T_X$ relation is the most"
"Note that in the standard case, where the background redshift is identified with the observed one, fp(ap) is simply replaced by the growthrate f(a)=2m and measuring the Hubble fluctuations would yield a direct measurement of f.","Note that in the standard case, where the background redshift is identified with the observed one, $f_\cD\left(a_{\cD}\right)$ is simply replaced by the growthrate $f\left(a\right)=\frac{\rmd \ln D\left(a\right)}{\rmd \ln a}$, and measuring the Hubble fluctuations would yield a direct measurement of $f$."
" In the real world, where the redshift captures the structure on the way from the source to us, it is not directly the background redshift."," In the real world, where the redshift captures the structure on the way from the source to us, it is not directly the background redshift."
" One would rather measure the modified ""growth rate” f(ap)."," One would rather measure the modified ""growth rate"" $f_\cD \left(a_{\cD}\right)$."
" The difference between these two quantities is small in our range of validity for fp(ap), however (corrections of linear order in the perturbations)."," The difference between these two quantities is small in our range of validity for $f_\cD\left(a_{\cD}\right)$, however (corrections of linear order in the perturbations)."
" For the first time, we brought together the well-established perturbative approach to incorporate inhomogeneities in Friedman-Lemaittre models (the theory of cosmological perturbations) and the ideas of cosmological backreaction and cosmic averaging in the Buchert formalism."," For the first time, we brought together the well–established perturbative approach to incorporate inhomogeneities in Friedman-Lemaîttre models (the theory of cosmological perturbations) and the ideas of cosmological backreaction and cosmic averaging in the Buchert formalism."
" Focusing on observations of the large-scale structure of the Universe at late times, we showed that the cosmic variance of cosmological parameters is in fact the leading order contribution of cosmological averaging."," Focusing on observations of the large–scale structure of the Universe at late times, we showed that the cosmic variance of cosmological parameters is in fact the leading order contribution of cosmological averaging."
" We studied volume averages, their expected means, and variances of the cosmological parameters Ho,QR,Qn,Qn (Qa is of higher order in perturbation theory)."," We studied volume averages, their expected means, and variances of the cosmological parameters $H_0, \Omega_\CR, \Omega_m,
\Omega_\Lambda$ $\Omega_\CQ$ is of higher order in perturbation theory)."
 Our central result is summarized in (47))-(51))., Our central result is summarized in \ref{eq:sigH}) \ref{eq:sigOQ}) ).
 Our extension of the backreaction study of Li&Schwarz(2008) to the ACDM case enabled us to study fluctuations for a wider class of cosmological models., Our extension of the backreaction study of \cite{li:scale} to the $\Lambda$ CDM case enabled us to study fluctuations for a wider class of cosmological models.
" We were able to confirm for the fluctuations in the Hubble rate that our results in comoving synchronous gauge agree with those found in Poisson gauge (Umehetal.,2010)..", We were able to confirm for the fluctuations in the Hubble rate that our results in comoving synchronous gauge agree with those found in Poisson gauge \citep{clarkson:hubble}.
" The use of general window functions allowed us to consider realistic survey geometries in detail and to calculate the fluctuations in the matter density, empirically found in the SDSS data by Driver&Robotham(2010),, directly from the underlying DM power spectrum."," The use of general window functions allowed us to consider realistic survey geometries in detail and to calculate the fluctuations in the matter density, empirically found in the SDSS data by \cite{driver:cosvar}, directly from the underlying DM power spectrum."
" Converting this information into curvature fluctuations, we found that regions of 540!Mpc diameter may still have a curvature parameter of ~0.01, even if the background curvature vanishes exactly."," Converting this information into curvature fluctuations, we found that regions of $540 h^{-1} \mathrm{Mpc}$ diameter may still have a curvature parameter of $\sim 0.01$, even if the background curvature vanishes exactly."
 We found that cosmic variance is a limiting factor even for surveys of the size of the SDSS survey., We found that cosmic variance is a limiting factor even for surveys of the size of the SDSS survey.
 A volume-limited sample up to a redshift of 0.5 was able to constrain the local curvature to 0.1 per cent., A volume–limited sample up to a redshift of $0.5$ was able to constrain the local curvature to $0.1$ per cent.
" Finally, we investigated the distortions of the local distance to a given redshift and found that it is less affected by the fluctuations of the local cosmic parameters than one mightexpect."," Finally, we investigated the distortions of the local distance to a given redshift and found that it is less affected by the fluctuations of the local cosmic parameters than one mightexpect."
" The distance measure Dy, used in BAO studies, is accurate to 0.2 percent for samples ranging up to z«0.35."," The distance measure $D_{V}$, used in BAO studies, is accurate to $0.2$ percent for samples ranging up to $z\approx0.35$."
" This means that BAO studies are not limited by cosmic variance, but by problems such as sampling variance and insufficient volume, as discussed in section 7.."," This means that BAO studies are not limited by cosmic variance, but by problems such as sampling variance and insufficient volume, as discussed in section \ref{sec:Variances-on-BAO}."
 In a next step one should incorporate the second-order effects into the expected means of the cosmological parameters., In a next step one should incorporate the second–order effects into the expected means of the cosmological parameters.
" There are no second-order corrections to the variances, as argued in section 2.."," There are no second–order corrections to the variances, as argued in section \ref{sec:Inhomogeneity-and-expansion}."
 Therefore a complete second-order treatment seems feasible., Therefore a complete second–order treatment seems feasible.
 The limitation of our approach comes from the fact that Buchert's formalism relies on spatial averaging., The limitation of our approach comes from the fact that Buchert's formalism relies on spatial averaging.
" Averaging on the light cone would be more appropriate (Gasperiniet thus the study in this work has been restricted to redshifts «1, where we expect light cone effects to play a subdominant role."," Averaging on the light cone would be more appropriate \citep{lightcone}, , thus the study in this work has been restricted to redshifts $\ll 1$ , where we expect light cone effects to play a subdominant role."
wwe have considered τι=0.,we have considered $\tau_{\rm h}=0$.
" We now show this ellect in Fieure 3. where we [ave fixed the eritical velocity my to 0.01. and calculated the dimensionless density and temperature for two cdilferent. botary lavers. with 7,=0.001 and 0.0001."," We now show this effect in Figure 3, where we have fixed the critical velocity $\tau_{\rm m}$ to 0.01, and calculated the dimensionless density and temperature for two different boundary layers, with $\tau_{\rm h} = 0.001$ and $0.0001$."
 Phey roth fulfill the constraint that zy;<my., They both fulfill the constraint that $\tau_{\rm h} < \tau_{\rm m}$.
 As shown. the terminated temperature depends mainly on my.," As shown, the terminated temperature depends mainly on $\tau_{\rm m}$."
 “Phe heat input fron he white dwarf mocifies only the asvniptotic properties of t10 base temperature a larger 7y will result in a slightly thicker “isothermal” [aver at the base. thus reducing the temperaare upstLCA anc causing the density to reach the terminated rlateau value further upstream.," The heat input from the white dwarf modifies only the asymptotic properties of the base temperature — a larger $\tau_{\rm h}$ will result in a slightly thicker “isothermal” layer at the base, thus reducing the temperature upstream and causing the density to reach the terminated plateau value further upstream."
 Figure 4 shows the dimensionless local bremsstrahlung emissivity jie (= 070677) as a function of the height £ in the region., Figure 4 shows the dimensionless local bremsstrahlung emissivity $j_{\rm br}$ $= \zeta^2 \theta^{1/2}$ ) as a function of the height $\xi$ in the post-shock region.
 The emissivity reaches a terminated value deermined by my. Wwvereas Chat from the conventional formulation ends to infinity.," The emissivity reaches a terminated value determined by $\tau_{\rm m}$, whereas that from the conventional formulation tends to infinity."
 Again the bremsstrahlung emissivity is nhodilied slightly by the input of heat from the white dwarf when TjL0 with increased emission for larger values of τμ., Again the bremsstrahlung emissivity is modified slightly by the input of heat from the white dwarf when $\tau_{\rm h} > 0$ with increased emission for larger values of $\tau_{\rm h}$.
 This is in fact the consequence of an implicit assumption that the white-dwarl energy Dux is much weaker than energy. [ux receased by accretion material., This is in fact the consequence of an implicit assumption that the white-dwarf energy flux is much weaker than energy flux released by accretion material.
55433 is the first known example of a Galactic relativistic jel source. and thus the forerunner of modern microquasar astrophivsies.,"	SS433 is the first known example of a Galactic relativistic jet source, and thus the forerunner of modern microquasar astrophysics."
 The optical spectrum of this object shows a number of strong. broad emission lines of (he Balmer and. [lel series. as well as several lines al unusual wavelengths.," The optical spectrum of this object shows a number of strong, broad emission lines of the Balmer and HeI series, as well as several lines at unusual wavelengths."
 These latter have been idenüfied as red/blue-shifted Balmer and Hel emission from collimated jets with intrinsic velocities of v£z0.260 1979).., These latter have been identified as red/blue-shifted Balmer and HeI emission from collimated jets with intrinsic velocities of $v \simeq 0.26c$ \citep{AbellMargon}.
" Furthermore. the Doppler shifts of these features change will ime in a cosinusoidal manner. leading to the label of ""moving lines”."," Furthermore, the Doppler shifts of these features change with time in a cosinusoidal manner, leading to the label of “moving lines”."
 This behavior is now widely accepted to be a svinplom of precession of (he jet axis in $5433 on a timescale of ~164 days (Margon1984)., This behavior is now widely accepted to be a symptom of precession of the jet axis in SS433 on a timescale of $\sim 164$ days \citep{Margon84}.
. Early studies of the precession in 55422 indicated possible instabilities or crits in the precessional clock (Andersonetal.1983)... whieh could give considerable insight into the accretion processes which must provide the precessional torque.," 	Early studies of the precession in SS433 indicated possible instabilities or drifts in the precessional clock \citep{Anderson}, which could give considerable insight into the accretion processes which must provide the precessional torque."
 However. \lareonson(1989) reviewed ten vears of $5432. (ming data and concluded Chat while significant," However, \citet{MargonAnderson} reviewed ten years of SS433 timing data and concluded that while significant"
"The thermodynamic trajectories of neutrino-driven outflows are obtained using a semi-analytic, spherically symmetric, general relativistic model of neutrino-driven winds.","The thermodynamic trajectories of neutrino-driven outflows are obtained using a semi-analytic, spherically symmetric, general relativistic model of neutrino-driven winds."
 This model has been developed in previous r- (Wanajoetal.2001;Wanajo2007) and vp-process (Wanajo2006) studies.," This model has been developed in previous r-process \citep{Wana2001,Wana2007} and $\nu$ p-process \citep{Wana2006}
 studies."
" Here, we describe several modifications added to the previous version."," Here, we describe several modifications added to the previous version."
" The equation of state for ions and arbitrarily degenerate, arbitrarily relativistic(ideal gas)electrons and positrons is taken from Timmes&Swesty(2000)."," The equation of state for ions (ideal gas) and arbitrarily degenerate, arbitrarily relativistic electrons and positrons is taken from \citet{Timm2000}."
". The root-mean-square averaged energies of neutrinos are taken to be 12, 14, and 14 MeV, for electron, anti-electron, and the other types of neutrinos, respectively,in light of a recent self-consistently exploding model of a 9M star (Kitaura,Janka,&Hillebrandt2006;2010).."," The root-mean-square averaged energies of neutrinos are taken to be 12, 14, and 14 MeV, for electron, anti-electron, and the other types of neutrinos, respectively,in light of a recent self-consistently exploding model of a $9 M_\odot$ star \citep{Kita2006, Hued2010, Muel2010}. ."
" These values are consistent with other recent studies for more massive progenitors (Fischeretal.2010), but substantially smaller than those taken in previous works (e.g.,12,22,and34MeVinWanajoetal.2001)."," These values are consistent with other recent studies for more massive progenitors \citep{Fisc2010}, , but substantially smaller than those taken in previous works \citep[e.g., 12, 22, and 34~MeV in ][]{Wana2001}."
". The mass ejection rate M at the neutrino sphere is determined such that the outflow becomes supersonic (i.e., wind)) through the sonic point."," The mass ejection rate $\dot M$ at the neutrino sphere is determined such that the outflow becomes supersonic (i.e., ) through the sonic point."
 The neutron star mass Λάπς is taken to be 1.4Mo for our standard model.," The neutron star mass $M_\mathrm{ns}$ is taken to be $1.4\, M_\odot$ for our standard model."
" The radius of the neutrino sphere is assumed to be R,(L,)=(Ryo—Ryi)(Ly/Lyo)+Ry as a function of the neutrino luminosity L, (taken to be the same for all the flavors), where Ryo=30km, R,,=10km, and [νο=109263.98x10°ergss7!."," The radius of the neutrino sphere is assumed to be $R_\nu (L_\nu) = (R_{\nu 0} - R_{\nu 1}) (L_\nu/L_{\nu 0}) +
R_{\nu 1}$ as a function of the neutrino luminosity $L_\nu$ (taken to be the same for all the flavors), where $R_{\nu 0} = 30\,
\mathrm{km}$, $R_{\nu 1} = 10\, \mathrm{km}$, and $L_{\nu 0} =
10^{52.6} = 3.98 \times 10^{52} \, \mathrm{ergs\ s}^{-1}$."
" This roughly mimics the evolution of R, in recent hydrodynamic simulations Burasetal.2006;Ar-conesetal. 2007)."," This roughly mimics the evolution of $R_\nu$ in recent hydrodynamic simulations \citep[e.g.,][]{Bura2006, Arco2007}."
". The wind(e.g., solution is obtained with L,=1x10?ergs~! (R,=12.5 km) for the standard model."," The wind solution is obtained with $L_\nu = 1 \times 10^{52}\, \mathrm{erg\ s}^{-1}$ $R_\nu = 12.5$ km) for the standard model."
" 'The time variations of radius r from the center, density p, and temperature T' for the standard model are shown in Figure 1 (black line)."," The time variations of radius $r$ from the center, density $\rho$, and temperature $T$ for the standard model are shown in Figure 1 (black line)."
" The time variations of r, p, and T after the wind- by the preceding supernova ejecta are calculated asfollows."," The time variations of $r$, $\rho$, and $T$ after the wind-termination by the preceding supernova ejecta are calculated asfollows."
" This phase is governed by the evolution of the preceding slowly outgoing ejecta, independent of the wind solution."," This phase is governed by the evolution of the preceding slowly outgoing ejecta, independent of the wind solution."
" In light of recent hydrodynamical calculations Arconesetal.2007),, we assume the time evolution (e.g.,of the outgoing ejecta to be pxt? and Tοςt~?/3, where t is the post-bounce time."," In light of recent hydrodynamical calculations \citep[e.g.,][]{Arco2007}, we assume the time evolution of the outgoing ejecta to be $\rho \propto t^{-2}$ and $T \propto
t^{-2/3}$, where $t$ is the post-bounce time."
" With these relations, we have for t>twt, where fa, Uwt, Twt, Pwt, and Tyare the time, velocity, radius, density, and temperature, respectively, just after the wind-termination."," With these relations, we have for $t > t_\mathrm{wt}$, where $t_\mathrm{wt}$, $u_\mathrm{wt}$, $r_\mathrm{wt}$ , $\rho_\mathrm{wt}$ , and $T_\mathrm{wt}$are the time, velocity, radius, density, and temperature, respectively, just after the wind-termination."
 Equation (7) represents the time variation of velocity after the wind-termination., Equation (7) represents the time variation of velocity after the wind-termination.
" In case that ryt is larger than that at the sonic point, rs, the Rankine-Hugoniot shock-jump conditions are applied at ryt to obtain t, Pwt, and Ty (see,e.g.,Arconesetal.2007;Kuroda, 2008). "," In case that $r_\mathrm{wt}$ is larger than that at the sonic point, $r_\mathrm{s}$, the Rankine-Hugoniot shock-jump conditions are applied at $r_\mathrm{wt}$ to obtain $u_\mathrm{wt}$, $\rho_\mathrm{wt}$, and $T_\mathrm{wt}$ \citep[see,
e.g.,][]{Arco2007,Kuro2008}. ."
"Equations(6) and (7) are obtained from equation (4) with the steady-state condition, ie., r?pu= constant (seePanov&Janka 2009).."," Equations(6) and (7) are obtained from equation (4) with the steady-state condition, i.e., $r^2 \rho u =$ constant \citep[see][]{Pano2009}. ."
" Note that equations (6) and (7) gives r(t)ct and u(t)=constantfort>> t4.In order to obtain t in equations (4)-(7) for a given trajectory with L,, the time evolution of L, at the neutrino sphereis assumed to be [L,(t),m,= Lyo(t/to)|, where t>to=0.2 8"," Note that equations (6) and (7) gives $r(t) \propto t$ and $u(t) = \mathrm{constant}$for$t \gg
t_\mathrm{wt}$ .In order to obtain $t$ in equations (4)-(7) for a given trajectory with $L_\nu$ , the time evolution of $L_\nu$ at the neutrino sphereis assumed to be $[L_\nu(t)]_{r = R_\nu} = L_{\nu 0}
(t/t_0)^{-1}$ , where $t > t_0 = 0.2$ s"
EGMF strength.,EGMF strength.
 Fig. 3((, Fig. \ref{fig:fitExample}( (
"a) shows a sample fit for a= 1.5, Eg=25 TeV, l=10, and B=3x10719 Gauss.","a) shows a sample fit for $\alpha=1.5$ , $E_0=25$ TeV, $\Gamma=10$, and $B=3\times 10^{-16}$ Gauss."
" Because we cannot exclude the existence of additional components contributing to the blazar emission (e.g.,Bóttcheretal.2008) and modifying the intrinsic spectrum of Eq. 9,,"," Because we cannot exclude the existence of additional components contributing to the blazar emission \citep[e.g.,][]{Bottcher2008} and modifying the intrinsic spectrum of Eq. \ref{eq:boost},"
" we also fit a broken power law with an index Qpreak below 80 GeV, as shown for instance in Fig. 3(("," we also fit a broken power law with an index $\alpha_\text{break}$ below 80 GeV, as shown for instance in Fig. \ref{fig:fitExample}( ("
b).,b).
 We find that the additional free parameter does not greatly affect our constraints on the EGMF., We find that the additional free parameter does not greatly affect our constraints on the EGMF.
" For each EGMF strength B, we find the best-fit model of the combined cascade and intrinsic emission for a wide range of cutoff energies Eo€(0.1 TeV, 100 TeV), allowing α and Qpreak to vary between 2.5 and the physically motivated constraint of 1.5 (e.g.,Malkov&O'CDrury2001;Aharonianetal. and fixing Τ to a typical value of 10."," 	For each EGMF strength $B$, we find the best-fit model of the combined cascade and intrinsic emission for a wide range of cutoff energies $E_0\in(0.1$ TeV, 100 TeV), allowing $\alpha$ and $\alpha_\text{break}$ to vary between 2.5 and the physically motivated constraint of 1.5 \citep[e.g.,][]{Malkov2001,HESS_2006} and fixing $\Gamma$ to a typical value of 10."
 We plot the x?2006) value of the best-fit model as a function of B in Fig., We plot the $\chi^2$ value of the best-fit model as a function of $B$ in Fig.
" 4 for blazar livetimes from 1 year to 10° years, and for the unlimited case."," \ref{fig:time_cuts} for blazar livetimes from 1 year to $10^6$ years, and for the unlimited case."
" At low B, the x? values converge because the cascade arrives promptly and at small angles."," At low $B$, the $\chi^2$ values converge because the cascade arrives promptly and at small angles."
" As B increases, the arrival angles and times of the cascade begin to spread out, diminishing the observed emission and providing a better fit to the observed data."," As $B$ increases, the arrival angles and times of the cascade begin to spread out, diminishing the observed emission and providing a better fit to the observed data."
 'The convergence of the curves in Fig., 	The convergence of the curves in Fig.
 4 to the infinite-livetime curve is easily understood., \ref{fig:time_cuts} to the infinite-livetime curve is easily understood.
 Combining Eq., Combining Eq.
 7 with Eq., \ref{eq:angle} with Eq.
" 8 and assuming small angles, we can translate a cut on the instrument angle 0, into a time cut AT;: For example, source photons at 1 TeV, which have a mean free path of L'~400 Mpc, will produce cascade photons of energy E=0.8 GeV, for which an angular cut of 6.z1° is appropriate for the LAT."," \ref{eq:time} and assuming small angles, we can translate a cut on the instrument angle $\theta_c$ into a time cut $\Delta T_c$: 	For example, source photons at 1 TeV, which have a mean free path of $L'\approx400$ Mpc, will produce cascade photons of energy $E\approx0.8$ GeV, for which an angular cut of $\theta_c\approx1^\circ$ is appropriate for the LAT."
" At the distance of RGB J0710--591 (L=500 Mpc), this translates into a time cut of AT,zz6x104 years."," At the distance of RGB J0710+591 $L\approx500$ Mpc), this translates into a time cut of $\Delta T_c\approx6\times10^4$ years."
 This is close the the livetime of ~10° years at which the X? curves in Fig., This is close the the livetime of $\sim10^5$ years at which the $\chi^2$ curves in Fig.
" 4 begin to converge to the unlimited-livetime case, becoming nearly indistinguishable at ~109 years."," \ref{fig:time_cuts} begin to converge to the unlimited-livetime case, becoming nearly indistinguishable at $\sim10^6$ years."
" If the blazar livetime is smaller, the livetime cut outweighs the angle cut, and the position of the x? curve depends on the blazar livetime."," If the blazar livetime is smaller, the livetime cut outweighs the angle cut, and the position of the $\chi^2$ curve depends on the blazar livetime."
" For longer livetimes, the angle cut becomes more constraining and the curves converge to the unlimited-livetime case."," For longer livetimes, the angle cut becomes more constraining and the curves converge to the unlimited-livetime case."
 We constrain the EGMF strength for a given livetime by finding the point at which x? exceeds its minimum value by Ax? for each curve in Fig. 4.., 	We constrain the EGMF strength for a given livetime by finding the point at which $\chi^2$ exceeds its minimum value by $\Delta\chi^2$ for each curve in Fig. \ref{fig:time_cuts}.
" Two sample confidence levels and 95%)), corresponding toAx? values of 2.72 and 3.84 (seee.g.,James2006) are indicated in the figure."," Two sample confidence levels and ), corresponding to$\Delta\chi^2$ values of 2.72 and 3.84 \citep[see e.g.,][]{James2006} are indicated in the figure."
 The limits derived from these confidence levels are shown in Fig., The limits derived from these confidence levels are shown in Fig.
 5 asa function of the blazar livetime., \ref{fig:confidence_limits} as a 	function of the blazar livetime.
" For livetimes below ~10* years, the limit on the EGMF strength scales with the blazar livetime"," For livetimes below $\sim10^4$ years, the limit on the EGMF strength scales with the blazar livetime"
contamination of οr photographic plate catalogue are argo.,contamination of our photographic plate catalogue are large.
 Qur star-galaxy separation software was inipoerfect snce out of the 110 objects for which redshufts were ueasured only 305 were galaxies., Our star-galaxy separation software was imperfect since out of the 410 objects for which redshufts were measured only 305 were galaxies.
 Such a high degree of contamination of the catalogue is at least partly due to he fac that we preferred to have a galaxy catalogue as colplete as possible. iu order not to “iniss” galaxies.," Such a high degree of contamination of the catalogue is at least partly due to the fact that we preferred to have a galaxy catalogue as complete as possible, in order not to “miss” galaxies."
 Our catalogue includes these 305 galaxy spectra. plus those oxeviouxlv published by Prous et al. (," Our catalogue includes these 305 galaxy spectra, plus those previously published by Proust et al. ("
1987). Quintana irez (1990) and Maliunuth et al. (,"1987), Quintana rez (1990) and Malumuth et al. ("
1992). aud some galaxies roni the CfA redshift catalogue (IIuchra et al.,"1992), and some galaxies from the CfA redshift catalogue (Huchra et al."
 1992). reaching a total of 166 ealaxy vedshifts after climinating objects observed twice.," 1992), reaching a total of 466 galaxy redshifts after eliminating objects observed twice."
 The positious of the objects for which we gathered reliable spectra (either from our observations or frou he literature] are shown in Fig. L.., The positions of the objects for which we gathered reliable spectra (either from our observations or from the literature) are shown in Fig. \ref{xy}.
 These positions are relative to the cluster center taken to be the position of the maxim N-ray enüssiou (Pislar 1908): Aoggu.y=p33737.97.699000—.—E15 E77.," These positions are relative to the cluster center taken to be the position of the maximum X-ray emission (Pislar 1998): $\alpha_{2000.0} = 04^h33^{m}37.9^s, \delta_{2000.0} = -13^\circ
15'47''$ ."
 This center is within To arcsec of the position of the cD ealaxy. a distance which is sinaller than the ROSAT PSPC pixel size. aud we will herefore consider hereafter that both positions coiucido.," This center is within 7 arcsec of the position of the cD galaxy, a distance which is smaller than the ROSAT PSPC pixel size, and we will therefore consider hereafter that both positions coincide."
 The spectra were reduced using the IRAF software., The spectra were reduced using the IRAF software.
 The fraanes were bias corrected iu the usual wax., The frames were bias corrected in the usual way.
 Velocities were nieasured by cross-correlating the observed spectra with different templates: a spectra of ALS Usinedly provided bv J. Perea) at a velocity of 200 kins 1. and stellar spectra of the standard stars WD 21331 and ΠΟ 18381. which were each observed every might during our 1991 run.," Velocities were measured by cross-correlating the observed spectra with different templates: a spectrum of M31 (kindly provided by J. Perea) at a velocity of $-300$ km $^{-1}$, and stellar spectra of the standard stars HD 24331 and HD 48381, which were each observed every night during our 1994 run."
 The cross-correlation techuique is that described by Toury Davis (1979) aud plemented in the NCSAO task of the RVSAO package in IRAF (surtz et al., The cross-correlation technique is that described by Tonry Davis (1979) and implemented in the XCSAO task of the RVSAO package in IRAF (Kurtz et al.
 1991)., 1991).
 The errors on the velocities derived from absorption lines are eiven automatically by this task., The errors on the velocities derived from absorption lines are given automatically by this task.
 The positious of cussion lines. when present. were measured by fitting cach liue with a gaussiau.," The positions of emission lines, when present, were measured by fitting each line with a gaussian."
 All the redshifts were measured by the same person (F.D.) iu a homogeneous wax., All the redshifts were measured by the same person (F.D.) in a homogeneous way.
 Recdshifts of msufBcieut quality were discarded G.c. those with a Tour Davis paraicter siualler than 2.0. except for three ealaxics with respective Toury Davis parameters of 1.5. 1.6 and 1.9 where the absorption lines seemed to be well enough defined for these redshifts to be kept in the final catalogue).," Redshifts of insufficient quality were discarded (i.e. those with a Tonry Davis parameter smaller than 2.0, except for three galaxies with respective Tonry Davis parameters of 1.5, 1.6 and 1.9 where the absorption lines seemed to be well enough defined for these redshifts to be kept in the final catalogue)."
 For ealaxies with absorption lines. two velocity standard stars from the Maurice et al. (," For galaxies with absorption lines, two velocity standard stars from the Maurice et al. ("
198D) list were observed cach welt in order to check the intrinsic quality of our velocity measurements.,1984) list were observed each night in order to check the intrinsic quality of our velocity measurements.
 The errors. derived by cross-correlating the star spectra to the spectruu of N31. range(from wieght to nieht) from +16 to £23 lan | for ΠΟ 21331. and from +17 to £Vkans ο for IID Ls83sl.," The errors, derived by cross-correlating the star spectra to the spectrum of M31, range(from night to night) from $\pm$ 16 to $\pm 23$ km $^{-1}$ for HD 24331, and from $\pm$ 17 to $\pm 43$ km $^{-1}$ for HD 48381."
 The mean internal error ou velocities derived from the in the mean wavelength calibration is 66 kms 1, The mean internal error on velocities derived from the in the mean wavelength calibration is 66 km $^{-1}$.
 For Cluission line measurements. the ClYTOlYS Oll velocities Were estimated. from the dispersion. of the velocities derived frou the various emission Ines present.," For emission line measurements, the errors on velocities were estimated from the dispersion of the velocities derived from the various emission lines present."
 When only one emission Ine was present we averaged the cussion aud absorption line redshifts whenever possible: if no reliable absorption line redshift was available. we estimated the internal error on a single cussion line to be the iutrinsic value of 66 kins 1.," When only one emission line was present we averaged the emission and absorption line redshifts whenever possible; if no reliable absorption line redshift was available, we estimated the internal error on a single emission line to be the intrinsic value of 66 km $^{-1}$ ."
 The nunber of redshifts obtained frou emission lines is 85., The number of redshifts obtained from emission lines is 85.
main sequence stars of normal chemical composition and. as next step. generalized. to variable giant stars with large metal ceficiency. Clementinietal 2000)).,"main sequence stars of normal chemical composition and, as next step, generalized to variable giant stars with large metal deficiency, \citealt{clem1}) )."
 Of course. in comparison with a line bv line analysis of ligh-cispersion spectroscopy. a multicolour photometry is sensitive only to the general shape of the optical flux of the star.," Of course, in comparison with a line by line analysis of high-dispersion spectroscopy, a multicolour photometry is sensitive only to the general shape of the optical flux of the star."
 However. it is possible to find the model atmosphere ον reproducing the continuum Ilux in the visible wavelength interval.," However, it is possible to find the model atmosphere by reproducing the continuum flux in the visible wavelength interval."
 Our method. can provide a global parameter of he chemical composition in the atmosphere summarized as metallicity M]., Our method can provide a global parameter of the chemical composition in the atmosphere summarized as metallicity [M].
 This. M] is the most. suitable for our purposes. because it accounts for the elfect of all elements on the optical continuum. including those that do not have ines.," This [M] is the most suitable for our purposes, because it accounts for the effect of all elements on the optical continuum, including those that do not have lines."
 Of course. it might. cülfer from the overall metallicity of the star (used. in the theory of stellar structure. as he pulsation does not stir up the deepest. lavers where he nuclear reactions take place) or from the averaged metallicity derived from high dispersion spectroscopy.," Of course, it might differ from the overall metallicity of the star (used in the theory of stellar structure, as the pulsation does not stir up the deepest layers where the nuclear reactions take place) or from the averaged metallicity derived from high dispersion spectroscopy."
 The fields are at high Galactic. latitude: GSC 4868-31. b=|24°: V312 Ser. b=(45°.," The fields are at high Galactic latitude: GSC 4868-0831, $b= +24\degr$; V372 Ser, $b=+45\degr$."
 Their recdening is very small. ancl therefore the parameters (logq.. Z5. d. AME) reported. in Section 4 are lower limits at the same ime if M] is fHixed.," Their reddening is very small, and therefore the parameters $\log g_{\rm e}$, $T_{\rm e}$, $d$, ${\cal M}_{\rm a}$ ) reported in Section 4 are lower limits at the same time if [M] is fixed."
 The upper limits are. £(21)).023.0.085 (Schlegel.Finkbeiner.&Davies1998).. or (CB13«0.03.0.045. Burstein&Lleiles1982... respectively.," The upper limits are $E(B-V) < 0.023, 0.085$ \citep{schl1}, or $E(B-V) < 0.03, 0.045$ \citealt{burs1}, respectively."
 We attribute our smaller £(21) to two factors: our method measures the reddening of a point source and. the excess reddening must originate [rom a region bevond GSC 0831. V372 Ser. and the comparison stars.," We attribute our smaller $E(B-V)$ to two factors: our method measures the reddening of a point source and the excess reddening must originate from a region beyond GSC 4868-0831, V372 Ser, and the comparison stars."
 The reddening £(2V)=0.085 of V372 Ser derived from the DIRBE cülfers from our value above the 5a level., The reddening $E(B-V)=0.085$ of V372 Ser derived from the DIRBE differs from our value above the $5\sigma$ level.
 ‘Taking this large (P5.V) would result in an unacceptable increment. from MISo=60.2 to S2c0.4., Taking this large $E(B-V)$ would result in an unacceptable increment from $\langle \Delta T_{\rm e}\rangle_{249}=6.2\pm 0.2$ to $8.2\pm 0.4$.
 The parameters would increase to logg.=4.11. =1092 leacing to a mass above Ον. which cannot be reconciled with any actual theoretical knowledge about pulsating stars.," The parameters would increase to $\overline{\log g_{\rm e}}=4.11$, $\overline{T_{\rm e}}=7092$ leading to a mass above $5{\cal M}_\odot$ which cannot be reconciled with any actual theoretical knowledge about pulsating stars."
" ""Therefore. EG.V)0.006 can be ruled out."," Therefore, $E(B-V) > 0.006$ can be ruled out."
 Concerning GSCASGS-O831. an interesting result can be seen from Fig. 2:," Concerning GSC4868-0831, an interesting result can be seen from Fig. \ref{fig1}:"
 the phases satisfving and not satisfving C! Scoreeate clearly in the colour-colour. diagrams and completely different 7:(/) and logg.(1) are obtained for these phases ifa BV£e photometry only is used as input., the phases satisfying and not satisfying ${\rm C}^{\rm (I)}$ segregate clearly in the colour-colour diagrams and completely different $T_{\rm e}(t)$ and $\log g_{\rm e}(t)$ are obtained for these phases if a $BVI_C$ photometry only is used as input.
 Fig., Fig.
 2 demonstrates that reliable 2;.(/) and ο(10) can be determined only if colour indices containing ( are used in the Li. logg. domain of RAR stars.," \ref{fig1} demonstrates that reliable $T_{\rm e}(t)$ and $g_{\rm e}(t)$ can be determined only if colour indices containing $U$ are used in the $T_{\rm e}$, $\log g_{\rm e}$ domain of RR stars."
 The use of one colour index or DV£e photometry only is not sullicient and can be misleading. even if it is limited to determining 7; only.," The use of one colour index or $BVI_C$ photometry only is not sufficient and can be misleading, even if it is limited to determining $T_{\rm e}$ only."
 Furthermore. our. ( observations might reveal the subeiant character of GSC 4868-0831 which was not suspected previously (Wils&Otero2005... Wils.Llove&Bernhare 20060... Szezveich&Fabryvezky 2007)).," Furthermore, our $U$ observations might reveal the subgiant character of GSC 4868-0831 which was not suspected previously \citealt{wils3}, \citealt{wils2}, \citealt{szcz1}) )."
 An important conclusion has emerged that it is not. possible to determine the luminosity class of a pulsating star if only periods or period ratio are available., An important conclusion has emerged that it is not possible to determine the luminosity class of a pulsating star if only periods or period ratio are available.
 Multicolour observations. covering the ultraviolet. are needed to classify DM pulsators properly.," Multicolour observations, covering the ultraviolet, are needed to classify DM pulsators properly."
" We have to emphasize. that μα and Zi, are first approximations from σας. (11))", We have to emphasize that $L_{\rm eq}$ and $T_{\rm eq}$ are first approximations from Eqs. \ref{4.901}) )
 and (12)) because some elements of the averaging were obtained assuming QSAA in phases when C was violated., and \ref{4.900}) ) because some elements of the averaging were obtained assuming QSAA in phases when ${\rm C}^{\rm (I)}$ was violated.
 Qualitative considerations sugeest a positive correction to Z.(/) of QSAA in. the shocked. phases when excess radiation and dissipation exis from shock waves., Qualitative considerations suggest a positive correction to $T_{\rm e}(t)$ of QSAA in the shocked phases when excess radiation and dissipation exist from shock waves.
" Corrections emerging [rom a dyvnamica model atmosphere would not modify the main fundamenta parameters V, and d because they were determined. from phases when both quantitative conditions of the validity of QSAA were satisfied.", Corrections emerging from a dynamical model atmosphere would not modify the main fundamental parameters ${\cal M}_{\rm a}$ and $d$ because they were determined from phases when both quantitative conditions of the validity of QSAA were satisfied.
" ©, and d can be considered. as wel substantiated empirical data from the ATLAS static mode atmospheres plus some basic hyvdrodynamies.", ${\cal M}_{\rm a}$ and $d$ can be considered as well substantiated empirical data from the ATLAS static model atmospheres plus some basic hydrodynamics.
. Dvnamica model atmospheres are bevond the scope of this series of papers., Dynamical model atmospheres are beyond the scope of this series of papers.
 The sampling of the euasi-repetitive curves introduce negligible error. whieh can be estimated by comparing 7; and 0 from the fitted colour curves (Denkó&Bareza2009) with those from the No=529 observations of W372 Ser.," The sampling of the quasi-repetitive curves introduced negligible error, which can be estimated by comparing $T_{\rm e}$ and $\vartheta$ from the fitted colour curves \citep{barc1} with those from the $N=529$ observations of V372 Ser."
 We remark that d=(1145+73) pe. ΝΤ. are the results for. [LE(GB1)=0003.Al] if the UAA is applied. that is. if OefOr=0 is assumed in Iq. (4)) (," We remark that $d=(1145\pm 73)$ pc, ${\cal M}_{\rm a}=(0.83\pm .17){\cal M}_\odot$ are the results for $\{E(B-V)=0.003,[M]=-0.53\}$ if the UAA is applied, that is, if $\partial v/\partial r=0$ is assumed in Eq. \ref{4.100}) ) ("
and. as a consequence.) =ed and deft= Pd).,"and, as a consequence, $\vartheta=R/d$ and $\partial v/\partial t={\ddot\vartheta}d$ )."
 This distance and the change of Ad7=6.12ΕΙο7644.52 gives LAUS=42.8L. which put V372 Ser in à position Just at the lower limit of stable DAL pulsation (Szabo.Ixolláth&Buchler 2004)., This distance and the change of ${\cal M}_{\rm a}d^{-2}=6.12\pm .31 \rightarrow 7.64\pm .82$ gives $L_{\rm eq}^{(\rm UAA)}=42.8L_\odot$ which put V372 Ser in a position just at the lower limit of stable DM pulsation \citep{szab1}. .
. Of course. the physical input of the UA is much less than that of our extended hyclrocvnamic treatment represented by Eqs. (4)). (5)).," Of course, the physical input of the UAA is much less than that of our extended hydrodynamic treatment represented by Eqs. \ref{4.100}) ), \ref{4.101}) )."
 In spite of the better agreement. the data from UA must not be accepted. because the UX is a rigid. and. less realistic approximation in comparison with a compressible model atmosphere.," In spite of the better agreement, the data from UAA must not be accepted because the UAA is a rigid and less realistic approximation in comparison with a compressible model atmosphere."
 The positions of our stars and SU Dra in a theoretical HIRD. (ie. Leg. uu) ave plotted in Fig. 4..," The positions of our stars and SU Dra in a theoretical HRD, (i.e. $[L_{\rm eq},T_{\rm eq}]$ ) are plotted in Fig. \ref{fig4}."
 For orientation. the instability strip ancl the zero age horizontal. branch of the metal deficient (Al 1.6) globular cluster M3 (NGC 5272) are shown (SilvaAguierrectal.," For orientation, the instability strip and the zero age horizontal branch of the metal deficient $\approx -1.6$ ) globular cluster M3 (NGC 5272) are shown \citep{agui1}."
2008).. We emphasize that Log.£4] correspond to L ancl 15 of non-variable stars. and they can be directly compared with those from the theoretical studies on stellar structure. pulsation and evolution.," We emphasize that $[L_{\rm eq},T_{\rm eq}]$ correspond to $L$ and $T_{\rm e}$ of non-variable stars, and they can be directly compared with those from the theoretical studies on stellar structure, pulsation and evolution."
 The only source of error is the dillerence of T.) and (0) in IE5qs. CL. 12))," The only source of error is the difference of $T_{\rm e}(t)$ and $\vartheta(t)$ in Eqs. \ref{4.901}, \ref{4.900}) )"
 from QSAA and dynamical model atmospheres. respectively.," from QSAA and dynamical model atmospheres, respectively."
 An error. has not. been propagated into the position of the star by semi-empirical relations like ως)V) ete., An error has not been propagated into the position of the star by semi-empirical relations like $\overline{T_{\rm e}}$ $\overline{(B-V)}$ etc.
 The period ratios 2)/2% are =0.7443 and 0.7450 for V372 Ser and GSC 4868-0831. respectively.," The period ratios $P_1/P_0$ are $=0.7443$ and $0.7450$ for V372 Ser and GSC 4868-0831, respectively."
 These are in the canonical range of Rite stars. as ντι of V372 Ser is also.," These are in the canonical range of RRd stars, as ${\cal M}_{\rm a}$ of V372 Ser is also."
" The equilibrium elfective temperatures of both stars are in the 5630<Z4,TOSO Ix interval where DAL pulsation can be stable.", The equilibrium effective temperatures of both stars are in the $5630 < T_{\rm eq} < 7080$ K interval where DM pulsation can be stable.
 The results of our analysis confirm that W372 Ser is à true RRe star., The results of our analysis confirm that V372 Ser is a true RRd star.
 However. it is subluminous by a factor ofz2-3 in comparison with the pulsation models of νά stars (Szabo.Ixolláth&," However, it is subluminous by a factor of$\approx 2\mbox{-}3$ in comparison with the pulsation models of RRd stars \citep{szab1}. ."
Buehler 2004).. MB of GSC 4868-0831 exceeds somewhat the canonical mass of Id. stars: its subluminosity is even larger. z4-5 ," ${\cal M}_{\rm a}$ of GSC 4868-0831 exceeds somewhat the canonical mass of RRd stars; its subluminosity is even larger, $\approx 4\mbox{-}5$ "
set to 0.42ro and the outer radius is set to 4.2ro.,set to $0.42r_0$ and the outer radius is set to $4.2r_0$.
 This resolution allows us to identify details as small as 0.04ro., This resolution allows us to identify details as small as $0.04r_0$.
" Casoli Masset (2009) showed that the coorbital co-rotation torque exerted by the disk over a planet occurs in a scale of zs~0.01ro for m~2ma, while Masset et al. ("," Casoli Masset (2009) showed that the coorbital co-rotation torque exerted by the disk over a planet occurs in a scale of $x_s \sim 0.01 r_0$ for $m \sim 2 m_{\oplus}$, while Masset et al. ("
2006b) showed that the horse-shoe zone width in a two-dimentional disk scale as ας=1.164/q/h.,2006b) showed that the horse-shoe zone width in a two-dimentional disk scale as $x_s =1.16\sqrt{q/h}$.
" Consequently, zs should increase with the planetary mass."," Consequently, $x_s$ should increase with the planetary mass."
" For example if we are interested in the dynamical evolution of Saturn-type planets, we expect that our choice for the spatial resolution should be adequate."," For example if we are interested in the dynamical evolution of Saturn-type planets, we expect that our choice for the spatial resolution should be adequate."
" To test this idea, we performed several simulations with the same initial conditions but different spatial resolutions for the disk, and compared the dynamical evolution of the planets."," To test this idea, we performed several simulations with the same initial conditions but different spatial resolutions for the disk, and compared the dynamical evolution of the planets."
 An example is shown in Figure 13 for two Saturn-like planets initially in circular orbits far from the cavity edge and in non-resonant conditions., An example is shown in Figure \ref{fig2_p} for two Saturn-like planets initially in circular orbits far from the cavity edge and in non-resonant conditions.
" Both simulations show the same qualitative results, not only in the final outcome of the planets (both runs end in coorbital configurations) but also in the behavior of the eccentricities."," Both simulations show the same qualitative results, not only in the final outcome of the planets (both runs end in coorbital configurations) but also in the behavior of the eccentricities."
 The only noticeable change is a slight difference in the radial distance at which the planets are trapped., The only noticeable change is a slight difference in the radial distance at which the planets are trapped.
" However, this is not significant and could be explained due to a decrease of the co-rotational torque from a low resolution description of the horse-shoe region."," However, this is not significant and could be explained due to a decrease of the co-rotational torque from a low resolution description of the horse-shoe region."
" Following the predictions of our N-body runs, all our hydro-simulations were done with two equal-mass planets (m1= m2) in a high-density disk."," Following the predictions of our N-body runs, all our hydro-simulations were done with two equal-mass planets $m_1=m_2$ ) in a high-density disk."
 Initial conditions were varied both with respect to the initial separation between the planets and the initial eccentricities., Initial conditions were varied both with respect to the initial separation between the planets and the initial eccentricities.
 We also considered different planetary masses between one Earth-mass and one Saturn-mass., We also considered different planetary masses between one Earth-mass and one Saturn-mass.
" Under a wide variety of initial conditions, we found that the orbital evolution of two Saturn-mass planets ultimately lead to the formation of stable coorbital configurations."," Under a wide variety of initial conditions, we found that the orbital evolution of two Saturn-mass planets ultimately lead to the formation of stable coorbital configurations."
" An example is shown in Figure 14,, where both bodies are trapped in L4/ L5-type motion with libration of both the resonant angle 0=A2—A and the differencea in longitudes of pericenter Aw=cv»—σι."," An example is shown in Figure \ref{fig4_p}, where both bodies are trapped in $L_4/L_5$ -type motion with a libration of both the resonant angle $\sigma = \lambda_2 - \lambda_1$ and the difference in longitudes of pericenter $\Delta \varpi = \varpi_2 - \varpi_1$."
" To test the robustness and stability of the coorbital solutions to additional perturbations, we performed a new series of runs including additional six Earth-mass planets initially far from the cavity edge."," To test the robustness and stability of the coorbital solutions to additional perturbations, we performed a new series of runs including additional six Earth-mass planets initially far from the cavity edge."
" The idea behind this set of experiments is to analyze both the orbital evolution of these new bodies, as well as their effects on the two coorbitals."," The idea behind this set of experiments is to analyze both the orbital evolution of these new bodies, as well as their effects on the two coorbitals."
" Typical results are shown in Figure 15,, where each color curve corresponds to a different planet (see inlaid color code)."," Typical results are shown in Figure \ref{fig3_p}, where each color curve corresponds to a different planet (see inlaid color code)."
" Planets 1 and 2 (i.e. m1 and mo», respectively) are the original Saturn-mass bodies in coorbital motion, while planets 3 through 8 are the new Earth-size masses."," Planets 1 and 2 (i.e. $m_1$ and $m_2$, respectively) are the original Saturn-mass bodies in coorbital motion, while planets 3 through 8 are the new Earth-size masses."
" Since this type of simulation included smaller masses and a larger population, we employed a high-resolution description of the disk."," Since this type of simulation included smaller masses and a larger population, we employed a high-resolution description of the disk."
" We perceived no significant instability in the coorbitals, and their configuration remained unaffected by the new masses."," We perceived no significant instability in the coorbitals, and their configuration remained unaffected by the new masses."
" These, however, suffered several cases of mutual scattering and close encounters that caused the ejection of three of the masses."," These, however, suffered several cases of mutual scattering and close encounters that caused the ejection of three of the masses."
 One of the close encounters (between planets 4 and 6) led to a temporary trapping of planet 4 in the cavity edge but finally led to the formation of a new coorbital system with planet 8., One of the close encounters (between planets 4 and 6) led to a temporary trapping of planet 4 in the cavity edge but finally led to the formation of a new coorbital system with planet 8.
 The role of scattering between planets as a formation mechanism for, The role of scattering between planets as a formation mechanism for
"Although there is evidence that magnetic winds brake the spin of stars, the underlying physics that produces the magnetic fields is not yet fully understood.","Although there is evidence that magnetic winds brake the spin of stars, the underlying physics that produces the magnetic fields is not yet fully understood."
 There is a distinction between the magnetic properties of stars on the upper and lower main sequence., There is a distinction between the magnetic properties of stars on the upper and lower main sequence.
" In the case of the Sun and solar-like stars the combination of a convective envelope and rotation is thought to generate the observable magnetic fields by an a—Ω dynamo, the strength of which depends upon the stellar angular velocity."," In the case of the Sun and solar-like stars the combination of a convective envelope and rotation is thought to generate the observable magnetic fields by an $\alpha-\Omega$ dynamo, the strength of which depends upon the stellar angular velocity."
" However, in contrast, only a minority of early-type stars, such as the Ap and Bp stars have observable magnetic fields."," However, in contrast, only a minority of early-type stars, such as the Ap and Bp stars have observable magnetic fields."
 Many attempts have been made to explain these phenomena, Many attempts have been made to explain these phenomena
Galaxies located in the core regions of clusters can experience strong rani pressure stripping of their ISM by the ICM. (0 the extent that they lose most of their ISAT gas (ie. fy 0).,"Galaxies located in the core regions of clusters can experience strong ram pressure stripping of their ISM by the ICM, to the extent that they lose most of their ISM gas (i.e. $f_{g}\sim 0$ )."
 They are therefore unlikely to experience starbursts cuving cluster merging. because {μον are already. gas-poor prior (o cluster merging.," They are therefore unlikely to experience starbursts during cluster merging, because they are already gas-poor prior to cluster merging."
 Thus the present stud is likely to overestimate (he fractions of starburst galaxies to some extent., Thus the present study is likely to overestimate the fractions of starburst galaxies to some extent.
" We need to model properly the variation inf, in galaxies of different IIubble tvpe. and how that translates into a varying f; with radius [rom the cluster centre. in order to predict much more precisely the possible fractions of starburst and poststarburst galaxies in merging clusters in our future stucies."," We need to model properly the variation in $f_{\rm g}$ in galaxies of different Hubble type, and how that translates into a varying $f_{\rm g}$ with radius from the cluster centre, in order to predict much more precisely the possible fractions of starburst and poststarburst galaxies in merging clusters in our future studies."
 It should also be noted that it is not only the external pressure of the ICM in clusters that can trigger starbursts in gas-rich galaxies (e.g.. INronberger et al.," It should also be noted that it is not only the external pressure of the ICM in clusters that can trigger starbursts in gas-rich galaxies (e.g., Kronberger et al."
 2008). but also cluster tidal fields (Bekki 1999).," 2008), but also time-dependent cluster tidal fields (Bekki 1999)."
 It therefore appears inevitable (hat merging clusters are likely to have a larger Traction of starburst galaxies than non-merging clusters., It therefore appears inevitable that merging clusters are likely to have a larger fraction of starburst galaxies than non-merging clusters.
 It would be diffieult. however. for observational studies to determine whether tidal effects or increased external ICM. pressure is the main driver for such an increased starburst Iraction.," It would be difficult, however, for observational studies to determine whether tidal effects or increased external ICM pressure is the main driver for such an increased starburst fraction."
 Since the time-varving tidal fields in merging clusters can also transformstellar disks (Bekki 1999). whereas the external pressure of the ICM is unlikely to do so. the morphological properties of each individual starburst galaxy will provide important clues as to which of the above (wo effects are responsible for (heir origin.," Since the time-varying tidal fields in merging clusters can also transform disks (Bekki 1999), whereas the external pressure of the ICM is unlikely to do so, the morphological properties of each individual starburst galaxy will provide important clues as to which of the above two effects are responsible for their origin."
 Recent and ongoing photometric and spectroscopic observations of galaxy. properties in a large sample of clusters with and without cold Ironts will soon reveal (he number fractions ol starburst and poststarburst galaxies ancl their spatial distributions ancl kinematics in these possibly merging and non-mergineg clusters (e.g.. IDwang Lee 2009: Ma et al.," Recent and ongoing photometric and spectroscopic observations of galaxy properties in a large sample of clusters with and without cold fronts will soon reveal the number fractions of starburst and poststarburst galaxies and their spatial distributions and kinematics in these possibly merging and non-merging clusters (e.g., Hwang Lee 2009; Ma et al."
 2010: Owers et al., 2010; Owers et al.
 2010)., 2010).
 These statistical studies will enable us to compare the simulated. kinematics (e.g.. line-o[-sight. velocity distributions) of starburst galaxies with (hose observed. thereby allowing more robust conclusions to be drawn as to whether the star formation histories of cluster member galaxies are dramatically changed by cluster merging.," These statistical studies will enable us to compare the simulated kinematics (e.g., line-of-sight velocity distributions) of starburst galaxies with those observed, thereby allowing more robust conclusions to be drawn as to whether the star formation histories of cluster member galaxies are dramatically changed by cluster merging."
 The observed peculiar spatial distributions (e.g.. rine-like structures) of poststarburst galaxies in clusters with substructures (e.g.. the Coma cluster: Pogelanti et al.," The observed peculiar spatial distributions (e.g., ring-like structures) of poststarburst galaxies in clusters with substructures (e.g., the Coma cluster; Poggianti et al."
 1999) will place strong. constraints on (he mass-ratios and radial velocities of merging clusters., 1999) will place strong constraints on the mass-ratios and radial velocities of merging clusters.
 KB. MSO and WJC all acknowledge the financial support of the Australian Research Council throughout the course of this work.," KB, MSO and WJC all acknowledge the financial support of the Australian Research Council throughout the course of this work."
 Numerical computations reported here were carried out both on the GRAPE svstem at the University of Western Australia and on those kindly made available bv the Center lor computational astrophvsies (CICA) of the National Astronomical Observatory of Japan., Numerical computations reported here were carried out both on the GRAPE system at the University of Western Australia and on those kindly made available by the Center for computational astrophysics (CfCA) of the National Astronomical Observatory of Japan.
"The correlations between the peak value of the total Galactic coalescence rate Rpeak and the model parameters (e.g. (the cut-off luminosity Ly, and the power index p) seem to be similar to those of DNSs we observed in paper I. As we found for the DNS svstems. values strongly depend on a pulsar luminosity function rather than a spatial distribution of pulsars in the Galaxy. in other words. Values rapidly increase as the fraction of [aint pulsars increases.","The correlations between the peak value of the total Galactic coalescence rate ${\cal R}_{\rm peak}$ and the model parameters (e.g. the cut-off luminosity $L_{\rm min}$ and the power index $p$ ) seem to be similar to those of DNSs we observed in paper I. As we found for the DNS systems, values strongly depend on a pulsar luminosity function rather than a spatial distribution of pulsars in the Galaxy, in other words, values rapidly increase as the fraction of faint pulsars increases."
 Close binaries consisting of compact objects (e.g. NS-WD binaries) are suggested as important GW sources in a [frequency range below 1 mllz., Close binaries consisting of compact objects (e.g. $-$ WD binaries) are suggested as important GW sources in a frequency range below 1 mHz.
 In this range. due to the large number of sources. LISA would not be able to resolve each source within a given Lrequency band.," In this range, due to the large number of sources, LISA would not be able to resolve each source within a given frequency band."
" Hence the Galactic binaries are expected (o establish a confusion noise level (or ""background) dominated by WD-WD binaries2001).", Hence the Galactic binaries are expected to establish a confusion noise level (or “background”) dominated by $-$ WD binaries.
 In this work. we consider the contribution from binaries (o the predicted confusion noise level.," In this work, we consider the contribution from binaries to the predicted confusion noise level."
 Using our results from the previous section. we calculate the amplitude of GW signals [roin binaries in the nearby Universe and compare it with the LISA sensitivitycurve.," Using our results from the previous section, we calculate the amplitude of GW signals from binaries in the nearby Universe and compare it with the LISA sensitivity."
 In this work. we assume that the tree observed svstems represent the whole population of binaries in our Galaxy.," In this work, we assume that the three observed systems represent the whole population of binaries in our Galaxy."
 We calculate the characteristic strain amplitude of GWs (5.) from binaries using the results given by Phinney (2001) lor cireular binaries., We calculate the characteristic strain amplitude of GWs $h_{\rm c}$ ) from binaries using the results given by Phinney (2001) for circular binaries.
 In general. binaries with an eccenlricily e emit GWs at [requencies f/—np. ie. Che n-th harmonic of the orbital Lrequency v.," In general, binaries with an eccentricity $e$ emit GWs at frequencies $f=n\nu$, i.e. the $n$ -th harmonic of the orbital frequency $\nu$."
 In the case of circular binaries. 7=2 due to the orbital symmetry and the quadrupole nature of GWs.," In the case of circular binaries, $n=2$ due to the orbital symmetry and the quadrupole nature of GWs."
 The eccentricity of PSRs J0751-2-1807 and J1757—5322 are e~LO and ~10© respectively.," The eccentricity of PSRs J0751+1807 and $-$ 5322 are $e\sim10^{-4}$ and $\sim 10^{-6}$, respectively."
 Hence. it is sale to consider them as circular binaries.," Hence, it is safe to consider them as circular binaries."
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular ,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular b,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular bi,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular bin,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular bina,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular binar,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular binary,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular binary!,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular binary!.,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
PSR J1141—6545 has an appreciable eccentricitv (e= 0.17). but for simplicity. we consiler only (he η=2 harmonic as if it were a circular binary!..,"PSR $-$ 6545 has an appreciable eccentricity $e=0.17$ ), but for simplicity, we consider only the $n=2$ harmonic as if it were a circular"
region contributes most of the inverse Compton flux. while the svuchrotvon emission comes mainlv from (he emission region as a whole.,"region contributes most of the inverse Compton flux, while the synchrotron emission comes mainly from the emission region as a whole."
 Within the injection region. the amplitude between consecutive Compton orders is ↕⋅∪∏↖≺↽↔↴⊔∡∖↽↶↵↓−↕↓↕↓↓⊤⊳↓⋅↴∿↴↶↴↴↴⋯⋯↿∎≼↜≺∢↓⋅≼↲≺⇂∿∖⋗∏⋝↻∏↥⋟∖⊽∐⇂≼↲⊔∐↲↕∐∙↿≼↲≺∢⊔∪∐↕⋅≼↲," Within the injection region, the amplitude between consecutive Compton orders is roughly $\gamma_{\rm min}^2 \tau_{\rm T} \sim \gamma_{\rm min}^{2(1-\alpha)}$ (cf."
↖≺↽↔↴↕∪∐⊔∐⋟∖⊽≀↧↴∐↓↕↽≻∐, eq. \ref{eq:1.9}] ]).
"⋯≺⇂≼↲↕⋟∖⊽⋟∖⇁∐↓≀↧↴∐≼↲↕⋅ :ÓUna) qur −⋅− − − − − bv a factor. R/r, due to the smaller column density of electrons.", Outside the injection region this amplitude is smaller by a factor $R/r_{\rm o}$ due to the smaller column density of electrons.
" The small value of 5, expected in LSPs causes the inverse Compton scattering also outside the injection region to be dominated by the low energv end of the electron distribution.", The small value of $\gamma_{\rm c}$ expected in LSPs causes the inverse Compton scattering also outside the injection region to be dominated by the low energy end of the electron distribution.
 The minimum in the spectral energv distribution. observed in LSPs is (wpically a [actor LO below the peak of the svnehrotvon component., The minimum in the spectral energy distribution observed in LSPs is typically a factor $10$ below the peak of the synchrotron component.
 In order to make MIC-Amnodels compatible with such minima. σαι9“r/RXo10U ds required. where n is the Compton order corresponding to the frequency of (he minimum.," In order to make MIC-models compatible with such minima, $\gamma_{\rm min}^{2(1-\alpha)} r_{\rm o} / R \lesssim 10^{-1/n}$ is required, where $n$ is the Compton order corresponding to the frequency of the minimum."
" Since μμ10 is needed for a smooth X-ray. spectrum. the strongest constraint on r,/ 1s obtained by using +yu,2210. which implies »ez2 (taking the mininmnm to occur at a frequency LOM Lz and a svichrotron peak frequency at LOM LIT)."," Since $\gamma_{\rm min} \lesssim 10$ is needed for a smooth X-ray spectrum, the strongest constraint on $r_{\rm o} / R$ is obtained by using $\gamma_{\rm min} \approx 10$, which implies $n\approx 2$ (taking the minimum to occur at a frequency $10^{17}$ Hz and a synchrotron peak frequency at $10^{13}$ Hz)."
" With a=0.62009b).. this shows that /r,,220 eives rise lo an X-ray minimuni deep enough to be consistent wilh observations."," With $\alpha \approx 0.6$, this shows that $R/r_{\rm o} \gtrsim 20$ gives rise to an X-ray minimum deep enough to be consistent with observations."
 The requirement for a much lareer value of e./ey inside the injection region than outside constrains the mechanism whereby the escaping electrons are dispersed (throughout. the svnchrotron emission volume., The requirement for a much larger value of $\epsilon_{\rm e}/\epsilon_{\rm B}$ inside the injection region than outside constrains the mechanism whereby the escaping electrons are dispersed throughout the synchrotron emission volume.
 Adiabatic expansion is unlikelv since the value of ej as well as (he radiative efficiency. would decrease., Adiabatic expansion is unlikely since the value of $\epsilon_{\rm B}$ as well as the radiative efficiency would decrease.
 A constant value of B and no adiabatie losses of (he electrons. as in (he above example. are consistent wilh scattering possibly in connection with streaming.," A constant value of $B$ and no adiabatic losses of the electrons, as in the above example, are consistent with scattering possibly in connection with streaming."
 In addition. (he acceleration of electrons could occur in regions with low values of D. suggesting that the injected. electrons may derive (heir οποιον [rom magnetic reconnection2010).," In addition, the acceleration of electrons could occur in regions with low values of $B$, suggesting that the injected electrons may derive their energy from magnetic reconnection."
. Bot of these processes may (ake place simultaneously and. hence. the transition from the svnchirotron to the inverse Compton component in the spectral energy distribution is sensitive to the combined effects of the acceleration and the subsequent spreading of the electrons through the whole emission volunie.," Both of these processes may take place simultaneously and, hence, the transition from the synchrotron to the inverse Compton component in the spectral energy distribution is sensitive to the combined effects of the acceleration and the subsequent spreading of the electrons through the whole emission volume."
 In a MlIC-scenario. spectral properties and [ας variations are «uite different [rom those in a source where the high frequency emission is mainiv due to a single scattering.," In a MIC-scenario, spectral properties and flux variations are quite different from those in a source where the high frequency emission is mainly due to a single scattering."
" As discussed above. the actual source properties need not differ too much between the (wo cases due to the sensitivitv ol 5, to the density of velativistic electrons. in particular. but also to the magnetic field strength/densitv of external photons."," As discussed above, the actual source properties need not differ too much between the two cases due to the sensitivity of $\gamma_{\rm c}$ to the density of relativistic electrons, in particular, but also to the magnetic field strength/density of external photons."
 Furthermore. (his sensitivity becomes more pronounced for lower values of 5. (cf.," Furthermore, this sensitivity becomes more pronounced for lower values of $\gamma_{\rm c}$ (cf."
 eq. [15]])., eq. \ref{eq:1.16}] ]).
 Even for a continuous distribution of intrinsic source parameters one max (hen expect a bi-mocel distribution. of observed properties. although transition cases should be observed.," Even for a continuous distribution of intrinsic source parameters one may then expect a bi-model distribution of observed properties, although transition cases should be observed."
 The first group could be adequately described by a standard svinclirotron/Conmpton model. in which multiple inverse Compton," The first group could be adequately described by a standard synchrotron/Compton model, in which multiple inverse Compton"
"Using the positions and determined distances to the clusters in our sample, we can investigate their distribution in the Galactic Plane.","Using the positions and determined distances to the clusters in our sample, we can investigate their distribution in the Galactic Plane."
 In the left panel of refdistrxy we show the distribution of all clusters in the X- plane., In the left panel of \\ref{distrxy} we show the distribution of all clusters in the X-Y plane.
" As blue triangles we plot the known globular clusters, while the open clusters are plotted as red symbols (large circles for clusters with ages above GGyr, small dots for clusters with ages below GGyr)."," As blue triangles we plot the known globular clusters, while the open clusters are plotted as red symbols (large circles for clusters with ages above Gyr, small dots for clusters with ages below Gyr)."
 The plot assumes a distance of the Sun to the Galactic Centre of kkpc., The plot assumes a distance of the Sun to the Galactic Centre of kpc.
 As expected one finds most of the globular clusters concentrated towards the Galactic Centre., As expected one finds most of the globular clusters concentrated towards the Galactic Centre.
 The open clusters are mostly found near the Suns position., The open clusters are mostly found near the Suns position.
" This is, however, a simple selection effect, as the sample contains naturally only clusters near enough to be visible in 2MASS data."," This is, however, a simple selection effect, as the sample contains naturally only clusters near enough to be visible in 2MASS data."
" Two details of the spatial distribution of the open clusters are, however, worth discussing in a bit more detail: 1) A histogram of the distribution of Galactic Longitude values of the open clusters in our sample (see refhistl)) reveals two regions with a smaller than average number of clusters."," Two details of the spatial distribution of the open clusters are, however, worth discussing in a bit more detail: i) A histogram of the distribution of Galactic Longitude values of the open clusters in our sample (see \\ref{histl}) ) reveals two regions with a smaller than average number of clusters."
 This is a) the region near the Galactic Centre (+ 660° away) and b) the Galactic Longitude range 120°<I« 1180*.," This is a) the region near the Galactic Centre $\pm$ $^\circ$ away) and b) the Galactic Longitude range $^\circ <
l <$ $^\circ$."
 In the case of a) there are two reasons for this., In the case of a) there are two reasons for this.
" Firstly, the high star density towards the Galactic Centre prevents the detection and identification of stellar clusters in this area (a selection effect when establishing the sample) and secondly the fact that there are indeed fewer old stellar clusters closer than the Sun to the Galactic Centre (see below)."," Firstly, the high star density towards the Galactic Centre prevents the detection and identification of stellar clusters in this area (a selection effect when establishing the sample) and secondly the fact that there are indeed fewer old stellar clusters closer than the Sun to the Galactic Centre (see below)."
" In the case of b), there seems to be no obvious reason for the paucity of old open clusters in this region compared to the same longitude range on the opposite side of the Galactic Anticentre (i.e. the region 180°«1 <2240°)."," In the case of b), there seems to be no obvious reason for the paucity of old open clusters in this region compared to the same longitude range on the opposite side of the Galactic Anticentre (i.e. the region $^\circ < l
<$ $^\circ$ )."
" One explanation could be one or several large, high extinction molecular clouds in this direction, preventing the detection of clusters."," One explanation could be one or several large, high extinction molecular clouds in this direction, preventing the detection of clusters."
" However, in the all sky extinction maps from Rowles Froebrich (2009) there is no indication of such clouds."," However, in the all sky extinction maps from Rowles Froebrich \cite{2009MNRAS.395.1640R} there is no indication of such clouds."
" Furthermore, the clusters which are detected in this longitude range, cover all distances between 1 and 7kkpc homogeneously, as well as extinction values between 0.0 and mmag Ax."," Furthermore, the clusters which are detected in this longitude range, cover all distances between 1 and kpc homogeneously, as well as extinction values between 0.0 and mag $A_K$."
" Finally, as can be seen in refhistl,, the effect is more pronounced for clusters with ages above GGyr."," Finally, as can be seen in \\ref{histl}, the effect is more pronounced for clusters with ages above Gyr."
" Hence, this region either suffers from an unknown selection effect, or indeed there are fewer than normal old open clusters present in this part of the Galaxy."," Hence, this region either suffers from an unknown selection effect, or indeed there are fewer than normal old open clusters present in this part of the Galaxy."
" ii) While our sample clearly contains a large number of clusters within kkpc from the Sun, there is a clear paucity of objects with distances to the Galactic Centre less than that of the Sun."," ii) While our sample clearly contains a large number of clusters within kpc from the Sun, there is a clear paucity of objects with distances to the Galactic Centre less than that of the Sun."
 This is a fact already noticed by Friel (1995) and replicated since then., This is a fact already noticed by Friel \cite{1995ARA&A..33..381F} and replicated since then.
 In our sample only of the open clusters with ages above GGyr are closer than 5kkpc to the Galactic Centre (only open clusters are further away than the Sun from the Galactic, In our sample only of the open clusters with ages above Gyr are closer than kpc to the Galactic Centre (only open clusters are further away than the Sun from the Galactic
the IIubble Space Telescope GIST) (Sharaetal.1996:Ixniggeοἱ2002).,"the Hubble Space Telescope (HST) \citep{sha96,kni02}."
. Recently. Chandra X-ray Observatory and HIST. observers claimed the detections of dozens of cataclvsmie variable (CV) candidates in the cores of several globular clusters.," Recently, Chandra X-ray Observatory and HST observers claimed the detections of dozens of cataclysmic variable (CV) candidates in the cores of several globular clusters."
 These candidates are X-ray bright in Chandra observations (Grindlayetal.2001) or UV-bright in IST observations (Ixniggeοἱal.2002) and clustered in the cores of their host clusters., These candidates are X-ray bright in Chandra observations \citep{gri01} or UV-bright in HST observations \citep{kni02} and clustered in the cores of their host clusters.
 ILowever. only zero or one of these UV or X-ray sources per cluster has vel exhibited the degree of variability (2 magnitude or larger outbursts) that originally led to labeling these variables as “cataclysmic.”," However, only zero or one of these UV or X-ray sources per cluster has yet exhibited the degree of variability (2 magnitude or larger outbursts) that originally led to labeling these variables as “cataclysmic.”"
 About hall of all the CVs known in the solar neighborhood exhibit large amplitiucle oulbursts — the Dwarf Novae., About half of all the CVs known in the solar neighborhood exhibit large amplitude outbursts — the Dwarf Novae.
 In the absence of spectroscopic confirmation (very challenging to obtain) and/or large amplitude outbursts. the N-rav aud UV-brieht. CV candidates in globular clusters may remain just that: candidates. for the loreseeable future.," In the absence of spectroscopic confirmation (very challenging to obtain) and/or large amplitude outbursts, the X-ray and UV-bright CV candidates in globular clusters may remain just that: candidates, for the foreseeable future."
 It is (hus important to fully substantiate everv claim of lavee amplitude variability in globular cluster stars., It is thus important to fully substantiate every claim of large amplitude variability in globular cluster stars.
 Only a lew invehtilable. spectrographically confirmed CVs (e.g. Grindlayetal.(1995):Edimonelsetal. (1999))). and cataclysmies confirmed by slitless lar-UlV spectroscopy (Ixniggeetal.2003:Ixnigge 2004).. (including 2 good candidates in Edmondsetal. (2003))) are known in globular elusters.," Only a few irrefutable, spectrographically confirmed CVs (e.g. \citet{gri95,edm99}) ), and cataclysmics confirmed by slitless far-UV spectroscopy \citep{kni03,kni04}, , (including 2 good candidates in \citet{edm03}) ) are known in globular clusters."
 A recent addition to this exclusive club of erupting CVs is the variable labelled V1 in the dense globular cluster MIS (Charles et al., A recent addition to this exclusive club of erupting CVs is the variable labelled V1 in the dense globular cluster M15 (Charles et al.
 2002. hereafter C02).," 2002, hereafter C02)."
 A15 is one of the densest elobular clusters in the Galaxy., M15 is one of the densest globular clusters in the Galaxy.
 It has probably undergone core Collapse ancl almost certainly. contains a significant amount of non-Iuminous matter in its core (Daumgzucdtiοἱal.2003)., It has probably undergone core collapse and almost certainly contains a significant amount of non-luminous matter in its core \citep{bau03}.
. The cluster hosts the remarkable low mass N-rav. binary (LMXD) 4U21274-119. whose optical counterpart ((Auriereetal.1984). is the eclipsing binary ACTI (Ilovaiskvetal.1987).," The cluster hosts the remarkable low mass X-ray binary (LMXB) 4U2127+119, whose optical counterpart \citep{aur84} is the eclipsing binary AC211 \citep{ilo87}."
. MID has been extensively imaged by LST. with which observers have discovered asharp central density cusp (Guhbathakurtaetal. 1996)..," M15 has been extensively imaged by HST, with which observers have discovered asharp central density cusp \citep{guh96}, ,"
with those in previous works (c.g. Veilleux ct al.,"with those in previous works (e.g., Veilleux et al."
 1999: Ixewley et al., 1999; Kewley et al.
 2001: Goto 2005: Cao et al., 2001; Goto 2005; Cao et al.
 2006: Hou et al., 2006; Hou et al.
 2009: Yuan et al., 2009; Yuan et al.
 2010)., 2010).
 In Fig., In Fig.
 5.r we present infrared luminosity.28.15 60 jan colour. and. LPAS 1tYO yam colour distributions of the combined. sample (432 |ινα) including this stucly. lim Sanders (1998). and llou et al. (," 5, we present infrared luminosity, $-$ 60 $\mu$ m colour, and $-$ 100 $\mu$ m colour distributions of the combined sample (432 ULIRGs) including this study, Kim Sanders (1998), and Hou et al. ("
2009).,2009).
 Dased on the Ixolmogorov-Smirnov (IX-S) test. we find that. the infrared Luminosity clistributicin οἱ broad-line ACGN ULIBCSs is significantly cillerent from trose Of non-AGN and narrow-line AGN ULIRGs. while those of non-AGN and narrow-line AGN WLIRGs are not so dillerent.," Based on the Kolmogorov-Smirnov (K-S) test, we find that the infrared luminosity distribution of broad-line AGN ULIRGs is significantly different from those of non-AGN and narrow-line AGN ULIRGs, while those of non-AGN and narrow-line AGN ULIRGs are not so different."
 We also find that the three types of ULIRGSs are sienificantly different from each other inLRAS 60 pm colour distribution but they are indistinguishable in 442.15 100 pm colour distribution., We also find that the three types of ULIRGs are significantly different from each other in $-$ 60 $\mu$ m colour distribution but they are indistinguishable in $-$ 100 $\mu$ m colour distribution.
 InZRAS 60 jun colour distribution. the significant cilference between AGN and non-AGN ULLBCGs strengthens that micd-infrared colour is a good indicator of AGN activity in infrared. galaxies (e.g... «e Grijp et al.," In $-$ 60 $\mu$ m colour distribution, the significant difference between AGN and non-AGN ULIRGs strengthens that mid-infrared colour is a good indicator of AGN activity in infrared galaxies (e.g., de Grijp et al."
 1985: Nell Llutchings 1992)., 1985; Neff Hutchings 1992).
 Phe little differences in28.15 100 sam colour distribution reinforce hat star formation clominates the emission of AGN in the fizr-infrared regime (e.g. Elbaz et al.," The little differences in $-$ 100 $\mu$ m colour distribution reinforce that star formation dominates the emission of AGN in the far-infrared regime (e.g., Elbaz et al."
 2010: Llatziminaoglou e al., 2010; Hatziminaoglou et al.
 2010: LIwane et al., 2010; Hwang et al.
 2010b: Shao et al., 2010b; Shao et al.
 2010)., 2010).
 On the «Xher hand. narrow-line ACN ULIRGSs have in general lower infrared luminosity ancl cooler184525. 60 pin colour than sroad-line AGN ULILIIBCGs do.," On the other hand, narrow-line AGN ULIRGs have in general lower infrared luminosity and cooler $-$ 60 $\mu$ m colour than broad-line AGN ULIRGs do."
 If we assume that the infrared emission in galaxies is emitted isotropically. and. does not depend on viewing angle (e.g.. Alulehaev et al.," If we assume that the infrared emission in galaxies is emitted isotropically, and does not depend on viewing angle (e.g., Mulchaey et al."
 1994: Schartmann et al., 1994; Schartmann et al.
 2008: Gandhi et al., 2008; Gandhi et al.
 2009). the dillerence between narrow-line AGN ancl broad-line AGN ULIRGs seems not to be compatible with the predictions of the orientation-dependent unification moclel of AGNs in which narrow-line ancl broacd-line ACGNs are intrinsically same objects observed from. cillerent angles (Antonucci 1993).," 2009), the difference between narrow-line AGN and broad-line AGN ULIRGs seems not to be compatible with the predictions of the orientation-dependent unification model of AGNs in which narrow-line and broad-line AGNs are intrinsically same objects observed from different angles (Antonucci 1993)."
 This conflict can be explained if narrow-line AGN ULIIHCGs host a central engine deeply buried by extended. dusty regions of star formation as proposed. by several authors (e.g. Genzel et al 1998: Gerssen et al.," This conflict can be explained if narrow-line AGN ULIRGs host a central engine deeply buried by extended, dusty regions of star formation as proposed by several authors (e.g., Genzel et al 1998; Gerssen et al."
 2004)., 2004).
 In Fig., In Fig.
 6. we show the relations between optical {line ratio or line width) and infrared. (infrared luminosity.LPS 60 [an colour or.28.15 100. pam colour) properties of the combined ULEIRG sample.," 6, we show the relations between optical line ratio or line width) and infrared (infrared luminosity, $-$ 60 $\mu$ m colour or $-$ 100 $\mu$ m colour) properties of the combined ULIRG sample."
 Broad-line ACN ULIRGs are not plotted: because their line widths and fluxes are not properly measured in manycases?., Broad-line AGN ULIRGs are not plotted because their line widths and fluxes are not properly measured in many.
. Since the spectral types are determined. by the optical emission lines. narrow-line ACGNs anc non-AGNs are well separated. along the ν-axis.," Since the spectral types are determined by the optical emission lines, narrow-line AGNs and non-AGNs are well separated along the y-axis."
 As expected. clear correlations are found. only in the middle panels (see the values of Spearman's rank correlation ," As expected, clear correlations are found only in the middle panels (see the values of Spearman's rank correlation coefficients)."
We checked these relations using various infrared coellicients).colours from the all-sky survey point source but we could not draw any meaningful⋠⋅ results due to the small number of data points., We checked these relations using various infrared colours from the all-sky survey point source but we could not draw any meaningful results due to the small number of data points.
 We studied: optical spectral properties of 115. southern ULIRGs using the spectra. obtained from our οΓιο observations. 2dCRS. and GAGS.," We studied optical spectral properties of 115 southern ULIRGs using the spectra obtained from our CTIO observations, 2dFGRS, and 6dFGS."
 For ULIBCs with numeasurement. we classified them in. the. standard diagnostic diagrams.," For ULIRGs with measurement, we classified them in the standard diagnostic diagrams."
 We classified the other ULIRGs using the niA5007 /LE7 line ratio against 111]A5007 line width diagram with an empirically determined. criterion., We classified the other ULIRGs using the $\lambda5007$ $\beta$ line ratio against $\lambda5007$ line width diagram with an empirically determined criterion.
 Main results are summarized below., Main results are summarized below.
 We thank the ςΓΕ stalf for their help on the observations., We thank the CTIO staff for their help on the observations.
 We also thank L. €. Hou for providing the line width data of the SDSS ULIRGs., We also thank L. G. Hou for providing the line width data of the SDSS ULIRGs.
 We are egrateful to anonvmous referee whose commoents helped to improve the original manuscript., We are grateful to anonymous referee whose comments helped to improve the original manuscript.
 This work was supported by Micd-career Research Program through NRE erant funded by the ALES'T (No.2010-0013875)., This work was supported by Mid-career Research Program through NRF grant funded by the MEST (No.2010-0013875).
 LESAL acknowledges the support ofthe Centre National d'I2tudes Spatiales (CNIZS).S.Cον isa member of the Dedicated Rescarchers for Extragalactic AstronoMy (DREAAL) team in korea Astronomy and Space Science Institute (ΝΑΡ)., H.S.H acknowledges the support ofthe Centre National d'Etudes Spatiales (CNES).S.C.K is a member of the Dedicated Researchers for Extragalactic AstronoMy (DREAM) team in Korea Astronomy and Space Science Institute (KASI).
separation P=Yn.raa and velocity Po—fugaoPea. the orbital energy is and the angular momentuni is With these. we computed the orbital eccentricitv. pericentric distance (from host halo center). and apocentric distance Orbital circularity is defined as the ratio of the orbit angular momentum to that of the cireular orbit with the same energy and can also be related to eccentricity In the point particle two-bocky approximation. the above quantities are all conserved throughout the orbit. so they do not depend on the separation of the satellite from the host halo.,"separation $\mathbf{r} = \mathbf{r}_{\rm host} - \mathbf{r}_{\rm sat}$ and velocity $\dot{\mathbf{r}} = \dot{\mathbf{r}}_{\rm host} - \dot{\mathbf{r}}_{\rm sat}$, the orbital energy is and the angular momentum is With these, we computed the orbital eccentricity, pericentric distance (from host halo center), and apocentric distance Orbital circularity is defined as the ratio of the orbit angular momentum to that of the circular orbit with the same energy and can also be related to eccentricity In the point particle two-body approximation, the above quantities are all conserved throughout the orbit, so they do not depend on the separation of the satellite from the host halo."
 Two of the above quantities (for example. circularity plus pericenter) are sulficient to uniquely describe an orbit.," Two of the above quantities (for example, circularity plus pericenter) are sufficient to uniquely describe an orbit."
 An alternate description can be given by the satellites radial. Vi. plus tangential Vo. velocity. though they need to be defined at a given radius. which here will be at the host halo virial racius.," An alternate description can be given by the satellite's radial, $\vr$ , plus tangential, $\vt$, velocity, though they need to be defined at a given radius, which here will be at the host halo virial radius."
 To clo this. we used the conservation of orbital energy anc angular momentum to evolve the velocities from the satellite’s measured. location to where its center crosses the host halo virial radius. A4. which was derived from the halo Fok mass aad concentration assuming a spherical NEW (7). density. profile.," To do this, we used the conservation of orbital energy and angular momentum to evolve the velocities from the satellite's measured location to where its center crosses the host halo virial radius, $\rvir$, which was derived from the halo FoF mass and concentration assuming a spherical NFW \citep{NFW96} density profile."
 llereon. all satellite halo distances and: velocities. are sealed to the host halo. virial raclius ancl virial circular velocity. Var=CMta.," Hereon, all satellite halo distances and velocities are scaled to the host halo virial radius and virial circular velocity, $\vvir = \sqrt{G\mvir/\rvir}$."
 The aim of this work is to examine the orbital parameter distributions of satellite halos as they first. pass through the virial radius of a larger host halo., The aim of this work is to examine the orbital parameter distributions of satellite halos as they first pass through the virial radius of a larger host halo.
 However. because of the finite time spacing of simulation outputs. satellite halos were identified at a variety of distances. [rom the host halo virial radius prior to merging.," However, because of the finite time spacing of simulation outputs, satellite halos were identified at a variety of distances from the host halo virial radius prior to merging."
 This presents some cdillicultv in. directly measuring orbital parameter distributions. since satellite halos can experience mass loss (or growth). dynamical friction. and tidal forces. all of which alter the orbital parameters as satellites traverse significant clistances.," This presents some difficulty in directly measuring orbital parameter distributions, since satellite halos can experience mass loss (or growth), dynamical friction, and tidal forces, all of which alter the orbital parameters as satellites traverse significant distances."
 To cireumyent this issue of finite time resolution. we estimated: satellite orbital parameter distributions bv selecting only satellite halos that were within a small separation of the host halo virial radius at the output prior to merging (asdoneby777).," To circumvent this issue of finite time resolution, we estimated satellite orbital parameter distributions by selecting only satellite halos that were within a small separation of the host halo virial radius at the output prior to merging \citep[as done by][]{VitKlyKra02,Ben05,WanJinMao05}."
 Specifically. we found satellite halos that were about to merge with a larger host halo aud whose edge (given by its own virial radius) was in range dod|Ar)Rey of the virial radius of the host We used Ar=0.25. finding no significant dillerence using smaller values.," Specifically, we found satellite halos that were about to merge with a larger host halo and whose edge (given by its own virial radius) was in range $[1,1+\Delta r]\,\rvir$ of the virial radius of the host We used $\Delta r = 0.25$, finding no significant difference using smaller values."
 Since more radial orbits spend less time within the eiven radial shell. to properly estimate the overall orbita distributions one must scale the orbital counts by. the crossing time within the shell.," Since more radial orbits spend less time within the given radial shell, to properly estimate the overall orbital distributions one must scale the orbital counts by the crossing time within the shell."
" Specifically. we sealed the counts by Adsl... where A? is the output time interva and foros. is the orbital crossing time for a satellite halo cdge to orbit [rom a separation of Ar to 0."" Since this method of estimating satellite orbita distributions considers only satellites close to the location of interest. it is little-alfected by mass stripping. dvnamica friction. and tidal forces. and the error on the estimatec orbital parameters is expected. to be less than 104 (2).."," Specifically, we scaled the counts by $\Delta t/t_{\rm cross}$, where $\Delta t$ is the output time interval and $t_{\rm cross}$ is the orbital crossing time for a satellite halo edge to orbit from a separation of $\Delta r$ to $0$ Since this method of estimating satellite orbital distributions considers only satellites close to the location of interest, it is little-affected by mass stripping, dynamical friction, and tidal forces, and the error on the estimated orbital parameters is expected to be less than $10\%$ \citep{Ben05}."
 There are. however. limitations to the point particle two- approximation. which ignores multi-body interactions and extended. triaxial halo profiles. (see. the Appendix).," There are, however, limitations to the point particle two-body approximation, which ignores multi-body interactions and extended, triaxial halo profiles (see the Appendix)."
 Some satellites (up to 20% at z=0) had their rji or Moye Within the radial selection shell. so these objects did not have a well-defined crossing time even though they did merge with the host halo in the next output.," Some satellites (up to $20\%$ at $z=0$ ) had their $\rperi$ or $\rapo$ within the radial selection shell, so these objects did not have a well-defined crossing time even though they did merge with the host halo in the next output."
 This population was dominated by satellites on highly circular orbits. such that Mae Was never bevond the outer radial shell.," This population was dominated by satellites on highly circular orbits, such that $\rapo$ was never beyond the outer radial shell."
 la weighting their counts for orbital distributions. these satellites were given a shell crossing time equal to the longest well-defined crossing time satellite in the eiven output. which amounts to a minimal weighting without ciscarding them.," In weighting their counts for orbital distributions, these satellites were given a shell crossing time equal to the longest well-defined crossing time satellite in the given output, which amounts to a minimal weighting without discarding them."
" Additionally. some merging satellites were on formally unbouncl orbits. though this represents less than 2!vs of satellites regardless of mass or Finally. we also examined how using Afooueri, instead of FoF massinfluences the We found that the averages of the orbital clistributions from the two methods are within error. and thus the specific"," Additionally, some merging satellites were on formally unbound orbits, though this represents less than $2\%$ of satellites regardless of mass or Finally, we also examined how using $M_{200{\rm crit}}$ instead of FoF massinfluences the We found that the averages of the orbital distributions from the two methods are within error, and thus the specific"
"If we were to ignore all the refinements in TIT and TV the same formulac hold provided thatp,, is replaced by (22)po Whenever we consider neutrou stars.","If we were to ignore all the refinements in III and IV the same formulae hold provided that$\rho_o$ is replaced by $\left ({m_n \over
m_e}\right )^3\rho_o$ whenever we consider neutron stars."
 Thus the radii are smaller. by the factor e. than the white cawart of the same mass with y=1.," Thus the radii are smaller, by the factor ${m_e \over m_n}$, than the white dwarf of the same mass with $\mu_e=1$."
 The Chandrasekhar limit is almost wichanged at 5.8242..., The Chandrasekhar limit is almost unchanged at $5.82 M_\odot$.
 Tu practice effects IIT aud IV. reduce this to between two and 3.6 solar masses. sce Shapiro Teukolsky.," In practice effects III and IV reduce this to between two and 3.6 solar masses, see Shapiro Teukolsky."
 These basic points were understood by Baade Zwicky (1931) writing soon after the neutron was discovered., These basic points were understood by Baade Zwicky (1934) writing soon after the neutron was discovered.
 Later they took spectra of the Crab Pulsar without realising that it pulsed thirty times a second. but Baade could not interpret the continuous spectrum as. in spite of Schott’s (1912) work. Svuchotrou radiation was not known in astronomy at that time.," Later they took spectra of the Crab Pulsar without realising that it pulsed thirty times a second, but Baade could not interpret the continuous spectrum as, in spite of Schott's (1912) work, Synchotron radiation was not known in astronomy at that time."
 0—].dSd£lat£ =O.," $\theta = 1, \ \ {d \theta \over d \xi} = 0 \ {\rm at}\ \xi = 0$ ."
 Let the edge be at£2 £i., Let the edge be at$\xi = \xi_1$ .
" Define the “mass” within © bxµίς)=Ih0""? d£.", Define the “mass” within $\xi$ by$\mu(\xi)=\int^\xi_0\theta^n\xi^2d\xi$ .
Observations anc modelling of the cireumnuclear regions (CNRs) of barred. Ingalaxies In]give important. clues about the nature of active galactic nuclei (AGN). (cireum)nuclear starbursts and the eas dynamics in the central kpe (c.g. Ixnapen et al.,"Observations and modelling of the circumnuclear regions (CNRs) of barred galaxies give important clues about the nature of active galactic nuclei (AGN), (circum)nuclear starbursts and the gas dynamics in the central kpc (e.g. Knapen et al."
 1995a.b: Buta Combes 1996: Elmecercen et al.," 1995a,b; Buta Combes 1996; Elmegreen et al."
 1997: Quillen et al., 1997; Quillen et al.
 1999: Reean Mulchaey. 1990: sec also reviews bv Ixnapen 1999 and Shlosman 1990)., 1999; Regan Mulchaey 1999; see also reviews by Knapen 1999 and Shlosman 1999).
 The presence of non-axisvnimetrics. such as oval distortions or bars. in the gravitational potential of a disc galaxy can lead to the infall of gascous material. [rom ealactocentrie radii of a [ew kpe. into the central kpe.," The presence of non-axisymmetries, such as oval distortions or bars, in the gravitational potential of a disc galaxy can lead to the infall of gaseous material, from galactocentric radii of a few kpc, into the central kpc."
 The nonaxisvmmetry facilitates the removal of angular momentum from the gas. anc directs it into the CNR. or I .i ≻∪⊳∖⊳∖⊀↓∣⋡⇂∖⇁↥⇂↥⋖⊾⊔⋯∼↓⋖⋅⊔⊳∖↿↔↓↕," The nonaxisymmetry facilitates the removal of angular momentum from the gas, and directs it into the CNR or possibly the nucleus (Shlosman et al."
↓∪⊳∖⊔↓⊳⋯∢⊾↿⊳↧↓⊳↓≤⋗≤⊓∶⇀∖↿↓↕⊳⋯⊳↧⊳∖⊳∖∪⊔↓⊳↧‘ ow‘ 1992:00. hyPhinneymu 1994)., 1990; Athanassoula 1992; Phinney 1994).
"o ThiThis gasE fuelling""uelli processOCOÓ is] expectedοςLO to induce morphological signatures in the bars anc CNIts."," This gas fuelling process is expected to induce morphological signatures in the bars and CNRs,"
"there is no ionization equilibrium for temperatures lower than 10"" Ik. the ionization state of the gas at Z7«c10"" Ix is obtained by solving the time-dependent ionization equations. for all ionic. species of LL. He. €. N. O. Ne. Ale. Si. 5. Ar and Fe.","there is no ionization equilibrium for temperatures lower than $10^6$ K, the ionization state of the gas at $T<10^6$ K is obtained by solving the time-dependent ionization equations, for all ionic species of H, He, C, N, O, Ne, Mg, Si, S, Ar and Fe."
" We adopt a non-equilibrium cooling function [or temperatures lower than 10"" Ix. since the recombination time of important ions is longer than the cooling time at these temperatures."," We adopt a non-equilibrium cooling function for temperatures lower than $10^6$ K, since the recombination time of important ions is longer than the cooling time at these temperatures."
 The cooling function and the coefficients of collisional ionization. recombination and charge exchange of the ionization equations are all calculated with the atomic database of the photoionization code AANGABA (Ciruenwald Viegas 1992).," The cooling function and the coefficients of collisional ionization, recombination and charge exchange of the ionization equations are all calculated with the atomic database of the photoionization code AANGABA (Gruenwald Viegas 1992)."
 The adopted abundances are solar (Cirevesse Anclers 1989)., The adopted abundances are solar (Grevesse Anders 1989).
" Since the version of AANGABA used in JE did not include molecules. their cooling function was valid only for Z7zο10"" Ix. In the present extension of JE calculations to lower temperatures. we use the equilibrium cooling functions. and. molecular fractions down to 7=LOO Ix. given bv a purely collisional model (Lepp et al."," Since the version of AANGABA used in JF did not include molecules, their cooling function was valid only for $T \ga 5\times 10^3$ K. In the present extension of JF calculations to lower temperatures, we use the equilibrium cooling functions and molecular fractions down to $T=100$ K, given by a purely collisional model (Lepp et al."
 1985. and references therein).," 1985, and references therein)."
 The clouds are modelled: as slabs shrinking in the direction transverse to that of the (which is the radial direction towards the cluster centre. since we assume spherical geometry for the Iow)).," The clouds are modelled as slabs shrinking in the direction transverse to that of the (which is the radial direction towards the cluster centre, since we assume spherical geometry for the )."
 The faces of the slab are in the plane of the magnetic field. so that the magnetic field. is perpendicular to the direction. of compression.," The faces of the slab are in the plane of the magnetic field, so that the magnetic field is perpendicular to the direction of compression."
 The initial density perturbations are characterized by an amplitude ct and a Length. scale L. following ορ=.isn(2xe/L)/(2x.riL). where .r is the direction. of the compression (expansion) of the perturbation.," The initial density perturbations are characterized by an amplitude $A$ and a length scale $L$, following $\delta\rho/\rho=A\,sin(2\pi\,x/L)/ (2\pi\,x/L)$, where $x$ is the direction of the compression (expansion) of the perturbation."
 We have also assumed that the perturbations are isobaric and nonlinear (.1 = 1)., We have also assumed that the perturbations are isobaric and nonlinear $A$ = 1).
 Vhe slab. geometry is justified. by the fact that the observed. line. emission is filamentary. which suggests that the perturbations are sheetlike rather than spherical.," The slab geometry is justified by the fact that the observed line emission is filamentary, which suggests that the perturbations are sheetlike rather than spherical."
 We start. to follow the evolution of the perturbations from the nonlinear stage in view of the uncertainties about processes suppressing the erowth of thermal instabilities in cooling Iows., We start to follow the evolution of the perturbations from the nonlinear stage in view of the uncertainties about processes suppressing the growth of thermal instabilities in cooling flows.
 In this work. we consider two models. representing perturbations evolving in the inner and in the outer parts of theow.," In this work, we consider two models, representing perturbations evolving in the inner and in the outer parts of the."
". In model A. representing the oof a cooling condensationin the inner aat a radius r—LO kpc. the unperturbed ng=0.1 cm and Z=10"" K are assumed."," In model A, representing the of a cooling condensationin the inner at a radius $r=10$ kpc, the unperturbed $n_H=0.1$ $^{-3}$ and $T=10^7$ K are assumed."
 Lowas fixed at 1 kpe [or all the models., $L$ was fixed at 1 kpc for all the models.
 TFhis length scale is suggested. for instance. by the spatial Huctuations in the velocity of optical filaments resolved in nearby cooling Lows (Lleckman et al.," This length scale is suggested, for instance, by the spatial fluctuations in the velocity of optical filaments resolved in nearby cooling flows (Heckman et al."
 1989)., 1989).
" We fixed 3=Pegf/D,,,OL for the unperturbed medium. a representative value for the range of 7=0.011: expected in the central 10 kpe of cooling Hows."," We fixed $\beta=P_B/P_{gas}=0.1$ for the unperturbed medium, a representative value for the range of $\beta=0.01-1$, expected in the central 10 kpc of cooling flows."
 We have considered that AIR. proceeds at an cllicicney AZ.=0.02., We have considered that MR proceeds at an efficiency $M_e=0.02$.
" The ALR heating is turned on only for 3>3,=1.", The MR heating is turned on only for $\beta > \beta_{on}=1$.
 The value of ἐν used in eq. (, The value of $l_{MR}$ used in eq. (
2) is derived assuming that the magnetic field sulfered a 1D compression as the gas condensed from the intracluster medium to form the condensations.,2) is derived assuming that the magnetic field suffered a 1D compression as the gas condensed from the intracluster medium to form the condensations.
" The smoothness of the radio images of radio haloes in clusters imply a correlation length. of the magnetic field 7,zz15 kpe (Tribble 1993).", The smoothness of the radio images of radio haloes in clusters imply a correlation length of the magnetic field $l_c \la 15$ kpc (Tribble 1993).
 During a 1D compression. the quantity fog is conserved.," During a 1D compression, the quantity $l_c n_H$ is conserved."
 Assuming £c;=10 kpe and με=107 emο for the intracluster medium. faye—£L within the condensations is obtained [from the local ng.," Assuming $l_{c,ICM}= 10$ kpc and $n_{H,ICM}=10^{-3}$ $^{-3}$ for the intracluster medium, $l_{MR}\equiv l_c $ within the condensations is obtained from the local $n_H$."
" Model DB. describing the outer ((r=100 kpc) has A4,=0.02 and 3%,=1. and unperturbed £=10 kpe ny=510?. and Yo=7101 WK. ‘Lhe radial variation of ng and T assumed in models A ancl B follows X-ray spectroscopic studies and image deprojection analysis which allow to derive temperature. gradients and the density. runs with radius."," Model B, describing the outer $r=100$ kpc) has $M_e=0.02$ and $\beta_{on}=1$, and unperturbed $L=10$ kpc, $n_H=5\times 10^{-3}$, and $T=7\times 10^{7}$ K. The radial variation of $n_H$ and $T$ assumed in models A and B follows X-ray spectroscopic studies and image deprojection analysis which allow to derive temperature gradients and the density runs with radius."
" The radial cependence of the density impliec » models A and D. ng=510""(rfl00kpc)L07C. ds consistent the averaged racial profile of ny=(4644088)-10.""(e/100kpe)LITPEUGU found by White. Jones Forman (1997) in their sample of arge (AL>50 vro 1) lows detected with theOBSERVATORY."," The radial dependence of the density implied by models A and B, $n_H =5\times 10^{-3} (r/100\; {\rm kpc})^{-1.30}$, is consistent the averaged radial profile of $n_H =(4.64 \pm 0.88) \times 10^{-3} (r/100\; {\rm kpc})^{-1.26\pm 0.19}$ found by White, Jones Forman (1997) in their sample of large $\dot M > 50$ $^{-1}$ ) s detected with the."
" With respect to the temperature gradient. he cooling How region LOSrX100 kpe separates the inner How.. where the gas temperature approaches he virial temperature of the central (tvpical σ=300 km or 1=6.6«10"" I) from the general ICM. with a temperature roughly equal to. the"," With respect to the temperature gradient, the cooling flow region $10 \la r \la 100$ kpc separates the inner , where the gas temperature approaches the virial temperature of the central (typical $\sigma=300$ km $^{-1}$ or $T=6.6\times 10^6$ K) from the general ICM, with a temperature roughly equal to the"
"at high-z black hole accretion is associated with the delayed onset of AGB stars, one naturally expects a correlation between Eddington ratio and nitrogen abundance, as observed by us.","at $z$ black hole accretion is associated with the delayed onset of AGB stars, one naturally expects a correlation between Eddington ratio and nitrogen abundance, as observed by us."
" Note that, in this picture, the correlation between Eddington ratio and emission-line flux ratios involving suggest that the timescale of AGN feeding are shorter than the timescale of the nitrogen enrichment by AGB stars (<10? years)."," Note that, in this picture, the correlation between Eddington ratio and emission-line flux ratios involving suggest that the timescale of AGN feeding are shorter than the timescale of the nitrogen enrichment by AGB stars $< 10^8$ years)."
" The relation between L/Zgqq and nitrogen abundance is associated to black hole accretion and star formation, i.e. on phenomena occurring on relatively short time scales (~105 years)."," The relation between $L/L_{\rm Edd}$ and nitrogen abundance is associated to black hole accretion and star formation, i.e. on phenomena occurring on relatively short time scales $\sim 10^8$ years)."
" Conversely, the relation between black hole mass and metallicity is more fundamental and representative of the global evolution of these systems, since it connects physical quantities integrated over the whole formation history of black holes and galaxies."," Conversely, the relation between black hole mass and metallicity is more fundamental and representative of the global evolution of these systems, since it connects physical quantities integrated over the whole formation history of black holes and galaxies."
" Within this context it is interesting to note that NLSIs, which are characterized by very high Eddington ratios, have highτν ratio (Shemmer et al."," Within this context it is interesting to note that NLS1s, which are characterized by very high Eddington ratios, have high ratio (Shemmer et al."
" 2004), in line with the trend found above for luminous high-z quasars and suggesting that NLS1s have undergone vigorous star formation (Nagao et al."," 2004), in line with the trend found above for luminous $z$ quasars and suggesting that NLS1s have undergone vigorous star formation (Nagao et al."
 2002)., 2002).
" It should be noted that previous studies have found that quasar hosts are characterized by vigorous star formation (e.g., Maiolino et al."," It should be noted that previous studies have found that quasar hosts are characterized by vigorous star formation (e.g., Maiolino et al."
 2007a; Netzer et al., 2007a; Netzer et al.
 2007; Lutz et al., 2007; Lutz et al.
 2008)., 2008).
" Our finding of a correlation between black hole accretion rate and nitrogen abundance suggests that at earlier times, by about 10? years, star formation was even higher."," Our finding of a correlation between black hole accretion rate and nitrogen abundance suggests that at earlier times, by about $^8$ years, star formation was even higher."
 We have investigated the relationship between the metallicity and SMBH mass or Eddington ratio by producing composite SDSS quasar spectra with 33 subsamples divided in intervals of black hole mass and Eddington ratio at 2.3«z3.0., We have investigated the relationship between the metallicity and SMBH mass or Eddington ratio by producing composite SDSS quasar spectra with 33 subsamples divided in intervals of black hole mass and Eddington ratio at $2.3 < z < 3.0$.
 In each of these spectra we measure emission-line flux ratios that are sensitive to the metallicity of the BLR gas., In each of these spectra we measure emission-line flux ratios that are sensitive to the metallicity of the BLR gas.
" We have investigated the Mgu-Zpig dependence or the L/Lgaa-Zgim dependence by comparing the emission-line flux ratios of the composite spectra by fixing the Eddington ratio or black hole mass, respectively."," We have investigated the $M_{\rm BH}$ $Z_{\rm BLR}$ dependence or the $L/L_{\rm Edd}$ $Z_{\rm BLR}$ dependence by comparing the emission-line flux ratios of the composite spectra by fixing the Eddington ratio or black hole mass, respectively."
" We have found the following results: By assuming that the Μμμ- Μιωι relation applies at high-redshift, the Mpgu-Zpig relation is likely a consequence of the Μιωι-Ζυιι relation in galaxies."," We have found the following results: By assuming that the $M_{\rm BH}$ $M_{\rm bul}$ relation applies at high-redshift, the $M_{\rm BH}$ $Z_{\rm BLR}$ relation is likely a consequence of the $M_{\rm bul}$ $Z_{\rm bul}$ relation in galaxies."
 The lack of redshift evolution of the LAGN-Zgirm relation does not necessarily imply a lack of metallicity evolution in AGN host galaxies., The lack of redshift evolution of the $L_{\rm AGN}$ $Z_{\rm BLR}$ relation does not necessarily imply a lack of metallicity evolution in AGN host galaxies.
" Indeed, as already suggested in previous works (e.g., Juarez et al."," Indeed, as already suggested in previous works (e.g., Juarez et al."
" 2009), a combination of selection effects and of the co-evolution between black hole and galaxies causes quasars to be detected only once their host galaxy is already chemically evolved at any epoch, regardless of redshift."," 2009), a combination of selection effects and of the co-evolution between black hole and galaxies causes quasars to be detected only once their host galaxy is already chemically evolved at any epoch, regardless of redshift."
distribution as a uniform dise of diameter 2 kpe in the centre of the bubble. comparable with recent observations of compact starbursts (222)...,"distribution as a uniform disc of diameter $2$ kpc in the centre of the bubble, comparable with recent observations of compact hyper-starbursts \citep{caseyetal09,walteretal09,maiolinoetal07}."
 Taking the value of 6.3+1.3 mJy for the SMM J222174+0015 flux density at 850 jjm (2) we estimate a flux density of 9.7 mJy at 400 GHz. and plot contours of the convolved brightness temperature distribution in the central region of Figure 2..," Taking the value of $6.3 \pm 1.3$ mJy for the SMM J222174+0015 flux density at 850 $\mu$ m \citep{solomonvandenbout05} we estimate a flux density of 9.7 mJy at 400 GHz, and plot contours of the convolved brightness temperature distribution in the central region of Figure \ref{fig:shockplot}."
 The brightness temperature greyscale is here removed subject to a simple cut of 7;<0.4 mK showing that. thanks to the angular resolution of ALMA. the extended bubble is discernible beyond the strong-but-localized dust emission at its centre.," The brightness temperature greyscale is here removed subject to a simple cut of $\Delta T_b \le 0.4$ mK showing that, thanks to the angular resolution of ALMA, the extended bubble is discernible beyond the strong-but-localized dust emission at its centre."
 In practice. a more sophisticated removal of SMG dust contamination will be aided by he possibility of ultra-high resolution mapping of the bright central dust (ALMA can reach FWHM = 0.011 aresec at 400 GHz). and by the characteristic signature of the tSZ effect on CMB photons: a maximum spectral excess/decrement around 385/144. GHz and null crossover at 218 GHz (2).," In practice, a more sophisticated removal of SMG dust contamination will be aided by the possibility of ultra-high resolution mapping of the bright central dust (ALMA can reach FWHM = 0.011 arcsec at 400 GHz), and by the characteristic signature of the tSZ effect on CMB photons: a maximum spectral excess/decrement around 385/144 GHz and null crossover at 218 GHz \citep{carlstrometal02}."
. Multiple observations will require ‘urther telescope time. however.," Multiple observations will require further telescope time, however."
 In Section ?? it was noted that relativistic corrections to the SZ ure not included in the determination of y., In Section \ref{sect:sz} it was noted that relativistic corrections to the tSZ are not included in the determination of $y$.
 Using equation(8). the temperature of the gas in the shock boundary is 75=|myARS (4g).which implies a temperature72 10K or all the objects listed in Table |..," Using equation, the temperature of the gas in the shock boundary is $T_2 = 3 (m_e + m_p) \dot{R}^2_2 / (4 k_{\textrm{B}}) $ , which implies a temperature of $T_2 \gtrsim 10^7$ K for all the objects listed in Table \ref{tab:targets}."
 Calculating the ofreduction in SZ intensity using the first and second order relativistic terms described by ?. yields corrections of =3% for the hyper-starbursts isted., Calculating the reduction in tSZ intensity using the first and second order relativistic terms described by \citet{challinorlasenby98} yields corrections of $\simeq 3 \%$ for the hyper-starbursts listed.
 The effect is therefore negligible for the current model. but relevant for more detailed calculations.," The effect is therefore negligible for the current model, but relevant for more detailed calculations."
 The discussion so far has assumed adiabatic bubble expansion. but depends upon losses due to radiative cooling being small within the hyper-starburst age /.," The discussion so far has assumed adiabatic bubble expansion, but depends upon losses due to radiative cooling being small within the hyper-starburst age $t$."
 At the shock boundary temperatures of 15m 10'K the gas will be fully ionized but not sufficiently energetic for electron-positron annihilation to contribute to radiative processes: thermal bremsstrahlung and inverse-Compton cooling by CMB photons will dominate radiative losses (this latter being the net cooling effect of the tSZ itself)., At the shock boundary temperatures of $T_2 \gtrsim 10^7$ K the gas will be fully ionized but not sufficiently energetic for electron-positron annihilation to contribute to radiative processes: thermal bremsstrahlung and inverse-Compton cooling by CMB photons will dominate radiative losses (this latter being the net cooling effect of the tSZ itself).
" The radiative cooling will be greatest at the shock boundary #2. where the combination of plasma temperature 75 and density po are greatest. and so we calculate timescales for radiative losses in this region to explore the validity of the adiabatic expansion approximation,"," The radiative cooling will be greatest at the shock boundary $R_2$, where the combination of plasma temperature $T_2$ and density $\rho_2$ are greatest, and so we calculate timescales for radiative losses in this region to explore the validity of the adiabatic expansion approximation."
 We consider first the effect of inverse-Compton scattering of CMB photons., We consider first the effect of inverse-Compton scattering of CMB photons.
 In the relevant non-relativistic limit. the net energy loss rate for a single electron in the shock plasma is given by (see. e.g. 2)) where σι is the Thomson scattering cross- esp is the Stefan-Boltzmann constant. Zev is the CMB photontemperature at that epoch. and ο is the velocity of the," In the relevant non-relativistic limit, the net energy loss rate for a single electron in the shock plasma is given by (see, e.g., \citealp{longair92}) ) where $\sigma_{\textrm{T}}$ is the Thomson scattering cross-section, $\sigma_{\textrm{SB}}$ is the Stefan-Boltzmann constant, $T_{\textrm{CMB}}$ is the CMB photontemperature at that epoch, and $v_e$ is the velocity of the"
Pontetal.(2004) investigated the influence of age in the WY. (7) versus [Fe/II) relationship from a theoretical point of view.,\citet{pont04} investigated the influence of age in the $W'_V$ $W'_I$ ) versus [Fe/H] relationship from a theoretical point of view.
" Thev used the theoretical calculations of CaT equivalent widths for different values of loggy. T,yy and metallicity caleulated by Jorgensen together with the Padova stellar evolution models (Girardietal.2002)."," They used the theoretical calculations of CaT equivalent widths for different values of $\log g$, $_{eff}$ and metallicity calculated by \citet{jcj92} together with the Padova stellar evolution models \citep{girardi02}."
. They concluded that (he variation of W with age for a fixed metallicity would be negligible [or clusters older than 4 Gvr., They concluded that the variation of $W'$ with age for a fixed metallicity would be negligible for clusters older than 4 Gyr.
 ILowever. this was not the case for the vounger clusters.," However, this was not the case for the younger clusters."
 This is observed clearly in Figure 15 by Pontetal. (2004).., This is observed clearly in Figure 15 by \citet{pont04}. .
 For a given metallicity. the sequences in (he Aly οσα and Mj-XC« planes are separated as a function of their ages for clusters vounger (han ~4 Gyr.," For a given metallicity, the sequences in the $_V$ $\Sigma Ca$ and $_I$ $\Sigma
Ca$ planes are separated as a function of their ages for clusters younger than $\sim$ 4 Gyr."
 According to this ealeulation. for the same metallicity. W decreases with age.," According to this calculation, for the same metallicity, $W'$ decreases with age."
 Thus. metallicities for clusters vounger than 4 Gyr. caleulated [rom calibrations computed Irom old stars. will be underestimated.," Thus, metallicities for clusters younger than 4 Gyr, calculated from calibrations computed from old stars, will be underestimated."
 This age dependence is more important in ihe My. XC plane than in the M; “Ca one., This age dependence is more important in the $_V$ $\Sigma Ca$ plane than in the $_I$ $\Sigma Ca$ one.
 This means that Wy) would be less sensitive to age than Wy., This means that $W'_I$ would be less sensitive to age than $W'_V$.
 Using the Jorgensenetal.(1992) models and the DaSTI stellar evolution models (Dietrinfernietal.2004).. we have estimated the expected W differences as a function ol age.," Using the \citet{jcj92} models and the BaSTI stellar evolution models \citep{pie04}, we have estimated the expected $W'$ differences as a function of age."
" From these caleulations. lor (wo clusters with the same metallicity and age 10.5 and 0.6 Gyr respectively, the voungest cluster Ht would be approximately 0.7 lower than that of the oldest one."," From these calculations, for two clusters with the same metallicity and age 10.5 and 0.6 Gyr respectively, the youngest cluster $W'_V$ would be approximately 0.7 lower than that of the oldest one."
 This implies that the metallicity obtained lor voung clusters using this calibration would be 0.25 dex more metal-poor than (he actual metallicity., This implies that the metallicity obtained for young clusters using this calibration would be 0.25 dex more metal-poor than the actual metallicity.
 In the case of Vj. the difference would be 0.4À.. so the metallicity obtained for young clusters wotld be 0.15 dex more metal-poor than the actual one.," In the case of $W'_I$, the difference would be 0.4, so the metallicity obtained for young clusters would be 0.15 dex more metal-poor than the actual one."
 As we can see in Figure 15 by (2004).. the difference would be similar for different metallicities.," As we can see in Figure 15 by \citet{pont04}, the difference would be similar for different metallicities."
 From our data. we confirm (that the influence of age is weak.," From our data, we confirm that the influence of age is weak."
 In Figure 13. we plot WI versus age for clusters with —0.17< ζω< 40.07., In Figure \ref{agetest} we plot $W'_I$ versus age for clusters with $-0.17\leq$ $_{CG97}\leq$ +0.07.
 We have selected. this range because il contains clusters with a wide range of ages and is small enough for the metallicity differences to be within the uncertainties., We have selected this range because it contains clusters with a wide range of ages and is small enough for the metallicity differences to be within the uncertainties.
 We can see that clusters with ages vounger than 5 αντ (NGC 2141. NGC 2682. NGC 6819 and NGC 7789) have similar WW than the oldest one (NGC! 6528).," We can see that clusters with ages younger than 5 Gyr (NGC 2141, NGC 2682, NGC 6819 and NGC 7789) have similar $W'_I$ than the oldest one (NGC 6528)."
 There are only two clusters that deviate wiclely [rom the behaviour of the others., There are only two clusters that deviate widely from the behaviour of the others.
 One of these is the voungest cluster. NGC 6705. which has a larger W7 than the oldest clusters.," One of these is the youngest cluster, NGC 6705, which has a larger $W'_I$ than the oldest clusters."
 This is contrary to the theoretical prediction that it should be smaller., This is contrary to the theoretical prediction that it should be smaller.
 However. we have to take into account that differences of 0.5 iim VW mean differences of —0.1 dex in |Fe/II].," However, we have to take into account that differences of 0.5 in $W'_I$ mean differences of $\sim$ 0.1 dex in [Fe/H]."
 So the observed variations are similar to the uncertainty in the determination of [Fe/II]., So the observed variations are similar to the uncertainty in the determination of [Fe/H].
 Our dataare not accurate enough to detect the influence of age because the uneertainty in the metallicity determination of clusters is similar to the expected variations due to age., Our dataare not accurate enough to detect the influence of age because the uncertainty in the metallicity determination of clusters is similar to the expected variations due to age.
"unrealistic upper cut-off in the number of considered iron-group ionization stages causes an artificial over-population of the highest ionization stage, and thus, affects its lines and the flux level.","unrealistic upper cut-off in the number of considered iron-group ionization stages causes an artificial over-population of the highest ionization stage, and thus, affects its lines and the flux level."
" InFigurel,, the impact of iron-group opacities on the astrophysical flux is demonstrated."," In, the impact of iron-group opacities on the astrophysical flux is demonstrated."
" For all values ofT.g,, the necessary ionization stages of all atoms are determined in advance by test calculations."," For all values of, the necessary ionization stages of all atoms are determined in advance by test calculations."
" The selection criterion is, that at least the ion(e."," The selection criterion is, that at least the ion."
"g.. in the case of species X), X""* that is dominant in the line-forming region has to be included together with the neighboring two,"," in the case of species X), $^{\mathrm{n}+}$ that is dominant in the line-forming region has to be included together with the neighboring two,."
"ie.. ΧΑ9 and X@+)+,", $^{\mathrm{(n-1)}+}$ and $^{\mathrm{(n+1)}+}$.
" 'Thus, the model atoms contain in general three to five ionization stages."," Thus, the model atoms contain in general three to five ionization stages."
"E.g.. in the case of our generic group model atom, is dominant at and [cm sec~?]2).."," in the case of our generic iron-group model atom, is dominant at and [cm $^{-2}$ ]."
 We therefore —selected the ionization stages —XVIII., We therefore selected the ionization stages –.
 Statistics of the model atoms are shown inTable1., Statistics of the model atoms are shown in.
". In total, 228 atomic levels are treated in NLTE, 360 additional levels in LTE, and 349 individual line transitions are considered."," In total, 228 atomic levels are treated in NLTE, 360 additional levels in LTE, and 349 individual line transitions are considered."
 A first grid of models is composed of H+He+C+N+0O with solar abundance ratios (7?) within and a fixed surface gravity of(Figure3)., A first grid of models is composed of H+He+C+N+O with solar abundance ratios \citep{AEA09} within and a fixed surface gravity of.
. We note that all synthetic energy distributions in our model grids described here are available at in (SEDs)Virtual Observatory VO)) compliant form from the service provided by 0.01/'TrSpectra.jsp?the /www.g-vo.org)) as well as for the use withdocs/xanadu/xspec., We note that all synthetic energy distributions (SEDs) in our model grids described here are available at in Virtual Observatory ) compliant form from the service provided by the ) as well as for the use with.
. We started with the RGS spectra of 2003 April(Table2).., We started with the RGS spectra of 2003 April.
" The data were reduced with the ESA Science Analysis (SAS) software, version 5.3.3, using the latest calibration files"," The data were reduced with the ESA Science Analysis (SAS) software, version 5.3.3, using the latest calibration files"
expand.,expand.
 This will alter the importance of the gas potential on the stars meaning that initially cool clusters are more able to survive gas expulsion. whilst initially warm clusters are less likely to.," This will alter the importance of the gas potential on the stars meaning that initially cool clusters are more able to survive gas expulsion, whilst initially warm clusters are less likely to."
 This can remove much of the dependence of survival after gas expulsion from the importance of the initial true SEI ancl place it instead on the importance of the initial dynamical state of the stars (see also ?:: 72))., This can remove much of the dependence of survival after gas expulsion from the importance of the initial true SFE and place it instead on the importance of the initial dynamical state of the stars (see also \citealp{Verschueren1989}; ; \citealp{Goodwin2009ApSS}) ).
 We have also shown that there is à very significant scatter due to both the intrinsic dillerences between (statistically the same) clusters (κου ?7)). and on the exact virial ratio at the onset of gas expulsion.," We have also shown that there is a very significant scatter due to both the intrinsic differences between (statistically the same) clusters (see \citealp{Allison2010}) ), and on the exact virial ratio at the onset of gas expulsion."
" There is. significant observational and theoretical evidence that the initial distributions of stars are not smooth. nor are theyvirialisecl (see 2: ?: Pir 75 Pi 7:5 Tu 2 2 and references in all of these papers) which is a natural consequence of eravoturbulent star formation (see e.g. ο, ο 7)."," There is significant observational and theoretical evidence that the initial distributions of stars are not smooth, nor are theyvirialised (see \citealp{Elmegreen2001}; ; \citealp{Bate2003}; ; \citealp{Bonnell2003}; ; \citealp{Bertout2006}; \citealp{Allen2007}; \citealp{Kraus2008}; \citealp{Gutermuth2009}; \citealp{Clarke2010}; \citealp{Bressert2010} and references in all of these papers) which is a natural consequence of gravoturbulent star formation (see e.g. \citealp{Elmegreen2004}; \citealp{McKee2007}; \citealp{Bergin2007}; \citealp{Clarke2010}) )."
 Indeed. evidence points towards sub-virial initial conditions for stars (seo ? and references therein).," Indeed, evidence points towards sub-virial initial conditions for stars (see \cite{Allison2010} and references therein)."
 We therefore argue that our initial conditions are far more realistic than those of a smooth. relaxed cluster as generally used. before (e.g. 2: 2: 7))," We therefore argue that our initial conditions are far more realistic than those of a smooth, relaxed cluster as generally used before (e.g. \citealp{Goodwin1997a}; \citealp{Goodwin2006}; \citealp{Baumgardt2007}) )."
 7? conduct a large parameter study. investigating the καν parameters| controlling the final bound. fraction of justers that have undergone gas expulsion using smooth uxd spherical initial stellar distributions., \cite{Proszkow2009} conduct a large parameter study investigating the key parameters controlling the final bound fraction of clusters that have undergone gas expulsion using smooth and spherical initial stellar distributions.
 They. find a clear rend with SEI although they. too. use sub-virial initial gacllar distributions.," They find a clear trend with SFE although they, too, use sub-virial initial stellar distributions."
 We note that we also see a trend. for increasing bound fraction with increasing SEI (Figure 7)). out that it is highly scattered.," We note that we also see a trend for increasing bound fraction with increasing SFE (Figure \ref{fboundsfe}) ), but that it is highly scattered."
 The source of this scatter is re use of elumpy initial stellar distribution., The source of this scatter is the use of clumpy initial stellar distribution.
 Therefore the gareneth of the star formation ellicienev as a predictor for 1. survival of a cluster to mass [oss is severely weakened with the use of far more realistic initially clumipy stellar clistributions., Therefore the strength of the star formation efficiency as a predictor for the survival of a cluster to mass loss is severely weakened with the use of far more realistic initially clumpy stellar distributions.
 lt has often previously heen assumed that for a cluster to survive it must have had a high. SPE and so the small numbers of clusters we see must be a high-SEE tail to the SEE distribution (see e.g. 2))., It has often previously been assumed that for a cluster to survive it must have had a high SFE and so the small numbers of clusters we see must be a high-SFE tail to the SFE distribution (see e.g. \citealp{Parmentier2008}) ).
 However. we have shown that some low-SEL clusters can survive if they are ‘luck’ chough to have the right initial conditions.," However, we have shown that some low-SFE clusters can survive if they are `lucky' enough to have the right initial conditions."
" Indeed. the low survival rates found for voung clusters of only ~LO per cent (2)) may be better explained as these being the few clusters with the ""right, initial conditionsthan being an extremely high-SELE (240 or 50 per cent) tail of star formation."," Indeed, the low survival rates found for young clusters of only $\sim
10$ per cent \citealp{Lada2003}) ) may be better explained as these being the few clusters with the `right' initial conditionsthan being an extremely high-SFE $>40$ or $50$ per cent) tail of star formation."
 We note that our simulations are highly idealised., We note that our simulations are highly idealised.
 Stars are equal mass. when ? have demonstrated. that. rapid mass segregation can occur with a more realistic initial mass function.," Stars are equal mass, when \cite{Allison2009b} have demonstrated that rapid mass segregation can occur with a more realistic initial mass function."
 We do not include any primordial binaries. although their presence can clearly influence the dynamics of a cluster (?)).," We do not include any primordial binaries, although their presence can clearly influence the dynamics of a cluster \citealp{Goodman1989}) )."
 The simulations of ? show that changes in binary fraction. ancl scattering between stars in the stellar outllow (that can result in binary hardening) can increase the final bound. fraction of a cluster.," The simulations of \cite{Kroupa2001} show that changes in binary fraction, and scattering between stars in the stellar outflow (that can result in binary hardening) can increase the final bound fraction of a cluster."
 In 7. the binary [raction does not change significantly during the cluster formation process as a result of formation in an initially high stellar density environment - within sub clumps., In \cite{Moeckel2010} the binary fraction does not change significantly during the cluster formation process as a result of formation in an initially high stellar density environment - within sub clumps.
 However their. initial conditions are limited to statistics of one. and we have further demonstrated that. clumpy initial conditions can result in highly stochastic behaviour.," However their initial conditions are limited to statistics of one, and we have further demonstrated that clumpy initial conditions can result in highly stochastic behaviour."
 Furthermore we assume instantaneous gas removal. although ? demonstrate that a slower rate of gas removal can result in a higher cluster survival rate.," Furthermore we assume instantaneous gas removal, although \cite{Baumgardt2007} demonstrate that a slower rate of gas removal can result in a higher cluster survival rate."
 The length. of the embedded: phase is fixed at 3 Myr in our simulations although this could) vary. depending on the nature of the gas removal mechanism., The length of the embedded phase is fixed at 3 Myr in our simulations although this could vary depending on the nature of the gas removal mechanism.
 Despite these simplifications. we argue that the idealised nature of the simulations has enabled us to more clearly test. the implications of an initially sub-virial. and clumpy stellar distribution.," Despite these simplifications, we argue that the idealised nature of the simulations has enabled us to more clearly test the implications of an initially sub-virial, and clumpy stellar distribution."
 We defer a less idealised study to a later paper., We defer a less idealised study to a later paper.
 Our simulations do have an obvious problem. however. in that we use a smooth. static background potential for the eas.," Our simulations do have an obvious problem, however, in that we use a smooth, static background potential for the gas."
 Phis has two main problems., This has two main problems.
 Firstly. the initial gas aud stellar distribution do not match. despite the fact that our stars are assumed to have formed from this background gas.," Firstly, the initial gas and stellar distribution do not match, despite the fact that our stars are assumed to have formed from this background gas."
 secondly. the gas is not able to respond to the motion of the stars. C," Secondly, the gas is not able to respond to the motion of the stars. ("
X third. but less important. problem. is that we assume that all of the stars form instantancously.},"A third, but less important problem, is that we assume that all of the stars form instantaneously.)"
 The importance of both problems comes down to how well the motions of the gas ancl the stars ave coupled., The importance of both problems comes down to how well the motions of the gas and the stars are coupled.
 Lf conditions are such that both stars and (at least a significant fraction of the) gas move together in the potential then we would expect both to collapse or expand together ancl the LSF and true star formation cllicicney to remain roughly constant., If conditions are such that both stars and (at least a significant fraction of the) gas move together in the potential then we would expect both to collapse or expand together and the $LSF$ and true star formation efficiency to remain roughly constant.
 Lowever. if the bulk of the gas does not notice the stars because it is not involved in their formation and. the relative gravitational inlluence of the stars is small. then our approximations should be roughly correct.," However, if the bulk of the gas does not notice the stars because it is not involved in their formation and the relative gravitational influence of the stars is small, then our approximations should be roughly correct."
 We would argue that at low true star formation elliciencies that the bulk of the gas would be uncoupled from the stars., We would argue that at low true star formation efficiencies that the bulk of the gas would be uncoupled from the stars.
 We are working on more detailed simulations with a live background potential which we will present in future papers., We are working on more detailed simulations with a live background potential which we will present in future papers.
 We perform N-bock simulations. of sub-structured.. non-equilibrium SOOAL. clusters of No=1000 equal-mass stars in a static background potential.," We perform $N$ -body simulations of sub-structured, non-equilibrium $500 M_\odot$ clusters of $N=1000$ equal-mass stars in a static background potential."
 Phe mass of gas is varied to simulate star formation efficiencies (SEIS) of 20 to 40 per cent., The mass of gas is varied to simulate star formation efficiencies (SFEs) of 20 to 40 per cent.
 After 3 Myr of dynamical evolution. the potential is instantancously removed to model the ellect of gas expulsion from the cluster.," After 3 Myr of dynamical evolution, the potential is instantaneously removed to model the effect of gas expulsion from the cluster."
 Previous work with initially smooth and. equilibrium clusters has shown that there is a critical SEIS for the survival of fat least part of) the cluster of ~30 per cent (e.g. 2:2? and references therein)., Previous work with initially smooth and equilibrium clusters has shown that there is a critical SFE for the survival of (at least part of) the cluster of $\sim 30$ per cent (e.g. \cite{Goodwin2006}; \cite{Baumgardt2007} and references therein).
 However. it has also been known that it isthe conditions of gas expulsionthat areimportant ininfluencing the evolution of the star cluster following gas expulsion (?.. 7)).," However, it has also been known that it isthe conditions of gas expulsionthat areimportant ininfluencing the evolution of the star cluster following gas expulsion \citealp{Verschueren1989}, , \citealp{Goodwin2009ApSS}) )."
 Our key results may be summarised as follows., Our key results may be summarised as follows.
Galaxy groups are key systems in advancing our understanding of structure formation and evolution.,Galaxy groups are key systems in advancing our understanding of structure formation and evolution.
 They contain the majority of galaxies in the universe. and are precursors to the most massive structures. ie. clusters. giving them cosmological importance.," They contain the majority of galaxies in the universe, and are precursors to the most massive structures, i.e. clusters, giving them cosmological importance."
 However. they show departures from the scaling relations obeyed by galaxy clusters indicating that they are not simply scaled-down versions of clusters.," However, they show departures from the scaling relations obeyed by galaxy clusters indicating that they are not simply scaled-down versions of clusters."
 Theoretical or computational models (e.g. Navarro.Frenk&White (1995))) based on simple gravitational collapse and shock heating would lead to self-similar structure of the inter-galactic medium (IGM) in clusters and groups., Theoretical or computational models (e.g. \citet{nfw95}) ) based on simple gravitational collapse and shock heating would lead to self-similar structure of the inter-galactic medium (IGM) in clusters and groups.
 This in turn implies scaling relations between the global properties of clusters: Lyx7%. Lx xcaland MxTy? where Lx is the total X-ray luminosity. Tx the gas temperature. c the velocity dispersion of cluster galaxies and A is the total gravitational mass of the cluster.," This in turn implies scaling relations between the global properties of clusters: $L_X\propto T_X^2$, $L_X\propto \sigma^4$ and $M\propto T_X^{3/2}$, where $L_X$ is the total X-ray luminosity, $T_X$ the gas temperature, $\sigma$ the velocity dispersion of cluster galaxies and $M$ is the total gravitational mass of the cluster."
 While some studies tind that the observed properties of the most massive clusters are close to the above relations (Allen&Fabian1998:Xue.Jin&Wu 2001). the self-similar model clearly breaks down in smaller systems. with observations indicating lower than expectec luminosities for a given temperature or velocity dispersion.," While some studies find that the observed properties of the most massive clusters are close to the above relations \citep{allen98,xuejinwu01}, the self-similar model clearly breaks down in smaller systems, with observations indicating lower than expected luminosities for a given temperature or velocity dispersion."
 White. found Ly-x77 for clusters observec with the X-ray Observatory.," \citet{wjf97}
 found $L_X\propto T^3$ for clusters observed with the X-ray Observatory."
 For galaxy groups. Mulchaey&Zabludoff(1998). found Lyx77. consistent with clusters. while Helsdon&Ponman(2000) and Xue&Wu(2000). founc much steeper relations. LxxT.," For galaxy groups, \citet{mz98} found $L_X
\propto T^3$, consistent with clusters, while \citet{helsdon00}
 and \citet{xue00} found much steeper relations, $L_X\propto T^5$."
 Galaxy groups are rapidly evolving and diverse systems. anc many are nof virialised (eg. Rasmussenetal. (20062).," Galaxy groups are rapidly evolving and diverse systems, and many are not virialised (eg. \citet{jesper06}) )."
 Thus studying a sample of well-characterised galaxy groups. in terms of their stellar properties and IGM. might help us to understand some of the observed diversity in group properties.," Thus studying a sample of well-characterised galaxy groups, in terms of their stellar properties and IGM, might help us to understand some of the observed diversity in group properties."
 Jonesetal.(2003) studied a flux-limited sample of old galaxy, \citet{jones03} studied a flux-limited sample of old galaxy
FT South and SMARTS epochs (~55298 MJD).The jet faded over the transition and our last REM observations (~55303 MJD). while the other component (probably the disc) stayed about the same.,"FT South and SMARTS epochs $\sim$ 55298 MJD).The jet faded over the transition and our last REM observations $\sim$ 55303 MJD), while the other component (probably the disc) stayed about the same."
 The H-band faded by a factor of 10 whereas the V-band faded by a factor >2.5., The $H$ -band faded by a factor of 10 whereas the $V$ -band faded by a factor $>$ 2.5.
 The jet contribution moved to lower frequencies: data taken on ~ 55298 MJD still have the jet dominating in the H-band in NIR. but no longer in the optical.," The jet contribution moved to lower frequencies: data taken on $\sim$ 55298 MJD still have the jet dominating in the $H$ -band in NIR, but no longer in the optical."
 The last SED ts bluer according to Fig. 7.., The last SED is bluer according to Fig. \ref{sed4}.
 This is exactly what one can expect when a source undergoes a LHS to HSS spectral transition: the colour changed and thermal processes started to dominate., This is exactly what one can expect when a source undergoes a LHS to HSS spectral transition: the colour changed and thermal processes started to dominate.
 During the rising LHS. Gandhi et al. (," During the rising LHS, Gandhi et al. ("
2011. A&AA. submitted) found that the jet spectrum was highly variable in the mid-IR.,"2011, A, submitted) found that the jet spectrum was highly variable in the mid-IR."
 The break between optically thick (self-absorbed) and optically thin synchrotron emission in the jet spectrum was found to vary between ~3.6 and 22 jun on timescales of minutes-hours., The break between optically thick (self-absorbed) and optically thin synchrotron emission in the jet spectrum was found to vary between $\sim 3.6$ and 22 $\mu m$ on timescales of minutes–hours.
 Here. in Fig. 7..," Here, in Fig. \ref{sed4},"
 we are witnessing this jet component fading over timescales of We produced broader-band SEDs in two distinct spectral states (hard and soft). composed of simultaneous radio (when available). (de-reddened) UV/optical/NIR and unabsorbed ray/soft y-ray data.," we are witnessing this jet component fading over timescales of We produced broader-band SEDs in two distinct spectral states (hard and soft), composed of simultaneous radio (when available), (de-reddened) UV/optical/NIR and unabsorbed X-ray/soft $\gamma$ -ray data."
 They are shown in Figures 8 and 9.., They are shown in Figures \ref{sed1bisfit} and \ref{sed2}.
 In Fig. 8..," In Fig. \ref{sed1bisfit},"
 we plot the data obtained during the first TToO (55259.9-S5261.1 MJD. Rev. 902)," we plot the data obtained during the first ToO (55259.9–55261.1 MJD, Rev. 902)."
 Fig., Fig.
 8 is remarkably similar to the Fig., \ref{sed1bisfit} is remarkably similar to the Fig.
 2 of ?.. focusing on GX 339-4 jet signatures during LHS.," 2 of \cite{corbfen02}, focusing on GX $-$ 4 jet signatures during LHS."
 One can interpret the results with simple power laws. but note that this is a very rough approach and not a physical model.," One can interpret the results with simple power laws, but note that this is a very rough approach and not a physical model."
 An extrapolation of the radio data up to the NIR/optical clearly does not fit the data., An extrapolation of the radio data up to the NIR/optical clearly does not fit the data.
 Excess NIR/optical emission is observed., Excess NIR/optical emission is observed.
 Possible sources of the residual emission are the disc. the irradiated face of the companion star (as the companion cannot contribute much. see ??)) and the jets.," Possible sources of the residual emission are the disc, the irradiated face of the companion star (as the companion cannot contribute much, see \citealt{Shab01,Hynes04}) ) and the jets."
 Similarly. several power law components with distinct slopes are needed to fit our SEDs.," Similarly, several power law components with distinct slopes are needed to fit our SEDs."
 For example. using a simple. double power law fit to the data in Fig.," For example, using a simple, double power law fit to the data in Fig."
 7 gives indices of —0.9+0.4 for the NIR range.," \ref{sed4} gives indices of $-0.9
\pm 0.4$ for the NIR range."
 We fixed the optical component (where the dise 1s assumed to dominate) to 1.7 which is the value we found fitting a single power law to our SED around 55303 MJD. where the jet contribution was absent or still negligible (thus assuming that. at that epoch. we only have contribution from the disc).," We fixed the optical component (where the disc is assumed to dominate) to 1.7 which is the value we found fitting a single power law to our SED around 55303 MJD, where the jet contribution was absent or still negligible (thus assuming that, at that epoch, we only have contribution from the disc)."
 For this SED in the soft state (Fig. 9)).," For this SED in the soft state (Fig. \ref{sed2}) ),"
 we observed important spectral changes both in the dise and hot medium components and found that a simple single power law model was enough to fit the UV/optical/NIR data (Fig. 7)).," we observed important spectral changes both in the disc and hot medium components and found that a simple single power law model was enough to fit the UV/optical/NIR data (Fig. \ref{sed4}) ),"
 with an index of 1.7+0.2., with an index of $1.7 \pm 0.2$.
 This is indicative of the fading jet component in339-4.. as previously discussed.," This is indicative of the fading jet component in, as previously discussed."
 The NIR/optical ESO/ISAAC spectroscopy taken at that time will be commented on in a forthcoming paper (Rahoui et al..," The NIR/optical ESO/ISAAC spectroscopy taken at that time will be commented on in a forthcoming paper (Rahoui et al.,"
 After the LHS. the radio spectral index became typical of optically thin synchrotron radiation. probably as a result of freely expanding plasma blobs previously ejected (see.e.g..?)..," After the LHS, the radio spectral index became typical of optically thin synchrotron radiation, probably as a result of freely expanding plasma blobs previously ejected \citep[see, e.g.,][]{Fender:2004}."
 This suggests that multiple ejection events took place during the outburst of aand then interacted with the interstellar medium ?.., This suggests that multiple ejection events took place during the outburst of and then interacted with the interstellar medium \citet{corb10b}.
 This could potentially result in re-acceleration of particles up to very high energies., This could potentially result in re-acceleration of particles up to very high energies.
 In general. the shape of our SEDs during the LHS are similar to the ones of the transient LMXB XTE J1118+480 (?2?):: in its 2000 outburst. the SED from radio to X-rays has been explained as a combination of synchrotron radiation from a jet and a truncated optically thick dise. whereas models assuming advection dominated accretion flows alone underestimated the optical and IR fluxes (?.. and references therein).," In general, the shape of our SEDs during the LHS are similar to the ones of the transient LMXB XTE $+$ 480 \citep{Chaty:2003,Zurita:2006}: in its 2000 outburst, the SED from radio to X-rays has been explained as a combination of synchrotron radiation from a jet and a truncated optically thick disc, whereas models assuming advection dominated accretion flows alone underestimated the optical and IR fluxes \citealt{Zurita:2006}, and references therein)."
 In 2005. discrepancies observed between the optical and IR SEDs of XTE J1118+480 suggested that the IR was dominated possibly by a jet whereas the optical was dominated by dise emission.," In 2005, discrepancies observed between the optical and IR SEDs of XTE $+$ 480 suggested that the IR was dominated possibly by a jet whereas the optical was dominated by disc emission."
 Power law fits to optical SEDs have also been performed for other BH XTs in outburst (?).., Power law fits to optical SEDs have also been performed for other BH XTs in outburst \citep{Hynes:2005}.
" All optical SEDs exhibit quasi power law spectra. with a ranging from 0.5-1.5. all steeper than that expected for a viscously heated. multi-temperature dise: S,ev!5 (which differs from our value taken closer to the peak)."," All optical SEDs exhibit quasi power law spectra, with $\alpha$ ranging from 0.5–1.5, all steeper than that expected for a viscously heated, multi-temperature disc: $S_\nu\propto\nu^{1/3}$ (which differs from our value taken closer to the peak)."
" Also. the authors found for these BHs that the UV/optical/X-ray data - when detected could be fitted with a simple black-body model of an aceretio disc heated by internal viscosity and. X-ray irradiation, but the inner radius could not be well constrained."," Also, the authors found for these BHs that the UV/optical/X-ray data - when detected - could be fitted with a simple black-body model of an accretion disc heated by internal viscosity and X-ray irradiation, but the inner radius could not be well constrained."
 They concludec that the flat-spectrum synchrotron emission may be important in the IR and optical in this source., They concluded that the flat-spectrum synchrotron emission may be important in the IR and optical in this source.
 However. they did not exclude the alternative explanation that the IR excess coulc ccome from the cool outer disc.," However, they did not exclude the alternative explanation that the IR excess could come from the cool outer disc."
 More recently. ?. even showed in the colour-magnitude diagram of XTE J1550—564 (Fig.," More recently, \citet{Russ2011} even showed in the colour-magnitude diagram of XTE $-$ 564 (Fig."
 | of their paper) that a can change dramatically over state transitions., 1 of their paper) that $\alpha$ can change dramatically over state transitions.
 In our observations of ppresented here. the contribution of the radio to synchrotron emission up to the NIR/optical is important. and we saw it fading over the transition (see Sect. 4.1).," In our observations of presented here, the contribution of the radio to synchrotron emission up to the NIR/optical is important, and we saw it fading over the transition (see Sect. \ref{inter}) )."
 However. another component. for example from the cooling dise and/or the irradiated companion. and/or an irradiated dise. (22) might be necessary to account for the NIR/optical excess observed.," However, another component, for example from the cooling disc and/or the irradiated companion, and/or an irradiated disc \citep{vanpar1994,Hynes:2005} might be necessary to account for the NIR/optical excess observed."
friction process (Navakshin.2006)..,friction process \citep{Nayakshin06a}.
 Phe dise swells as time progresses. as can be seen in the two edge-on views of the disc from simulation 52 shown in Fig. 5..," The disc swells as time progresses, as can be seen in the two edge-on views of the disc from simulation S2 shown in Fig. \ref{fig:edgeon} ."
 The gaseous disc can thus become thicker than it would have been on its own., The gaseous disc can thus become thicker than it would have been on its own.
 Llowever. depending on conditions (i.e. cooling parameter 3. initial total clise mass. mass spectrum of stars). the two clises can be coupled or decoupled.," However, depending on conditions (i.e. cooling parameter $\beta$, initial total disc mass, mass spectrum of stars), the two discs can be coupled or decoupled."
 In the latter case the stellar disc has a larger gcometrical thickness than the gaseous clisc., In the latter case the stellar disc has a larger geometrical thickness than the gaseous disc.
 In such a case the rate at which stars heat the disc is not trivially calculated., In such a case the rate at which stars heat the disc is not trivially calculated.
 This is especially so if stellar radiation. winds and supernovae are taken into account. as stellar energy release above the disc is much less effective in disc heating than it is inside the disc.," This is especially so if stellar radiation, winds and supernovae are taken into account, as stellar energy release above the disc is much less effective in disc heating than it is inside the disc."
 We expect that. following an increase in disc elfective temperature and geometrical thickness with time. the vertically averaged: density of the cise must. crop.," We expect that, following an increase in disc effective temperature and geometrical thickness with time, the vertically averaged density of the disc must drop."
 The disc may then evolve into a non sell-eravitating state as ( increases above unity (see equation 1))., The disc may then evolve into a non self-gravitating state as $Q$ increases above unity (see equation \ref{q}) ).
 Further disc fragmentation should. cease., Further disc fragmentation should cease.
 This cllect can be noticed in the right hand panel of Fig. 2.., This effect can be noticed in the right hand panel of Fig. \ref{fig:fig1}.
 Near the inner edge of the disc. there are stars but no high density clunips or filaments implving that the cise is no longer fragmenting in that region.," Near the inner edge of the disc, there are stars but no high density clumps or filaments implying that the disc is no longer fragmenting in that region."
 Figure 6. shows the racdius-integrated fragmentation rate of the disc in the simulations with 3=0.3 (run SI. thick solid curve) 3—2 (82. dashed) and 3= (83. thin solid).," Figure \ref{fig:fragmrate} shows the radius-integrated fragmentation rate of the disc in the simulations with $\beta=0.3$ (run S1, thick solid curve), $\beta=2$ (S2, dashed) and $\beta=3$ (S3, thin solid)."
 The fragmentation rate is defined as the total mass of first cores createcl per unit time., The fragmentation rate is defined as the total mass of first cores created per unit time.
 Not surprisingly. the shorter the cooling time (smaller 3). the more rapid is disc fragmentation.," Not surprisingly, the shorter the cooling time (smaller $\beta$ ), the more rapid is disc fragmentation."
 Lhe more vigorous disc fragmentation explains why there are more stars and dense bound gas clumps in Fig., The more vigorous disc fragmentation explains why there are more stars and dense bound gas clumps in Fig.
 2 than in Fig., \ref{fig:fig1} than in Fig.
 3. at the same time (/—15)., \ref{fig:fig2} at the same time $t=75$ ).
 One can also see that the expectation of a decrease of the fragmentation rate with time is borne out., One can also see that the expectation of a decrease of the fragmentation rate with time is borne out.
 At the peaks of the respective curves. most of the cise mass is still in gascous form.," At the peaks of the respective curves, most of the disc mass is still in gaseous form."
 Therefore. the decline in the fragmentation rate with time is indeed. mostly a consequence of a change in the disc state (higher Q-parameter) rather than due to the disc running out of gas.," Therefore, the decline in the fragmentation rate with time is indeed mostly a consequence of a change in the disc state (higher $Q$ -parameter) rather than due to the disc running out of gas."
 ]t is interesting to compare the distribution functions of stellar masses from our simulations., It is interesting to compare the distribution functions of stellar masses from our simulations.
" Phe end state of our simulations. i.c. when the majority of gas is turned. into stars. corresponds to the ""initial mass function"" (IME) of a stellar population."," The end state of our simulations, i.e., when the majority of gas is turned into stars, corresponds to the “initial mass function” (IMF) of a stellar population."
 Pherefore we use this name to refer to our mass distributions., Therefore we use this name to refer to our mass distributions.
 Figure 7. shows the AIF of stars ormed in the three simulations S183., Figure \ref{fig:imf23} shows the IMF of stars formed in the three simulations S1–S3.
 Table 1 lists the irst two moments of the distribution. ic. the average mass. Ali}. and (y.," Table 1 lists the first two moments of the distribution, i.e., the average mass, $\left<M_*\right>$, and $\left<M_*^2\right>^{1/2}$."
 τμ ds clear from. both the ligure and he table. that the longer the cooling time. the more top-weavy (or. equivalently. bottom-light) is the resulting ME.," It is clear from both the figure and the table, that the longer the cooling time, the more top-heavy (or, equivalently, bottom-light) is the resulting IMF."
 This outcome is not surprising., This outcome is not surprising.
 As we saw in Section ??.. ragmentation is fastest for the smallest values of the cooling xwameter 2.," As we saw in Section \ref{sec:quenching}, fragmentation is fastest for the smallest values of the cooling parameter $\beta$."
 Further. fragmentation stalls in our models when the stars heat up the disc above Q=1 (see Fig. 6)).," Further, fragmentation stalls in our models when the stars heat up the disc above $Q=1$ (see Fig. \ref{fig:fragmrate}) )."
 At this point. disc [fragmentation stops but accretion onto stars continues.," At this point, disc fragmentation stops but accretion onto stars continues."
 The average mass of a star reached by the time the gas supply is exhausted: is roughly inversely proportional to the number of stars at the time the disc fragmentation stalls., The average mass of a star reached by the time the gas supply is exhausted is roughly inversely proportional to the number of stars at the time the disc fragmentation stalls.
 As we find many more stars in the tests, As we find many more stars in the tests
"in the outer regions of the disce. while the nuclear regions show strongest values for D,, 4000.","in the outer regions of the disc, while the nuclear regions show strongest values for $_n$ 4000."
 Their findings are consistent with the idea that the star formation in passive late-type galaxies ceased a few Gyr ago., Their findings are consistent with the idea that the star formation in passive late-type galaxies ceased a few Gyr ago.
 Hence. the blue passive galaxies can be considered to be the progenitors of their red counterparts. in which star formation has been recently shut-off: these galaxies will eventually aequire redder colours and. eartype morphologies. over a period that depends upon their environment (Gotoetal.2003c).," Hence, the blue passive galaxies can be considered to be the progenitors of their red counterparts, in which star formation has been recently shut-off: these galaxies will eventually acquire redder colours and early-type morphologies, over a period that depends upon their environment \citep{goto03c}."
. have recently identified. characterized and studied the evolution of galaxies which. though morphologically classified as E/SOs. lie on the blue sequence in the colour-stellar mass space. and suggest that these systems almost always have a bluer outer disk. and that they are a population in transition. potentially evolving on to the red sequence. or have (re)formed a disk and are about to fall back in the class of late-types.," have recently identified, characterized and studied the evolution of galaxies which, though morphologically classified as E/S0s, lie on the blue sequence in the colour-stellar mass space, and suggest that these systems almost always have a bluer outer disk, and that they are a population in transition, potentially evolving on to the red sequence, or have (re)formed a disk and are about to fall back in the class of late-types."
 Even more interesting are red galaxies with signs of active star formation in their fibre spectra., Even more interesting are red galaxies with signs of active star formation in their fibre spectra.
 We find that of the 539 galaxies in this class have significant emission lines on the basis of which their SSFR has been evaluated., We find that of the 539 galaxies in this class have significant emission lines on the basis of which their SSFR has been evaluated.
 These red star-forming galaxies are made up of at least two kinds of galaxies (chosen from simple limits in colour and SSER)., These red star-forming galaxies are made up of at least two kinds of galaxies (chosen from simple limits in colour and SSFR).
 One kind appears to be very similar in photometric properties to the red sequence galaxies but with spectroscopically derived properties consistent with their late-type counterparts. while the other class has both photometricand spectroscopic properties similar to the red sequence galaxies. yet has a high SSFR.," One kind appears to be very similar in photometric properties to the red sequence galaxies but with spectroscopically derived properties consistent with their late-type counterparts, while the other class has both photometric spectroscopic properties similar to the red sequence galaxies, yet has a high SSFR."
 The origin of the latter class is not clear., The origin of the latter class is not clear.
 It is likely that these galaxies are currently experiencing a starburst. which has started very recently (0.5 Gyr).," It is likely that these galaxies are currently experiencing a starburst, which has started very recently $\lesssim$ 0.5 Gyr)."
" Stellar populations resulting from such a starburst would be detected in spectral indices such as H,, EW (leading to high values of SFR/M*). but will not dominate the D,, 4000 and the Hs EW for another —0.5-1 Gyr."," Stellar populations resulting from such a starburst would be detected in spectral indices such as $_\alpha$ EW (leading to high values of $^*$ ), but will not dominate the $_n$ 4000 and the $_\delta$ EW for another $\sim$ 0.5–1 Gyr."
 The former sub- is most likely to be bright cluster galaxies (BCGs). with nuclear star-formation linked with the AGN (seeGoto2006:Bild-felletal.2008:Reichard 2009).," The former sub-population is most likely to be bright cluster galaxies (BCGs), with nuclear star-formation linked with the AGN \citep[see][]{goto06,bildfell,reichard}."
. In the following we turn our attention to this sub-population., In the following we turn our attention to this sub-population.
 The hypothesis that AGN and star formation activity in the core of a galaxy may be linked is supported by the fact that ~30% of the red star-forming galaxies with emission lines are classitied as AG on the BPT diagram., The hypothesis that AGN and star formation activity in the core of a galaxy may be linked is supported by the fact that $\sim$ of the red star-forming galaxies with emission lines are classified as AGN on the BPT diagram.
 The fact that of the emission-line red sequence galaxies (defined according to Fig. 23) , The fact that of the emission-line red sequence galaxies (defined according to Fig. \ref{ssf-gr}) )
are also classitiec as AGN does not contradict this scenario because either (1) the rec sequence AGN galaxies in general have high stellar mass. and so can have low SSFR even though they have significantly high SFR. or (i) they do not have enough cold gas to form stars. because of the presence of an active nuclei.," are also classified as AGN does not contradict this scenario because either (i) the red sequence AGN galaxies in general have high stellar mass, and so can have low SSFR even though they have significantly high SFR, or (ii) they do not have enough cold gas to form stars, because of the presence of an active nuclei."
 A significant fraction of giant ellipticals show evidence of ongoing star formation or signs of recent («2 Gyr) star formation in their cores. particularly in environments where the likelihood of recent mergers is high (e.g..MeDermidetal.2006:Nolaneal.2006:Nolan.Raychaudhury.&Kabán 2007).," A significant fraction of giant ellipticals show evidence of ongoing star formation or signs of recent $<$ 2 Gyr) star formation in their cores, particularly in environments where the likelihood of recent mergers is high \citep[e.g.,][]{mcd06,nolan06,nolan07}. ."
. For a sample of BCGs in 48 X-ray luminous clusters. Bildfelletal.(2008). find tha of the BCGs have colour profiles (g+) that turn bluer toward the centre.," For a sample of BCGs in 48 X-ray luminous clusters, \citet{bildfell} find that of the BCGs have colour profiles $(g\! -\! r)$ that turn bluer toward the centre."
 Our radial distribution of galaxies in this category (Fig. 33. ," Our radial distribution of galaxies in this category (Fig. \ref{dist}) ),"
compared to the distribution of blue star-forming galaxies. is consistent with this (see refsec:radial}).," compared to the distribution of blue star-forming galaxies, is consistent with this (see \\ref{sec:radial}) )."
" Elsewhere. Gallazzietal.(2009) analyse an extensive dataset. covering UV to IR SEDs for galaxies in the Abell 901/902 cluster pair, to show that ~40¢° of the star-forming galaxies residing in intermediate tohigh density environments have optically red colours."," Elsewhere, \citet{gallazzi09} analyse an extensive dataset, covering UV to IR SEDs for galaxies in the Abell 901/902 cluster pair, to show that $\sim$ of the star-forming galaxies residing in intermediate tohigh density environments have optically red colours."
 They show that these galaxies are not starbursts. and," They show that these galaxies are not starbursts, and"
Several studies (2???) as well as the cust maps of ? have shown the reddening of U Sco to beinthe range k(BY)= 0.36.,"Several studies \citep{barlow,amores,burstein} as well as the dust maps of \cite{schlegel} have shown the reddening of U Sco to be inthe range $E(B-V) = 0.09 - 0.36$ ."
 The ratios of line Ηχος were estimated. and compared to the theoretical values derived from ?.., The ratios of line fluxes were estimated and compared to the theoretical values derived from \cite{hummer}.
 We use he optical spectra taken on day S.S1 at LT and day. 9.43 at CTIO (Figs | and 2). and the Ht spectrum taken on clay 9.43 at NT (Fig 4) to estimate the reddening. as by this ime the [ux ratios are converging on case B values.," We use the optical spectra taken on day 8.81 at LT and day 9.43 at CTIO (Figs 1 and 2), and the IR spectrum taken on day 9.43 at NTT (Fig 4) to estimate the reddening, as by this time the flux ratios are converging on case B values."
 We use the extinction law ancl assumption of /?=3.1 of ?.., We use the extinction law and assumption of $R=3.1$ of \cite{howarth}.
 Using the ratios IL2/Dac. 113/Da. Hz /Pac. ancl Hz Da he reddening was found to be in the range L(V)=LO0.29 with a mean of (D.V)=0.14250.12. consistent with the previous studies.," Using the ratios $\beta$ $\epsilon$, $\beta$ $\beta$, $\gamma$ $\epsilon$, and $\gamma$ $\beta$ the reddening was found to be in the range $E(B-V) = 0.0 - 0.29$ with a mean of $E(B-V) = 0.14\pm0.12$, consistent with the previous studies."
 Although U Sco is at a distance of kkpe. low reddening is consistent with both the line of sight leaving the plane of the galaxy and the system being ata height of z= 4.5kkpe above the galactic plane (?)..," Although U Sco is at a distance of kpc, low reddening is consistent with both the line of sight leaving the plane of the galaxy and the system being at a height of $z=4.5$ kpc above the galactic plane \citep{schaeferlong}."
 bor the spectra used in this work we adopt. £06132202 as à good compromise between our own value and previous estates., For the spectra used in this work we adopt $E(B-V) = 0.2$ as a good compromise between our own value and previous estimates.
 The helium abundance of U Sco was caleulatecl usingLL.Let. anc recombination lines.," The helium abundance of U Sco was calculated using, and recombination lines."
 We use the line Duxes on days S.81 (LE). 09.93-11.93. (ΑΛΛΟ). and 9.433 (NPL) to estimate the helium abundance. as by this time the line ratios are converging on case D values.," We use the line fluxes on days 8.81 (LT), 9.93-11.93 (SAAO), and 9.43 (NTT) to estimate the helium abundance, as by this time the line ratios are converging on case B values."
 The laree Ha/1L7 ratio throughout the time coverage we have available shows that Lla should not be used in our abundance analysis due to optical depth ellects., The large $\alpha$ $\beta$ ratio throughout the time coverage we have available shows that $\alpha$ should not be used in our abundance analysis due to optical depth effects.
 Abundance analyses of helium are complicated by the moetastability of the lowest triplet level of Lei. 278. which," Abundance analyses of helium are complicated by the metastability of the lowest triplet level of , $^3$ S, which"
The VLA total intensity iso-contours of WNB 173446407. WNB 182946911. and WNB 18514570 at 1.4 and GGHz are respectively shown in top left and right panels of 11 to 3.,"The VLA total intensity iso-contours of WNB 1734+6407, WNB 1829+6911, and WNB 1851+570 at 1.4 and GHz are respectively shown in top left and right panels of 1 to 3."
 The GGHz contours are overlaid to the optical image of the red Digital Sky Survey (DSS2). while the GGHz contours are overlaid to the spectral index image between the two radio frequencies.," The GHz contours are overlaid to the optical image of the red Digital Sky Survey (DSS2), while the GHz contours are overlaid to the spectral index image between the two radio frequencies."
 We recover from the VLA archive all the useful data that are available for the radio sources B2 012033 and B2 1610229., We recover from the VLA archive all the useful data that are available for the radio sources B2 0120+33 and B2 1610+29.
 A summary of relevant parameters of the images produced is listed in 33., A summary of relevant parameters of the images produced is listed in 3.
 The source B2 0120433 has a useful 322MMHz B-array observation in the VLA archive. (project. ACOSAS)., The source B2 0120+33 has a useful MHz B-array observation in the VLA archive (project AC0845).
 The observation was acquired in line mode with 15 spectral channels in a bandwidth of 6 MHz per each of the two IPs., The observation was acquired in line mode with 15 spectral channels in a bandwidth of 6 MHz per each of the two IFs.
 The source 3C 48 was used as amplitude. phase. and bandpass calibrator.," The source 3C 48 was used as amplitude, phase, and bandpass calibrator."
 In the final imaging the data were averaged to 6 channels and mapped using a wide-field imaging technique. which corrects for distortions in the image caused by the non-coplanarity of the VLA over a wide field of view.," In the final imaging the data were averaged to 6 channels and mapped using a wide-field imaging technique, which corrects for distortions in the image caused by the non-coplanarity of the VLA over a wide field of view."
" A set of overlapping fields was used to cover an area of about 2.5°x2.5"" around the radio source.", A set of overlapping fields was used to cover an area of about $2.5\degr\times2.5\degr$ around the radio source.
 All these fields were included in CLEAN and used for several loops of phase self-calibration., All these fields were included in CLEAN and used for several loops of phase self-calibration.
 The central frequency of the final images ts 324 MHz., The central frequency of the final images is 324 MHz.
 At GGHz we produced a B+C image at a resolution of about10’., At GHz we produced a B+C image at a resolution of about.
 This image is shown in top-left panel of 44., This image is shown in top-left panel of 4.
 The resolution of the C array GGHz image alone well matches that of the MMHz image and the two have been used to derive the spectral index image shown in top-right panel of Fig.44., The resolution of the C array GHz image alone well matches that of the MHz image and the two have been used to derive the spectral index image shown in top-right panel of 4.
 It is worthwhile to mention that flux densities of both the MMHz and the GGHz images are perfectly consistent with that of the WENSS and the NVSS., It is worthwhile to mention that flux densities of both the MHz and the GHz images are perfectly consistent with that of the WENSS and the NVSS.
 At GGHz we were able to combine the B and C array., At GHz we were able to combine the B and C array.
 Only the core and the brightest part of the lobes are detected in this image that can be used however to derive an integrated flux density for the radio source., Only the core and the brightest part of the lobes are detected in this image that can be used however to derive an integrated flux density for the radio source.
 At GGHz we recover a D array image in which only the core of the radio galaxy is visible., At GHz we recover a D array image in which only the core of the radio galaxy is visible.
 The source B2 1610+29 has an angular size of about two aremin and has a very relaxed morphology which is completely resolved out by the A configuration of the VLA at 1.4 GHz., The source B2 1610+29 has an angular size of about two arcmin and has a very relaxed morphology which is completely resolved out by the A configuration of the VLA at 1.4 GHz.
 Thus. we opted to combine just the B and C array at this frequency. obtaining a final resolution of 55xS75 and a rms noise level of 70 jJy/beam.," Thus, we opted to combine just the B and C array at this frequency, obtaining a final resolution of $5\farcs5 \times 5\farcs5$ and a rms noise level of 70 $\mu$ Jy/beam."
 This is the highest resolution image we present for this source and it is shown as 150-contours in top-left panel of 55., This is the highest resolution image we present for this source and it is shown as iso-contours in top-left panel of 5.
 The flux density of our B+C image is consistent. within the errors. with the NVSS flux.," The flux density of our B+C image is consistent, within the errors, with the NVSS flux."
 Therefore. we are confident to have recovered most of flux density from the extended source structure.," Therefore, we are confident to have recovered most of flux density from the extended source structure."
 We also produced images at 1.4 and GGHz in C and D configuration. respectively.," We also produced images at 1.4 and GHz in C and D configuration, respectively."
" These images. which have been restored with the same resolution of 17""x17"" and corrected from the primary beam attenuation. are characterized by roughly the same uv-coverage and have been used to produce the spectral index image shown in the top-right panel of 55."," These images, which have been restored with the same resolution of $17\arcsec \times 17\arcsec$ and corrected from the primary beam attenuation, are characterized by roughly the same uv-coverage and have been used to produce the spectral index image shown in the top-right panel of 5."
 Finally. for the purposes of the analysis of the source integrated spectrum. we reduced separately the two IFs of the L band observation obtaining a measure of the source flux densities at 1452 and MMHz.," Finally, for the purposes of the analysis of the source integrated spectrum, we reduced separately the two IFs of the L band observation obtaining a measure of the source flux densities at 1452 and MHz."
 The VLA integrated flux density and spectral index for the individual source components are listed in 44., The VLA integrated flux density and spectral index for the individual source components are listed in 4.
 Characterized by two relaxed lobes lacking hot-spots. the radio morphology of WNB1I734+6407 resembles that of B2 0924-30. which is considered the prototype of fossil radio galaxies (Cordey 1987. but see also Jamrozy et al.," Characterized by two relaxed lobes lacking hot-spots, the radio morphology of WNB1734+6407 resembles that of B2 0924+30, which is considered the prototype of fossil radio galaxies (Cordey 1987, but see also Jamrozy et al."
 2004 for more recent data on this source)., 2004 for more recent data on this source).
 The radio spectrum of the fossil lobes of WNB 173446407 is so steep that their surface brightness at 6 cm is below the sensitivity level of our observations., The radio spectrum of the fossil lobes of WNB 1734+6407 is so steep that their surface brightness at 6 cm is below the sensitivity level of our observations.
 Thus. we can only place lower limits on the spectral index which result in «>2.7 and α>2.3 for the northern and southern lobe. respectively.," Thus, we can only place lower limits on the spectral index which result in $\alpha>2.7$ and $\alpha>2.3$ for the northern and southern lobe, respectively."
 We also observe a slightly extended component coincident with the galaxy center., We also observe a slightly extended component coincident with the galaxy center.
 Also this feature has a quite steep spectral index. α=1.3. and its nature remains unclear.," Also this feature has a quite steep spectral index, $\alpha=1.3$, and its nature remains unclear."
 In the VLA A-array image (not shown) this feature is elongated in N-S direction with an extent of about 7 kpe x 14 kpe. resembling a core-jet morphology.," In the VLA A-array image (not shown) this feature is elongated in N-S direction with an extent of about 7 kpc $\times$ 14 kpc, resembling a core-jet morphology."
 The radio appearance of WNB1829+691] is virtually identical to that of 3C 338. a nearby radio source associated with the central dominant galaxy in the cooling flow cluster Abell 2199 (Jones Preston 2001).," The radio appearance of WNB1829+6911 is virtually identical to that of 3C 338, a nearby radio source associated with the central dominant galaxy in the cooling flow cluster Abell 2199 (Jones Preston 2001)."
 In both sources we observe the presence of fossil plasma remaining from a previous activity, In both sources we observe the presence of fossil plasma remaining from a previous activity
We notice in refhurstfiux that at the lowest flux levels 260 shows an increase in burst rate with increasing persisteut flux.,We notice in \\ref{burstflux} that at the lowest flux levels $-$ 260 shows an increase in burst rate with increasing persistent flux.
 When 26 reaches 1.7«LO? ergenn2s|. the burst rate drops by a factor of 5 iux at higher x levels the burst rate slowly decreases.," When $-$ 26 reaches $1.7\times10^{-9}$ $\ergcms$, the burst rate drops by a factor of 5 and at higher flux levels the burst rate slowly decreases."
 129. 0 and 21 all show only an increasing bus rate. while 536 is the oulv source that shows a decreasing burst rate.," $-$ 429, $-$ 0 and $-$ 24 all show only an increasing burst rate, while $-$ 536 is the only source that shows a decreasing burst rate."
 Lh CX3311 and 30 show a drop by a factor of &5 in burst rate over a πα rauge of persistent flux.," $-$ 44, $+$ 1 and $-$ 30 show a drop by a factor of $\simeq$ 5 in burst rate over a small range of persistent flux."
 291 is the only source for which no clear trend is visible. aud the burst rate stavs constant over the total observed. fiux range.," $-$ 294 is the only source for which no clear trend is visible, and the burst rate stays constant over the total observed flux range."
 This source traces out the lowest fluxes within our salple., This source traces out the lowest fluxes within our sample.
 Several sources are known to show quasi-periodic burst recurrence times duriug certain periods., Several sources are known to show quasi-periodic burst recurrence times during certain periods.
 The best example is 21. in 1996-1997 it exhibited a burst every ~6 hours. aud the burst wait times were constant within a few nünutes for long periods of tine (Uboertini οἳ al.," The best example is $-$ 24, in 1996-1997 it exhibited a burst every $\simeq$ 6 hours, and the burst wait times were constant within a few minutes for long periods of time (Ubertini et al."
 1999: Coechi et al., 1999; Cocchi et al.
 20015)., 2001b).
 Iu the left panel of refwaiting we plotted the wait time as a function of the osersisteut flux for 21., In the left panel of \\ref{waiting} we plotted the wait time as a function of the persistent flux for $-$ 24.
 Most wait times appear ο follow a straight line at the bottom of the figure., Most wait times appear to follow a straight line at the bottom of the figure.
 A second gear trend cau clearly be distinguished above this iue (ancl evel wo more above that)., A second linear trend can clearly be distinguished above this line (and even two more above that).
 Civen the fact BeppoSAN has a 96-minute low-carth orbit. it is xobable hat bursts are missed during data gaps and that iultiples of the burst wait times are observed.," Given the fact that BeppoSAX has a 96-minute low-earth orbit, it is probable that bursts are missed during data gaps and that multiples of the burst wait times are observed."
 We checked bursts with long wait ines within one observation aud find that or all of them the previous burst may have occurred during an earth occultation or South Atlantic Anomaly oiswage., We checked bursts with long wait times within one observation and find that for all of them the previous burst may have occurred during an earth occultation or South Atlantic Anomaly passage.
 From refwaiting we see that the wait time between the bursts decreases linearly with increasing persistent flux., From \\ref{waiting} we see that the wait time between the bursts decreases linearly with increasing persistent flux.
 We formed a least-squares fit on the bursts where the, We performed a least-squares fit on the bursts where the
CB2 simulations grow to higher masses than in the SB3 simulation.,CB2 simulations grow to higher masses than in the SB3 simulation.
" As a consequence, the scale height of the dust is lower in these simulations, as heavier particles are more difficult to stir up by turbulence."," As a consequence, the scale height of the dust is lower in these simulations, as heavier particles are more difficult to stir up by turbulence."
" We illustrate the mass distribution at t=1, 10, 100, 316, 10?, 104 yr in Fig."," We illustrate the mass distribution at $t=1$, 10, 100, 316, $10^3$, $10^4$ yr in Fig."
 7 for the CB1 model., \ref{fig:mass_br_1d-4_or} for the CB1 model.
 The particle evolution has two phases in these simulations., The particle evolution has two phases in these simulations.
" The first 1000 yr are identical for the SB3 and CB1/CB2 simulations, respectively (see also the first five snapshots of Figs."," The first 1000 yr are identical for the SB3 and CB1/CB2 simulations, respectively (see also the first five snapshots of Figs."
 6 and 7))., \ref{fig:mass_simpl_1d-4_or} and \ref{fig:mass_br_1d-4_or}) ).
" In this phase, particles start sedimenting, and the rain-out particles reach the midplane."," In this phase, particles start sedimenting, and the rain-out particles reach the midplane."
 The dust evolution in the SB3 model halts at this point as only bouncing collisions happen., The dust evolution in the SB3 model halts at this point as only bouncing collisions happen.
" However, during the second phase of the CB1 and CB2 simulations, particles can grow to higher masses because there are areas in the parameter space that is favorable for growth somewhat beyond the bouncing barrier (see Paper I and Paper II for a detailed explanation)."," However, during the second phase of the CB1 and CB2 simulations, particles can grow to higher masses because there are areas in the parameter space that is favorable for growth somewhat beyond the bouncing barrier (see Paper I and Paper II for a detailed explanation)."
" Due to this growth, particles reach 107? g in mass for the CB1 and CB2 simulations."," Due to this growth, particles reach $10^{-2}$ g in mass for the CB1 and CB2 simulations."
 The time evolution of the enlargement parameter is quite different in the two cases., The time evolution of the enlargement parameter is quite different in the two cases.
 Fluffy aggregates are produced with enlargement parameter V=10° when the Okuzumi porosity model is used., Fluffy aggregates are produced with enlargement parameter $\Psi = 10^3$ when the Okuzumi porosity model is used.
" However, these aggregates are strongly compacted by bouncing and their final enlargement parameter is between 2 and 20."," However, these aggregates are strongly compacted by bouncing and their final enlargement parameter is between 2 and 20."
" When the Ormel porosity model is considered, the enlargement parameter never exceeds 20 and by the end of the simulation the enlargement parameter is between 2 and 10."," When the Ormel porosity model is considered, the enlargement parameter never exceeds 20 and by the end of the simulation the enlargement parameter is between 2 and 10."
" As we see, the enlargement parameter evolved through very different ways, but the final enlargement parameters are not so much different for the CB1 and CB2 simulations."," As we see, the enlargement parameter evolved through very different ways, but the final enlargement parameters are not so much different for the CB1 and CB2 simulations."
 Figure 8 illustrates the collision history of the CB1 simulation., Figure \ref{fig:coll_hist_br_1d-4} illustrates the collision history of the CB1 simulation.
 If we compare this figure to Figs., If we compare this figure to Figs.
" 5, 7 and 10 of Paper II, we see that the features are more smeared out in Fig."," 5, 7 and 10 of Paper II, we see that the features are more smeared out in Fig."
 8 than in the other figures., \ref{fig:coll_hist_br_1d-4} than in the other figures.
 As we simulate, As we simulate
Integrating the second of (6)) gives A). where ut=5524 contains the initial conditions. and with F(\) given by (A8)). with ay=0assumed in (AI0)).,"Integrating the second of \ref{LAE3}) ) gives ), where $u'_0=\gamma'_0\beta'_0$ contains the initial conditions, and with $F(\chi)$ given by \ref{LAE10a}) ), with $a_0=0$assumed in \ref{LAE6}) )."
 The solution (2.3)) implies a Lorentz [actor ου and a velocity Gin unitsof ο) ΩΣ)," The solution \ref{LAE4}) ) implies a Lorentz factor , and a velocity (in unitsof $c$ ) ."
"explain the dashed curve, with the difference that the neutron Fermi momentum is the highest in this case.","explain the dashed curve, with the difference that the neutron Fermi momentum is the highest in this case."
" These effects are of course especially pronounced when kj is much larger than k7, see Fig."," These effects are of course especially pronounced when $k_F^n$ is much larger than $k_F^p$, see Fig."
 4., 4.
" In summary, strong variations with energy of the proton and neutron mean free path can be seen, as well as large differences between the two, in a rather narrow region around the Fermi “thresholds” for pp and nn scatterings."," In summary, strong variations with energy of the proton and neutron mean free path can be seen, as well as large differences between the two, in a rather narrow region around the Fermi “thresholds"" for $pp$ and $nn$ scatterings."
" As energy increases, however, the mean free path becomes essentially insensitive to isospin asymmetry."," As energy increases, however, the mean free path becomes essentially insensitive to isospin asymmetry."
 This is in agreement with the conclusions of Ref. [3].., This is in agreement with the conclusions of Ref. \cite{Jiang07}.
" Finally, we show in Table 2 some of the in-medium cross sections used for the present calculations of the mean free path in asymmetric matter."," Finally, we show in Table 2 some of the in-medium cross sections used for the present calculations of the mean free path in asymmetric matter."
 The pp and nn cross sections are quite similar to each other except in the low energy region where one may be sizable while the other is still suppressed., The $pp$ and $nn$ cross sections are quite similar to each other except in the low energy region where one may be sizable while the other is still suppressed.
 We presented predictions of the mean free path for protons and neutrons in isospin symmetric or asymmetric matter based on microscpic predictions of in-medium cross sections., We presented predictions of the mean free path for protons and neutrons in isospin symmetric or asymmetric matter based on microscpic predictions of in-medium cross sections.
 The mean free path in exotic matter is a fundamentally important quantity which finds applications in diverse areas including radiobiology., The mean free path in exotic matter is a fundamentally important quantity which finds applications in diverse areas including radiobiology.
" As it appears reasonable, very low-energy protons and neutrons can have dramatically different propagation properties in strongly asymmetric matter."," As it appears reasonable, very low-energy protons and neutrons can have dramatically different propagation properties in strongly asymmetric matter."
 Our conclusion is that an experimental signature of sensitivity of in-medium scattering to isospin asymmetry may be sought by probing highly asymmetric matter with energies close to the proton and neutron Fermi surfaces., Our conclusion is that an experimental signature of sensitivity of in-medium scattering to isospin asymmetry may be sought by probing highly asymmetric matter with energies close to the proton and neutron Fermi surfaces.
" Otherwise, isospin asymmetry has only a very minor impact on the mean free path."," Otherwise, isospin asymmetry has only a very minor impact on the mean free path."
 We recall that our baseline calculation of the cross sections is a microscopic one., We recall that our baseline calculation of the cross sections is a microscopic one.
 The assumptions we made in this paper concerning kinematics and sharpness of the Pauli operator simply have the purpose to make the discussion more transparent and can be improved or removed depending on the specific needs of potential users and, The assumptions we made in this paper concerning kinematics and sharpness of the Pauli operator simply have the purpose to make the discussion more transparent and can be improved or removed depending on the specific needs of potential users and
"of momentum from both feedback methods tends, slowly with time, to the same expected limit.","of momentum from both feedback methods tends, slowly with time, to the same expected limit."
" On the opposite, a higher conversion rate (α=4) produces an excess of thermal energy, creating too few (thermal case) or converting too much (kinetic case) momentum."," On the opposite, a higher conversion rate $\alpha=4$ ) produces an excess of thermal energy, creating too few (thermal case) or converting too much (kinetic case) momentum."
 With this higher o the convergence to Sedov's energy partition happens again slowly with time., With this higher $\alpha$ the convergence to Sedov's energy partition happens again slowly with time.
 We confirm here that the stability of the numerical scheme depends on the artificial viscosity We also demonstrate that the value of a=2 reproduces best Sedov's test through a faster numerical convergence., We confirm here that the stability of the numerical scheme depends on the artificial viscosity We also demonstrate that the value of $\alpha=2$ reproduces best Sedov's test through a faster numerical convergence.
 This is in agreement with the result of Monaghan(1992) for the choice of α when considering shock fronts., This is in agreement with the result of \cite{Monaghan1992} for the choice of $\alpha$ when considering shock fronts.
" However, we cannot yet advise for this somewhat higher value than typically used in cosmological simulations since, in the artificial viscosity formalism ofGADGET-2,, the a parameter does not differentiate shocks from shear flows."," However, we cannot yet advise for this somewhat higher value than typically used in cosmological simulations since, in the artificial viscosity formalism of, the $\alpha$ parameter does not differentiate shocks from shear flows."
" In the latter, a too high conversion rate would affect dramatically the velocity field at the contact surface."," In the latter, a too high conversion rate would affect dramatically the velocity field at the contact surface."
" In this specific matter of the artificial viscosity in SPH simulations, we assess that more work is required before obtaining a more general criterion."," In this specific matter of the artificial viscosity in SPH simulations, we assess that more work is required before obtaining a more general criterion."
" In simulations with a large dynamical range (e.g. astrophysical problems of galaxy and structure formation), the required time scales can span several orders of magnitude."," In simulations with a large dynamical range (e.g. astrophysical problems of galaxy and structure formation), the required time scales can span several orders of magnitude."
" In state-of-the-art cosmological simulations, up to several billions of particles are integrated over time."," In state-of-the-art cosmological simulations, up to several billions of particles are integrated over time."
 The use of constant and global time integration steps is then prohibitive in terms of computational costs., The use of constant and global time integration steps is then prohibitive in terms of computational costs.
 The individual time-step integration scheme (firstintroducedbyAarseth1963;Makino1991) allows particles to be integrated on time-steps which are functions of the local state.," The individual time-step integration scheme \citep[first introduced by][]{Aarseth1963,Makino1991} allows particles to be integrated on time-steps which are functions of the local state."
 SM09 demonstrated that an accurate description of feedback processes which involve strong energy perturbations can only be achieved by ensuring fast information transfer when using individual time-steps., SM09 demonstrated that an accurate description of feedback processes which involve strong energy perturbations can only be achieved by ensuring fast information transfer when using individual time-steps.
 They proposed an innovative scheme in which the time-step length of neighbouring gas particles is constrained by a limiter., They proposed an innovative scheme in which the time-step length of neighbouring gas particles is constrained by a limiter.
 SM09 validated their time-step limiter algorithm with Sedov’s blast wave test., SM09 validated their time-step limiter algorithm with Sedov's blast wave test.
 They also tested the effect of the time-step limiter with a SN-like explosion in a self-gravitating halo of cold gas., They also tested the effect of the time-step limiter with a SN-like explosion in a self-gravitating halo of cold gas.
 They showed that the conservation of energy and linear momentum agree remarkably well with those obtained with a more conservative (but computationally more expensive) global time-step scheme., They showed that the conservation of energy and linear momentum agree remarkably well with those obtained with a more conservative (but computationally more expensive) global time-step scheme.
" However, SM09 used initial conditions that included the explosion energy."," However, SM09 used initial conditions that included the explosion energy."
 This leads to setting the time-step of heated/kicked particles and limiting the time-step of their neighbours at the start of the simulation., This leads to setting the time-step of heated/kicked particles and limiting the time-step of their neighbours at the start of the simulation.
" Therefore, the integration was correctly performed over the appropriate time-step from the beginning."," Therefore, the integration was correctly performed over the appropriate time-step from the beginning."
" In the most general case, e.g. in cosmological simulations where the feedback events from SN or AGN activity do not affect only active particles, this would not happen."," In the most general case, e.g. in cosmological simulations where the feedback events from SN or AGN activity do not affect only active particles, this would not happen."
 We specifically design our tests to consider the energy input after the simulation has been running for several steps., We specifically design our tests to consider the energy input after the simulation has been running for several steps.
 This prevents any time-step adjustment before the integration of the dynamical and hydrodynamical equations is performed., This prevents any time-step adjustment before the integration of the dynamical and hydrodynamical equations is performed.
 We also state here that we do not limit the maximum size of the time-step., We also state here that we do not limit the maximum size of the time-step.
" This setup is justified by the aim of describing the more general case of energetic feedback, where individual time-steps reflect only the hydrodynamical state of particles."," This setup is justified by the aim of describing the more general case of energetic feedback, where individual time-steps reflect only the hydrodynamical state of particles."
" In order to study an asymmetric problem, we modify the case of the self-gravitating gas halo by off-setting the explosion."," In order to study an asymmetric problem, we modify the case of the self-gravitating gas halo by off-setting the explosion."
" In the following sections, we briefly present our method to maintain a high accuracy when considering strong feedback events in simulations using an individual time-step scheme."," In the following sections, we briefly present our method to maintain a high accuracy when considering strong feedback events in simulations using an individual time-step scheme."
 The reader interested in the implementation of this method should refer to the appendices for a detailed description., The reader interested in the implementation of this method should refer to the appendices for a detailed description.
" The results of the two sets of test simulations are then presented, focusing on the conditions needed to achieve the concordance between feedback methods."," The results of the two sets of test simulations are then presented, focusing on the conditions needed to achieve the concordance between feedback methods."
" In the context of the hierarchical time-stepping scheme presented in Appendix A,, it is important to note that if the energy content of a gas particle is suddenly increased, the particle itself will be informed as soon as it becomes active."," In the context of the hierarchical time-stepping scheme presented in Appendix \ref{app:leapfrog}, it is important to note that if the energy content of a gas particle is suddenly increased, the particle itself will be informed as soon as it becomes active."
" Within a few active time-steps, a fraction of the energy is efficiently converted from one form to the other (thermal to kinetic or versa))."," Within a few active time-steps, a fraction of the energy is efficiently converted from one form to the other (thermal to kinetic or )."
" If thermal energy is injected, it will be converted into momentum through strong hydro accelerations due to the large pressure gradient."," If thermal energy is injected, it will be converted into momentum through strong hydro accelerations due to the large pressure gradient."
" If heated particles and their neighbours do not adjust their time-steps, the integration would lead to very large velocities and an artificial excess of kinetic energy."," If heated particles and their neighbours do not adjust their time-steps, the integration would lead to very large velocities and an artificial excess of kinetic energy."
" In the other case, the injected kinetic energy is converted into thermal energy through shocks."," In the other case, the injected kinetic energy is converted into thermal energy through shocks."
" Eventually, the integration over a long time-step would lead to an excess of thermal energy, again violating energy conservation."," Eventually, the integration over a long time-step would lead to an excess of thermal energy, again violating energy conservation."
" In both scenarios, the violation of energy conservation is due to the fact that the particles do not react soon enough to the sudden change of their hydrodynamical state."," In both scenarios, the violation of energy conservation is due to the fact that the particles do not react soon enough to the sudden change of their hydrodynamical state."
" Even if the particle becomes active at the proper time, if the new energy state is not taken into account in defining the length of the next step, the integration following the energy increase will also be done over a too long time-step, leading again to non-conservation of energy."," Even if the particle becomes active at the proper time, if the new energy state is not taken into account in defining the length of the next step, the integration following the energy increase will also be done over a too long time-step, leading again to non-conservation of energy."
" In any case, it is therefore crucial to capture the initial stage of energy injection."," In any case, it is therefore crucial to capture the initial stage of energy injection."
 We emphasise here the problems that one may encounter when implementing feedback modules in SPH codes., We emphasise here the problems that one may encounter when implementing feedback modules in SPH codes.
 We show in Appendix D that the signal velocity and the acceleration are the key information needed to define the time-step., We show in Appendix \ref{app:dtcriteria} that the signal velocity and the acceleration are the key information needed to define the time-step.
 In the public release of both quantities are calculated during the computation of the local hydrodynamical forces., In the public release of both quantities are calculated during the computation of the local hydrodynamical forces.
" Consequently, any later change of the energetics of an active particle would not be taken into account during the calculation of the next time-step."," Consequently, any later change of the energetics of an active particle would not be taken into account during the calculation of the next time-step."
" On the other hand, injecting the energy before the hydro calculation would make the current time-step inconsistent with the new hydrodynamical state of the particle."," On the other hand, injecting the energy before the hydro calculation would make the current time-step inconsistent with the new hydrodynamical state of the particle."
 If the current, If the current
in several filters based on a few hours of observations with a 4-m class telescope. a non-detection in a single red optical filter withAST.. ~3—4o0 detections in two near-infrared filters. and possibly observations with Spitzer/IRAC.,"in several filters based on a few hours of observations with a 4-m class telescope, a non-detection in a single red optical filter with, ${\sim}3{-}4{\sigma}$ detections in two near-infrared filters, and possibly observations with /IRAC."
 Specifically. we investigate the degeneracy between z7-7 interpretations of such datasets with lower redshift alternatives. and how such degeneracies might be broken.," Specifically, we investigate the degeneracy between $z{\gs}7$ interpretations of such datasets with lower redshift alternatives, and how such degeneracies might be broken."
 We use Version [.1 of (Bolzonella et 22000) to fit standard Bruzual Charlot (1993) single stellar population models (Burst. E. Sa. Se. Im) to the photometric data.," We use Version 1.1 of (Bolzonella et 2000) to fit standard Bruzual Charlot (1993) single stellar population models (Burst, E, Sa, Sc, Im) to the photometric data."
 We assume a Calzetti et ((2000) extinction law. allow dust extinction in #11916 to lie. in the range O<Ay<4. adopt E(B—V)=0.03 for extinction within the Milky Way (Schlegel et 11998). and use Madau’s (1995) preseription for absorption by the. inter-galactic medium.," We assume a Calzetti et (2000) extinction law, allow dust extinction in 1916 to lie in the range $0{\le}A_V{\le}4$, adopt $E(B-V)=0.03$ for extinction within the Milky Way (Schlegel et 1998), and use Madau's (1995) prescription for absorption by the inter-galactic medium."
 In each case described below. if we fit the models to all available photometric information across the full redshift range (Oxzz ΤΟ. then we obtain \7<1 (all  values are quoted per degree of freedom) for all redshifts beyond z-7.," In each case described below, if we fit the models to all available photometric information across the full redshift range $0{\le}z{\le}11$ ), then we obtain ${\chi}^2{\ls}1$ (all ${\chi}^2$ values are quoted per degree of freedom) for all redshifts beyond $z{\simeq}7$."
" This is because the models are defined to have zero flux short-ward of the Lyman limit (A,,,20.09124/m)."," This is because the models are defined to have zero flux short-ward of the Lyman limit ${\lambda}_{\rm
rest}{=}0.0912{\mu}{\rm m}$ )."
" At «ο. the models also have negligible flux short-ward of Lyman-o. (À,,20.1215jm) due to the Lyman forests and Gunn Peterson absorption."," At $z{\gs}6$, the models also have negligible flux short-ward of ${\alpha}$ ${\lambda}_{\rm
rest}{=}0.1215{\mu}{\rm m}$ ) due to the Lyman forests and Gunn Peterson absorption."
 These two spectral features are progressively redshifted long-ward of the observed optical filters at high redshift. rendering the shorter wavelength detection limits irrelevant to the fits.," These two spectral features are progressively redshifted long-ward of the observed optical filters at high redshift, rendering the shorter wavelength detection limits irrelevant to the fits."
 The goodness of fit is therefore systematically over-estimated at high redshift if all photometric information 1s included in the fit across the full redshift range., The goodness of fit is therefore systematically over-estimated at high redshift if all photometric information is included in the fit across the full redshift range.
 We therefore fit the models to the data in a series of redshift chunks — only the photometric information that Hes longward of red-shifted Lyman-o 1s considered in each redshift chunk., We therefore fit the models to the data in a series of redshift chunks – only the photometric information that lies longward of red-shifted ${\alpha}$ is considered in each redshift chunk.
 The results described below are insensitive to whether the Lyman limit ts substitutec for Lyman-a., The results described below are insensitive to whether the Lyman limit is substituted for ${\alpha}$ .
 First. we fit the models to PO4Zs VRIJHK-band ground-based photometry tthe VAR7J-band detection limits anc the H K-band detections listed in column 4 of Table 1)).," First, we fit the models to P04's $VRIJHK$ -band ground-based photometry the $VRIJ$ -band detection limits and the $HK$ -band detections listed in column 4 of Table \ref{table:imaging}) )."
 Note that ΡΟΗ used different sized apertures for the near-infrarec (1.57) and optical (0.6) photometry respectively., Note that P04 used different sized apertures for the near-infrared $1.5''$ ) and optical $0.6''$ ) photometry respectively.
 We find two acceptable solutions: zz2.6 and 5-10 (Fig. 2)).," We find two acceptable solutions: $z{\simeq}2.6$ and $z{\simeq}10$ (Fig. \ref{fig:photoz}) ),"
 neither of which require any dust-obscuration within #11916. the latter having the lower formal 4 value.," neither of which require any dust-obscuration within 1916, the latter having the lower formal ${\chi}^2$ value."
 We also scale PO4's ground-based optical non-detections to that appropriate for à consistent photometric aperture of three times the seeing disk at all wavelengths (V:226.2. R226.3. F724.7).," We also scale P04's ground-based optical non-detections to that appropriate for a consistent photometric aperture of three times the seeing disk at all wavelengths $V{\ge}26.2$, $R{\ge}26.3$, $I{\ge}24.7$ )."
 This marginally improves the goodness of fit of the z—2.6 solution. and is otherwise indistinguishable from the original fit.," This marginally improves the goodness of fit of the $z{\simeq}2.6$ solution, and is otherwise indistinguishable from the original fit."
 We retain the matched photometric apertures for the remainder of this analysis., We retain the matched photometric apertures for the remainder of this analysis.
 To improve the discrimination between low- and high-redshift solutions. we add the sensitive AST//WFPC2 detection threshold of Απο. (Table 1)) to the photometric constraints.," To improve the discrimination between low- and high-redshift solutions, we add the sensitive /WFPC2 detection threshold of $R_{702}{\ge}27.0$ (Table \ref{table:imaging}) ) to the photometric constraints."
 We fit the synthetic SEDs to the combined VARR;o:7JHK dataset., We fit the synthetic SEDs to the combined $VRR_{702}IJHK$ dataset.
 The goodness of fit of the 72.59 solution is marginally worse relative to the VRIJHK-band analysis. because the most stringent optical non-detection comes from. the Rso-band.," The goodness of fit of the $z{\simeq}2.59$ solution is marginally worse relative to the $VRIJHK$ -band analysis, because the most stringent optical non-detection comes from the $R_{702}$ -band."
 However. in general. the 47 as a function of redshift is indistinguishable between this fit and the VR//HK-band fits.," However, in general, the ${\chi}^2$ as a function of redshift is indistinguishable between this fit and the $VRIJHK$ -band fits."
 This is probably because the \- is dominated by the non-detections in six out of the eight observed filters., This is probably because the ${\chi}^2$ is dominated by the non-detections in six out of the eight observed filters.
 To test this we re-fit the models. limiting the data to just the Ryq2./HK-bands. tthe most sensitive non-detection (R72). the reddest non-detection (J) and the two detections (HA).," To test this we re-fit the models, limiting the data to just the $R_{702}JHK$ -bands, the most sensitive non-detection $R_{702}$ ), the reddest non-detection $J$ ) and the two detections $HK$ )."
 The impact of the sensitive detection limit from the data is now clearly evident n the new  distribution. as shown in Fig. 2..," The impact of the sensitive detection limit from the data is now clearly evident in the new ${\chi}^2$ distribution, as shown in Fig. \ref{fig:photoz}."
 The only acceptable fit to the Ryo2/AK-band data is at zz10., The only acceptable fit to the $R_{702}JHK$ -band data is at $z{\simeq}10$.
 We show the best fit SEDs at z22.6 and -ΞΙ0. based on the AR;7HK-band data in the right panel of Fig. 2..," We show the best fit SEDs at $z{=}2.6$ and $z{=}10$, based on the $R_{702}JHK$ -band data in the right panel of Fig. \ref{fig:photoz}. ."
 This demonstrates the potential power of IRAC photometry to. further. discriminate between low and high-redshift interpretations of candidate high-redshift galaxies., This demonstrates the potential power of IRAC photometry to further discriminate between low and high-redshift interpretations of candidate high-redshift galaxies.
" The SED of low redshift solutions z-2.6) is red at As,7-3—5,/m. eand the SED of high redshift solutions (z7:10) is blue at similar wavelengths."," The SED of low redshift solutions $z{\simeq}2.6$ ) is red at ${\lambda}_{\rm
 obs}{\sim}3{-}5{\mu}{\rm m}$, and the SED of high redshift solutions $z{\simeq}10$ ) is blue at similar wavelengths."
" We add the detection limits obtained at Àa,23.6j/mm and μπι to the ΑποK-band data and re-fit the spectral templates.", We add the detection limits obtained at ${\lambda}_{\rm obs}{=}3.6{\mu}{\rm m}$ and $4.5{\mu}{\rm m}$ to the $R_{702}JHK$ -band data and re-fit the spectral templates.
 The result is shown in the left panel of Fig., The result is shown in the left panel of Fig.
2 —the reduced X of the z—2.6 solution is now in excess of 2. and the only redshifts which achieve an acceptable fit to the data are 9.5zzz101.5.," \ref{fig:photoz} – the reduced ${\chi}^2$ of the $z{\simeq}2.6$ solution is now in excess of 2, and the only redshifts which achieve an acceptable fit to the data are $9.5{\ls}z{\ls}11.5$."
 In summary. a single sensitive non-detection derived from imaging. combined with deep ground-based near-infrared imaging and 72.400 second integrations with Spitzer//IAC at 3.6 and 4.5j/m can break the degeneracy between low and high-redshift interpretations of candidate 227 stellar systems (see also Εσαπ et 22004: Eyles et 22005).," In summary, a single sensitive non-detection derived from imaging, combined with deep ground-based near-infrared imaging and ${\sim}2,400$ second integrations with /IRAC at $3.6$ and $4.5{\mu}{\rm m}$ can break the degeneracy between low and high-redshift interpretations of candidate $z{\gs}7$ stellar systems (see also Egami et 2004; Eyles et 2005)."
 Whilst this is not an exhaustive study. it demonstrates that despite the demise of #11916. searches for gravitationally magnified galaxies at extremely high redshifts using ground-based near-infrared data remain feasible when combined with sensitive space-based optical and mid-infrared data.," Whilst this is not an exhaustive study, it demonstrates that despite the demise of 1916, searches for gravitationally magnified galaxies at extremely high redshifts using ground-based near-infrared data remain feasible when combined with sensitive space-based optical and mid-infrared data."
 We have analyzed new and archival observations of the <=10 galaxy #11916 behind the foreground galaxy cluster Jens 11835 (z20.25) spanning 0.955Αρ5m., We have analyzed new and archival observations of the $z{=}10$ galaxy 1916 behind the foreground galaxy cluster lens 1835 $z{=}0.25$ ) spanning $0.35{\le}{\lambda}_{\rm obs}{\le}5{\mu}{\rm m}$.
 Statistically significant flux is not detected in. any of these data. including our independent re-analysis of PO4's discovery H- and K-band data.," Statistically significant flux is not detected in any of these data, including our independent re-analysis of P04's discovery $H$ - and $K$ -band data."
 The 30 detection thresholds are: F(XsS0445m)24.5.107 Ρο. spectral resolution element. R503::727.0. H225.0. K25.0. F(3.6um)z0.75jJy. and. F(d.5jym)z0.754Jy.. where. the photometric limits are calculated in apertures with a diameter 3« that of the seeing disk.," The $3{\sigma}$ detection thresholds are: $F({\lambda}_{\rm obs}{=}0.44{\mu}{\rm m}){\gs}4.5{\times}10^{-19}{\rm
 erg\,s^{-1}\,cm^{-2}}$ per spectral resolution element, $R_{702}{\ge}27.0$, $H{\ge}25.0$, $K{\ge}25.0$, $F(3.6{\mu}{\rm m}){\le}0.75{\mu}{\rm Jy}$ and $F(4.5{\mu}{\rm m}){\le}0.75{\mu}{\rm Jy}$, where the photometric limits are calculated in apertures with a diameter $3{\times}$ that of the seeing disk."
 Combining these results with those of Bremer et ((2004) and Weatherley et ((2004). we are therefore forced by the data to conclude that there is no statistically sound evidence for the existence of 211916.," Combining these results with those of Bremer et (2004) and Weatherley et (2004), we are therefore forced by the data to conclude that there is no statistically sound evidence for the existence of 1916."
 We also show that inefficient bad-pixel rejection and issues relating to the calculation of the background noise can broadly account for the differences between PO4's near-infrared photometry and our own using the same data., We also show that inefficient bad-pixel rejection and issues relating to the calculation of the background noise can broadly account for the differences between P04's near-infrared photometry and our own using the same data.
 The balance of probability is therefore that 411916 was a false detection in. PO4s analysis., The balance of probability is therefore that 1916 was a false detection in P04's analysis.
 From a broader perspective. the demise of #11916 warnsof the hazards of operating close to the detection threshold of deep ground-based near-infrared data.," From a broader perspective, the demise of 1916 warnsof the hazards of operating close to the detection threshold of deep ground-based near-infrared data."
 The need for gravitational magnification to boost the observed flux of faint (L:O0.1 L) galaxies at z—06-8and populations of galaxies at still higher redshifts (see refintro)) is undiminished by our results on #11916.," The need for gravitational magnification to boost the observed flux of faint $L{\ls}0.1\,L^\star$ ) galaxies at $z{\sim}6$ –8and populations of galaxies at still higher redshifts (see \\ref{intro}) ) is undiminished by our results on 1916."
 However an important issue is whether it is feasible to find such galaxies using ground-based facilities., However an important issue is whether it is feasible to find such galaxies using ground-based facilities.
 We explore. this, We explore this
"scale according to where we asstuue au ideal equation of state (106. Py=Πρ) aud take the total pressure to be P7, plus that of the equipartition maguctic field.","scale according to where we assume an ideal equation of state (i.e., $P_g=R_g\rho T/\mu$ ) and take the total pressure to be $P_g$ plus that of the equipartition magnetic field."
 Here. Ry is the gas constant aud pois the 116221 molecular weight.," Here, $R_g$ is the gas constant and $\mu$ is the mean molecular weight."
 It is straightforward to iutegrate Equation (3) for a power-law temperature distribution. which gives where and the pressure scale height is For the gas to be confined within the core of the cluster. we must have where vis of order onc.," It is straightforward to integrate Equation (3) for a power-law temperature distribution, which gives where and the pressure scale height is For the gas to be confined within the core of the cluster, we must have where $\eta$ is of order one."
 This places an upper But on the temperature. since clearly a hotter gas will have a greater extension.," This places an upper limit on the temperature, since clearly a hotter gas will have a greater extension."
" More specifically. we see from Equation (7) that For example. when o=Ll. we see that Ty=5«105,7. IK. Tn principle. the gas cau therefore attain the temperature (25&107 IS) required to produce the observed GIIz (thermal) svuchrotron radiation from Ser A* ifyze3 (see Melia 1991)."," More specifically, we see from Equation (7) that For example, when $\alpha=-1$, we see that $T_0=5\times 10^8\eta^2$ K. In principle, the gas can therefore attain the temperature $\simeq 5\times 10^9$ K) required to produce the observed GHz (thermal) synchrotron radiation from Sgr A* if $\eta\simeq 3$ (see Melia 1994)."
 In addition. we see that for a captured gas mass Mayas~10 SAT. the equipartition magnetic field cau reach an intensity of ~10 eauss. or ercater.," In addition, we see that for a captured gas mass $M_{gas}\sim 10^{-3}\;M_\odot$ , the equipartition magnetic field can reach an intensity of $\sim 10$ gauss, or greater."
 This is low compared with the value (~ huudreds of gauss) used in the best fits by Molia (1991) and Naravan. Yi. aud Alahadevan (1995). but is certainly within the range where an attempt to fit the spectrum with this coufiguratiou of trapped eas isduterestine.," This is low compared with the value $\sim$ hundreds of gauss) used in the best fits by Melia (1994) and Narayan, Yi, and Mahadevan (1995), but is certainly within the range where an attempt to fit the spectrum with this configuration of trapped gas isinteresting."
atmosphere of this star.,atmosphere of this star.
 A plot of the EWs of C II A 4267 line versus the mass transfer rates for ten stars points out an existence of a relationship between these quantities., A plot of the EWs of C II $\lambda$ 4267 line versus the mass transfer rates for ten stars points out an existence of a relationship between these quantities.
 The higher mass transfer rate corresponds to the smaller EWs., The higher mass transfer rate corresponds to the smaller EWs.
" This may be taken as mixing on the surface layers of the accreted star which leads to lower C and higher N abundance (de Greve and Cugier, 1989)."," This may be taken as mixing on the surface layers of the accreted star which leads to lower C and higher N abundance (de Greve and Cugier, 1989)."
 Logarithmic abundances relative to the standard stars and the Sun clearly confirm this expectation., Logarithmic abundances relative to the standard stars and the Sun clearly confirm this expectation.
" Average [Nc/N:«] are estimated to be -0.82 and -0.54 relative to the Sun and the standard stars, respectively."," Average $ [N_{C} /N_{tot}] $ are estimated to be -0.82 and -0.54 relative to the Sun and the standard stars, respectively."
" Specially, abundance ratios will give us important information about evolutionary histories and current status of the interacting Algols."," Specially, abundance ratios will give us important information about evolutionary histories and current status of the interacting Algols."
" However, it should be noted that measuring the isotopic ratios in the mass-losing secondaries is difficult due to their faintness, contributing a few per cent to the total light."," However, it should be noted that measuring the isotopic ratios in the mass-losing secondaries is difficult due to their faintness, contributing a few per cent to the total light."
 The authors acknowledge generous allotments of observing time at Asiago Observatory (Italy) and TUBITAK National Observatory (TUG) of Turkey., The authors acknowledge generous allotments of observing time at Asiago Observatory (Italy) and TUBITAK National Observatory (TUG) of Turkey.
 We also wish to thank the Turkish Scientific and Technical Research Council for supporting this work through grant No., We also wish to thank the Turkish Scientific and Technical Research Council for supporting this work through grant No.
 109T708 and TUBBITTAK National Observatory (TUG) for a partial support in using RTT150 with project number 1OARTT150-, 109T708 and TÜBBİTTAK National Observatory (TUG) for a partial support in using RTT150 with project number 10ARTT150-493-0.
From equations (15)) and (51)) a formula for the mass of the galaxy is found An application of the previous formula is reported in Figure (15)) . where the mass ol galaxies as a function of the absolute magnitude in the five bands of SDSS is drawn.,"From equations \ref{massa}) ) and \ref{mstar}) ) a formula for the mass of the galaxy is found An application of the previous formula is reported in Figure \ref{masse}) ), where the mass of galaxies as a function of the absolute magnitude in the five bands of SDSS is drawn."
 In this Figure Vi... is different lor each selected. band. and equal to the value suggestedin equation (16) of Blantonetal.(2001 )..," In this Figure ${M}_{bol,\sun}$ is different for each selected band and equal to the value suggestedin equation (16) of \cite{blanton}. ."
to measure the SAIBIT mass. spin aud the location of the X-ray flare.,"to measure the SMBH mass, spin and the location of the X-ray flare."
 By stucdving such transfer fictions. it is found that a characteristic signature of rapidly rotating black holes is a cred-tail’ ou the traustfer function.," By studying such transfer functions, it is found that a characteristic signature of rapidly rotating black holes is a `red-tail' on the transfer function."
 This feature corresponds to lighly redshifted aud delaved line eimissiou that originates from au mvardly moving rug of illuination/fluorescence that asviiptoticallv freezes at the horizon (see Revuolds et al., This feature corresponds to highly redshifted and delayed line emission that originates from an inwardly moving ring of illumination/fluorescence that asymptotically freezes at the horizon (see Reynolds et al.
 1999 for a discussion of this feature)., 1999 for a discussion of this feature).
 The primary observational difficulty in characterizing iron line reverberation will be obtaining the required signal-to-noise., The primary observational difficulty in characterizing iron line reverberation will be obtaining the required signal-to-noise.
 One nimmst be able to measure an iron line profile on a timescale of than~GMfe=5OOAMs. where we have normalized. to ainass of 105AL...," One must be able to measure an iron line profile on a timescale of $t_{\rm
reverb}\sim GM/c^3\approx 500 M_8\,{\rm s}$, where we have normalized to a mass of $10^8\,{\rm M}_\odot$."
 This requires an iustriuneut such asConstellation-X., This requires an instrument such as.
 Figure 3 shows thatConstellation-X can indeed detect reverberation from a bright ACN with a 1uass of 105AL...," Figure 3 shows that can indeed detect reverberation from a bright AGN with a mass of $10^8\,{\rm
M}_\odot$."
 Furthermore. the sieuatures of black hole spin aay well be witlin reach ofC'onstellution-X (Young Reynolds 1999).," Furthermore, the signatures of black hole spin may well be within reach of (Young Reynolds 1999)."
 Although these simulations lake the somewhat artificial assumption that the N-rav flare is lustantaneous aud located on the axis of the syste. it provides encouragement that reverberation sienatures nav be observable in the foresecable future.," Although these simulations make the somewhat artificial assumption that the X-ray flare is instantaneous and located on the axis of the system, it provides encouragement that reverberation signatures may be observable in the foreseeable future."
 Of course. the occurrence of multiple. overlapping flares will also hamper the interpretation of iron line reverberation.," Of course, the occurrence of multiple, overlapping flares will also hamper the interpretation of iron line reverberation."
 The best way to discutangle these flares is still the subject of current work., The best way to disentangle these flares is still the subject of current work.
 Towever.Constellution-X mav have the required signal to noise to allow the direct fitting of multiple transter functions to real data (seo Young Reynolds 1999).," However, may have the required signal to noise to allow the direct fitting of multiple transfer functions to real data (see Young Reynolds 1999)."
 I will eud by briefly discussing an exciting idea which will allow us to image the central regions of nearby ACN with sufficieut angular resolution to probe structure on scales sinaller than the size of the eveut horizou., I will end by briefly discussing an exciting idea which will allow us to image the central regions of nearby AGN with sufficient angular resolution to probe structure on scales smaller than the size of the event horizon.
 By combining diffraction limited X-ray optics with the interferometric technologies that are currently being developed for the Space Interferometer Mission (110. it is within our technological reach to construct au X-ray interferometer capable of achieving sub-auicroarcsecoud resolution (this concept las become known asAZANTAL the Micro-arcsee Νταν Tuterterometer Mission: sce Zia.yous).," By combining diffraction limited X-ray optics with the interferometric technologies that are currently being developed for the Space Interferometer Mission ), it is within our technological reach to construct an X-ray interferometer capable of achieving sub-microarcsecond resolution (this concept has become known as, the Micro-arcsec X-ray Interferometer Mission; see )."
 As well as the obvious appeal of directly imaging an accreting black hole. au observatory capable of achieving 0.14arcsec would vicld major scicutific return.," As well as the obvious appeal of directly imaging an accreting black hole, an observatory capable of achieving $0.1\mu\,{\rm arcsec}$ would yield major scientific return."
 The ecoletry of the N-ray source (aud the spatial nature of the N-vay flares) would be open to direct dmaging studies., The geometry of the X-ray source (and the spatial nature of the X-ray flares) would be open to direct imaging studies.
 N-rav activity or fluorescence frou within the radius of inaveinal stability could be easily sccu (this region would be well resolved)., X-ray activity or fluorescence from within the radius of marginal stability could be easily seen (this region would be well resolved).
" We nuelit also expect there to be X-ray cussion from the base of the jet im the region where the magnetic field couples to the black hole spin via the Dlaudford-Zuajek process,", We might also expect there to be X-ray emission from the base of the jet in the region where the magnetic field couples to the black hole spin via the Blandford-Znajek process.
 Such emission could be imaged. thereby providing the first look at these exotic physical mechanisius at work.," Such emission could be imaged, thereby providing the first look at these exotic physical mechanisms at work."
 If au interferometer can be constructed with sufficient effective area. we will be able to use the fluorescent iron line to male detailed velocity maps across the image.," If an interferometer can be constructed with sufficient effective area, we will be able to use the fluorescent iron line to make detailed velocity maps across the image."
 These velocity fields would provide direct constraints of the black hole mass aud spin. and iaplicitlv provide a stringent test of strong field. General Relativity.," These velocity fields would provide direct constraints of the black hole mass and spin, and implicitly provide a stringent test of strong field General Relativity."
Galaxy clusters have been known ancl studied. for many vears. and the racial dependence of cluster temperatures is becoming a testing ground for models of structure formation and for our understanding of gas dvnamics.,"Galaxy clusters have been known and studied for many years, and the radial dependence of cluster temperatures is becoming a testing ground for models of structure formation and for our understanding of gas dynamics."
 Galaxy. clusters are not Iso-thermal., Galaxy clusters are not iso-thermal.
 Instead. the emerging temperature profile is one where the temperature is approximately Lat or increases from the centre to some characteristic radius. and then decreases again for larger raclii.," Instead, the emerging temperature profile is one where the temperature is approximately flat or increases from the centre to some characteristic radius, and then decreases again for larger radii."
 The central temperature. clecrement has been much discussed. and. the possibility. of cooling flows has been explained in exeellent reviews (Fabian1994:Donahue 2003).," The central temperature decrement has been much discussed and the possibility of cooling flows has been explained in excellent reviews \cite{fabian94,donahue03}."
. Many. clusters are well Gt with a power law Zr in the very central region (Voigt&Fabian 2003).. with a slope. r. between 0.15 and 0.45.," Many clusters are well fit with a power law $ T \sim r^\tau$ in the very central region \cite{voigt03}, , with a slope, $\tau$, between $0.15$ and $0.45$."
 Numerical simulations are only now beginning to see this central decrement (see Motl et al., Numerical simulations are only now beginning to see this central decrement (see Motl et al.
 (2003). for references)., \shortcite{motl03} for references).
 The outer. temperature decrease is well established both observationally (Alarkevitehetal.1998:DePointecouteau2003) and numerically (see Lin et al.," The outer temperature decrease is well established both observationally \cite{markevitch98,degrandi02,kaastra03,arnaud} and numerically (see Lin et al."
 (2003) and references therein)., \shortcite{lin03} and references therein).
 The simulations have even started to provide estimates of the outer temperature. decrement which are well fit with power laws in fair agreement with observations., The simulations have even started to provide estimates of the outer temperature decrement which are well fit with power laws in fair agreement with observations.
 Such non-trivial temperature profiles also alfect derived. cosmological parameters like the. Hubble parameter (Battistellietal.2003:Lin.Norman 2003).," Such non-trivial temperature profiles also affect derived cosmological parameters like the Hubble parameter \cite{battistelli,lin03}. ."
. On the other hand. observations of the (SZ) effect are becoming increasingly accurate (Laroqueetal.2002:DePetrisBattis- 2003).. however. most analyses of SZ observational data are made under the simplifving assumption of iso-thermality.," On the other hand, observations of the (SZ) effect are becoming increasingly accurate \cite{laroque02,coma,battistelli}, however, most analyses of SZ observational data are made under the simplifying assumption of iso-thermality."
 Several groups have considered. the ability of [uture SZ observations to extract cluster parameters (Ixnox. 2004).. but always under the assumption of iso-thermality.," Several groups have considered the ability of future SZ observations to extract cluster parameters \cite{knox,AHL}, but always under the assumption of iso-thermality."
 We therefore set out to study the elect on the SZ derived parameters of non-trivial cluster temperature profiles., We therefore set out to study the effect on the SZ derived parameters of non-trivial cluster temperature profiles.
 In section 2 we show that the relevant quantities for SZ observations are. Compton-averaged. in. particular. we emphasize the difference between the temperatures derived through X-ray and SZ observation For clusters which are not iso-thermal.," In section \ref{sec:sz} we show that the relevant quantities for SZ observations are Compton-averaged, in particular we emphasize the difference between the temperatures derived through X-ray and SZ observation for clusters which are not iso-thermal."
" In section 3. we use observed. ancl simulated cluster profiles to show that the systematic shift in the derived. peculiar. velocity. ¢,. is of the order 1020%."," In section \ref{sec:cluster} we use observed and simulated cluster profiles to show that the systematic shift in the derived peculiar velocity, $v_p$, is of the order $10-20\%$."
 Finally we comment on the possibility of ceprojection of SZ observed spectra in section 4.., Finally we comment on the possibility of deprojection of SZ observed spectra in section \ref{sec:onion}.
" The SZ elfect is dominated by the Compton parameter sor fn Hyd. where er is the Thomson cross section. n..Z; ancl n, arenumber density. temperature ancl mass of the electrons. and the integral is along the line of sight."," The SZ effect is dominated by the Compton parameter y = _T n_e dl, where $\sigma_T$ is the Thomson cross section, $n_e,T_e$ and $m_e$ arenumber density, temperature and mass of the electrons, and the integral is along the line of sight."
 If we extract the radial dependenceof the parameters, If we extract the radial dependenceof the parameters
Including the SiO emission places constraints on the initial repartition of silicon-bearing material but is of little help in discriminating between the models which fit the H» rotational excitation diagram.,Including the SiO emission places constraints on the initial repartition of silicon-bearing material but is of little help in discriminating between the models which fit the $_2$ rotational excitation diagram.
 We have validated our approach to calculating the water line intensities. which uses the results of shock model as input to a separate radiative transfer code. based on the LVG approximation.," We have validated our approach to calculating the water line intensities, which uses the results of shock model as input to a separate radiative transfer code, based on the LVG approximation."
 To this end. we provide a full comparison between the statistical-equilibrium and the steady-state approximations. based on reference shock models of the stationary. C-type. and non-stationary. CJ-type.," To this end, we provide a full comparison between the statistical-equilibrium and the steady-state approximations, based on reference shock models of the stationary, C-type, and non-stationary, CJ-type."
 We make predictions of water emission line intensities for the SiO knot in the BHR71 outflow., We make predictions of water emission line intensities for the SiO knot in the BHR71 outflow.
 As we expect emission by water to coincide with that by SiO and H». we provide the results of the LVG calculations of the water line intensities for models that fit satisfactorily both the SiO and H» Given that the excitation conditions are relatively uniform in the internal parts of the outflow. we consider that our caleulations can be used to compare with observations of water throughout this region.," As we expect emission by water to coincide with that by SiO and $_2$, we provide the results of the LVG calculations of the water line intensities for models that fit satisfactorily both the SiO and $_2$ Given that the excitation conditions are relatively uniform in the internal parts of the outflow, we consider that our calculations can be used to compare with observations of water throughout this region."
 We shall discuss in a future publication. ground-based observations of CO transitions. undertaken with. the APEX facility. that could bring further insights into. the physical conditions in the BHR71 outflow.," We shall discuss in a future publication ground-based observations of CO transitions, undertaken with the APEX facility, that could bring further insights into the physical conditions in the BHR71 outflow."
 Complementary observations of sulphur-bearing molecules and other shock tracers. such as HNCO. are planned.," Complementary observations of sulphur-bearing molecules and other shock tracers, such as HNCO, are planned."
 One of the aims of these studies is to investigate the suitability of “chemically active outflows’. such as BHR71. for observations with the new generation of instruments. such as ALMA.," One of the aims of these studies is to investigate the suitability of chemically active outflows', such as BHR71, for observations with the new generation of instruments, such as ALMA."
Alost of the analyses of Cosmic Microwave. Dackground (CAIB) data focus on the measurement of the power spectrum of temperature fluctuations.,Most of the analyses of Cosmic Microwave Background (CMB) data focus on the measurement of the power spectrum of temperature fluctuations.
 Information. on this second. order moment is. crucial to determine the fundamental parameters of the cosmological mocel corresponding to our universe., Information on this second order moment is crucial to determine the fundamental parameters of the cosmological model corresponding to our universe.
 However. this determination relies on the Gaussian. hypothesis for the temperature distribution.," However, this determination relies on the Gaussian hypothesis for the temperature distribution."
 Establishing the statistical character of. the CAIB fluctuations will provide crucial evidence about the ohvsical origin of the primordial density. Ductuations in he carly universe., Establishing the statistical character of the CMB fluctuations will provide crucial evidence about the physical origin of the primordial density fluctuations in the early universe.
 Simple inflationary moceLg edic a Gaussian. homogeneous. anc isotropic random fiel or the temperature Huctuations.," Simple inflationary models predict a Gaussian, homogeneous and isotropic random field for the temperature fluctuations."
 On the contrary. non-standard inflation ancl cosmic defects. eencricalA oreclic non-Gaussian random fields.," On the contrary, non-standard inflation and cosmic defects generically predict non-Gaussian random fields."
 Reeent CAIB observations by Boomerang. DASI and ALANIALA-1 (Netterfield et al.," Recent CMB observations by Boomerang, DASI and MAXIMA-1 (Netterfield et al."
 2001. uavke et al.," 2001, Pryke et al."
 2001. Stompor et al.," 2001, Stompor et al."
 2001) have establishec or the first time the presence of multiple acoustic peaks in he CALB power spectrum., 2001) have established for the first time the presence of multiple acoustic peaks in the CMB power spectrum.
 As à consequence cosmic defects cannot be the dominant. source of density. perturbations in he universe., As a consequence cosmic defects cannot be the dominant source of density perturbations in the universe.
 Even if they are present as a sub-dominant component confirmation of its existence will be best made by appropriate techniques searching for non-Gaussian features in the CMD maps., Even if they are present as a sub-dominant component confirmation of its existence will be best made by appropriate techniques searching for non-Gaussian features in the CMB maps.
 Since a random. field. can departure [rom a Gaussian one in many dillerent. ways there is not à unique wav to detect and characterise deviations from Gaussianityv., Since a random field can departure from a Gaussian one in many different ways there is not a unique way to detect and characterise deviations from Gaussianity.
 Thus. depending on the kind. of features one is looking for some specific methods will prove to be more ellicient than others.," Thus, depending on the kind of features one is looking for some specific methods will prove to be more efficient than others."
 Efficient methods are able to extract relevant information on the non-Gaussian nature of the data which is otherwise hidden in the temperature Ductuation maps., Efficient methods are able to extract relevant information on the non-Gaussian nature of the data which is otherwise hidden in the temperature fluctuation maps.
 A large number of methods have been already proposed. to search. for non-CGaussianity in CMD maps., A large number of methods have been already proposed to search for non-Gaussianity in CMB maps.
 The methods can be grouped by the spaces (real. Fourier...) in whieh they act.," The methods can be grouped by the spaces (real, Fourier,...) in which they act."
 In real space. standard. quantities used. are the cumulants which contain information on the I-pdf only.," In real space, standard quantities used are the cumulants which contain information on the 1-pdf only."
 Information on the n-pdf£ can be obtained through the Edgeworth expansion (Contaldi ct al., Information on the n-pdf can be obtained through the Edgeworth expansion (Contaldi et al.
 2000) or alternative expansions with a proper normalization (Rocha et al., 2000) or alternative expansions with a proper normalization (Rocha et al.
 2000)., 2000).
 Other quantities focus on topological and gcomietric statistics. c.g. Minkowski functionals implemented. on the sphere (Schmalzing anc Gorski 1998): statistics of excursion sets. e.g. characteristics of peaks (Darreiro. nez-CGonzallez ancl Sanz 2001). extrema correlation function (Lleavens anc Gupta. 2001): also. ecometrical Jaracteristies of polarisation have already been investigated vasclsky and Novikov 1998).," Other quantities focus on topological and geometric statistics, e.g. Minkowski functionals implemented on the sphere (Schmalzing and Gorski 1998); statistics of excursion sets, e.g. characteristics of peaks (Barreiro, nez-Gonzállez and Sanz 2001), extrema correlation function (Heavens and Gupta 2001); also geometrical characteristics of polarisation have already been investigated (Naselsky and Novikov 1998)."
 Alultifractal analysis and HHroughness have been applied. to the COBE-DAIR. data —Diego οἱ al., Multifractal analysis and roughness have been applied to the COBE-DMR data (Diego et al.
 1998. Mollerach et al.," 1998, Mollerach et al."
 1999)., 1999).
 In Fourier space. 10 bispectrum has been applied. in. several occasions to analyse the COBE-DAIR data (see e.g. Ferreira ct al.," In Fourier space, the bispectrum has been applied in several occasions to analyse the COBE-DMR data (see e.g. Ferreira et al."
 1998) as well as an extension to include. possible correlations among multipoles (Magueijo 2000)., 1998) as well as an extension to include possible correlations among multipoles (Magueijo 2000).
 An alternative approach is to work in eigen space. extracting the eigenmocdes from a principal component analysis.," An alternative approach is to work in eigen space, extracting the eigenmodes from a principal component analysis."
 This approach has been taken by Bromley and Tegmark (1999) for the CODE-DAMIDIt data and by Wu et al. (, This approach has been taken by Bromley and Tegmark (1999) for the COBE-DMR data and by Wu et al. (
2001) for the ALANIALA-1 data.,2001) for the MAXIMA-1 data.
 In spite of all this effort there is not any strong evidence of deviations, In spite of all this effort there is not any strong evidence of deviations
We performed Monte Carlo simulation of conal emission of radio pulsars according to procedure described in items 1-16 listed above.,We performed Monte Carlo simulation of conal emission of radio pulsars according to procedure described in items 1–16 listed above.
" We checked 160 (80 for circular and 80 for elliptical beams) different combinations of trial functions for parent distribution periods P, inclination angles a and opening angles pio."," We checked 160 (80 for circular and 80 for elliptical beams) different combinations of trial functions for parent distribution periods $P$, inclination angles $\alpha$ and opening angles $\rho_{10}$."
 For each combination of the above distribution functions we searched the two-dimensional grid of free parameters of the period distribution function., For each combination of the above distribution functions we searched the two-dimensional grid of free parameters of the period distribution function.
 It gave the total number of 641360 single simulation runs., It gave the total number of $641\:360$ single simulation runs.
" In all solutions we identified ""detectable"" cases in which all probabilities (of P, α and Wio) were greater than 0.005 (0.5%)) and occurrence of interpulses were at the observed"," In all solutions we identified ”detectable” cases in which all probabilities (of $P$, $\alpha$ and $W_{10}$ ) were greater than 0.005 ) and occurrence of interpulses were at the observed"
"dust halo hypothesis with respect to steady diffuse emission, as expected, for example, in a supernova remnant.","dust halo hypothesis with respect to steady diffuse emission, as expected, for example, in a supernova remnant."
" However, the data quality of the outburst observations is not good enough to detect significant variability in the extended emission at the level expected for a dust scattering halo."," However, the data quality of the outburst observations is not good enough to detect significant variability in the extended emission at the level expected for a dust scattering halo."
" For the spectroscopy (performed with the 12.6 fitting package; Arnaud 1996)) we concentrate first on the sspectra which, owing to the EPIC instrument high throughput and long observing time, are those with the best statistical quality."," For the spectroscopy (performed with the 12.6 fitting package; \citealt{arnaud96}) ) we concentrate first on the spectra which, owing to the EPIC instrument high throughput and long observing time, are those with the best statistical quality."
 We fit the spectra from the three observations simultaneously with the hydrogen column density tied between all data sets., We fit the spectra from the three observations simultaneously with the hydrogen column density tied between all data sets.
" Photons having energies below 2 keV and above 10 keV were ignored, owing to the very few counts from1833-0832."," Photons having energies below 2 keV and above 10 keV were ignored, owing to the very few counts from."
". This resulted in about 6100+80, 5800+90, and 4000+70 net EPIC-pn counts in the three observations in chronological order."," This resulted in about $6100\pm80$, $5800\pm90$, and $4000\pm70$ net EPIC-pn counts in the three observations in chronological order."
 The abundances used are those of Anders&Grevesse(1989) and photoelectric absorption cross-sections are from Balucinska-Church&McCammon(1992)., The abundances used are those of \citet{anders89} and photoelectric absorption cross-sections are from \citet{balucinska92}.
". Using the first twelve Swift//XRT observations in PC mode, Gó£üsetal.(2010a) could not discriminate between a power-law and a blackbody as the model that better describes the spectrum of the SGR persistent emission."," Using the first twelve /XRT observations in PC mode, \citet{gogus10short} could not discriminate between a power-law and a blackbody as the model that better describes the spectrum of the SGR persistent emission."
" A fit of the ppn and MOS data with an absorbed power law yields a relatively high x? value [x2=1.21 for 673 degrees of freedom (dof)] and structured residuals, while a better fit o—1.05 for 673 dof) is obtained using a blackbody model (corrected for the absorption)."," A fit of the pn and MOS data with an absorbed power law yields a relatively high $\chi^2$ value $\chi^2_\nu=1.21$ for 673 degrees of freedom (dof)] and structured residuals, while a better fit $\chi^2_\nu=1.05$ for 673 dof) is obtained using a blackbody model (corrected for the absorption)."
 No additional spectral components are statistically required., No additional spectral components are statistically required.
 The results of this simultaneous modelling are presented in Table (the values of the spectral parameters were not significantly different when each observation was fit separately)., The results of this simultaneous modelling are presented in Table (the values of the spectral parameters were not significantly different when each observation was fit separately).
" The best-fitting hydrogen column density is (10.4+0.2)x10?? cm? (here and in the following uncertainties are at lo confidence level, unless otherwise noted) and the blackbody temperature is consistent within the errors with being constant at kT’~1.2 keV, while the observed flux slightly (but significantly, see Table 2)) decreased from Fc3.9x10? tto 3.1x107? (in the 2-10 keV band) during? the ~20 days that separate the first and last ppointings."," The best-fitting hydrogen column density is $(10.4\pm0.2)\times10^{22}$ $^{-2}$ (here and in the following uncertainties are at $\sigma$ confidence level, unless otherwise noted) and the blackbody temperature is consistent within the errors with being constant at $kT\simeq1.2$ keV, while the observed flux slightly (but significantly, see Table \ref{fits}) ) decreased from $F\simeq3.9\times10^{-12}$ to $3.1\times10^{-12}$ (in the 2–10 keV band) during the $\sim$ 20 days that separate the first and last pointings."
" Assuming an arbitrary distance of 10 kpc, the corresponding mmaximum and minimum luminosities are ~9x10? aand ~7x10?4 ((Table 2))."," Assuming an arbitrary distance of 10 kpc, the corresponding maximum and minimum luminosities are $\sim$$9\times10^{34}$ and $\sim$$7\times10^{34}$ (Table \ref{fits}) )."
 For each oobservation we performed phase-resolved spectroscopy by extracting the spectra for different selections of phase intervals., For each observation we performed phase-resolved spectroscopy by extracting the spectra for different selections of phase intervals.
" No significant variations with phase were detected, all the spectra being consistent with the model and parameters of the phase-averaged spectrum, simply re-scaled in normalisation."," No significant variations with phase were detected, all the spectra being consistent with the model and parameters of the phase-averaged spectrum, simply re-scaled in normalisation."
" Most magnetars, especially when in outburst, exhibit more complex spectra, usually fit by the superposition of two/three blackbodies or a blackbody and a high-energy power law (e.g. Mereghettietal.2005;Esposito2009;Bernardini2009,2011;Israeletal.2007a;Enoto2009;Rea 2009))."," Most magnetars, especially when in outburst, exhibit more complex spectra, usually fit by the superposition of two/three blackbodies or a blackbody and a high-energy power law (e.g. \citealt{mte05short,esposito09short,bernardini09short,bernardini11short,icd07,enoto09short,rea09short}) )."
 The spectrum of lis reminiscent of that reported by Espositoetal.(2010a) for 00418+5729 during its 2009 outburst., The spectrum of is reminiscent of that reported by \citet{esposito10short} for 0418+5729 during its 2009 outburst.
" However, at the time of the analysis reported in Espositoetal.(2010a),, only low-counts-statistics aand sspectra were available for SGR00418+5729."," However, at the time of the analysis reported in \citet{esposito10short}, only low-counts-statistics and spectra were available for 0418+5729."
" Higher quality ddata, now public, clearly show that a more complex model is required for the emission of In order to achieve better statistics and higher signal to noise, we merged the data from the first two oobservations (since there is no evidence for variations in the flux and spectrum of the source between them) and accumulated a combined spectrum which is presented in Fig. 2.."," Higher quality data, now public, clearly show that a more complex model is required for the emission of In order to achieve better statistics and higher signal to noise, we merged the data from the first two observations (since there is no evidence for variations in the flux and spectrum of the source between them) and accumulated a combined spectrum which is presented in Fig. \ref{spec}."
" Again, the single-blackbody model provides an excellent fit (x2=1.00 for 352 dof), and no additional components are required."," Again, the single-blackbody model provides an excellent fit $\chi^2_\nu=1.00$ for 352 dof), and no additional components are required."
" By including in the model a power-law component with photon index fixed to 3 (see e.g. Reaetal. 2009)), we can set 3c upper limits of 22596 and 823096 on the contribution of this component to the total observed and unabsorbed fluxes, respectively."," By including in the model a power-law component with photon index fixed to 3 (see e.g. \citealt{rea09short}) ), we can set $3\sigma$ upper limits of $\approx$ and $\approx$ on the contribution of this component to the total observed and unabsorbed fluxes, respectively."
 This shows that the relatively low flux and high absorption of mmake even our deep oobservations not very sensitive to the presence of a second spectral component in this SGR., This shows that the relatively low flux and high absorption of make even our deep observations not very sensitive to the presence of a second spectral component in this SGR.
 We repeated the phase-resolved spectral analysis on the combined dataset., We repeated the phase-resolved spectral analysis on the combined dataset.
" Once more, we see no evidence"," Once more, we see no evidence"
very low mass trausfer rates <1033 xy4 (Prialuik&Ikovetz 1995).,very low mass transfer rates $<10^{-11}$ $^{-1}$ \citep{PK95}.
. For mass transfer rates above this hreshold. we interpolateover the results of Prialuik&I&ovetz(1995). to obtain the mass accretion efficiencies. and we account for optically thick winds at high mass rausfer rates (Ilachisuetal.1999.seealsoBelcezvu-aileddescription)...," For mass transfer rates above this threshold, we interpolateover the results of \citet{PK95} to obtain the mass accretion efficiencies, and we account for optically thick winds at high mass transfer rates \citep[][see also Belczynski, Bulik \& Ruiter (2005)
 Section 2.3 for a more detailed description]{HKN99}."
" The ouly difference between the accumulation effcieucies in this work aud those of Delezvuskietal.(2005) is that here we have additionally included an updated prescription for accretion of ivdrogeu-rich matter frou, Nomotootal.(2007).. in which fully efficient accumulation is oulv achieved for a vorv narrow range of nass transfer rates. and is also dependent upon the white dwarf accretor mass (Nomotoetal.2007.sceequations5& 6).."," The only difference between the accumulation efficiencies in this work and those of \citet{BBR05} is that here we have additionally included an updated prescription for accretion of hydrogen-rich matter from \citet{Nom07}, in which fully efficient accumulation is only achieved for a very narrow range of mass transfer rates, and is also dependent upon the white dwarf accretor mass \citep[][see equations 5 \& 6]{Nom07}."
 Accretion of heliuu-rich matter is treated in the sale fashion as in Belezvuskietal.(2005): accretion prescriptions are adopted from Kato&Iachisu(1999)., Accretion of helium-rich matter is treated in the same fashion as in \citet{BBR05}: accretion prescriptions are adopted from \citet{KH99}.
.. However this study contains one major difference which affects our results: though in this work we do allow for the formation of sub-Chancdrasclhar mass SNe Ia. we do uot include these binaries as potential SN Ia progenitors here.," However this study contains one major difference which affects our results: though in this work we do allow for the formation of sub-Chandrasekhar mass SNe Ia, we do not include these binaries as potential SN Ia progenitors here."
 In Belezvuskietal.(2005).. a large fraction (61% )) of the Type Ia SNe which contributed to the presented delay times iu the standard model were in factChandrauschhar SNe Ta. in which the accumulation of 0.1 AL. of We-vich material on the WD surface could lead to an edee-lit detonation and subsequeut SN Ta (Woosley&Weaver1991:katoIIachisu1999).," In \citet{BBR05}, a large fraction $61$ ) of the Type Ia SNe which contributed to the presented delay times in the standard model were in fact SNe Ia, in which the accumulation of $\sim 0.1$ $_{\odot}$ of He-rich material on the WD surface could lead to an edge-lit detonation and subsequent SN Ia \citep{WW94,KH99}."
. The criteria used here for defining SNe Ia progenitors arising from different formation chanucls are different from the work of Delezvusiàietal.(2005)., The criteria used here for defining SNe Ia progenitors arising from different formation channels are different from the work of \citet{BBR05}.
. Iu this work. we only consider accreting WDs which have obtained a duas of d.d AL. as potential SNe Ta. We iatke the distinction between SN Ia progenitors with CO white dwarfs accreting from non-degenerate lyvclroee-vich companions (SDS). aud white dwarfs accreting from helimm-vich companions (AM CV).," In this work, we only consider accreting WDs which have obtained a mass of 1.4 $_{\odot}$ as potential SNe Ia. We make the distinction between SN Ia progenitors with CO white dwarfs accreting from non-degenerate hydrogen-rich companions (SDS), and white dwarfs accreting from helium-rich companions (AM CVn)."
 We note that Te WDs never reach the Chandrascklar mass in our sinulatious with the adopted accumulation plysics. aud we have assumedthat oxveeu-uceon-nmaeguesmuu WDs collapse to forma a neutron star (accretion induced collapse) upon reaching the Chandrasekhar mass. rather than producing a SN Ia (e...Mivajretal.1980).," We note that He WDs never reach the Chandrasekhar mass in our simulations with the adopted accumulation physics, and we have assumedthat oxygen-neon-magnesium WDs collapse to form a neutron star (accretion induced collapse) upon reaching the Chandrasekhar mass, rather than producing a SN Ia \citep[e.g.,][]{Miy80}."
. For the SDS aud AAT CVn cases; we record the binary às à SN Ia ouce the accretiug WD has reached the Chandrasekhar mass.," For the SDS and AM CVn cases, we record the binary as a SN Ia once the accreting WD has reached the Chandrasekhar mass."
" The SDS may occur when a WD accretes matter via RLOF from any bydrogen-vich companion (οιο, mail sequence or evolved star)."," The SDS may occur when a WD accretes matter via RLOF from any hydrogen-rich companion (e.g., main sequence or evolved star)."
 I£ the WD acciuulates enough lwdrogeu on its surface such that steady Durning cau occur. the WD can increase in mass up to 1.1 M. carbon ignites explosively in the WD centre aud the result is a SN Ia. In the AM CVn scenario. we assume the result is a SN Ta if the CO WD reaches 1.1 M. via stable RLOF from) a helimm-rich companion.," If the WD accumulates enough hydrogen on its surface such that steady burning can occur, the WD can increase in mass up to 1.4 $_{\odot}$, carbon ignites explosively in the WD centre and the result is a SN Ia. In the AM CVn scenario, we assume the result is a SN Ia if the CO WD reaches 1.4 $_{\odot}$ via stable RLOF from a helium-rich companion."
 We allow for different types of heliuu-rich. donors in the AM. ΟΝα scenario: helium stars (non-degenerate stars burning helium iu the core or m a shell which have been stripped of their outer hydrogen envelope). helm white dwarts. and bybrid white dwarfs (CO-rich core. leliumerich mantle).," We allow for different types of helium-rich donors in the AM CVn scenario: helium stars (non-degenerate stars burning helium in the core or in a shell which have been stripped of their outer hydrogen envelope), helium white dwarfs, and hybrid white dwarfs (CO-rich core, helium-rich mantle)."
 We adopt two contrasting star formation rates (SFR): an dmstantaneous burst at f=0 (elliptical galaxw) and a coustaut SER for 10 Cr (spiral galaxv)., We adopt two contrasting star formation rates (SFR): an instantaneous burst at $t=0$ (elliptical galaxy) and a constant SFR for 10 Gyr (spiral galaxy).
 Both populations are then evolved up to 15 Car., Both populations are then evolved up to 15 Gyr.
" The mass formed in stars in both cases is the same: 6&101"" AL... which corresponds to the stellar 1iass in the Milky Wav (MW:Klvpiunetal.2002)."," The mass formed in stars in both cases is the same: $6 \times 10^{10}$ $_{\odot}$, which corresponds to the stellar mass in the Milky Way \citep[MW;][]{KZS02}."
. In cach population we adopt a binary fraction of 50% (2/3 stars in dnuaries). though this fraction may be overestimating he binary population :unong low-mass stars (Lada2006).. aud. underestimating the binary fraction among uassive stars (Il&obuluickv&Fryer2007).," In each population we adopt a binary fraction of $50\%$ (2/3 stars in binaries), though this fraction may be overestimating the binary population among low-mass stars \citep{Lad06}, and underestimating the binary fraction among massive stars \citep{KF07}."
. All stars are evolved with solar-like metallicity (Z=0.02)., All stars are evolved with solar-like metallicity $Z=0.02$ ).
 While he initial distributions representative of the phvsical characteristics of ZAMS binaries. aud the correct wav in which to treat communon envelope evolution and nagnetic braking are all somewhat uucertaim. our clioices Or Various distribution fuuctions are constrained by available observations (see below).," While the initial distributions representative of the physical characteristics of ZAMS binaries, and the correct way in which to treat common envelope evolution and magnetic braking are all somewhat uncertain, our choices for various distribution functions are constrained by available observations (see below)."
 The magnetic braking xescription which we adopt is that of Ivanova&Taam(2003).. which is based on a two-component coronal nodel (seeSection3.2ofBelezvuskietal. 2008)..," The magnetic braking prescription which we adopt is that of \citet{IT03}, which is based on a two-component coronal model \citep[see Section 3.2 of][]{Bel08}. ."
 ZAMS inasses (Afzayis) span the ecutive mass range: LOS.150AL...," ZAMS masses $M_{\rm ZAMS}$ ) span the entire mass range: $0.08-150
\msun$ ."
 Sinele stars aud binary primarics are drawn froma 3-component broken power law initial mass πιοΊο (Ixroupaetal.1993).. and secondary masses are obtained from a flat mass ratio distribution (Alazehetal.(1992).. ¢=sccondary /primary). which is the canonical choice among population svuthesis studies (Litetal.m106).," Single stars and binary primaries are drawn from a 3-component broken power law initial mass function \citep{KTG93}, and secondary masses are obtained from a flat mass ratio distribution \citet{Maz92}, $q=$ secondary/primary), which is the canonical choice among population synthesis studies \citep{LYH06}."
. Towever. given the observational selection effects. je true mass ratio distribution amoug ZAAIS binarics remains unknown. though it is likely dependent upon stellar mass (Trimble1990:Duqueunov&Mayor1991).," However, given the observational selection effects, the true mass ratio distribution among ZAMS binaries remains unknown, though it is likely dependent upon stellar mass \citep[][]{Tri90,DM91}."
. It has been suggested that the mass ratio distribution among local spectroscopic field binaries as well as vouug carly-type stars is peaked near unity (Fisheretal.2005:Ikobuluickv&Frver 2007).," It has been suggested that the mass ratio distribution among local spectroscopic field binaries as well as young early-type stars is peaked near unity \citep{FSS05,KF07}."
. Initial orbital separations m our caleulatious span a wide range up to LO’ Rand are drawn from a distribution xο (Abt1983).. which lias been found to be representative of the local population of binarics (Lépine&Bougiorno2007).," Initial orbital separations in our calculations span a wide range up to $10^{5}$ $_{\odot}$ and are drawn from a distribution $\propto 1/a$ \citep{Abt83}, which has been found to be representative of the local population of binaries \citep{LB07}."
. Initial eccentrieities are drawn from a distribution Wile)=2c (Duqueunoy&Mavor1991)., Initial eccentricities are drawn from a distribution $\Psi(e)=2e$ \citep{DM91}.
. Close binaries (and potential SN Ia progenitors) are believed to. encounter a. phase. of conunon: envelope evolution. aud this remains one of the most poorlv-nuderstood phases in stellar astroplivsics.," Close binaries (and potential SN Ia progenitors) are believed to encounter a phase of common envelope evolution, and this remains one of the most poorly-understood phases in stellar astrophysics."
 For this reason. we present the delay times aud rates of SNe Ia for wo different CE paraimcterizations.," For this reason, we present the delay times and rates of SNe Ia for two different CE parameterizations."
 Some comparison vetween different prescriptions for CE evolution have )een tested against observations for local double white chwarts (Nelemanus&Tout2005)., Some comparison between different prescriptions for CE evolution have been tested against observations for local double white dwarfs \citep{NT05}.
. It is uuclear at this »m:t how the common envelope phase should best be reated in population svuthesis studies; aud curreutly detailed models are not sophisticated enough toexplore he parameter space in detail (Ricker&Taam2008).," It is unclear at this point how the common envelope phase should best be treated in population synthesis studies, and currently detailed models are not sophisticated enough toexplore the parameter space in detail \citep{RT08}."
. For this work. we choose the o prescription for CE evolution (Webbink1981).. in which the orbital energv of the binary is used to uubind the couuuon euvelope from the svstem.," For this work, we choose the $\alpha$ ' prescription for CE evolution \citep{Web84}, in which the orbital energy of the binary is used to unbind the common envelope from the system."
 We clioose two different. values for the paralucterization of common envelope removal efficiency acq., We choose two different values for the parameterization of common envelope removal efficiency $\alpha_{\rm CE}$ .
 For our standard model (Model 1) we choose Oop\A= d. where A ds a function of the donor euvelope structure. and is of order unitv (seoalso 2006)..," For our standard model (Model 1) we choose $\alpha_{\rm CE} \times \lambda = 1$ , where $\lambda$ is a function of the donor envelope structure, and is of order unity \citep[see also][]{vVP06}. ."
 As an alternative. we additionally iuelude a model with decreasedcommon euveloperemoval cficiency in which ace.A=0.5 ," As an alternative, we additionally include a model with decreasedcommon enveloperemoval efficiency in which $\alpha_{\rm CE} \times \lambda = 0.5$ "
The measured shear at each source plane becomes. with a noise teri now added. Now 6) is the mass overdensity averaged through the shell of width Ay: Ay is a (actor that couverts this overdensity into the convergence measured on a leus plane at infinity.,"The measured shear at each source plane becomes, with a noise term now added, Now $m_\ell(\bft)$ is the mass overdensity averaged through the shell of width $\Delta \chi$; $k_\ell$ is a factor that converts this overdensity into the convergence measured on a lens plane at infinity."
 The shear measurement noise is cletermined by the areal density (4 of sliape-mieasurable galaxies ou shell s. aud σ.220.32005).," The shear measurement noise is determined by the areal density $n_s$ of shape-measurable galaxies on shell $s$, and $\sigma_\gamma\approx 0.3$."
 The true mass fluctuation (6) is not directly observable., The true mass fluctuation $m_\ell(\bft)$ is not directly observable.
 But we assume that we deduce [rom the galaxy. distribution some estimate of the mass overceusity 6) that is correlated. with the true distribution., But we assume that we deduce from the galaxy distribution some estimate of the mass overdensity $g_\ell(\bft)$ that is correlated with the true distribution.
 This proxy. field need not be the galaxy density itself. but rather the result of any algorithin applied to the observed galaxy. fekl.e.g. involving the assigniuent of lialos to ealaxies. groups. and clusters.," This proxy field need not be the galaxy density itself, but rather the result of any algorithm applied to the observed galaxy field, involving the assignment of halos to galaxies, groups, and clusters."
 The correlation coefficient of the true auc proxy inass fields is where the angle brackets deuote averaging over direction0., The correlation coefficient of the true and proxy mass fields is where the angle brackets denote averaging over direction.
 The r; are not kuown but are assumed positive., The $r_i$ are not known but are assumed positive.
 We assume that the shells are thick enough to be uuncorrelated., We assume that the shells are thick enough to be uncorrelated.
 We also clefine for each shell an uuknown bias via Note that this is the bias of the mass withrespect to theestimated mass., We also define for each shell an unknown bias via Note that this is the bias of the mass withrespect to the mass.
 The more conventional bias 6 of the estimated(galaxy) mass to the true mass is. for Caussiau-cistributecd g anda. related to this via The shear noise df is assumed to have correlations with neither the true nor the proxy 1uass distribution.," The more conventional bias $b$ of the estimated(galaxy) mass to the true mass is, for Gaussian-distributed $g$ and $m$, related to this via The shear noise $\delta\kappa$ is assumed to have correlations with neither the true nor the proxy mass distribution."
 With these definitions we could extract curvature information from our observed shear data as follows: we cross-correlate each couvergeuce field &; with each foreground proxy field gy to give ai observable quantity γι, With these definitions we could extract curvature information from our observed shear data as follows: we cross-correlate each convergence field $\kappa_s$ with each foreground proxy field $g_\ell$ to give an observable quantity $X_{s\ell}$ .
 From the expectation values are, From the expectation values are
The curve hits the j£ axis al 7=0 when ji=jt.,The curve hits the $\mu$ axis at $T=0$ when $\mu=\mu_0$.
 The chemical potential at (he critical point is fl —jt hence ," The chemical potential at the critical point is _c = _0, hence f_1(t) = n_c _0."
It is apparent [rom the above expression For (1) that this whole parameterization is only good when T«Ti. otherwise /1(0) would become imaginary.," It is apparent from the above expression for $f_1(t)$ that this whole parameterization is only good when $T < T_0$, otherwise $f_1(t)$ would become imaginary."
 Knowing this. it is straightforward to derive a simple formula for the latent heat per unit volume.," Knowing this, it is straightforward to derive a simple formula for the latent heat per unit volume."
 Making use of phase equilibrium. AP=0. THT). pn-—fly gq ου Seti) = which of course is valid onlv lor /«0.," Making use of phase equilibrium, $\Delta P = 0$, $T = T_x$, $\mu = \mu_x$, $\eta_g = - \eta_l$, one finds (t) = which of course is valid only for $t < 0$."
 Theindependent constants max be taken as; aδνfaepto.pte.dons. plus the functions fo(/) and. f5(/).," Theindependent constants may be taken as: $a_-, b_-, f_{\sigma}, \mu_0, \mu_c, T_c, n_c$, plus the functions $\bar{f}_0(t)$ and $\bar{f}_2(t)$."
 This parameterization captures the critical exponents and amplitudes which are universal., This parameterization captures the critical exponents and amplitudes which are universal.
 To proceed further we require more information on (he equation of state of QCD away [rom the critical point., To proceed further we require more information on the equation of state of QCD away from the critical point.
 For future reference it may be noted that one could add more terms to the free energy. without changing the basic picture., For future reference it may be noted that one could add more terms to the free energy without changing the basic picture.
 For example. one could add fiU)! and fA).," For example, one could add $f_4(t) \eta^4$ and $f_8(t) \eta^8$."
 However. f;0) must vanish al /=0 so as not to allect the critical behavior.," However, $f_4(t)$ must vanish at $t=0$ so as not to affect the critical behavior."
" The coefficient. f/;(/) need not vanishal /=0 but it becomes irrelevant compared to the dominant term f,|77|7 as η— 0. as do all powers of 7 greater than σ."," The coefficient $f_8(t)$ need not vanishat $t=0$ but it becomes irrelevant compared to the dominant term $f_{\sigma} |\eta|^{\sigma}$ as $\eta \rightarrow 0$ , as do all powers of $\eta$ greater than $\sigma$ ."
"the disk aud wind is B,,.. the azimuthal field streneth at the disk surface. which sets the torque exerted by the wind on the disk (BPs2).","the disk and wind is $B_{\phi,{\rm s}}$, the azimuthal field strength at the disk surface, which sets the torque exerted by the wind on the disk (BP82)."
" As with the surface fiek inclination angle. the value of DB,,, is determined b processes acting outside the disk: specifically. the value of D,,. is determined by the coustraünt that the wine solution pass znoothlv through the Alfvéóun point."," As with the surface field inclination angle, the value of $B_{\phi,{\rm s}}$ is determined by processes acting outside the disk; specifically, the value of $B_{\phi,{\rm s}}$ is determined by the constraint that the wind solution pass smoothly through the Alfvénn point."
 Accretion flows in a weaklv ionized disk will drag In nmaeguetie flux. opposing the natural tendency of the amaeguetic fux to diffuse outwards.," Accretion flows in a weakly ionized disk will drag in magnetic flux, opposing the natural tendency of the magnetic flux to diffuse outwards."
 Studies of accretion. disks that sled angular momentum due to AIITD. turbulence developed from the maguetorotationa instability CMBRI. Balbus&Hawley 19913) suggest tha the magnetic diffusivity aud the effective. viscosity of such disks are comparable (Camu&(ναι2009:Lesur&Lougaretti 2009).. aerecing with previousο work by Lubowetal.(199la) which suggested that the magnetic diffusivity of AIRactive acerction disks nay be too large o allow cfiicicut inward drageiug of maguctic flux.," Studies of accretion disks that shed angular momentum due to MHD turbulence developed from the magnetorotational instability (MRI, \citealt{bh91}) ) suggest that the magnetic diffusivity and the effective viscosity of such disks are comparable \citep{gg09, ll09}, , agreeing with previous work by \citet{lpp94} which suggested that the magnetic diffusivity of MRI-active accretion disks may be too large to allow efficient inward dragging of magnetic flux."
 This xobleuni can be overcome in au MBlL-active disk threaded oa large-scale field that removes angular ποιοται roni surface livers in which the MBI i$ suppressed (Rothstem&Lovelace2008:etal.2009).," This problem can be overcome in an MRI-active disk threaded by a large-scale field that removes angular momentum from surface layers in which the MRI is suppressed \citep{rl08, lrb09}."
. More generally. this problem disappears if the angular ποιοτι transport mechanisin is decoupled from the uechanisia that provides the magnetic diffusivitv so hat the diffusivity is much smaller than the effective viscosity of the disk.," More generally, this problem disappears if the angular momentum transport mechanism is decoupled from the mechanism that provides the magnetic diffusivity so that the diffusivity is much smaller than the effective viscosity of the disk."
 This occurs naturallv in a disk. roni Which angular momentum is removed through a uaenetocentrifusal wind. allowing magnetic flux to be drageed in and distributed over a wiud-driviug disk.," This occurs naturally in a disk from which angular momentum is removed through a magnetocentrifugal wind, allowing magnetic flux to be dragged in and distributed over a wind-driving disk."
 As noted bv OLOI. the evolution of the maguctic Hux through the disk determines whether magnetic configurations capable of launchiug outflows can be uaintained over the accretion timescale.," As noted by OL01, the evolution of the magnetic flux through the disk determines whether magnetic configurations capable of launching outflows can be maintained over the accretion timescale."
" Previous studies of wind-diiving disks have not treated the effects of naenetic flux evolutiou iu a selfconsistent nuiumner: iu a local model. OLOL caleulated the rate of flux migration eiven D,. aud δι. which were free parameters,"," Previous studies of wind-driving disks have not treated the effects of magnetic flux evolution in a self-consistent manner: in a local model, OL01 calculated the rate of flux migration given $B_{r,{\rm s}}$ and $B_{\phi,{\rm s}}$ , which were free parameters."
" The οσα] models of Campbell(2003.2005). expressed B.,.. and the surface field inclination angle iu terms of unconstrained paraicters that describe the structure of he disk and the processes operating within it."," The local models of \citet{ca03, ca05} expressed $B_{\phi,{\rm s}}$ and the surface field inclination angle in terms of unconstrained parameters that describe the structure of the disk and the processes operating within it."
" WIN93 also ciuploved a local model iu which they determined D,,. sclfconsisteutly using the Alfvéun point constraint. nit left the surface feld iuclination angle uncoustraimed. describing the magnetic flux mieration speed using the ree parainueter ep."," WK93 also employed a local model in which they determined $B_{\phi,{\rm s}}$ self-consistently using the Alfvénn point constraint, but left the surface field inclination angle unconstrained, describing the magnetic flux migration speed using the free parameter $\epsilon_{\rm B}$."
" The elobal models of L9G aud E97 also solved for B,,. using the Alfvénn point constraint nit left the surface field inclination angle aud maguetic Hux στα unconstrained. csscutially setting ευ=0."," The global models of L96 and F97 also solved for $B_{\phi,{\rm s}}$ using the Alfvénn point constraint but left the surface field inclination angle and magnetic flux migration unconstrained, essentially setting $\epsilon_{\rm B} = 0$."
 For an overviewο of recent work ou scuuanalytic models and complementary wuunerical smulatious of protostellar jets. see the article by I&ónigl&Sahucron(2011).," For an overview of recent work on semianalytic models and complementary numerical simulations of protostellar jets, see the article by \citet{ks11}."
". I present here a new, scnuanalytic. radially self-similar. elobal model that matches disk solutions with BPs2- wind solutions. adopting an approach similar to hat of L96 but including for the first timea fully selt-consistent calculation of the effects of magnetic flux wieration and the maguetie field structure above the disk: ep. DB... Bo. aud D,./D.. ave no longer free xuwanieters but are determined by plysical coustraiuts iuposed at the disk surface and at critical poiuts of the outflow."," I present here a new, semianalytic, radially self-similar, global model that matches disk solutions with BP82-type wind solutions, adopting an approach similar to that of L96 but including for the first timea fully self-consistent calculation of the effects of magnetic flux migration and the magnetic field structure above the disk: $\epsilon_{\rm B}$, $B_{r,{\rm s}}$, $B_{\phi, {\rm s}}$ and $B_{r,{\rm s}}/B_{z,{\rm s}}$ are no longer free parameters but are determined by physical constraints imposed at the disk surface and at critical points of the outflow."
 Iu addition. the equations describing the disk are ormulated im terms of the conductivity tensor. allowing or fully ecneral non-ideal disk models.," In addition, the equations describing the disk are formulated in terms of the conductivity tensor, allowing for fully general non-ideal disk models."
 The plan of the oper is as follows., The plan of the paper is as follows.
 In Section ?7. I review the equatious describing a (uearly) steady-state; axisviunetric. weaklv ionized disk threaded by a imagnuetie field. aud I show row the system reduces to a set of ordinary differential equations under the radial sclfsimilarity asswuption.," In Section \ref{form} I review the equations describing a (nearly) steady-state, axisymmetric, weakly ionized disk threaded by a magnetic field, and I show how the system reduces to a set of ordinary differential equations under the radial self-similarity assumption."
 Iu Section ??7. T discuss coustraints from the selt-consisteut reatinent of the field-line inchuation at the top of the disk and from matching disk and wind solutions., In Section \ref{paramconstr} I discuss constraints from the self-consistent treatment of the field-line inclination at the top of the disk and from matching disk and wind solutions.
 Iu Section 77. T present a representative matched disk/wiud solution. aud in Section 77. I sumunarize the study aud outline plans to iuprove on the curent model.," In Section \ref{res} I present a representative matched disk/wind solution, and in Section \ref{concs} I summarize the study and outline plans to improve on the current model."
 I consider a weakly ionized. geometrically thin. vertically isothermal axisviunetric disk in nem-Iseplerian rotation around a protostar of mass AL that is threaded by an ordered. large-scale imuagnetic field.," I consider a weakly ionized, geometrically thin, vertically isothermal, axisymmetric disk in near-Keplerian rotation around a protostar of mass $M$ that is threaded by an ordered, large-scale magnetic field."
 I examine the general. non-kdeal ΑΠΟ case (Section ??)). expressing the matterfield interaction using the conductivity tensor formalisin. aud seek solutions that are steady on timescales shorter than the accretion time (Section ?7?7)).," I examine the general, non-ideal MHD case (Section \ref{casegen}) ), expressing the matter–field interaction using the conductivity tensor formalism, and seek solutions that are steady on timescales shorter than the accretion time (Section \ref{goveqgen}) )."
 I follow L96 aud FO? in imakine the simplitving assumption of radial seltsnailaritv to produce a elobal disk/wind model (Section ??)). aud I derive the midplane values of various plivsical quantities (Section ??)).," I follow L96 and F97 in making the simplifying assumption of radial self-similarity to produce a global disk/wind model (Section \ref{selfgen}) ), and I derive the midplane values of various physical quantities (Section \ref{midgen}) )."
 Iu Section ?? I specialize to the case of the aiubipolar conductivity regime., In Section \ref{caseAD} I specialize to the case of the ambipolar conductivity regime.
 T once again derive the relevaut sclsiuiar disk equations (Section 27)) and the midplane values for the variables of iutegration (Section ?7))., I once again derive the relevant self-similar disk equations (Section \ref{selfAD}) ) and the midplane values for the variables of integration (Section \ref{midAD}) ).
 The interiors of protostellar disks have a low ionization action. uecessitating the use of non-ddeal ΑΠΌ to escribe the behavior of the disk material iu the presence of a large-scale magnetic field that is assumed to thread he disk.," The interiors of protostellar disks have a low ionization fraction, necessitating the use of non-ideal MHD to describe the behavior of the disk material in the presence of a large-scale magnetic field that is assumed to thread the disk."
 The uon-ideal matter-field interaction can be escribed using a iulti-cuid approach or by introducimg he conductivity tensor (os.Nouiel&Salineron2011).. which eucapsulates the respouse of the dominant neutral Huid to the presence of the maguetic field.," The non-ideal matter-field interaction can be described using a multi-fluid approach or by introducing the conductivity tensor \citep[e.g.,][]{ks11}, which encapsulates the response of the dominant neutral fluid to the presence of the magnetic field."
 The effects ft the maenetic field are communicated to the neutrals w the development of chargeneutral dift velocities ue to the Loreutz forces ou charged species. aud the consequeut svstematic charge.neutral drag forces.," The effects of the magnetic field are communicated to the neutrals by the development of charge–neutral drift velocities due to the Lorentz forces on charged species, and the consequent systematic charge–neutral drag forces."
" Under he conductivity tensor formalisii. this is expressed iu the ecucralized version of Olini's law. where E; is the electric field ina frame comoving with the dominant neutral fluid. aud Ee. are the components of Ee parallel aud perpendiculayEy to the magnetic field. respectively, and ois the conductivity tensor."," Under the conductivity tensor formalism, this is expressed in the generalized version of Ohm's law, where $E$ $_{\rm c}$ is the electric field in a frame comoving with the dominant neutral fluid, $E$ $_{\rm c\parallel}$ and $E$ $_{\rm c\perp}$ are the components of $E$ $_{\rm c}$ parallel and perpendicular to the magnetic field, respectively, and $\sigma$ is the conductivity tensor."
 The colmponcnts of gare the Olun conductivity o. the Tall conductivity oq. and the Pedersen conductivity op.," The components of $\sigma$are the Ohm conductivity $\sigma_{\rm O}$ , the Hall conductivity $\sigma_{\rm H}$, and the Pedersen conductivity $\sigma_{\rm P}$."
 For a giveu charged species (denoted by subscript 7j). the product of. the evrofrequency. (Enn)ZjeB and the neutral-charge moieutuii exchaugetime 1/(5;p) gives a nueasure of the coupling of the charged species to the," For a given charged species (denoted by subscript $j$ '), the product of the gyrofrequency $\left(\frac{|Z_{j}|e|B|}{m_{j}c}\right)$ and the neutral-charge momentum exchangetime $1/\left(\gamma_{j}\rho\right)$ gives a measure of the coupling of the charged species to the"
"where Q is the effective cross section of the inass transfer throat at the Lagrangian point L4. Q=1.9x4107TT,hyD cnr.U where Ty is. the surfaceB Τοράνο ofn the secondary aud Pu, the orbital period iu hours: ὃς is the sound speed. and prj the density at £4. which. in the case of an isothermal atmosphere. cam be expressed as: where pg is the deusity calculated at a reference level. HR the secoudary radius (defined by this reference position). Ry, the Roche lobe radius and Jf the scale height.","where $Q$ is the effective cross section of the mass transfer throat at the Lagrangian point $L_1$, $Q = 1.9 \times 10^{17} T_4 P_{\rm hr}^2$ $^2$, where $T_4$ is the surface temperature of the secondary and $P_{\rm hr}$ the orbital period in hours; $c_s$ is the sound speed, and $ \rho_{L1}$ the density at $L_1$, which, in the case of an isothermal atmosphere, can be expressed as: where $\rho_0$ is the density calculated at a reference level, $R$ the secondary radius (defined by this reference position), $R_{L1}$ the Roche lobe radius and $H$ the scale height."
" Iu quiescence. the mass transfer rate is low. of order of 107το ο 5,"," In quiescence, the mass transfer rate is low, of order of $10^{15} - 10^{16}$ g $^{-1}$."
" At short orbital periods aud correspondingly low secondarys temperatures, both Q and e. are sul."," At short orbital periods and correspondingly low secondary's temperatures, both $Q$ and $c_s$ are small."
 If one assuned that pg correspouds to the Xiotospherie density. which for a low ass main sequence star is of the order of 10/7 oο 5m. one would fud AL.~10067Πε v 1," If one assumed that $\rho_0$ corresponds to the photospheric density, which for a low mass main sequence star is of the order of $10^{-5}$ g $^{-3}$, one would find $\dot{M}_{\rm tr}
\sim 10^{16} e^{(R-R_{L1})/H}$ g $^{-1}$."
" This in ftum would mean that (RoRey)fff is μαμα so that A. varies as 262TU. aud therefore. is. uot very sensitiveMp to Πα.τοι,"," This in turn would mean that $(R-R_{L1})/H$ is small so that $\dot{M}_{\rm tr}$ varies as $T_4^{3/2}$, and therefore is not very sensitive to illumination."
 Iu the mradiated case. however. the reference density pp does not correspond to the photosphlieric density but to the density at the base of the isothermal atiuosphliere. which extends much deeper than the unilliuninated plotosphere when the atimosphere is idhuninated with an mradiation eniperature excecdine 101 K. For those high iuuination fluxes. the outer lavers of the star are affected on a hermal time scale (seconds) at least down to a poiut where the unperturbed temperature equals the inradiation eni])orature.," In the irradiated case, however, the reference density $\rho_0$ does not correspond to the photospheric density but to the density at the base of the isothermal atmosphere, which extends much deeper than the unilluminated photosphere when the atmosphere is illuminated with an irradiation temperature exceeding $^4$ K. For those high illumination fluxes, the outer layers of the star are affected on a thermal time scale (seconds) at least down to a point where the unperturbed temperature equals the irradiation temperature."
 This point is the base of the isotlermal avers in the iluuinated case. aud from models of very ow mss stars (Dorian ct al. L989)).," This point is the base of the isothermal layers in the illuminated case, and from models of very low mass stars (Dorman et al. \cite{dnc89}) ),"
" one gets py~10.7> e cnn7.m with resulting very high mass transfer rates. exceeding 1025 est,"," one gets $\rho_0 \sim 10^{-3}$ g $^{-3}$, with resulting very high mass transfer rates, exceeding $^{18}$ g $^{-1}$."
 The latter estimate d8 al upper Dui. as it docs not take iuto account the fact that a fraction of the secondary is shielded from irradiation by the accretion disc: partial shiclding docs not suppress the cuhancement of mnass trausfer. since circulation at the surface of the star prevents the existence of large temperature eradieuts. but reduces it in à complex wav which we are not attempting to describe here.," The latter estimate is an upper limit, as it does not take into account the fact that a fraction of the secondary is shielded from irradiation by the accretion disc; partial shielding does not suppress the enhancement of mass transfer, since circulation at the surface of the star prevents the existence of large temperature gradients, but reduces it in a complex way which we are not attempting to describe here."
" Observations show. however. that the increase cau be quite significant: the mass trausfer rate rises by a factor 2 im aud (Sinak 1995)). whereas Vost (1983)) fuds that inναι, the bright spot luminosity increases by a factor ~ 15 curving maxima aud decline of outbursts close to a superoutburst. which le attributed to a corresponding mnerease nm nlass transfer uuder the effect of illumination."," Observations show, however, that the increase can be quite significant; the mass transfer rate rises by a factor 2 in and (Smak \cite{s95}) ), whereas Vogt \cite{v83}) ) finds that in, the bright spot luminosity increases by a factor $\sim$ 15 during maximum and decline of outbursts close to a superoutburst, which he attributed to a corresponding increase in mass transfer under the effect of illumination."
 Although the evidence is not very strone. there are some indications that the hot spot brightcnine. aud heuce the mass transfer inercase. is cdelaved with respect to the eruption bv a dav or two: this could be either the response time ofthe secondary (Simals 1995)). or the thermal inertia of the white dwarf (ouly an equatorial belt is instantaneously heated by accretiou. so irradiation of the secondary is delaved).," Although the evidence is not very strong, there are some indications that the hot spot brightening, and hence the mass transfer increase, is delayed with respect to the eruption by a day or two; this could be either the response time of the secondary (Smak \cite{s95}) ), or the thermal inertia of the white dwarf (only an equatorial belt is instantaneously heated by accretion, so irradiation of the secondary is delayed)."
 There is evideuce that in dwarf novae accretion disces aro truncated. as incicated by cussion line profiles (Menunikeut Arenas 1998)). or the detection of a sjenificaut quiesceunt X-ray and UV flux (Lasota 1996)).," There is evidence that in dwarf novae accretion discs are truncated, as indicated by emission line profiles (Mennikent Arenas \cite{ma98}) ), or the detection of a significant quiescent X-ray and UV flux (Lasota \cite{l96}) )."
" Tf this is due to the presence of a magnetic field. the ner disc radius rg, is a siuple function of the massaccretion rate onto the white dwarf: where My is the mass accretion rate m units of D es 4) AL, is the white cawarf mass and piggy is the ος moment in units of 1079. Cem?."," If this is due to the presence of a magnetic field, the inner disc radius $r_{\rm in}$ is a simple function of the massaccretion rate onto the white dwarf: where $\dot{M}_{15}$ is the mass accretion rate in units of $^{15}$ g $^{-1}$, $M_1$ is the white dwarf mass and $\mu_{30}$ is the magnetic moment in units of $^{30}$ $^3$."
" The value of 1 should be such as to allow 57g,=Ry in outbursts where 415 the primary radius. as most DNs do not show then colereut pulsations."," The value of $\mu_{30}$ should be such as to allow $r_{\rm in} = R_1$ in outbursts, where $R_1$ is the primary radius, as most DNs do not show then coherent pulsations."
 Iu quiescence. however. cohereut oscillatious are observed (Patterson et al. 1998)).," In quiescence, however, coherent oscillations are observed (Patterson et al. \cite{prkm98}) ),"
 and for example in μου© 50 (Lasota ct al. 1999))., and for example in $\mu_{30} \approx$ 50 (Lasota et al. \cite{lkc99}) ).
 Au inner hole iu the disc can be also due to evaporation., An inner hole in the disc can be also due to evaporation.
 The plivsics οἳ evaporation 1 poorly understood but several models were proposed (Mover Alever-Iofiueister 199L. Liu et al. 1997..," The physics of evaporation is poorly understood but several models were proposed (Meyer Meyer-Hofmeister \cite{mm94}, Liu et al. \cite{lmm97},"
 Kato Nakamura 1998.. Shaviv ct al. 1999)).," Kato Nakamura \cite{kn98}, Shaviv et al. \cite{sww99}) )."
 The evaporation rates X du the disc are. however. quite uncertain.," The evaporation rates $\dot{\Sigma}$ in the disc are, however, quite uncertain."
 Evaporation is normally accounted for by introducing the additional tex X iu the mass conservation equation., Evaporation is normally accounted for by introducing the additional term $\dot{\Sigma}$ in the mass conservation equation.
 However. because evaporation is expected to iucrease owards the accreting body and since the local mass ransfer rate in the disc increases sharply with radius during quiescence. he effects of evaporation are iniportaut essentiallv very close to the dise immer edge. aud can be reated assunnuime that the disc mner radius is a function of the accretion rate outo the white dwarf. just as in the case of the formation of a magnetosphere.," However, because evaporation is expected to increase towards the accreting body, and since the local mass transfer rate in the disc increases sharply with radius during quiescence, the effects of evaporation are important essentially very close to the disc inner edge, and can be treated assuming that the disc inner radius is a function of the accretion rate onto the white dwarf, just as in the case of the formation of a magnetosphere."
" To first order. he effect of evaporation is to create a hole in quiescence which increases the recurrence time: what matters is thus he immer disc radius in quiescence aud not the detailed Way ry, Varies with M."," To first order, the effect of evaporation is to create a hole in quiescence which increases the recurrence time; what matters is thus the inner disc radius in quiescence and not the detailed way $r_{\rm in}$ varies with $\dot{M}_{\rm acc}$."
 We shall. therefore use equation (7)) in all cases. Using Hag asa free. wncoustrained parameter that merely describes the size of the hole generated in the disc by either the presence of as magnetic field. or by evaporation.," We shall therefore use equation \ref{eq:rin}) ) in all cases, using $\mu_{30}$ as a free, unconstrained parameter that merely describes the size of the hole generated in the disc by either the presence of as magnetic field, or by evaporation."
" Iu the following. we discuss the iufluence of cach individual effect mentioned above: our reference situation is that of a system with a 1.0 M. primary. whose radius is 5ονqusI0? cn: the mass trauster. rate is. Bstole st, and the average outer disc radius is 1.8«1019 cin."," In the following, we discuss the influence of each individual effect mentioned above; our reference situation is that of a system with a 1.0 $_\odot$ primary, whose radius is $5 \times 10^8$ cm; the mass transfer rate is $3 \times
10^{16}$ g $^{-1}$ , and the average outer disc radius is $1.8 \times 10^{10}$ cm."
 These parameters are typical of short period dwarf novae with lnassive primaries., These parameters are typical of short period dwarf novae with massive primaries.
 The light curve corresponding to the, The light curve corresponding to the
"model £57«d by assuming. the underlying. 3D overdensity. field. has a correlation. function. £""«d given. by Note that we are using £' to represent the real-space dimensional correlation function.",model $\xi_p^h$ by assuming the underlying 3D overdensity field has a correlation function $\xi^h$ given by Note that we are using $\xi^r$ to represent the real-space 3-dimensional correlation function.
 As can be seen in Fig. 10..," As can be seen in Fig. \ref{fig:autoABC},"
" this model is well-matched to the measured £,.", this model is well-matched to the measured $\xi_p$.
 The justification for this model is that the multiplicative boost to the projected density, The justification for this model is that the multiplicative boost to the projected density
small anisotropies observed in the angular distribution (Mészárosetal.. 1999)); short- and long- GRBs in BATSE catalog actually show two different angular. distributions Vavreketal..2008.Balazs 2009)). but the connection to the hypothesis of the instrumental selection effect is not still completely clear.,"small anisotropies observed in the angular distribution \cite{mes99}) ); short- and long- GRBs in BATSE catalog actually show two different angular distributions \cite{vav08,bal09}) ), but the connection to the hypothesis of the instrumental selection effect is not still completely clear."
 In this paper we characterize the angular distributions of the two classes of GRBs with multifractal analysis from a selected sample of BATSE and Swift observations. taking also in consideration the redshift relativistic effects and compare the results obtained from the two catalogues.," In this paper we characterize the angular distributions of the two classes of GRBs with multifractal analysis from a selected sample of BATSE and Swift observations, taking also in consideration the redshift relativistic effects and compare the results obtained from the two catalogues."
 In Chapter 2 are described the mathematical basis of the method., In Chapter 2 are described the mathematical basis of the method.
 In Chapter 3 are estimated the fractal/multifractal dimensions by determining the distribution moments of the multifractal spectrum from the second up to the tenth order and then draw the conclusions., In Chapter 3 are estimated the fractal/multifractal dimensions by determining the distribution moments of the multifractal spectrum from the second up to the tenth order and then draw the conclusions.
 Being GRBs associated to galaxies. they should trace the angular distribution of their host galaxies at distances slightly larger than those estimated with classical Supernovae. te. distances the order the Gpe (Paezyfíski.1986.Usov&Chibisov. 1975)). across which inhomogeneity in the distribution of luminous matter should be averaged out from the Mpe scale.," Being GRBs associated to galaxies, they should trace the angular distribution of their host galaxies at distances slightly larger than those estimated with classical Supernovae, i.e. distances the order the Gpc \cite{pac86,uso75}) ), across which inhomogeneity in the distribution of luminous matter should be averaged out from the Mpc scale."
 Mandelbrot and Peebles (1980) first used fractal geometry to deseribe the angular and spatial distribution of galaxies., Mandelbrot and Peebles (1980) first used fractal geometry to describe the angular and spatial distribution of galaxies.
" In their work they used Lévvy-Rayleigh random paths. Le. infinite-variance. stable. generalized random walks in which the step lengths are described by a tailed probability distribution,"," In their work they used Lévvy-Rayleigh random paths, i.e. infinite-variance, stable, generalized random walks in which the step lengths are described by a tailed probability distribution."
 Galaxies are placed at the steps of a random walk and each galaxy is randomly connected with another in the close proximity. like mimicking the random motior of a fly in the air passing through each given point.," Galaxies are placed at the steps of a random walk and each galaxy is randomly connected with another in the close proximity, like mimicking the random motion of a fly in the air passing through each given point."
 The power-law constraint of the motion gives the fractal exponent of the random path and rules the distribution of the jump lengths and the direction of each jump ts taken isotropically at random (Martinez. 2002))., The power-law constraint of the motion gives the fractal exponent of the random path and rules the distribution of the jump lengths and the direction of each jump is taken isotropically at random \cite{mar02}) ).
 Galaxy surveys show that the distribution of the luminous matter in the universe is more complicated than that of a single fractal., Galaxy surveys show that the distribution of the luminous matter in the universe is more complicated than that of a single fractal.
 The distribution of galaxies has instead multifractal properties (Pietronero.1987.Célérier&Thieberger.2001 )). with a tendency to cluster. following the well known peculiar structures made by patterns of voids and filaments (Kurokawaetal.. 1999)).," The distribution of galaxies has instead multifractal properties \cite{pie87,cel01}) ), with a tendency to cluster following the well known peculiar structures made by patterns of voids and filaments \cite{kur99}) )."
 The CfA survey of nearby galaxies. for example. presents scaling properties with a correlation dimension in the interval Ds.=1.3—1.," The CfA survey of nearby galaxies, for example, presents scaling properties with a correlation dimension in the interval $D_{2}\simeq 1.3-1.4$."
 The two-point correlation function and the power spectrum analysis indicate that the distribution of galaxies at very large scales becomes homogeneous and isotropic like the X-ray background emitted by active galactic nuclei (AGNs). in agreement with the Cosmological Principle (Peebles. 1993)).," The two-point correlation function and the power spectrum analysis indicate that the distribution of galaxies at very large scales becomes homogeneous and isotropic like the X-ray background emitted by active galactic nuclei (AGNs), in agreement with the Cosmological Principle \cite{pee93}) )."
 Multifractal scaling analysis revealed the presence of aggregated structures also in. some samples of galaxies with distances larger than r~20 Mpe., Multifractal scaling analysis revealed the presence of aggregated structures also in some samples of galaxies with distances larger than $r\sim 20$ Mpc.
 At scales larger than 30 Mpe. the angular distribution of luminous. matter tends to homogeneity.," At scales larger than $30$ Mpc, the angular distribution of luminous matter tends to homogeneity."
 According to Pietronero.1987. the distribution of luminous matter has multifractal properties with fractal dimensions varying from dΞ1.23. for the nearby galaxies. up to the value α=2 at very high redshifts.," According to \cite{pie87} the distribution of luminous matter has multifractal properties with fractal dimensions varying from $d=1.23$, for the nearby galaxies, up to the value $d\simeq 2$ at very high redshifts."
 Usually. multifractal distributions are present when a structure has different fractal dimensions on different parts of the support: in other words. when spatial correlations are present and change the geometrical shape of the distribution at different scales (Falconer.1990).," Usually, multifractal distributions are present when a structure has different fractal dimensions on different parts of the support; in other words, when spatial correlations are present and change the geometrical shape of the distribution at different scales \cite{fal90}) )."
 Those distributions cannot be adequately described by a geometrical support with a single fractal dimension. but require instead a whole spectrum of dimensions.," Those distributions cannot be adequately described by a geometrical support with a single fractal dimension, but require instead a whole spectrum of dimensions."
" To characterize the angular distribution of GRBs in terms of multifractals. we use the moments that estimates the fractality of the distribution by calculating the multifractalspectrum of generalized dimensions D, in a given range (Falconer.1990.Feder. 1988))."," To characterize the angular distribution of GRBs in terms of multifractals, we use the ” that estimates the fractality of the distribution by calculating the multifractalspectrum of generalized dimensions $D_q$ in a given range \cite{fal90,fed88}) )."
 Consider N GRBs in a box of size L divided into cells of size r., Consider $N$ GRBs in a box of size $L$ divided into cells of size $r$.
" The sample is said to have a non-null q—£/ moment £, on the scale r where 7; is the number of objects present in the i-th cell.", The sample is said to have a non-null $q-th$ moment $\xi_q$ on the scale $r$ where $n_i$ is the number of objects present in the i-th cell.
 Higher moments reflect the emergence of structures present in the denser regions., Higher moments reflect the emergence of structures present in the denser regions.
 From the scaling relation conducted on a two-dimensional section we obtain (Kurokawa et al., From the scaling relation conducted on a two-dimensional section we obtain (Kurokawa et al.
" 2001). where the coefficients D, belong to the spectrum of generalized dimensions."," 2001), where the coefficients $D_q$ belong to the spectrum of generalized dimensions."
" Do is the dimension. D, 1s thedimension. and D» is thedimension."," $D_0$ is the , $D_1$ is the, and $D_2$ is the."
 To estimate the multifractality of the distribution one has to determine each generalized dimension D—q as a function of the g-th moment., To estimate the multifractality of the distribution one has to determine each generalized dimension $D-q$ as a function of the $q$ -th moment.
" One gives evidence to the dependence of each of the fractal dimensions from the moments by drawing the q—D, plot. which is based on the Lipshitz-Hóllder exponent (Benassietal..1998.Goltz. 1997))."," One gives evidence to the dependence of each of the fractal dimensions from the moments by drawing the $q-D_q$ plot, which is based on the Lipshitz-Höllder exponent \cite{ben98, gol97}) )."
" In the q—D, plot one defines those simple fractal structures described by only one dimension. D,. and in the plot all the D,’s result equal to D,."," In the $q-D_q$ plot one defines those simple fractal structures described by only one dimension, $D_f$ , and in the plot all the $D_q$ 's result equal to $D_f$."
" In the general case of amultifractaL. instead. D, is usually decreasing for higher and higher values of the moments q: if g>q'. then D,x until converging to the asymptotic value D for a distribution of infinite objects."," In the general case of a, instead, $D_q$ is usually decreasing for higher and higher values of the moments $q$: if $q>q'$, then $D_q \leq D_{q'}$ until converging to the asymptotic value $D_\infty$ for a distribution of infinite objects."
 To characterise with a good precision the multifractal distribution of a sample of object in space one usually does not need to calculate the generalized dimensions up to the limit dictated by the number of points in the space., To characterise with a good precision the multifractal distribution of a sample of object in space one usually does not need to calculate the generalized dimensions up to the limit dictated by the number of points in the space.
 For our purposes a good estimation was obtained by taking as upper limit g= 10., For our purposes a good estimation was obtained by taking as upper limit $q=10$ .
 From Swift recent observations we selected a sample of 444 GRBswithout ambiguity of classification. 52 of which are classifiedas short- GRBs after the correction of Too from the relativistic effects (data updated at 2009-09-15).," From Swift recent observations we selected a sample of $444$ GRBswithout ambiguity of classification, $52$ of which are classifiedas short- GRBs after the correction of $T_{90}$ from the relativistic effects (data updated at 2009-09-15)."
 From BATSE, From BATSE
for assigned value of V. (e.. of c) it is determined by the value of the central dark matter density. Powe,"for assigned value of $V_{\infty}$ (i.e., of $\sigma$ ) it is determined by the value of the central dark matter density, $\rho_{DM}^{0}$."
 For our zero-thickness disk. we take the surface density to be proportional to the observed surface brightness p(R). Le. here the function 6 is the Dirac delta-function.," For our zero-thickness disk, we take the surface density to be proportional to the observed surface brightness $\mu(R)$, i.e.: here the function $\delta$ is the Dirac delta-function."
" The ratio M/L is the mass-to-light ratio of the stellar disk. which we take to be constant: for assigned photometric profile iR). it represents the only free dimensional scale of the stellar density distribution,"," The ratio $M/L$ is the mass-to-light ratio of the stellar disk, which we take to be constant; for assigned photometric profile $\mu(R)$, it represents the only free dimensional scale of the stellar density distribution."
 The model is constructed by solving the complete Poisson equation for the system. in which the matter densities are defined by the distributions (7)) and (11): This is a nonlinear differential equation for the total gravitational potential pj. which must be solved with the appropriate boundary conditions: Eq. (8)).," The model is constructed by solving the complete Poisson equation for the system, in which the matter densities are defined by the distributions \ref{rhoiso}) ) and \ref{rhodisk}) ): This is a nonlinear differential equation for the total gravitational potential $\Phi_T$, which must be solved with the appropriate boundary conditions: Eq. \ref{zeropot}) ),"
 as the condition at the center. and Eq. (9)).," as the condition at the center, and Eq. \ref{asypot}) ),"
 for the asymptotic behavior at large radii., for the asymptotic behavior at large radii.
 The solution of the Poisson Eq. (12)).," The solution of the Poisson Eq. \ref{dimpois}) ),"
" which depends on the two free scales Phy and M/L. will allow us to calculate the rotation curve of the model V,,LR)=WROD;/OR)|.o. to be compared with the observed V(R)."," which depends on the two free scales $\rho_{DM}^{0}$ and $M/L$, will allow us to calculate the rotation curve of the model $V_{mod}(R)=\sqrt{R(\partial \Phi_T/\partial R)|_{z=0}}$, to be compared with the observed $V(R)$."
 Note that this method is intrinsically different from. the standard parametric decompositions of the mass models of spiral galaxies., Note that this method is intrinsically different from the standard parametric decompositions of the mass models of spiral galaxies.
 In. standard parametric analyses. the contribution of the dark matter halo to the total gravitational potential has a fixed shape (usually. spherically symmetric). which is independent of the presence and gravitational importance of the stellar disk and which can be changed only through a set of parameters a;. that is Φρη=Ppay(R.zia;," In standard parametric analyses, the contribution of the dark matter halo to the total gravitational potential has a fixed shape (usually, spherically symmetric), which is independent of the presence and gravitational importance of the stellar disk and which can be changed only through a set of parameters $a_i$, that is $\Phi_{DM}=\Phi_{DM}(R,z;a_i)$."
 The present self-consistent models have a fixed parametric profile for the shape of the gravitational potential of the dark matter component., The present self-consistent models have a fixed parametric profile for the shape of the gravitational potential of the dark matter component.
 Through Eq. (12)).," Through Eq. \ref{dimpois}) ),"
 the dark matter gravitational potential. and thus its contribution to the total rotation curve. depends implicitly also on the properties of the stellar disk (its surface brightness u(R) and its mass to light ratio). that is Dp;=Ppa(R.τρ:WOR).M/L).," the dark matter gravitational potential, and thus its contribution to the total rotation curve, depends implicitly also on the properties of the stellar disk (its surface brightness $\mu(R)$ and its mass to light ratio), that is $\Phi_{DM}=\Phi_{DM}(R,z; \rho_{DM}^{0}; \mu(R), M/L)$."
 In conclusion. for a given rotation curve V(R) and surface brightness of the stellar disk jj(R). the best-fit self-consistent disk-halo decomposition will correspond to the self-consistent model defined by the pair (oaMIL). together with the asymptotic velocity V4. that gives the best fit to the observed rotation curve.," In conclusion, for a given rotation curve $V(R)$ and surface brightness of the stellar disk $\mu(R)$, the best-fit self-consistent disk-halo decomposition will correspond to the self-consistent model defined by the pair $(\rho_{DM}^{0}, {M/
L})$, together with the asymptotic velocity $V_{\infty}$, that gives the best fit to the observed rotation curve."
" As units of length and of specific energy we will consider the scale length of the stellar disk /: and the square of the velocity dispersion c=(eV2,/2. respectively. so that we introduce the dimensionless coordinates and the dimensionless gravitational potential We separate the Poisson Eq. (12))"," As units of length and of specific energy we will consider the scale length of the stellar disk $h$ and the square of the velocity dispersion $\sigma^2=V_{\infty}^2/2$, respectively, so that we introduce the dimensionless coordinates and the dimensionless gravitational potential We separate the Poisson Eq. \ref{dimpois}) )"
 in the two Poisson equations for the two different mass components of the system., in the two Poisson equations for the two different mass components of the system.
 The dimensionless equation for the stellar disk reads Here we have introduced the normalized surface density. defined by ECR)2XXR)/X(0). and the parameter f which is the dimensionless version of the weight of the stellar disk as measured by the mass-to-light ratio M/L.," The dimensionless equation for the stellar disk reads Here we have introduced the normalized surface density, defined by $\hat\Sigma(R)\equiv{{\Sigma(R)}/{\Sigma(0)}}$, and the parameter $\beta$ which is the dimensionless version of the weight of the stellar disk as measured by the mass-to-light ratio $M/L$."
 To solve Eq. (15)).," To solve Eq. \ref{poisdisk}) ),"
 we may use the standard technique of Hankel transforms (see Toomre (1963))., we may use the standard technique of Hankel transforms (see Toomre (1963)).
 In this way the potential associated with any surface density profile (M/LiR) can be calculated exactly., In this way the potential associated with any surface density profile $(M/L)\mu(R)$ can be calculated exactly.
" Consistent with the condition (8)). we also impose that w,(0.0)=0."," Consistent with the condition \ref{zeropot}) ), we also impose that $\psi_D(0,0)=0$."
 In particular. for purely exponential disks. X=exp(—&). the associated gravitational potential is given by where Jo follows the standard notation of Bessel functions.," In particular, for purely exponential disks, $\hat\Sigma=\exp(-\xi)$, the associated gravitational potential is given by where $J_0$ follows the standard notation of Bessel functions."
 In turn. the Poisson equation for the dark matter halo can be written as where the parameter a can be seen as the dimensionless expression of the central dark matter density scale: The parameter « measures the halo concentration.," In turn, the Poisson equation for the dark matter halo can be written as where the parameter $\alpha$ can be seen as the dimensionless expression of the central dark matter density scale: The parameter $\alpha$ measures the halo concentration."
 With an eye to the central region of the system. a measure of the halo mass relative to the the mass of the disk is given by the parameter a/p=οΣυ (see also Eq. (12))).," With an eye to the central region of the system, a measure of the halo mass relative to the the mass of the disk is given by the parameter $\alpha/\beta = \rho_{DM}^{0} h / \Sigma_0$ (see also Eq. \ref {dimpois}) ))."
 Thus the dimensional pair QuyMID has its dimensionless counterpart in (cr.8).," Thus the dimensional pair $(\rho_{DM}^{0}, {M/ L})$ has its dimensionless counterpart in $(\alpha, \beta)$."
 In the (a.8) plane. the straight lines defined by Bxa can be interpreted as lines of approximately given central density of the halo with respect to the mass of the disk: steep lines are those characterizing relatively light halos.," In the $(\alpha, \beta)$ plane, the straight lines defined by $\beta \propto \alpha$ can be interpreted as lines of approximately given central density of the halo with respect to the mass of the disk; steep lines are those characterizing relatively light halos."
" A better defined measure of the weight of the halo will be given in the next Section,", A better defined measure of the weight of the halo will be given in the next Section.
 The Poisson Eq. (18)), The Poisson Eq. \ref{poishalo}) )
 for the dark matter halo has to be solved numerically (see Appendix A)., for the dark matter halo has to be solved numerically (see Appendix A).
 Here the gravitational potential of the disk. calculated separately as described above. acts as an external potential that breaks the spherical symmetry and shapes the dark matter density distribution.," Here the gravitational potential of the disk, calculated separately as described above, acts as an external potential that breaks the spherical symmetry and shapes the dark matter density distribution."
" The relevant boundary conditions can. be obtained from. the boundary conditions imposed on the total gravitational potential: where the symbol vj, indicates the constant value approached (with spherical symmetry) by the dimensionless gravitational potential of the disk at large distances from the center of the system.", The relevant boundary conditions can be obtained from the boundary conditions imposed on the total gravitational potential: where the symbol $\psi_{D}^{\infty}$ indicates the constant value approached (with spherical symmetry) by the dimensionless gravitational potential of the disk at large distances from the center of the system.
The SuperCOSALOS machine. digitises the scanned photograph at. 10/7 resolution (0.67 aresec) in 15bit grey levels (further details of the hardware are given in Llambly. et al.,The SuperCOSMOS machine digitises the scanned photograph at $\mu$ m resolution (0.67 arcsec) in 15–bit grey levels (further details of the hardware are given in Hambly et al.
 L998)., 1998).
" Pixel values are in transmission units and are related to incident intensity via the usual transmission CE). density (2) and intensity (2) relationships on the the loglinear part of the photographic characteristic curve: where 5 is the gradient of the response in the logy,£ plane."," Pixel values are in transmission units and are related to incident intensity via the usual transmission $T$ ), density $D$ ) and intensity $I$ ) relationships on the the log–linear part of the photographic characteristic curve: where $\gamma$ is the gradient of the response in the $D-\log_{10}I$ plane."
 Each photograph is digitised. in ~2 hr. producing 2 Ghvte of data.," Each photograph is digitised in $\sim2$ hr, producing $\sim2$ Gbyte of data."
" The pixel map is ""blocked down in 64.64 pixel cells during scanning to produce a raw median sky map: all subsequent. processing is then carried out ollline.", The pixel map is `blocked down' in $64\times64$ pixel cells during scanning to produce a raw median sky map; all subsequent processing is then carried out offline.
 Object detection in the raw pixel map requires an estimate of the sky value at every point within the twodimensional image., Object detection in the raw pixel map requires an estimate of the sky value at every point within the two–dimensional image.
 Phe raw median values of the sky are first. filtered using a weighted median filter (Brownrige 1984) of the form applied. over filtering. scales of 3.2. and 1. blocks (c.g. Stobic 1986).," The raw median values of the sky are first filtered using a weighted median filter (Brownrigg 1984) of the form applied over filtering scales of 3,2 and 1 blocks (e.g. Stobie 1986)."
 This results in. variations in scale size greater than ~2 mm (or ~2 aremin) being retained as sky variations in the background array., This results in variations in scale size greater than $\sim2$ mm (or $\sim2$ arcmin) being retained as sky variations in the background array.
 ‘This is obviouslyμ a compromise between detection of low surface rightness objects at this scale or greater and accurate image »wameterisation in areas of real sky variation., This is obviously a compromise between detection of low surface brightness objects at this scale or greater and accurate image parameterisation in areas of real sky variation.
 Experiments with cillering scale lengths showed that the choice of raw rlocking of 0.64 mm squares coupled with iterative filtering scale sizes of 3.2 and 1 (e.g. Stobie 1986) is optimal in erms of minimising spurious detections due to. for example. ow level scattered light around bright stars. (see also MacGillivray Stobie 1984).," Experiments with differing scale lengths showed that the choice of raw blocking of 0.64 mm squares coupled with iterative filtering scale sizes of 3,2 and 1 (e.g. Stobie 1986) is optimal in terms of minimising spurious detections due to, for example, low level scattered light around bright stars (see also MacGillivray Stobie 1984)."
 Alost mocern Schmidt photographs have one or more intensity calibration “stepwedges’ exposed on their sides or corners (see. for example. “Tritton 1983).," Most modern Schmidt photographs have one or more intensity calibration `step–wedges' exposed on their sides or corners (see, for example, Tritton 1983)."
 In. principle. it is possible to measure the transmission of the various intensity spots and calibrate measured. transmission. into. incident relative intensity.," In principle, it is possible to measure the transmission of the various intensity spots and calibrate measured transmission into incident relative intensity."
 In. practice. however. there are a number of reasons why even rough stepwedge calibration is clifficult to achieve.," In practice, however, there are a number of reasons why even rough step–wedge calibration is difficult to achieve."
 Variations in emulsion sensitivity across the plate and nonuniform illumination of the wedge spots themselves can cause large errors., Variations in emulsion sensitivity across the plate and non–uniform illumination of the wedge spots themselves can cause large errors.
 Also. the fact that SuperCOSMOS scans bw illumination. of a strip of emulsion means that when imaging dense spots. the light levels recorded are badly allected by cülfracted light originating from relatively clear emulsion surrounding the spot.," Also, the fact that SuperCOSMOS scans by illumination of a strip of emulsion means that when imaging dense spots, the light levels recorded are badly affected by diffracted light originating from relatively clear emulsion surrounding the spot."
 Measurements of the 16spot KAPNO step wedges on modern UINSTU plates using a Macbeth cülfuse clensitometer show values of characteristic curve gradient ?2+10% (indeed. such a value is a criterion. for a survey plate and inclusion in the final UINSTU photographic skv atlases).," Measurements of the 16--spot KPNO step wedges on modern UKSTU plates using a Macbeth diffuse densitometer show values of characteristic curve gradient $\gamma\sim2\pm10$ (indeed, such a value is a criterion for a survey--grade plate and inclusion in the final UKSTU photographic sky atlases)."
 Llence. we use anessumced transmission.tointensity calibration with slope corresponding to 5=2 (older glass copies of the POSSE -7I7 survey are analysed using a dillerent. value due to their poorer contrast).," Hence, we use an transmission--to--intensity calibration with slope corresponding to $\gamma\approx2$ (older glass copies of the POSS–I `E' survey are analysed using a different value due to their poorer contrast)."
 This procedure is vindicated not least. bv. the. demonstration. in Section 2.3.1.. of measurements of the gradients of the instrumental magnitude scale with respect to CCD data. where we show that the resulting scales have close to unit eracdient with scatter at around the level.," This procedure is vindicated not least by the demonstration, in Section \ref{grads}, of measurements of the gradients of the instrumental magnitude scale with respect to CCD data, where we show that the resulting scales have close to unit gradient with scatter at around the level."
 Image detection requires thresholding above a local sky estimate., Image detection requires thresholding above a local sky estimate.
 The sss pixel data are thresholded at a constant percentage cut above local sky., The SSS pixel data are thresholded at a constant percentage cut above local sky.
" X. sky. noise. parameter COxky, is determined [rom a small (—LO10 arcmin) region of image data near the centre of cach plate.", A sky noise parameter $\sigma_{{\rm SKY}_0}$ is determined from a small $\sim10\times10$ arcmin) region of image data near the centre of each plate.
" The skv backerouncl value Isicy,, in this region is determined as described. in Beard. MacGillivray Thaniseh (1990) while the ey, estimator uses the relationship σον148.MADzagiy,. where Ανν, is the median of absolute deviations of the pixel values about the sky. estimate (the factor LAS is the sealing needed to convert. the MAD estimate for a Gaussian. distribution to the equivalent Gaussian sigma)."," The sky background value $_{{\rm SKY}_0}$ in this region is determined as described in Beard, MacGillivray Thanisch (1990) while the $\sigma_{{\rm SKY}_0}$ estimator uses the relationship $\sigma_{{\rm SKY}_0}=1.48\times{\rm MAD}_{{\rm SKY}_0}$ , where $_{{\rm SKY}_0}$ is the median of absolute deviations of the pixel values about the sky estimate (the factor 1.48 is the scaling needed to convert the MAD estimate for a Gaussian distribution to the equivalent Gaussian sigma)."
" This vields a ""percentage eut value PCUT (=100.2.3esny,/dsey,) such that at any part of the image where Zpungsg ds the threshold. intensity. {ων is he local sky estimate based on bilinear interpolation of he filtered sky map. and PCUT is typically for survey plates."," This yields a `percentage cut' value PCUT $=100\times2.3\sigma_{{\rm SKY}_0}/I_{{\rm SKY}_0}$ ) such that at any part of the image where $I_{\rm THRESH}$ is the threshold intensity, $I_{\rm SKY}$ is the local sky estimate based on bilinear interpolation of the filtered sky map, and PCUT is typically for survey plates."
 Note that this singleparameter multiplicative hresholding is designed to cope with the fact that the sky evel (and therefore the noise level) changes over the Schmidt ield due to vignetting. and à zeropoint sky intensity. value is chosen for the transmissiontointensitv calibration such hat the threshold. in terms of the number of & above sky. remains constant regardless of sky variations.," Note that this single–parameter multiplicative thresholding is designed to cope with the fact that the sky level (and therefore the noise level) changes over the Schmidt field due to vignetting, and a zeropoint sky intensity value is chosen for the transmission–to–intensity calibration such that the threshold, in terms of the number of $\sigma$ above sky, remains constant regardless of sky variations."
 Once the thresholded. pixel. set. has been. determine or a plate. pixel connectivity can be analysed ane image parameterisation can take place.," Once the thresholded pixel set has been determined for a plate, pixel connectivity can be analysed and image parameterisation can take place."
 Phe details of the COSMOS pixel analyser are contained in AlacCuillivray Stobie (1984). Stobie (1986) and. Beare et al. (," The details of the COSMOS pixel analyser are contained in MacGillivray Stobie (1984), Stobie (1986) and Beard et al. ("
1990) anc will not be repeated here: however we describe some details that are relevant to SuperCOSMOS survey data;,1990) and will not be repeated here; however we describe some details that are relevant to SuperCOSMOS survey data.
 Many of the techniques used are similar to those developed for the APM machine (c.g. Irvin 1985: Draper Eaton 1999)., Many of the techniques used are similar to those developed for the APM machine (e.g. Irwin 1985; Draper Eaton 1999).
 For image detection. Sfold. nearest.neighbour pixe," For image detection, 8–fold nearest–neighbour pixel"
episode [rom the primary.,episode from the primary.
 The secondary accretes a large fraction of the outflowing mass., The secondary accretes a large fraction of the outflowing mass.
 We now turn to discuss the light curves (Fig. 1))., We now turn to discuss the light curves (Fig. \ref{fig:lightcurves1}) ).
 The early light eurve of showed 3 peaks (μιαν οἱ al., The early light curve of showed 3 peaks (Munari et al.
 2002). and then had a rapid decay by 5 magnitudes in about 20 days. shorter (han the ~80davs duration of the high-brightness phase.," 2002), and then had a rapid decay by 5 magnitudes in about 20 days, shorter than the $\sim 80 \days$ duration of the high-brightness phase."
 It. is unknown if such peaks existed in NGC 300 OT (Bond et al., It is unknown if such peaks existed in NGC 300 OT (Bond et al.
 2009). but this may be due to a poor lime coverage of the early light curve.," 2009), but this may be due to a poor time coverage of the early light curve."
 However. in the decline phase. the light curve of NGC 300 OT resembles those of V838. Mon and 5 Car GE (Fig. 1)).," However, in the decline phase, the light curve of NGC 300 OT resembles those of V838 Mon and $\eta$ Car GE (Fig. \ref{fig:lightcurves1}) )."
 The similarity of NGC 300 OT to V838 Mon ancl few other ILOTs was already noticed by other authors (e.g. Derger et al., The similarity of NGC 300 OT to V838 Mon and few other ILOTs was already noticed by other authors (e.g. Berger et al.
 2009. S1iuth et al.," 2009, Smith et al."
 2009)., 2009).
 In the case of NGC 300 OT. 120 days were required for a decay by 5 magnitudes. a time longer than the high-brightness phase.," In the case of NGC 300 OT, 120 days were required for a decay by 5 magnitudes, a time longer than the high-brightness phase."
 For the GE of η Car. it took no less than ~18vr for the light curve to decay by 5 magnitudes (IIumnphreys et al.," For the GE of $\eta$ Car, it took no less than $\sim 18 \yr$ for the light curve to decay by 5 magnitudes (Humphreys et al."
 1999)., 1999).
 Remarkably. when we normalize the timescales of the three objects without changing je magnitude scale we find that the light curves look verv similar (Fig. 2)).," Remarkably, when we normalize the timescales of the three objects without changing the magnitude scale we find that the light curves look very similar (Fig. \ref{fig:lightcurves2}) )."
" A similar procedure is known as ""stretch correction” when applied to SN lisht curves.", A similar procedure is known as “stretch correction” when applied to SN light curves.
 The normalized V.ope of the decline behaves the same way for all three eruptions., The normalized slope of the decline behaves the same way for all three eruptions.
 We consider this a hint val all three. eruptions are governed by (he same physical mechanism., We consider this a hint that all three eruptions are governed by the same physical mechanism.
 For example. the residual energy. source is not as high (even when the bolometric light. curve is considered) as in radioactive decay in SNe.," For example, the residual energy source is not as high (even when the bolometric light curve is considered) as in radioactive decay in SNe."
 Weaker contribution is possible. e.g. [rom mass accretion ab a low rate (like in the scenario for V338 Mon: Tvlenca 2005).," Weaker contribution is possible, e.g. from mass accretion at a low rate (like in the scenario for V838 Mon; Tylenda 2005)."
 Our leading candidate lor the powering mechanism is dissipation of accretion energy., Our leading candidate for the powering mechanism is dissipation of accretion energy.
 We consider the similar slopes as supporting argument to NGC 300 OT beinga scaled-down version of the GE of Car., We consider the similar slopes as supporting argument to NGC 300 OT beinga scaled-down version of the GE of $\eta$ Car.
 The rapid decay in Iuminositv is determined by Che decrease in the supply of accreted mass., The rapid decay in luminosity is determined by the decrease in the supply of accreted mass.
 Later. the slower decay is dictated by the settling of the inner part of the inflated envelope on the accreting object.," Later, the slower decay is dictated by the settling of the inner part of the inflated envelope on the accreting object."
 In this later phase the Iuminosity is below the Edcdington Iuminositx., In this later phase the luminosity is below the Eddington luminosity.
 The Gime scale itself depends on the accreting object aud on (he structure of the inflated envelope., The time scale itself depends on the accreting object and on the structure of the inflated envelope.
 On the other hand. when we compare NGC 300 OT to SNe (Fig. 1))," On the other hand, when we compare NGC 300 OT to SNe (Fig. \ref{fig:lightcurves1}) )"
 we find that the V.ope of the decline phase of NGC 300 OT is steeper than the slope of the decline phase οἱ 5Ne., we find that the slope of the decline phase of NGC 300 OT is steeper than the slope of the decline phase of SNe.
 Even SN 20085 which is considered (he “twin” of NGC 300 OT (Prieto et al., Even SN 2008S which is considered the “twin” of NGC 300 OT (Prieto et al.
 2003) shows this cliserepancy its decay was more gradual than that of NGC 300 OT in the V-band., 2008) shows this discrepancy – its decay was more gradual than that of NGC 300 OT in the V-band.
 This is despite (he fact that its progenitor had properties similar to of NGC 300 OT. and the outburst duration was similar.," This is despite the fact that its progenitor had properties similar to of NGC 300 OT, and the outburst duration was similar."
 ISN 20085 is indeed à SN. as argued by Dotticella et al. (," If SN 2008S is indeed a SN, as argued by Botticella et al. ("
2009). ils kinetic energy is expected to be ~10—100 larger than the radiated energy.,"2009), its kinetic energy is expected to be $\sim 10-100$ larger than the radiated energy."
 In such a case SN 20085 would be located in the region of exploding stars in Figure 3.. consistent. with its elaimed SN nature.," In such a case SN 2008S would be located in the region of exploding stars in Figure \ref{fig:totEvst}, , consistent with its claimed SN nature."
where sf and yp awe free parameters. and Aq is a modified jewel function. which represents the projected density profile ofan edge-onexponential disc.,"where $A_d$ and $r_d$ are free parameters, and $K_1$ is a modified Bessel function, which represents the projected density profile of an edge-on exponential disc."
 With only rather noisy photometric profiles and no detailed surface photometry for our galaxies. we only fit a clise model at large raclius as fitting a bulge-disc model proves very much underconstrained and does not give reliable values for the parameters of the disc and bulge.," With only rather noisy photometric profiles and no detailed surface photometry for our galaxies, we only fit a disc model at large radius as fitting a bulge-disc model proves very much underconstrained and does not give reliable values for the parameters of the disc and bulge."
 Ehe best fit photometric models are shown in Fig., The best fit photometric models are shown in Fig.
 12 where the models are fitted independently on each side of the galaxy. ancl the derived photometric scale-leneths ry are given in Table 1..," \ref{plotphotom} where the models are fitted independently on each side of the galaxy, and the derived photometric scale-lengths $r_d$ are given in Table \ref{table1}."
 This quantity differs from. the photometric scale length usually estimated. by. fitting bulge-dise models ., This quantity differs from the photometric scale length usually estimated by fitting two-component bulge-disc models .
(Barnes&IIeruquist1991:Barnes1992) aud drive strong inflows of gas into the circumnucenr rmon. (CasolietIleruquist1996DiMatteoetal.2001 ).,"\citep{barnes91, barnes92} and drive strong inflows of gas into the circumnuclear region \citep{casoli91,braine93,
barnes02, barnes96, dimatteo07}."
. These inflows aro a result of strong tidal toreMos. Issipation. aud strong asviunetres in the stellar «istrilition (c.g.bars:seeDarues2002:Barnes&Teruquist1996) leading to efhcieut star formation (DiAatteC0OTal.2007).," These inflows are a result of strong tidal torques, dissipation, and strong asymmetries in the stellar distribution \citep[e.g., bars; see][]{barnes02, barnes96} leading to efficient star formation \citep{dimatteo07}."
. While jc dutall of eas in imnergers has 2001] {iscussed within ιο context of smaulatious iud )bservalous (CasolietIleruquist199(xDiMatteoetal.2007 ).. currently. there ias been onlv Linited quautitatix teupts to address he relatiouship between iufall. Ornation. aud the netallicity(e..Perezotal.2006:Rupkect2010).," While the infall of gas in mergers has been discussed within the context of simulations and observations \citep{casoli91,braine93, barnes02, barnes96, dimatteo07}, currently, there has been only limited quantitative attempts to address the relationship between infall, star formation, and the metallicity \citep[e.g.,][]{perez06, rupke10}."
. The trends we find iu the sinulations aro fascinating and explain a number of imporant picnonienon that rave bee1 observed iu interacting aud τιwreie galaxies., The trends we find in the simulations are fascinating and explain a number of important phenomenon that have been observed in interacting and merging galaxies.
 Aletallicitv changes cming the merger event is a powerful way of testing the tinune au streneth of the easons inflows axd ultimateVv in understanding the fate of post-merecr1iteractiug οἱalaxies., Metallicity changes during the merger event is a powerful way of testing the timing and strength of the gas inflows and ultimately in understanding the fate of post-merger/interacting galaxies.
 Overall. we fid broad agreenont wih the resiIts of our simulations or the evolulon of the metallictv and wha is observed.," Overall, we find broad agreement with the results of our simulations for the evolution of the metallicity and what is observed."
 We find for example that the net decrease i ithe metalicity is a few eutlis «fa dex. which is comparable to whatis observediu ucreing aud interacting ealaxies (INewleyctal.2006>Rupke2008:Michel 2010).," We find for example that the net decrease in the metallicity is a few tenths of a dex, which is comparable to what is observed in merging and interacting galaxies \citep{kewley06,rupke08, dansac08, alonso10}."
. Iu addition. as ds likely observed. our smnulaIOUS show that the peak of the star formation is close in tie to the ni inmiuiuthe emeunmunuclear metallicity (Rurieetal.20is:Rodrigues 2008).," In addition, as is likely observed, our simulations show that the peak of the star formation is close in time to the minimum in the circumnuclear metallicity \citep{rupke08,rodrigues08}."
. The trends coud he explained cutirely by the strength of the gas inflows of relatively metal-poor eas from the outer regions o the disks which both fuel star ornuation aud dilute the metal coutcn of the gas in aud around the nucleus., The trends could be explained entirely by the strength of the gas inflows of relatively metal-poor gas from the outer regions of the disks which both fuel star formation and dilute the metal content of the gas in and around the nucleus.
 The streneth of the inflows likely depeuds both on the initial eas coiteut of the clisks. how it is disributed. the characteristics of he orits twoueh which the Oogalaxies merec.Oo alc the nass ratios of the galaxies.," The strength of the inflows likely depends both on the initial gas content of the disks, how it is distributed, the characteristics of the orbits through which the galaxies merge, and the mass ratios of the galaxies."
 Tere we concentrated oi the nodelling a range of orital parameters and a 20% eas Tactions., Here we concentrated on the modelling a range of orbital parameters and a $\%$ gas fractions.
 S1bsequent papers will focus on varving the nass ratios. eas fraction. and eas distributions.," Subsequent papers will focus on varying the mass ratios, gas fraction, and gas distributions."
 The simulalon resuls suggest that the strongest star onuation axd dilution of the initial uuclear metallicity Qccnurs around )orcenutre passage of the merecr., The simulation results suggest that the strongest star formation and dilution of the initial nuclear metallicity occurs around pericentre passage of the merger.
" This effect ius been observed iu tvat the star formation at small exubit a tail to high Ue cquivalent widtIs NUN(ο,ο,.Woodsetal.2006:Bartonct2000)."," This effect has been observed in that the star formation at small separations exhibit a tail to high $\alpha$ equivalent widths \citep[e.g.,][]{woods06, barton00}."
. At this tie. we also expect the wiiwin in the metallicity to occur.," At this time, we also expect the minimum in the metallicity to occur."
 This expect reationship has been enmiplasised in Rewlexotuneal.(2006) based on a study of pairs of galaxies., This expect relationship has been emphasised in \citet{kewley06} based on a study of pairs of galaxies.
 T1ο of Kewevetal.(2006) is particularly appropriate for this type of study as it is a robust sample of pairs of luteractingC» C»galaxies and show Oogenerally a modest merease int jer star formation rates and modest likelv dilutious., The sample of \citet{kewley06} is particularly appropriate for this type of study as it is a robust sample of pairs of interacting galaxies and show generally a modest increase in their star formation rates and modest likely dilutions.
" For example. their nios intense starbursts.owlhüeh show sjenificaut offsets in iietallicitv. only have offsets of about 0,2-0.3 dex conrpared to non-interacting ealaxics with sinular B-baud absolute magnitudes."," For example, their most intense starbursts, which show significant offsets in metallicity, only have offsets of about 0.2-0.3 dex compared to non-interacting galaxies with similar B-band absolute magnitudes."
" This is the amount we fiud iu our siulatio15,", This is the amount we find in our simulations.
 We can explün the full range of cireunaunmclear metallicities observed m more violen iuergers im the local (Rupkectal.2008) or distant universe (Rodrigues.etal.2005)., We can explain the full range of circumnuclear metallicities observed in more violent mergers in the local \citep{rupke08} or distant universe \citep{rodrigues08}.
 Dv violeut mergers. we nean in this coutext ealaxies wuch aro luminous in the infrared and have star formation rates several to almost 100 times that of the Milkv. W:ay.," By violent mergers, we mean in this context galaxies which are luminous in the infrared and have star formation rates several to almost 100 times that of the Milky Way."
 These are t10 dost appropriate for coniparison because. as we have sLOW]. as the merecr aclvauices. he star formatioji rate will decTease and the metallicity iucicuses reaching vaues that are higher than the initial netallicitics.," These are the most appropriate for comparison because, as we have shown, as the merger advances, the star formation rate will decrease and the metallicity increases reaching values that are higher than the initial metallicities."
 To show that our models can reproduce the data. we have made a nronte-carlo simiation of t1ο metallicity distribution we nuelt obtain if we obxyved set of ealaxies near their ani star formajon rate and allow for a range in their properties auk the iutriusic scatter in the 1inass-1uetallicitv relationship," To show that our models can reproduce the data, we have made a monte-carlo simulation of the metallicity distribution we might obtain if we observed a set of galaxies near their maximum star formation rate and allow for a range in their properties and the intrinsic scatter in the mass-metallicity relationship."
 For such a coniparison saliple. we would naturaly draw on the most violent star-formüus galaxies in the local 1niverse which are niajor niergers.," For such a comparison sample, we would naturally draw on the most violent star-forming galaxies in the local universe which are major mergers."
 This is the poiut in otry simulations. where we eet the most significaut influx of materia iuto the cmeuunuclear regions aud near the peal of star formation.," This is the point in our simulations, where we get the most significant influx of material into the circumnuclear regions and near the peak of star formation."
 At tljs stage. we can dilute the circunaunclear eas with relativolv metal poor gas youn the outer cisk decreasing the «overall metallicity by about a factor of 2 xd relativelv bricf time (of-order au iuteral dynanical time of about ys vrs).," At this stage, we can dilute the circumnuclear gas with relatively metal poor gas from the outer disk decreasing the overall metallicity by about a factor of 2 and relatively brief time (of-order an internal dynamical time of about $^8$ yrs)."
 Thus. 1n our simpe moute-carlo éyproach to the inetallicity estimates. Ww(Coa low for a decrease of 0.1 dex.," Thus, in our simple monte-carlo approach to the metallicity estimates, we allow for a decrease of $\pm$ 0.1 dex."
 We aake this alowahce as we oulv have single motallicity eradicut in our all our mnodcls O.1 dex )., We make this allowance as we only have single metallicity gradient in our all our models $-$ 0.1 dex $^{-1}$ ).
 However. ietallicity eradieus dn spiral ealaxies range Toni about (Q0 to 0.924ex 1," However, metallicity gradients in spiral galaxies range from about 0.0 to $-$ 0.2 dex $^{-1}$."
 To allow for this veWwiation. we have assiguec araise in the metallicity alowed of 0.1 dex.," To allow for this variation, we have assigned a range in the metallicity allowed of 0.1 dex."
 We do not kirow the nature of the progenitors either., We do not know the nature of the progenitors either.
 Thev could lave a range of initial eircuuunnuclear metallicities aud thle Lass netallicity relationship las a πιοποσα scatter 0.15 dex (Trenoutietal.200 1)., They could have a range of initial circumnuclear metallicities and the mass metallicity relationship has a significant scatter 0.15 dex \citep{tremonti04}. .
. Thus we allow for an additional source of «utter of 0.15 dex OLeflect the scatter iu the Lass-uctalicitv relation., Thus we allow for an additional source of scatter of 0.15 dex to reflect the scatter in the mass-metallicity relation.
 Finally. there is also measurement uucertaiuties which are about 0.1 dex (Rupkeetal.2008).," Finally, there is also measurement uncertainties which are about 0.1 dex \citep{rupke08}."
. Wo asstue that cach source of scatte ο and that cach can be linearly combined., We assume that each source of scatter is Gaussian and that each can be linearly combined.
 This assumption Is approwmiate since (at the metallicities «observed here) the uncertaiutv in the metallicity is not correlated to the metallicity. our numerical simulations are completely incdependeut of the data. aud the scatter in the mass relatioushipis intriusic.," This assumption is appropriate since (at the metallicities observed here) the uncertainty in the metallicity is not correlated to the metallicity, our numerical simulations are completely independent of the data, and the scatter in the mass-metallicity relationshipis intrinsic."
 One such iunoute-carlo siulation is shown in Fig., One such monte-carlo simulation is shown in Fig.
 8 but all reproduce the data from Ruakeetal.(2008) reasonably well., \ref{fig:MCmassmetallicity} but all reproduce the data from \citet{rupke08} reasonably well.
 Uufortuuately.," Unfortunately,"
detailedad surface napping techuques.,d surface mapping techniques.
We can see that the brighter stars (for which the S/N ratio is higher) show a larger scatter in the abundances of tron than the fainter stars.,We can see that the brighter stars (for which the $S/N$ ratio is higher) show a larger scatter in the abundances of iron than the fainter stars.
 The difference 15 quite significant. considering the errors associated to the estimate of the scatter: in the bin -2«Me.«—]. with 271 stars. we obtain an ris scatter of 0.033£0.002 from the GIRAFFE observations. whereas for the brighter bin —4<«-3 (81 stars) we have 0.065+0.004.," The difference is quite significant, considering the errors associated to the estimate of the scatter: in the bin $-2< M_K^0 <-1$, with 271 stars, we obtain an $rms$ scatter of $0.033\pm 0.002$ from the GIRAFFE observations, whereas for the brighter bin $-4< M_K^0 <-3$ (81 stars) we have $0.065\pm 0.004$."
 This difference ts Me.statistically significant. at a level of more than 4c.," This difference is statistically significant, at a level of more than $\sigma$."
 The results from the UVES spectra are essentially consistent with those derived from GIRAFFE (although with larger statistical errors owing to the much smaller number of stars available) and reveal a clear trend of the rs as a function of the luminosity., The results from the UVES spectra are essentially consistent with those derived from GIRAFFE (although with larger statistical errors owing to the much smaller number of stars available) and reveal a clear trend of the $rms$ as a function of the luminosity.
 This trend can be explained with the reasonable assumption that the atmosphere of each star (among the more Juminous objects) has a distinct peculiarity or intrinsic time variability., This trend can be explained with the reasonable assumption that the atmosphere of each star (among the more luminous objects) has a distinct peculiarity or intrinsic time variability.
 Ita et al. (, Ita et al. (
2002) found that very bright RGB stars in. the LMC are indeed variables (likely pulsating).,2002) found that very bright RGB stars in the LMC are indeed variables (likely pulsating).
 However. the stars considered here are considerably fainter. so that we are more inclined to attribute this variability to the large sizes (and small numbers) of convective cells with respect. e.g.. to solarlogues?.," However, the stars considered here are considerably fainter, so that we are more inclined to attribute this variability to the large sizes (and small numbers) of convective cells with respect, e.g., to solar."
. As a consequence. there is an uncertainty (of the order of about 0.04-0.05 dex) in the Fe abundances derived from an instantaneous observation. since the convective elements of à red giant are so large that only à small number of them occupy the surface of the star at any time (Schwarzschild 1975; Gray et al.," As a consequence, there is an uncertainty (of the order of about 0.04-0.05 dex) in the Fe abundances derived from an instantaneous observation, since the convective elements of a red giant are so large that only a small number of them occupy the surface of the star at any time (Schwarzschild 1975; Gray et al."
 2008)., 2008).
 To account for this uncertainty we would need to assume that the temperature in the line formation region differs from star to stars. with an rary scatter of about 50 K. While Schwarzschild (1975) estimated in about 400 the number of convective cells in the atmosphere of a red giant. our sample is mostly formed by warmer stars.," To account for this uncertainty we would need to assume that the temperature in the line formation region differs from star to stars, with an $rms$ scatter of about 50 K. While Schwarzschild (1975) estimated in about 400 the number of convective cells in the atmosphere of a red giant, our sample is mostly formed by warmer stars."
 We then assume that the number of convective cells is larger. 1n our case. say about 1000.," We then assume that the number of convective cells is larger, in our case, say about 1000."
 Of those. only half are obviously visible and not all of them with the same weight. because of the limb darkening.," Of those, only half are obviously visible and not all of them with the same weight, because of the limb darkening."
 Thus. we would get random fluctuations of about a factor 20 smaller than those expected at the star surface between the hotter regions (the top of the ascending columns) and the cooler regions (descending columns). having temperature differences by about 1000 K. Hence. we could expect that random fluctuations of the temperature in the region of line formation in a giant are effectively of the order of 50 K. even at similar effective temperatures.," Thus, we would get random fluctuations of about a factor 20 smaller than those expected at the star surface between the hotter regions (the top of the ascending columns) and the cooler regions (descending columns), having temperature differences by about 1000 K. Hence, we could expect that random fluctuations of the temperature in the region of line formation in a giant are effectively of the order of 50 K, even at similar effective temperatures."
 Following the same line of thought. one could expect that the luminosity is not perfectly constant. but should show fluctuations up to 0.01 mag (over timescales of a few tens of days): the same would show up in fluctuations of radial velocity (over similar timescales and amplitude of about 200-300 m/s).," Following the same line of thought, one could expect that the luminosity is not perfectly constant, but should show fluctuations up to 0.01 mag (over timescales of a few tens of days); the same would show up in fluctuations of radial velocity (over similar timescales and amplitude of about 200-300 m/s)."
 These expectations seem to agree with observations of stars similar to those in our sample (the threshold being -2«My<-1.4. Cummings 1999; Carney et al.," These expectations seem to agree with observations of stars similar to those in our sample (the threshold being $-2<M_V<-1.4$, Cummings 1999; Carney et al."
 2008)., 2008).
 In summary. since the spread measured from brighter stars could include the intrinsic. noise due to the convection. and pulsation (both in radial and non-radial mode). the implication is that it is better to use fainter stars in order to obtain the star- scatter.," In summary, since the spread measured from brighter stars could include the intrinsic noise due to the convection and pulsation (both in radial and non-radial mode), the implication is that it is better to use fainter stars in order to obtain the star-to-star scatter."
 From these stars (practically restricted to those with GIRAFFE spectra) the typical ss in a cluster should be about 0.033 dex (see Tab. 3)).," From these stars (practically restricted to those with GIRAFFE spectra) the typical $rms$ in a cluster should be about 0.033 dex (see Tab. \ref{t:diffe}) ),"
 or less. since this scatter value includes also the measurement errors.," or less, since this scatter value includes also the measurement errors."
 Finally. we remark that the scatter seems to be larger for more metal-poor clusters.," Finally, we remark that the scatter seems to be larger for more metal-poor clusters."
 We speculate that this is in part due to possibly larger observational errors. but in part it could be due to an increased transparency of the atmosphere. hence pointing to a larger relevance of the 3d-structure.," We speculate that this is in part due to possibly larger observational errors, but in part it could be due to an increased transparency of the atmosphere, hence pointing to a larger relevance of the 3d-structure."
 An alternative explanation of the larger scatter observed at higher luminosities (in particular in the magnitude bin -4<Myx< —3) could be contamiation from asymptotic giant branch (AGB) stars., An alternative explanation of the larger scatter observed at higher luminosities (in particular in the magnitude bin $-4< M_K <-3$ ) could be contamination from asymptotic giant branch (AGB) stars.
 Were this the case. we would expect that the sequences are sufficiently well separated (and the contamination negligible) at fainter magnitudes.," Were this the case, we would expect that the sequences are sufficiently well separated (and the contamination negligible) at fainter magnitudes."
 On the other hand. the temperature difference between RGB and AGB ts small. in the more luminous bin.," On the other hand, the temperature difference between RGB and AGB is small, in the more luminous bin."
 However. the error would be larger where the two sequences are so close that cannot be distinguished. but separated enough to give a wrong temperature when the luminosity is used. as in the procedure adopted 1n our analysis (see Carretta et al.," However, the error would be larger where the two sequences are so close that cannot be distinguished, but separated enough to give a wrong temperature when the luminosity is used, as in the procedure adopted in our analysis (see Carretta et al."
 2009a for further details)., 2009a for further details).
 In. this scenario. we would expect a distribution essentially Gaussian in form at lower luminosities. where the contamination is negligible.," In this scenario, we would expect a distribution essentially Gaussian in form at lower luminosities, where the contamination is negligible."
 At higher magnitudes. the distribution should become asymmetric. with a tal toward lower abundances given by the possible interlopers AGB stars for which the effective temperature (hence the abundance) would be underestimated.," At higher magnitudes, the distribution should become asymmetric, with a tail toward lower abundances given by the possible interlopers AGB stars for which the effective temperature (hence the abundance) would be underestimated."
 This asymmetry shoulc be more pronounced in the bin 4«My<-3., This asymmetry should be more pronounced in the bin $-4< M_K <-3$.
 We checked this possible etfect by computing the residuals in Fe abundances derived from GIRAFFE spectra in different bins of magnitudes. again restricting to the clusters in. the metallicity range —1.4. <[Fe/H]<—0.7 dex.," We checked this possible effect by computing the residuals in Fe abundances derived from GIRAFFE spectra in different bins of magnitudes, again restricting to the clusters in the metallicity range $-1.4<$ $<-0.7$ dex."
 The histogram of the residuals in various magnitude bins are plotted i Fig., The histogram of the residuals in various magnitude bins are plotted in Fig.
 7 There seems to be no particular asymmetry in the range —4«My—3. where the distribution appears only slightly larger.," \ref{f:fek2} There seems to be no particular asymmetry in the range $-4< M_K <-3$, where the distribution appears only slightly larger."
 In conclusion. this test seems to support the idea that the larger spread 1s due to the stellar atmospheres and to the variability provoked by the photospherte convection and stellar pulsation.," In conclusion, this test seems to support the idea that the larger spread is due to the stellar atmospheres and to the variability provoked by the photospheric convection and stellar pulsation."
 Having at hand an unprecedented large sample of GC red giants with homogeneously measured metal abundances. a logic step is to re-address the issue of an accurate metallicity scale for," Having at hand an unprecedented large sample of GC red giants with homogeneously measured metal abundances, a logic step is to re-address the issue of an accurate metallicity scale for"
We thanks Autojella Natta aud Steven Beckwith for reading the i.anuscript aud providing helpfu sugeestions.,We thank Antonella Natta and Steven Beckwith for reading the manuscript and providing helpful suggestions.
 We fiuk the OVRO/CARMA staff aud the CARMA observers for their assistance iu obtaining the daa., We thank the OVRO/CARMA staff and the CARMA observers for their assistance in obtaining the data.
 Support for the Combines Array for Rescarch in Alillnueter Astrononiv construction was derived from the Cordon ak Betty Moore. Forncdation. the Ikeuneth T. ax Eileen L. Norris Fouudatiou. the Associates of the Califoruia Iustitute of Technology. the states of California. IEHlhuoi. and Alarvlaucd. aud the National Science Foundation.," Support for the Combined Array for Research in Millimeter Astronomy construction was derived from the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. Norris Foundation, the Associates of the California Institute of Technology, the states of California, Illinois, and Maryland, and the National Science Foundation."
 Ongoing CARATA development aud operaions are supported by the National Science Founation under a cooperative aegrecluent (eraut AST 055200). and by the CARMA partuer universities.," Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement (grant AST 08-38260), and by the CARMA partner universities."
 This work was performed in partf uier contract with the Jet Propulsion Laboratory (JPL) funded bv NASA through the Michclson Fellowship Program., This work was performed in part under contract with the Jet Propulsion Laboratory (JPL) funded by NASA through the Michelson Fellowship Program.
 JPL is managed for NASA x the California Iustitute of Technology., JPL is managed for NASA by the California Institute of Technology.
[rom the distance. Az’ sav. its position advances in each period of the LAEW: with (4) given by (2.3)).,"from the distance, $\Delta z'$ say, its position advances in each period of the LAEW: ), with $\beta'(\chi)$ given by \ref{LAE5}) )."
" The mean dift velocity is ONz//2z. which depends on 4j. and tay. and is zero for uj, 0. The 4-velocity constructed from the 3-velocity 0Nz//2zx is rather cumbersome."," The mean drift velocity is $\Omega'\Delta z'/2\pi$ , which depends on $u'_0$, and $u_{\rm max}$, and is zero for $u'_0=0$ The 4-velocity constructed from the 3-velocity $\Omega'\Delta z'/2\pi$ is rather cumbersome."
 Alternatively. one might considera drift 4-velocitv found by dividing the displacement (5.2)) by the proper time elapsed in a waveperiod. Av sav.," Alternatively, one might considera drift 4-velocity found by dividing the displacement \ref{DeltaZ}) ) by the proper time elapsed in a waveperiod, $\Delta\tau$ say."
One has,"One has ,"
spindown and predicted a distinct achromatic feature.,spindown and predicted a distinct achromatic feature.
 A similar model was discussed earlier by Dai Lu (1998) who considered the energy injection to the forward shock by a millisecond pulsar with much a weaker magnetic field., A similar model was discussed earlier by Dai Lu (1998) who considered the energy injection to the forward shock by a millisecond pulsar with much a weaker magnetic field.
 This model is one of the candidates to interpret the majority of the X-ray plateaus observed in many Swift GRB afterglows (Zhang et al., This model is one of the candidates to interpret the majority of the X-ray plateaus observed in many Swift GRB afterglows (Zhang et al.
 2006; Nousek et al., 2006; Nousek et al.
 2006)., 2006).
 This model does not invoke the internal dissipation of the magnetar wind., This model does not invoke the internal dissipation of the magnetar wind.
" On the other hand, if the magnetar wind indeed dissipates internally before hitting the blastwave, it is possible that it would generate an ""internal"" X-ray plateau whose X-ray luminosity tracks the spindown luminosity if the energy dissipation and radiation efficiency remain constant."," On the other hand, if the magnetar wind indeed dissipates internally before hitting the blastwave, it is possible that it would generate an ""internal"" X-ray plateau whose X-ray luminosity tracks the spindown luminosity if the energy dissipation and radiation efficiency remain constant."
" If the magnetar undergoes direct collapse into a black hole before spin down, then the X-ray plateau would be followed by a very steep decay."," If the magnetar undergoes direct collapse into a black hole before spin down, then the X-ray plateau would be followed by a very steep decay."
" This is the light-curve feature we investigate and hereafter we call this feature an ""internal plateau""."," This is the light-curve feature we investigate and hereafter we call this feature an ""internal plateau""."
" In the eera the early X-ray light curve, observed within the first few hours of the GRB, has been found to be complex (e.g. Nousek 22006; O’Brien 22006)."," In the era the early X-ray light curve, observed within the first few hours of the GRB, has been found to be complex (e.g. Nousek 2006; O'Brien 2006)."
" The so-called canonical X-ray light curve observed in a significant fraction of GRBs (Evans 22009) has a short period of fast decay, often with flares superimposed, which are usually over within the first hour after the burst."," The so-called canonical X-ray light curve observed in a significant fraction of GRBs (Evans 2009) has a short period of fast decay, often with flares superimposed, which are usually over within the first hour after the burst."
" This is followed by a shallower decay period lasting from a few hours up to a day with a temporal decay index a~0.5 (where the X-ray flux, fvxv~*t~% and B is the spectral index)."," This is followed by a shallower decay period lasting from a few hours up to a day with a temporal decay index $\alpha \sim 0.5$ (where the X-ray flux, $f_{\nu} \propto \nu^{-\beta}
t^{-\alpha}$ and $\beta$ is the spectral index)."
 After this X-ray plateau there is a smooth transition to a modest power law decay of a~1—1.5., After this X-ray plateau there is a smooth transition to a modest power law decay of $\alpha \sim 1-1.5$.
" Willingale ((2007) found that the X-ray light curve of most GRBs, including those of the canonical form, can be represented by two components - the prompt and afterglow — plus flares (Section 2)."," Willingale (2007) found that the X-ray light curve of most GRBs, including those of the canonical form, can be represented by two components – the prompt and afterglow — plus flares (Section 2)."
" However, we find that in a small minority of bursts a period of relatively constant emission (compared to the general lightcurve) followed by a steep decline can be identified which does not fit this phenomenological model."," However, we find that in a small minority of bursts a period of relatively constant emission (compared to the general lightcurve) followed by a steep decline can be identified which does not fit this phenomenological model."
 The observed feature instead resembles the proposed signature of a magnetar spin down., The observed feature instead resembles the proposed signature of a magnetar spin down.
" Troja ((2007) found such a feature dominates the X-ray light curve of GRB 070110 from z1,000—20,000 seconds in and proposed it was due to a spinning down millisecond pulsar."," Troja (2007) found such a feature dominates the X-ray light curve of GRB 070110 from $\approx 1,000 - 20,000$ seconds in and proposed it was due to a spinning down millisecond pulsar."
" Starling ((2008) found a similar, earlier feature in GRB 070616 ending at about 600 seconds."," Starling (2008) found a similar, earlier feature in GRB 070616 ending at about 600 seconds."
 Liang et al. (, Liang et al. (
2007) systematically analyzed a sample of X-ray plateaus and identified several more plateaus that are followed by decays with slopes steeper than a=—3.,2007) systematically analyzed a sample of X-ray plateaus and identified several more plateaus that are followed by decays with slopes steeper than $\alpha = -3$.
 It is the combination of a plateau followed by a steep decay which distinguishes these from the canonical behaviour., It is the combination of a plateau followed by a steep decay which distinguishes these from the canonical behaviour.
 We regard these objects as candidate internal plateaus., We regard these objects as candidate internal plateaus.
 We have conducted a systematic investigation of the GRB X-ray light curves observed by uup to the end of 2008., We have conducted a systematic investigation of the GRB X-ray light curves observed by up to the end of 2008.
 Using an automated fitting procedure in this paper we identify 10 bursts which may have an emission component powered by magnetar spin-down dominating the light curve for some period of time in the form of an internal plateau., Using an automated fitting procedure in this paper we identify 10 bursts which may have an emission component powered by magnetar spin-down dominating the light curve for some period of time in the form of an internal plateau.
 Assuming this internal plateau is caused by the spinning down of a magnetar we use its properties to constrain the magnetic field and initial spin period of the magnetar., Assuming this internal plateau is caused by the spinning down of a magnetar we use its properties to constrain the magnetic field and initial spin period of the magnetar.
 The criteria for selecting those GRBs with internal plateaus is discussed in Section 2., The criteria for selecting those GRBs with internal plateaus is discussed in Section 2.
 Section 3 compares the properties of the internal plateaus with the magnetar model and we discuss the implications in Section 4., Section 3 compares the properties of the internal plateaus with the magnetar model and we discuss the implications in Section 4.
 Over of GRB X-ray light curves are well described by a two component model with a prompt and an afterglow component as described in Willingale ((2007)., Over of GRB X-ray light curves are well described by a two component model with a prompt and an afterglow component as described in Willingale (2007).
 These components are described by an exponential that relaxes into a power law; this function can be expressed for the prompt, These components are described by an exponential that relaxes into a power law; this function can be expressed for the prompt
are very close to the largest outflow parameters ever measured (see e.g. Lóppez-Sepulcre et al.,are very close to the largest outflow parameters ever measured (see e.g. Lóppez-Sepulcre et al.
" 2009 and Wu et al. 2004)),"," \cite{lopsep} and Wu et al. \cite{wu04}) ),"
 corresponding to a YSO powering the outflow of at least ~10? Lo., corresponding to a YSO powering the outflow of at least $\sim10^5~L_\odot$ .
" More precisely, using Wu et al. (2004))"," More precisely, using Wu et al. \cite{wu04}) )"
" relationship between bolometric luminosity (1ο) of the powering source and outflow momentum rate, from the value in Table 3 one obtains [μοι=6x10°Lo."," relationship between bolometric luminosity $L_{\rm bol}$ ) of the powering source and outflow momentum rate, from the value in Table \ref{tout} one obtains $L_{\rm bol}=6\times10^6$."
. This is much higher than the luminosity estimated for the HMC (~10? Lo) by Osorio et, This is much higher than the luminosity estimated for the HMC $\sim10^5~L_\odot$ ) by Osorio et
value. AQ can be estimated as where A@ denotes the differential rotation In order to confirm eq. (49)).,"value, $\Delta\Omega$ can be estimated as where $\Delta \theta$ denotes the differential rotation In order to confirm eq. \ref{deltaomega}) ),"
 we conduct two sets of simulations. firstly. varving the value of turbulent. viscosity (vy). and secondly the amplitude of the A effect (Ag).," we conduct two sets of simulations, firstly varying the value of turbulent viscosity $\nu_\mathrm{0v}$ ), and secondly the amplitude of the $\Lambda$ effect $\Lambda_0$ )."
 Note that the setting for turbulent viscosity does not reflect à real situation. since the coefficients ol turbulent viscosity and the A effect should have a common value.," Note that the setting for turbulent viscosity does not reflect a real situation, since the coefficients of turbulent viscosity and the $\Lambda$ effect should have a common value."
 Nonetheless. (his is necessary for the purpose of our investigation.," Nonetheless, this is necessary for the purpose of our investigation."
 The simulation results are shown in Figures 12 and 13.., The simulation results are shown in Figures \ref{vis_vari} and \ref{lam_vari}.
" We obtain AQxr4, and AQxAR which are consistent witheq. (49))."," We obtain $\Delta \Omega \propto \nu_\mathrm{0v}^{-0.88}$ and $\Delta \Omega \propto \Lambda_0^{1.1}$, which are consistent witheq. \ref{deltaomega}) )."
 Fig., Fig.
 14 shows the results of the dependence of AQ on Q) (Cases 1-5)., \ref{angular_difference} shows the results of the dependence of $\Delta \Omega$ on $\Omega_0$ (Cases 1-5).
 Asterisks denote the difference al the surface between the equator ancl the pole. squares show the difference between the equator and the colatitude @=45°. and triangles are the difference between the equator and the colatitude @=60°.," Asterisks denote the difference at the surface between the equator and the pole, squares show the difference between the equator and the colatitude $\theta=45^\circ$, and triangles are the difference between the equator and the colatitude $\theta=60^\circ$."
 The difference at low latitudes (squares and triangles) monotonically increases with stellar angular velocity., The difference at low latitudes (squares and triangles) monotonically increases with stellar angular velocity.
 However this is not the case for angular velocity difference between the equator and the pole (asterisk)., However this is not the case for angular velocity difference between the equator and the pole (asterisk).
 As we discussed in 8?2?.. when stellar rotation velocity is large. the Tavlor-Proudman state is achieved. meaning the gradient of angular velocity at the surface concentrates in lower latitudes.," As we discussed in \ref{taylor}, when stellar rotation velocity is large, the Taylor-Proudman state is achieved, meaning the gradient of angular velocity at the surface concentrates in lower latitudes."
 Due to this concentration. A@ becomes smaller in Eq. (49))," Due to this concentration, $\Delta\theta$ becomes smaller in Eq. \ref{deltaomega}) )"
 with larger values of O4., with larger values of $\Omega_0$.
 Thus. AQ) does not show an explicit dependence on Ou.," Thus, $\Delta \Omega_0$ does not show an explicit dependence on $\Omega_0$."
 At low latitudes. A@ is fixed and the angular velocity dillerence increases with stellar angular velocity.," At low latitudes, $\Delta\theta$ is fixed and the angular velocity difference increases with stellar angular velocity."
 We obtain AQxὉμ (between the equator and the colatitude @=45°: squares) and AQxOt? (between the equator and the colatitude @= 60°: triangles)., We obtain $\Delta\Omega\propto\Omega_0^{0.43}$ (between the equator and the colatitude $\theta=45^\circ$: squares) and $\Delta\Omega\propto\Omega_0^{0.55}$ (between the equator and the colatitude $\theta=60^\circ$ : triangles).
" This indicates that AQ/Q, decreases with stellar angular velocity.", This indicates that $\Delta\Omega/\Omega_0 $ decreases with stellar angular velocity.
 These resulls are consistent wilh previous stellar observations 2005)..," These results are consistent with previous stellar observations \citep{1996ApJ...466..384D,2003A&A...398..647R,2005MNRAS.357L...1B}. ."
SiO and CS. being an order of magnitude less: the predicted LIC3N abundance is several orders of magnitude less than the upper limit set by Bachillerctal.(1997).,"SiO and CS, being an order of magnitude less; the predicted ${\rm HC_3N}$ abundance is several orders of magnitude less than the upper limit set by \citet{bachiller.et.al97}."
 However the nmiocels severely underpredict the abundancees of HON. LINC. CN and HCO.," However the models severely underpredict the abundances of HCN, HNC, CN and ${\rm HCO^{+}}$."
 This indicates that these molecules whose earlier clump abundances were comparable to the observed levels survive for longer than our simple mocel predicts., This indicates that these molecules whose earlier clump abundances were comparable to the observed levels survive for longer than our simple model predicts.
 The chemistry of the gas in the slow wind ejected [rom a star in the AGB phase is strongly allected by photoprocesses driven by the hardening radiation field of the central star. as the system ds transformed. into a PPN and then to a PN.," The chemistry of the gas in the slow wind ejected from a star in the AGB phase is strongly affected by photoprocesses driven by the hardening radiation field of the central star, as the system is transformed into a PPN and then to a PN."
 UW othe slow wind density distribution is smoothly monotonic. then the degradation product. molecules in the eas are confined to the PPN stage.," If the slow wind density distribution is smoothly monotonic, then the degradation product molecules in the gas are confined to the PPN stage."
 However. if the slow wind is clumpy. perhaps because of clumpiness in the precursor AGB atmosphere. then molecules in clumps are protected against photoclissociation until significantly later times. and may be present in the PN phase of evolution.," However, if the slow wind is clumpy, perhaps because of clumpiness in the precursor AGB atmosphere, then molecules in clumps are protected against photodissociation until significantly later times, and may be present in the PN phase of evolution."
 We have presented. Lists of molecules that may. represent tracers of such clumpiness., We have presented lists of molecules that may represent tracers of such clumpiness.
 Benzene survives the transition into the PN phase. suggesting that some long-lived molecules may last long enough to be injected into the general interstellar mecium.," Benzene survives the transition into the PN phase, suggesting that some long-lived molecules may last long enough to be injected into the general interstellar medium."
 A comparison with observations suggests that this model can. at the PPN and voung PN stage. correctly predict the abundance. of several molecular species. to within an order of magnitude.," A comparison with observations suggests that this model can, at the PPN and young PN stage, correctly predict the abundance of several molecular species to within an order of magnitude."
 At [ater stages. the model severely. underestimates the abundance of some short-chain molecules.," At later stages, the model severely underestimates the abundance of some short-chain molecules."
 Thus the shielding of molecules in elumps may be even more effective than as described in this simple moclel., Thus the shielding of molecules in clumps may be even more effective than as described in this simple model.
 ‘The original aim of this work was to chemically identify when clumps form in PNe., The original aim of this work was to chemically identify when clumps form in PNe.
 Ehe comparison of the results in the early stages with the data from CRLGIS and NGC 7027 supports the hypothesis that the clumps form at an early stage: if they formed. later out of low extinction gas. the abuncdances of the molecules would be significantly depleted compared with the values observed.," The comparison of the results in the early stages with the data from CRL618 and NGC 7027 supports the hypothesis that the clumps form at an early stage; if they formed later out of low extinction gas, the abundances of the molecules would be significantly depleted compared with the values observed."
 As noted. above. the evidence of the continuing presence of molecules in an evolved. PN like the Helix indicates that the protective ability of the clumps persists until late stages.," As noted above, the evidence of the continuing presence of molecules in an evolved PN like the Helix indicates that the protective ability of the clumps persists until late stages."
 Future work in this field should include a more realistic physical model for the different post-AGB phases to account for the survival of molecules in the eblumps., Future work in this field should include a more realistic physical model for the different post-AGB phases to account for the survival of molecules in the clumps.
 Lt should also be possible to mocel individual sources and thus carey out a much more detailed comparison with observations both from existing facilities and. also from forthcoming facilities such as ALMA., It should also be possible to model individual sources and thus carry out a much more detailed comparison with observations both from existing facilities and also from forthcoming facilities such as ALMA.
 We thank the referee. Pierre Cox. for a very constructive report that led to an improved. paper.," We thank the referee, Pierre Cox, for a very constructive report that led to an improved paper."
 We also thank John Dyson and Jonathan Rawlines for useful discussions., We also thank John Dyson and Jonathan Rawlings for useful discussions.
 MPH acknowledges the support of PPARC., MPR acknowledges the support of PPARC.
 SV thanks the Ealian Space Agency (ASL) and DAW thanks the Leverhulme Trust for financial support., SV thanks the Italian Space Agency (ASI) and DAW thanks the Leverhulme Trust for financial support.
 PC was a Fellow supported. by the European Commission MIU programme in. Astrophysical Chemistry while this work was carried out., PC was a Fellow supported by the European Commission TMR programme in Astrophysical Chemistry while this work was carried out.
The radio emission [rom star-forming galaxies is of great importance as an apparently unbiased tracer of the star-formation rate (SER). being neither obscured nor enhanced bv the presence. of dust. (e.g. Condon 1992).,"The radio emission from star-forming galaxies is of great importance as an apparently unbiased tracer of the star-formation rate (SFR), being neither obscured nor enhanced by the presence of dust (e.g. Condon 1992)."
 Αι sub-milliJansky. luxes. deep radio surveys detect. large number of star-forming galaxies. which overtake radio galaxies ancl other active galactic nuclei (AGN) as the most. numerous type of source.," At sub-milliJansky fluxes, deep radio surveys detect large number of star-forming galaxies, which overtake radio galaxies and other active galactic nuclei (AGN) as the most numerous type of source."
 Hence a deep ractio-selected source sample will clliciently select the most intensely star-forming galaxies in the survey volume., Hence a deep radio-selected source sample will efficiently select the most intensely star-forming galaxies in the survey volume.
 Llamuner et al. (, Hammer et al. (
1995) investigated the optical IDs. of a saniple of radio sources detected in a deep Very Large Array (VLA) survey of one of the Canada-Erance ecshift Survey fields. to a limit in integrated flux £(5.0Cllz)> l6j]y.,"1995) investigated the optical IDs of a sample of radio sources detected in a deep Very Large Array (VLA) survey of one of the Canada-France Redshift Survey fields, to a limit in integrated flux $F(5.0~\rm
GHz)\geq 16\mu Jy$ ."
 Spectroscopy. revealed these galaxies to be mixture of (i) giant ellipticals. presumably with ACN. (ii) star-forming disks at redshifts out to z=1.16. ancl (iii) lower redshift. high-excitation emission-line galaxies.," Spectroscopy revealed these galaxies to be mixture of (i) giant ellipticals, presumably with AGN, (ii) star-forming disks at redshifts out to $z=1.16$, and (iii) lower redshift, high-excitation emission-line galaxies."
 Winclhorst et al. (, Windhorst et al. (
1985) and. Richards et al. (,1985) and Richards et al. (
1998). obtained similar lindings | although with fewer AGN and more clisk galaxies al even fainter Bux limits.,1998) obtained similar findings – although with fewer AGN and more disk galaxies – at even fainter flux limits.
 Using HIST imaging. they estimated that at least 60 per cent of the radio-selected ealaxies were interacting.," Using HST imaging, they estimated that at least $\sim 60$ per cent of the radio-selected galaxies were interacting."
 In à previous paper (Roche. Lowenthal ancl Ixoo 2002:," In a previous paper (Roche, Lowenthal and Koo 2002;"
positive vertex deviation. similar to (hat found by Morenoetal.(1999). [or the GB stars alter eliminating the ones belonging to the Pleiades moving group.,"positive vertex deviation, similar to that found by \citet{Mor99} for the GB stars after eliminating the ones belonging to the Pleiades moving group."
" So when we work with samples (either of GB or LGD stars) in which a single moving group clearly dominates over the others the vertex deviation is positive aud close to /,~207. which is the value expected from the dynamic equations of the svstem for this age eroup."," So when we work with samples (either of GB or LGD stars) in which a single moving group clearly dominates over the others the vertex deviation is positive and close to $l_{v} \sim 20\degr$, which is the value expected from the dynamic equations of the system for this age group."
 The negative vertex deviation of the GB seems to be originated by (he relative position of the centroids of (wo moving eroups rather than by (he distribution of the residual velocities as a whole., The negative vertex deviation of the GB seems to be originated by the relative position of the centroids of two moving groups rather than by the distribution of the residual velocities as a whole.
 We must note that the analvsis of the voung stars -specially those belonging to the GB- based on the Schwarzschild distribution from which the velocity ellipsoid is tailorecl. is not the best suited to describe the reality of the svstem.," We must note that the analysis of the young stars -specially those belonging to the GB- based on the Schwarzschild distribution from which the velocity ellipsoid is tailored, is not the best suited to describe the reality of the system."
 The velocity field is dominated. as we have seen. by the presence of moving groups. ancl thus itis lar from the hypothesis of a homogeneous and stationary svstem.," The velocity field is dominated, as we have seen, by the presence of moving groups, and thus it is far from the hypothesis of a homogeneous and stationary system."
 The problem thus lies in that the classic analvsis of the velocity ellipsoid is applied to a set of moving eroups., The problem thus lies in that the classic analysis of the velocity ellipsoid is applied to a set of moving groups.
 Already Mihalas&Binney(1981) pointed out that (he cause of the vertex deviation was the existence of moving groups in the Galactic velocity field., Already \citet{Mih81} pointed out that the cause of the vertex deviation was the existence of moving groups in the Galactic velocity field.
 Bul even though the velocity ellipsoid does not strictly correspond to a physice reality in our case. (hat doesn’t invalidate our result. that is. the fundamental contribution of the GB to the negative vertex deviation for the O and D stars.," But even though the velocity ellipsoid does not strictly correspond to a physic reality in our case, that doesn't invalidate our result, that is, the fundamental contribution of the GB to the negative vertex deviation for the O and B stars."
 Thus. the different values of the vertex deviation found in the literature can be explained according to the dillerent proportions of GB stars present in the respective samples.," Thus, the different values of the vertex deviation found in the literature can be explained according to the different proportions of GB stars present in the respective samples."
 This translates the question about (he origin of the negative vertex deviation to the investigation ol the origin of the moving groups., This translates the question about the origin of the negative vertex deviation to the investigation of the origin of the moving groups.
 Although the latter is out of the scope of this paper. we want to stress that the more we study in detail the nature of the GB. the more we find indications that we must probe both its nature and origin as a set of moving eroups.," Although the latter is out of the scope of this paper, we want to stress that the more we study in detail the nature of the GB, the more we find indications that we must probe both its nature and origin as a set of moving groups."
 Another striking result is the inclination of the GD's velocity ellipsoid of about with respect to the CV. plane. although (his value does not have a great. statistical significance.," Another striking result is the inclination of the GB's velocity ellipsoid of about $10\degr \pm 12\degr$ with respect to the $UV$ plane, although this value does not have a great statistical significance."
 While the third axis (W) of the LGD ellipsoid merely shows an inclination of 27. the GB's ellipsoid inclination (107)) reminds that of the spatial svstem wilh respect {ο the Galactic plane (αμ~147 as we found in Paper I).," While the third axis $W'$ ) of the LGD ellipsoid merely shows an inclination of $2\degr$, the GB's ellipsoid inclination ) reminds that of the spatial system with respect to the Galactic plane $i_{GB} \sim 14\degr$ as we found in Paper I)."
 Some authors have pointed out that the GB could be oscillating as a whole around the Galactic plane (ComerónGrenier 2003).," Some authors have pointed out that the GB could be oscillating as a whole around the Galactic plane \citep{Com99,Per03}."
. But our result. does not discard a different interpretation in terms of a juxtaposition of moving groups: this links back to the idea (hat a descriptive parameter of the GB -such as this inclination of the minor axis of the ellipsoid- can be interpreted in terms of the relative position (in the velocity space) of two moving groups., But our result does not discard a different interpretation in terms of a juxtaposition of moving groups: this links back to the idea that a descriptive parameter of the GB -such as this inclination of the minor axis of the ellipsoid- can be interpreted in terms of the relative position (in the velocity space) of two moving groups.
to fit the observed SEDs (?)..,to fit the observed SEDs \citep{2007ApJS..169..328R}.
 Pherefore. we do not claim that the models presented here are a unique solution for the disk structure.," Therefore, we do not claim that the models presented here are a unique solution for the disk structure."
 Instead. they should be seen as exemplary fits to the SED.," Instead, they should be seen as exemplary fits to the SED."
 For the objects. considered. here. further complications are introduced. by the significant. variability in the optical and near-infrared. hands.," For the objects considered here, further complications are introduced by the significant variability in the optical and near-infrared bands."
 To facilitate the modeling. we fix a number of parameters.," To facilitate the modeling, we fix a number of parameters."
 Following ?.. the average interstellar extinction is assumed to be chy—0.5 mmag and the distance ppc.," Following \citet{2008AJ....135.1616S}, the average interstellar extinction is assumed to be $A_V=0.5$ mag and the distance pc."
 Since the near- and mid-infrared. SED is not. sensitive to global disk. parameters. we fix the total disk mass at 10AI. and the outer disk radius at AU.," Since the near- and mid-infrared SED is not sensitive to global disk parameters, we fix the total disk mass at $10^{-3}\,M_{\odot}$ and the outer disk radius at AU."
 σον! is a cluster. without evidence for deeply. embedded: sources. therefore we fit the SEDs without including an envelope.," $\sigma$ Ori is a cluster without evidence for deeply embedded sources, therefore we fit the SEDs without including an envelope."
 Additionally. we assume that accretion rates are faürlv low and disk heating is thus dominated. by radiation from the central. sources.," Additionally, we assume that accretion rates are fairly low and disk heating is thus dominated by radiation from the central sources."
 With these assumptions. the main [free parameters are the Daring power. the inclination. and the inner clisk radius (?)..," With these assumptions, the main free parameters are the flaring power, the inclination, and the inner disk radius \citep{2002ApJ...567.1183W}."
 To explore possiblereasons for the discrepancy in the 24;j0n lux. we match the SEDs of 2 and #333. two objects with similar effective temperature. with mocdels.," To explore possiblereasons for the discrepancy in the $24\,\mu m$ flux, we match the SEDs of 2 and 33, two objects with similar effective temperature, with models."
 We find that only one parameter. the [aring power. needs to be mocified to account for the difference.," We find that only one parameter, the flaring power, needs to be modified to account for the difference."
 This is illustrated in the two panels in Fig. δ.., This is illustrated in the two panels in Fig. \ref{f10}.
 Phe best fit of the optical/near- data is produced with Z;;=2900 KK for object #£22 ancl Why for object. $333., The best fit of the optical/near-infrared data is produced with $T_{\mathrm{eff}} = 2900$ K for object 2 and K for object 33.
 For the two disks we use essentially the same model parameters. with the exception of the Παπάς power.," For the two disks we use essentially the same model parameters, with the exception of the flaring power."
 We assume fy=0.025. 39=1.2 for 22 and fy=0.02 and 3%=l.l for 4333.," We assume $h_0=0.025$, $\beta=1.2$ for 2 and $h_0=0.02$ and $\beta=1.1$ for 33."
" At LAA distance from the star. this results in a scaleheight of 2, or objec 2 and LOR, for object 333."," At AU distance from the star, this results in a scaleheight of $\,R_\star$ for object 2 and $\,R_\star$ for object 33."
 As seen in the Mots. such a choice of parameters can nicely explain the ).6cledlex Hux dillerence at 244700 in these two SEDs.," As seen in the plots, such a choice of parameters can nicely explain the dex flux difference at $\mu m$ in these two SEDs."
 The arger scaleheight leads to a larger surface area that is irradiated by the stellar light and thus an increased. level of micd-infrared (lux., The larger scaleheight leads to a larger surface area that is irradiated by the stellar light and thus an increased level of mid-infrared flux.
 Similarly. the SED of source #443 can ος explained. with an enhanced Uaring power.," Similarly, the SED of source 43 can be explained with an enhanced flaring power."
 This provides additional constraints on the scenario oesented. in Sect., This provides additional constraints on the scenario presented in Sect.
 5.5. to explain the variability., \ref{comp} to explain the variability.
 While he photometric variations for 333 and 3xÉéll9 are best explained by hot spots. object 422 likely undergoes changes in cireumstellar extinction. caused by inhomogenitics at the inner disk edge (see Sect. 5.5)).," While the photometric variations for 33 and 19 are best explained by hot spots, object 2 likely undergoes changes in circumstellar extinction, caused by inhomogenities at the inner disk edge (see Sect. \ref{comp}) )."
 For this to have an cllect on the optical/near-infrared fluxes. the line of sight has to intercept parts of the accretion disk (Sect. 5.3)).," For this to have an effect on the optical/near-infrared fluxes, the line of sight has to intercept parts of the accretion disk (Sect. \ref{ext}) )."
 A disk with comparatively high degree of Haring and thus relatively large scaleheights would naturally comply with this prerequisite., A disk with comparatively high degree of flaring and thus relatively large scaleheights would naturally comply with this prerequisite.
 On the other hand. relatively small scaleheights favour the possibility. of having a unobstructed. line of sight towards the hot spots in the inner accretion zone. as itis the case for #119 and #333.," On the other hand, relatively small scaleheights favour the possibility of having a unobstructed line of sight towards the hot spots in the inner accretion zone, as it is the case for 19 and 33."
 Hence. the results from the SED mocleling fit into the interpretation given in Sect. 5.5..," Hence, the results from the SED modeling fit into the interpretation given in Sect. \ref{comp}."
 In Sect. 4.0.," In Sect. \ref{var},"
 we have identified three low-mass objects in a OOri showing strong variability with amplitudes exceeding mmag in the J-band. two of them with masses close to or below the substellar limit.," we have identified three low-mass objects in $\sigma$ Ori showing strong variability with amplitudes exceeding mag in the J-band, two of them with masses close to or below the substellar limit."
 Their lighteurves are not strictly periodic. instead. we see evidence for multiple periods with additional irregular contributions.," Their lightcurves are not strictly periodic, instead we see evidence for multiple periods with additional irregular contributions."
 A few more likely VLM objects in σ OOri exhibit small-scale photometric variations with amplitudes =0.1 mmag., A few more likely VLM objects in $\sigma$ Ori exhibit small-scale photometric variations with amplitudes $\la 0.1$ mag.
 This type of variability. is in line with the phenomenological classification of variability. in “Po Tauri stars. as given by 7: The low-level modulations can be classified as tvpe E variations. mostly seen in cliskless stars. while the large-scale. partly irregular modulations resemble the type IE. variability. predominantly seen in clisk-bcaring objects.," This type of variability is in line with the phenomenological classification of variability in T Tauri stars, as given by \citet{1994AJ....108.1906H}: The low-level modulations can be classified as 'type I' variations, mostly seen in diskless stars, while the large-scale, partly irregular modulations resemble the 'type II' variability, predominantly seen in disk-bearing objects."
 Thus. the variability characteristic of very. low mass stars and brown cdwarfs is completely analogous to the one seen in solar-mass T Tauri stars.," Thus, the variability characteristic of very low mass stars and brown dwarfs is completely analogous to the one seen in solar-mass T Tauri stars."
 For two highly variable objects. the colour variability is consistent with hot spots. formed. by the shocklront of he aceretion How.," For two highly variable objects, the colour variability is consistent with hot spots, formed by the shockfront of the accretion flow."
 Since hot spots are co-rotating with he objects. they are expected. to induce a periodic [ux modulation: additionally. changes in the [low properties can cause irregular variations.," Since hot spots are co-rotating with the objects, they are expected to induce a periodic flux modulation; additionally, changes in the flow properties can cause irregular variations."
 For object. 333. possibly a oxown dwarf. the best fit spot temperature is 0000 Why with a filling factor of 20%..," For object 33, possibly a brown dwarf, the best fit spot temperature is $\sim 6000$ K with a filling factor of $\ga 20$ ."
 While this value is similar o the constraints for hot spots in other brown. clwarls IxIx. 7. see also ?)). it dillers considerably from. assumptions mace in models for Ho. profiles for substellar objects (10000-12000Ix.2)..," While this value is similar to the constraints for hot spots in other brown dwarfs K, \citet{2008arXiv0801.3525H} see also \citet{2008MNRAS.tmp..853K}) ), it differs considerably from assumptions made in models for $\alpha$ profiles for substellar objects \citep[10000-12000\,K,][]{2003ApJ...592..266M}."
 Phis object also has an upper imit on the accretion rate of <5-10IA. ο (2). illustrating that large-scale. variability due to hotspots can » generated even at very low accretion rates.," This object also has an upper limit on the accretion rate of $<5 \cdot 10^{-11}\,M_{\odot}$ $^{-1}$ \citep{2008A&A...481..423G}, illustrating that large-scale variability due to hotspots can be generated even at very low accretion rates."
 For the third highly variable candidate. an object close o the substellar boundary. the dominant cause of variability is probably a change in the circumstellar extinction. bv ον~ 2mmae.," For the third highly variable candidate, an object close to the substellar boundary, the dominant cause of variability is probably a change in the circumstellar extinction by $A_V \sim 2$ mag."
 This may be related. to inhomosgenities in he inner disk at raciii AAU. which is close to the inner edge of the disk.," This may be related to inhomogenities in the inner disk at radii AU, which is close to the inner edge of the disk."
 SED modeling confirms an enhanced disk scaleheight for this object. which favours the detection of disk inhomogenities in lighteurves (Sect. 6)).," SED modeling confirms an enhanced disk scaleheight for this object, which favours the detection of disk inhomogenities in lightcurves (Sect. \ref{sed}) )."
 lt ds. well-established that a substantial number of accreting objects exhibits. high-level variability in optical ancl near-infrarecl bands (c.g.οο).," It is well-established that a substantial number of accreting objects exhibits high-level variability in optical and near-infrared bands \citep[e.g.][]{1994AJ....108.1906H,1995A&A...299...89B,2001AJ....121.3160C,2008A&A...485..155A}."
 Our results demonstate that this behaviour extends into the substellar regime., Our results demonstate that this behaviour extends into the substellar regime.
 The fraction of accreting. objects in the σοΟνὶ cluster is for IN3-MB5 stars and for objects later than ALS. based on the Lla emission (?).. ," The fraction of accreting objects in the $\sigma$ Ori cluster is for K3-M5 stars and for objects later than M5, based on the $\alpha$ emission \citep{2003AJ....126.2997B}. ."
Thus. we expect that the total sample of 135 objects published in SEO4 containsabout 15-20 accretors (accounting for contamination).," Thus, we expect that the total sample of 135 objects published in SE04 containsabout 15-20 accretors (accounting for contamination)."
 We find four sources with hieh-level variability confirmed in at least two independent monitoring campaigns 22," We find four sources with high-level variability confirmed in at least two independent monitoring campaigns 2, 19, 33,"
Our derivation is closely analogous to the Nakar et al (2002) orphan rate calculation. aud includes the approximation that GRB jets have a standard energy.,"Our derivation is closely analogous to the Nakar et al (2002) orphan rate calculation, and includes the approximation that GRB jets have a standard energy."
 Lu this model. the light curve of off axis orphan afterglows assumes a staucdard form.," In this model, the light curve of off axis orphan afterglows assumes a standard form."
" The masxinitun time fyget.) at which an orphan is detectable is then cleterminecl uniquely by the redshift z of the burst aud the maguitude limit i"" of the observatious."," The maximum time $t_{max}(z,m)$ at which an orphan is detectable is then determined uniquely by the redshift $z$ of the burst and the magnitude limit $m$ of the observations."
 For completeness. we reproduce Lere the result of inverting equatious 6-7 of NPGO?: or the case where vy.«r. aud or the case 4.>7.," For completeness, we reproduce here the result of inverting equations 6–7 of NPG02: for the case where $\nu_c < \nu$, and for the case $\nu_c > \nu$."
 Quantities in these equations are scaled to dimeusionless or ces unit. vali whose logarithius are indicated in the subscripts: p is the slope of the power law clistribution of euergies for electrous accelerated in the forward shock of the expanding GRB remuaut: e.4 aud εµ.ο are the fractious of iuterual eue‘oy eolng iuo relativistic electrois and magnetic fields a this shock. scaled to uuits of 0.] aud 0.01 respectively: my is the number density of the ambient inedium: Epis is he initial kinetic energy of he ejecta in units of 107ere. 141.> ds the observed [recteucvy in units of 10MTHy. 2 ix the GRB reshift. and dox is the luminosity cistaice in units of 107€Lu.," Quantities in these equations are scaled to dimensionless or cgs unit values whose logarithms are indicated in the subscripts: $p$ is the slope of the power law distribution of energies for electrons accelerated in the forward shock of the expanding GRB remnant; $\epsilon_{e,-1}$ and $\epsilon_{B,-2}$ are the fractions of internal energy going into relativistic electrons and magnetic fields at this shock, scaled to units of $0.1$ and $0.01$ respectively; $n_0$ is the number density of the ambient medium; $E_{0,50.7}$ is the initial kinetic energy of the ejecta in units of $10^{50.7} \erg$, $\nu_{14.7}$ is the observed frequency in units of $10^{14.7}\Hz$, $z$ is the GRB redshift, and $d_{L,28}$ is the luminosity distance in units of $10^{28} \cm$."
 The funeious go(p)=10PPPOy—0.98)[(p2/(1?t and g2)/(p—-1)p vary slowly with p for p=2.2 but becorje uucdefine Dpx2. sisince the assiüued physical 11odel for the radiatiig electron population gives a clivergent total electron energy if p<2.," The functions $g_0(p) =10^{-0.56 p} (p-0.98) [(p-2)/(p-1)]^{p-1}$ and $g_1(p) = 10^{-0.31 p} (p-0.04) [(p-2)/(p-1)]^{p-1}$ vary slowly with $p$ for $p \ga 2.2$ but become undefined for $p \le 2$, since the assumed physical model for the radiating electron population gives a divergent total electron energy if $p \le 2$."
 Aun interesting feature of fyyey is that it declines extremely slowly at high recdslifts.," An interesting feature of $t_{1,max}$ is that it declines extremely slowly at high redshifts."
 /παρσιΠΠ) is plotted for a few values of min figure 3..," $t_{1,max}(z,m)$ is plotted for a few values of $m$ in figure \ref{tmax}."
" Because our calculation of /TO as assuined the ""universa- lisht curve for a ofL-axis ]et. it is only valid for Manan>/jjj."," Because our calculation of $t_{1,max}$ has assumed the “universal” light curve for an off-axis jet, it is only valid for $t_{1,max} > \tjet$."
" We 1erelo'e overplot ljhobs for rest-frame fje,=0.τ day."," We therefore overplot $t_{jet,obs}$ for rest-frame $\tjet = 0.7$ day."
" Cousider first the distribtion of twn-on /,,, uuder the approximation. hat tle alierglow becomes visible when Opp.L/P.", Consider first the distribution of “turn-on $t_{on}$ under the approximation that the afterglow becomes visible when $\theta_{obs} = 1 / \Gamma$.
 In the late time evolution of a spreading QGRB jet. Px/E?> (Rhoacls 1997. 1999).," In the late time evolution of a spreading GRB jet, $\Gamma \propto t^{-1/2}$ (Rhoads 1997, 1999)."
" The clistribution of off-axis angles should go as (Needui,XOops.", The distribution of off-axis angles should go as $dN_{GRB} / d\theta_{obs} \propto \theta_{obs}$.
" Uwe consider : vsinele redshift au a single jet euergy. we then obtain Here Neva has uuits of events per day per unit volume. while dNcngChan is"" in. events per day> per unit volume."," If we consider a single redshift and a single jet energy, we then obtain Here $N_{GRB}$ has units of events per day per unit volume, while $dN_{GRB} \over dt_{on}$ is in events per $^2$ per unit volume."
 The factor of (1+=)1 in the rate is required to account for cosmological time dilatiou of the observed turn-on time., The factor of $(1+z)^{-1}$ in the rate is required to account for cosmological time dilation of the observed turn-on time.
"Receutly Falthanunar (Fallamar2006) conumnented on the unjustified use of the motion of magnetic ficld lines iu plasnia. ollowius Alivel. who vigorously warned against heuyustified use of concept ""Erozeu-iu niaeuetic field” iu his lae vears (Alfve11976).","Recently Falthammar \citep{b01} commented on the unjustified use of the motion of magnetic field lines in plasma, following Alfven, who vigorously warned against the unjustified use of concept “Frozen-in magnetic field"" in his late years \citep{b02}."
. However. tie scientists pav little attention to the sugeestions.," However, the scientists pay little attention to the suggestions."
 The concept of maenctic field line motion was applied exteusivev in space plasma plivsies ad solar owesies., The concept of magnetic field line motion was applied extensively in space plasma physics and solar physics.
 Magnelc LOCOmection becomes rot topic in explaining nagjetic activities., Magnetic reconnection becomes hot topic in explaining magnetic activities.
 The xoblem is the eeueralized Ohuis law of MIID theo:v which shows the couple )otween velocity of plasuτὰ and the magnetic field., The problem is the generalized Ohm's law of MHD theory which shows the couple between velocity of plasma and the magnetic field.
 NIID(Magnuetolivdrocvauandes) is the plysical-naticmaatical framework. that concerns the αναος of uaaenetic fields in elecvically couductiug fiuids. e.g.àY ii plaxuas aud licuil metals.," MHD(Magnetohydrodynamics) is the physical-mathematical framework that concerns the dynamics of magnetic fields in electrically conducting fluids, e.g. in plasmas and liquid metals."
 One of the LOS famous scholars associated with AID was the Swedish plivsicist Haunes Alfven. who," One of the most famous scholars associated with MHD was the Swedish physicist Hannes Alfven, who"
signal to noise to allow reliable local surface deusity measurements of the stream.,signal to noise to allow reliable local surface density measurements of the stream.
 Multiple image eravitational leusing by a galaxy of a backerounud quasars is another probe of sub-halos., Multiple image gravitational lensing by a galaxy of a background quasars is another probe of sub-halos.
 Mao&Schueider(1998) poiuted out that the anomalous flux ratios relative to a smooth poteutial model that accurately predicts the locations of the miages are the natural consequence of substructure in galaxy halos., \cite{MS:98} pointed out that the anomalous flux ratios relative to a smooth potential model that accurately predicts the locations of the images are the natural consequence of substructure in galaxy halos.
" A statistical analysis of available stroug leus svstems (Dalal&Kochauek2002:ItochanckDalal2001) assis NEW density profile sub-halo profiles fouud that about of the halo mass (with a verv large uucertaiuty) is in substructure,"," A statistical analysis of available strong lens systems \citep{DK:02,KD:04} assuming NFW density profile sub-halo profiles found that about of the halo mass (with a very large uncertainty) is in substructure."
 Leusing is very sensitive to the ceutral density profile which Springeletal.(2005) found to be close to an Eimasto profile. which is shallower thau the NFW profile.," Lensing is very sensitive to the central density profile which \citet{Aquarius} found to be close to an Einasto profile, which is shallower than the NFW profile."
 The theoretical situation has become somewhat unclear. siuce a projection of a sub-halo rich u-body galaxy simulation (Xuetal.2009.2010). fines nuage distortions generalle smaller than those seen iu strong leus systems.," The theoretical situation has become somewhat unclear, since a projection of a sub-halo rich n-body galaxy simulation \citep{Xu:09,Xu:10} finds image distortions generally smaller than those seen in strong lens systems."
 A better overall understanding aud tighter lits requires better data on more strong leus systems CIreu2010)., A better overall understanding and tighter limits requires better data on more strong lens systems \citep{Treu:10}.
. Iu this paper we analyze M31s NW star stream for the density variations that sub-halos are expected to duco., In this paper we analyze M31's NW star stream for the density variations that sub-halos are expected to induce.
 We first trace the mean ceuterline of the stream., We first trace the mean centerline of the stream.
 The Dhuuinositv function and total buuinositv of the main part of the stream are estimated as au indication of the progenitor system., The luminosity function and total luminosity of the main part of the stream are estimated as an indication of the progenitor system.
 The streams width and deviations from the ceuterline are measured., The stream's width and deviations from the centerline are measured.
 The iiecasuremoent of the density along the stream gives the data which we use to test whether the stream is smooth or luupy., The measurement of the density along the stream gives the data which we use to test whether the stream is smooth or lumpy.
 An illustrative simulation is used to show the expected scale aud degree of stream structure as the amount of substructure increases from the visible chwarts to the thousands of halosthat LCDM predicts., An illustrative simulation is used to show the expected scale and degree of stream structure as the amount of substructure increases from the visible dwarfs to the thousands of halosthat LCDM predicts.
 PAndAS has acquired y and / band images covering nearly 100 square degrees of M31-M325 with the CFIIT Meeaprime system., PAndAS has acquired $g$ and $i$ band images covering nearly 400 square degrees of M31-M33 with the CFHT Megaprime system.
 The images have a uals S/N=10 at /4p=21.5 mae., The images have a minimum S/N=10 at $i_{AB}\simeq 24.5$ mag.
 Details of data reduction are in MeConnachieetal.(2010). and al. (2011)., Details of data reduction are in \citet{McC:10} and \citet{Richardson:11}.
. We adopt the MeConnachieetal.(2005) distance of 78b kpc (DAL of 2LI7 mae) for M31. although we note a statistically consistent. but somewhat sinaller. Cepheid distance having a DAL 1.32+0.12 mag or 731 kpe (Vilardelletal.2007).," We adopt the \citet{McConnachie:05} distance of 785 kpc (DM of 24.47 mag) for M31, although we note a statistically consistent, but somewhat smaller, Cepheid distance having a DM $\pm0.12$ mag or 731 kpc \citep{VJR:07}."
. The adopted distance indicates that the survey reaches to M;~0 πας at biel signal-to-noise., The adopted distance indicates that the survey reaches to $_i\sim 0$ mag at high signal-to-noise.
 We use data up to and including Jamuary. 2010.," We use data up to and including January, 2010."
 huage extraction vields approximately 31 nmülliou stellar objects in the images., Image extraction yields approximately 31 million stellar objects in the images.
 Matching the colors and maguitudes of the stars to isochrones of varving metallicity at a fixed age of 12 Corr (Dotteretal.2008) identifies M31 Red Ciaut Branch (ROB) stars using the procedures of MeConnachieetal.(2010)., Matching the colors and magnitudes of the stars to isochrones of varying metallicity at a fixed age of 12 Gyr \citep{Dotter:08} identifies M31 Red Giant Branch (RGB) stars using the procedures of \citet{McC:10}.
. Since we are secius M31 through our own galaxy. it turus out that more than of the stars do not match RGB isochroucs at the adopted distance.," Since we are seeing M31 through our own galaxy, it turns out that more than of the stars do not match RGB isochrones at the adopted distance."
" These stars are not included im the map. although all stars above the completeness η, are useful for the measurement of ecometric biases of the nuages."," These stars are not included in the map, although all stars above the completeness limit are useful for the measurement of geometric biases of the images."
 The stars are individually extinction corrected with galactie dust maps (MeConnuachie.etal. 2010).., The stars are individually extinction corrected with galactic dust maps \citep{McC:10}. .
 We are left with about L3 imillion nominal red seiaut branch stars to 21.5 mag which have been assigned, We are left with about 4.3 million nominal red giant branch stars to 24.5 mag which have been assigned
"and expands, compressing the gas in a thin rim and generating a low density wake within.","and expands, compressing the gas in a thin rim and generating a low density wake within."
" To understand in more detail the nature of the resulting BH orbit, Figure 5 shows unsharp masked maps of the gravitational potential of the gas at the moment when the BH reaches the apocentre for the first time."," To understand in more detail the nature of the resulting BH orbit, Figure \ref{map_potential} shows unsharp masked maps of the gravitational potential of the gas at the moment when the BH reaches the apocentre for the first time."
 The unsharp masking has been performed by subtracting from the initial potential a version smoothed on a scale of ~600h~'pc.," The unsharp masking has been performed by subtracting from the initial potential a version smoothed on a scale of $\sim 600\,h^{-1} \, {\rm
 pc}$."
" As can be seen in the top panel, for the case without BH accretion, the high density wake exerts gravitational drag on the BH and the net gravitational force from the surrounding gas acting on the BH is directed prograde along the y-axis."," As can be seen in the top panel, for the case without BH accretion, the high density wake exerts gravitational drag on the BH and the net gravitational force from the surrounding gas acting on the BH is directed prograde along the $y$ -axis."
" In the case with BH accretion and feedback, the situation is more complex, as shown in the bottom panel."," In the case with BH accretion and feedback, the situation is more complex, as shown in the bottom panel."
" Here the BH is located within the dip associated with high density compressed gas, while above the BH a hot low density plume is expanding away from the BH."," Here the BH is located within the dip associated with high density compressed gas, while above the BH a hot low density plume is expanding away from the BH."
" The net gravitational force acting on the BH from this gas distribution is pointing retrograde along the y-axis, and thus the BH experiences a gravitational drag in the opposite direction than in the top panel."," The net gravitational force acting on the BH from this gas distribution is pointing retrograde along the $y$ -axis, and thus the BH experiences a gravitational drag in the opposite direction than in the top panel."
" Moreover, cold gas in the galactic disc, which approaches the hot gas plume due to the rotation of the galaxy, is compressed and transfers its angular momentum to the gas in the plume."," Moreover, cold gas in the galactic disc, which approaches the hot gas plume due to the rotation of the galaxy, is compressed and transfers its angular momentum to the gas in the plume."
 This effect further reduces the effective drag force onto the BH from gas accretion., This effect further reduces the effective drag force onto the BH from gas accretion.
" 'The BH describes a precessing elliptical trajectory, and with each passage the ellipses widen such that after five relatively close passages to the galactic centre (after about 3x10?yrs) the orbit starts to circularise, but with a BH that is counter-rotating, even though it started out with the same radial motion."," The BH describes a precessing elliptical trajectory, and with each passage the ellipses widen such that after five relatively close passages to the galactic centre (after about $3 \times 10^8\,$ yrs) the orbit starts to circularise, but with a BH that is counter-rotating, even though it started out with the same radial motion."
" Due to the feedback a ring of hot, low density gas forms, within which the BH orbits, partially decoupled from the rest of the galaxy."," Due to the feedback a ring of hot, low density gas forms, within which the BH orbits, partially decoupled from the rest of the galaxy."
" The BH trajectory is loop-like, with 2-500!pc excursions from the plane."," The BH trajectory is loop-like, with $\pm 500 \,h^{-1} \, {\rm pc}$ excursions from the plane."
" For the rest of the simulated time of ~1.7x10? yrs, the BH remains at a considerable distance from the galactic centre, in the range of 3.5—6.5h!kpc, even though ~7x10? yrs after the kick, the apocentric distance starts to decay, indicating that the BH will eventually return to the centre (see also Appendix A))."," For the rest of the simulated time of $\sim 1.7
\times 10^9\,$ yrs, the BH remains at a considerable distance from the galactic centre, in the range of $3.5 \,-\, 6.5\,h^{-1} \, {\rm kpc}$, even though $\sim 7 \times 10^8\,$ yrs after the kick, the apocentric distance starts to decay, indicating that the BH will eventually return to the centre (see also Appendix \ref{appen}) )."
 The case where the BH is kicked perpendicular to the plane of the galaxy is shown in panel (c) of Figure 4.., The case where the BH is kicked perpendicular to the plane of the galaxy is shown in panel (c) of Figure \ref{BHorbit}.
" During the first ~3x10 yrs, the BH describes about a dozen orbits before it returns to the galactic centre."," During the first $\sim 3 \times
10^8\,$ yrs, the BH describes about a dozen orbits before it returns to the galactic centre."
" In this case, the BH feedback does not significantly delay the return timescale because the BH accretes only during short passages through the galactic plane, due to the mostly radial orbit along the z-axis."," In this case, the BH feedback does not significantly delay the return timescale because the BH accretes only during short passages through the galactic plane, due to the mostly radial orbit along the $z$ -axis."
" Note also that the maximum distance reached by the BH at the first apocentre is lower than for the BH kicked in the plane of the galaxy, as a result of the larger gravitational force exerted from the disc of the galaxy in this direction (see panel on the right)."," Note also that the maximum distance reached by the BH at the first apocentre is lower than for the BH kicked in the plane of the galaxy, as a result of the larger gravitational force exerted from the disc of the galaxy in this direction (see panel on the right)."
" Finally, in the right-hand panel of Figure 4,, we also show a case where the kick velocity is much lower, equal to 375kms~' (blue line)."," Finally, in the right-hand panel of Figure \ref{BHorbit}, we also show a case where the kick velocity is much lower, equal to $375\, {\rm km\,s^{-1}}$ (blue line)."
" Here the BH briefly leaves the minimum of the potential and reaches a radial distance of 500A!pc, but then returns to the centre within 7x10° yrs."," Here the BH briefly leaves the minimum of the potential and reaches a radial distance of $500 \,h^{-1} \, {\rm pc}$, but then returns to the centre within $7 \times 10^6\,$ yrs."
" In Figure 6 we show the BH mass (top panel) and bolometric luminosity (middle panel) as a function of time, as well as the time evolution of the total star formation rate (SFR) of the host galaxy (bottom panel) for different cases."," In Figure \ref{mbh_iso} we show the BH mass (top panel) and bolometric luminosity (middle panel) as a function of time, as well as the time evolution of the total star formation rate (SFR) of the host galaxy (bottom panel) for different cases."
" The green lines are for the case where the BH does not experience a gravitational recoil, while the red and orange lines are for vxick=700kms! parallel and perpendicular to the plane of the galaxy, respectively."," The green lines are for the case where the BH does not experience a gravitational recoil, while the red and orange lines are for $v_{\rm kick}= 700\, {\rm km \, s^{-1}}$ parallel and perpendicular to the plane of the galaxy, respectively."
" From the top panel it can be seen that the BH which does not experience any recoil increases its mass moderately, by ~8.5x10”Μο over 2 Gyrs."," From the top panel it can be seen that the BH which does not experience any recoil increases its mass moderately, by $\sim 8.5 \times 10^7 \,h^{-1}{\rm M}_\odot\,$ over $2\,$ Gyrs."
" This is due to secular processes which gradually transport some gas towards the inner regions, fuelling BH accretion."," This is due to secular processes which gradually transport some gas towards the inner regions, fuelling BH accretion."
" The bolometric is initially around 10*°ergs”! and slowly decreases with time, but we note that after ~0.5 Gyrs the BH accretion falls below 0.01 of"," The bolometric is initially around $10^{45}\,{\rm
 erg\,s^{-1}}$ and slowly decreases with time, but we note that after $\sim
0.5\,$ Gyrs the BH accretion falls below $0.01$ of"
for a region of sizeRp to become supercritical.,for a region of size$\sim R_\mathrm{BE}$ to become supercritical.
 For the set of shock moclels we are studyiug (see Table 1. for inflow parameters). calculated radii aucl mass-to-flux ratios for Bounor-Ebert spheres uuder post-shock couditious without ambipolar diffusion are listed iu Table 2..," For the set of shock models we are studying (see Table \ref{thickness} for inflow parameters), calculated radii and mass-to-flux ratios for Bonnor-Ebert spheres under post-shock conditions without ambipolar diffusion are listed in Table \ref{Model}."
 Iu all cases. Epi is muuch smaller than 1. which meaus uo collapse at the BE scale could happen iu tlie post-shock region without significant ambipolar diffusion.," In all cases, $\Gamma_\mathrm{BE}$ is much smaller than $1$, which means no collapse at the BE scale could happen in the post-shock region without significant ambipolar diffusion."
 More generally. since Cpe~ος(μον)2<<p under GMC couditious. most post-shock regious are sullicieutly magnetized that gravitational collapse of cores would be prevented uuless ambipolar diffusion occurs.," More generally, since $\Gamma_\mathrm{BE} \sim c_s / \left(v_0 v_{\mathrm{A},0}\right)^{1/2} \ll 1$ under GMC conditions, most post-shock regions are sufficiently magnetized that gravitational collapse of low-mass cores would be prevented unless ambipolar diffusion occurs."
 Note that if p>rro. as is true in the candidate core material for transient C shocks. Rpye will be lower than the value in the table.," Note that if $\rho > r_f\rho_0$, as is true in the candidate core material for transient C shocks, $R_\mathrm{BE}$ will be lower than the value in the table."
 Current theoretical aud observational work suggests that shocks produced by supersoulte turbulence play a role in compressing gas to form prestellar cores., Current theoretical and observational work suggests that shocks produced by supersonic turbulence play a role in compressing gas to form prestellar cores.
 Our fiudiugs that the neutrals are compressed more than the magnetic field during the early stages of shock formation raise an interesting question: Is it. possible for a subcritical cloud. to [orm supercritical cores in shocks. which cau then gravitationally collapse promptly?," Our findings that the neutrals are compressed more than the magnetic field during the early stages of shock formation raise an interesting question: Is it possible for a subcritical cloud to form supercritical cores in shocks, which can then gravitationally collapse promptly?"
 For a prestellar core to collapse. the region must be deuse enough so that sell-egravity overcomes the magnetic support.," For a prestellar core to collapse, the region must be dense enough so that self-gravity overcomes the magnetic support."
 Although sel(-gravity is not included in the present mocels. we can make an initial assessinent of whether trausient. C shocks are likely to affect the ability of cores to collapse promptly alter they [oru," Although self-gravity is not included in the present models, we can make an initial assessment of whether transient C shocks are likely to affect the ability of cores to collapse promptly after they form."
 Iu the context of our convergiug flow test. we shall define “caucicdate core” material to be regious where the background has uniform »/By. so the time evolution of this candidate core material is easily caleulated.," In the context of our converging flow test, we shall define “candidate core"" material to be regions where the background has uniform $n/B_y$, so the time evolution of this candidate core material is easily calculated."
 Pliysically. this candidate core material correspouds to that in the ceutral peak ol nrBy as shown iu e.g. Fig.," Physically, this candidate core material corresponds to that in the central peak of $n/B_y$ as shown in e.g. Fig."
 L or Fig. 6.., \ref{cvg1} or Fig. \ref{CC4}.
 For steady shocks with compression factor ry produced by a two-sided converging flow with inflow speed cios rom both sides. the upstream speed in the shock frame is eg=cungasrp/Qr—1). al Ushock=Cinow/(ry— 1).," For steady shocks with compression factor $r_f$ produced by a two-sided converging flow with inflow speed $v_\mathrm{inflow}$ from both sides, the upstream speed in the shock frame is $v_0 = v_\mathrm{inflow} r_f/ \left(r_f-1\right)$, and $v_\mathrm{shock} = v_\mathrm{inflow}/\left(r_f -1\right)$ ."
" The rate at which the coltunn density grows for a steady shock is therefore where n—ny,+ng=0.6141 is assiunecd.", The rate at which the column density grows for a steady shock is therefore where $n = n_\mathrm{H_2} + n_\mathrm{He} = 0.6n_\mathrm{H}$ is assumed.
We do have clear evidence of Malm(quist bias in the present study.,We do have clear evidence of Malmquist bias in the present study.
 Figure 11 shows that compared to the huninosity functions for the five spectral classes. the luminosity distributions predicted. lor the GUSBAD sample are shifted. to considerably higher huminosities.," Figure 11 shows that compared to the luminosity functions for the five spectral classes, the luminosity distributions predicted for the GUSBAD sample are shifted to considerably higher luminosities."
 The Aalmequist biases for sp=1-5 are0.70.0.53.0.31.0.28 and 0.36. respectively. for an average οἱ 0.44.," The Malmquist biases for $sp = 1-5$ are$0.70, 0.53, 0.31, 0.28$ and $ 0.36$, respectively, for an average of $0.44$."
 This value is tantalizinglv close to the offset shown in eq. (, This value is tantalizingly close to the offset shown in eq. (
13).,13).
 However. our values are based on an assumed gaussian shape of (he spectral luminosity functions with a dispersion Of Moe.=0.5. designed to make the overall luminosity function reasonably smooth.," However, our values are based on an assumed gaussian shape of the spectral luminosity functions with a dispersion of $\sigma_{\log L} = 0.5$, designed to make the overall luminosity function reasonably smooth."
 As long as we do not have reliable information about the distribution of logLi; al given lin space. we will not be able to assess the ellect of Malimeuist bias accurately.," As long as we do not have reliable information about the distribution of $\log L_{\rm iso}$ at given in space, we will not be able to assess the effect of Malmquist bias accurately."
 We present in (his final section some commentary on (he approach used in this paper and on the results., We present in this final section some commentary on the approach used in this paper and on the results.
 1., 1.
" This work validates (he use of aas a cosmological distance indicator,", This work validates the use of as a cosmological distance indicator.
 It requires a complete sample of GRBs above a well defined flux limit., It requires a complete sample of GRBs above a well defined flux limit.
 The GUSBAD catalog with its uniform treatment of all sources is eminently useful for this purpose., The GUSBAD catalog with its uniform treatment of all sources is eminently useful for this purpose.
 The input in the form of [forced us to a cleductive approach where we used a given shape of the (spectral) Iuminosity Iunction and iterated its center Iuminosity and rest [rame uuntil the observed aad wwere fit., The input in the form of forced us to a deductive approach where we used a given shape of the (spectral) luminosity function and iterated its center luminosity and rest frame until the observed and were fit.
 An important advantage of the deductive method is that the derivation of source counts from a given Iuminosityv [unction can be carried out accurately., An important advantage of the deductive method is that the derivation of source counts from a given luminosity function can be carried out accurately.
 2., 2.
 This approach is in contrast to the inductive method generally used when redshilts are available., This approach is in contrast to the inductive method generally used when redshifts are available.
 This involves deriving densities in bins of redshift and Iuminositv. and correlating these with redshilt and/or huninosity to derive (heΙΟ., This involves deriving densities in bins of redshift and luminosity and correlating these with redshift and/or luminosity to derive the.
.. In such an approach. il is essential to check that the resulting huminosity Function produces Che input source counts and ecorrectly.," In such an approach, it is essential to check that the resulting luminosity function produces the input source counts and correctly."
 IE (his is not the case. predictions about. expected numbers of sources at. large redshift will be unrealistie (Schmidt.2004).," If this is not the case, predictions about expected numbers of sources at large redshift will be unrealistic \citep{sch04}."
. 3., 3.
" IEGRDs were not subject to cosmological evolution. the input in terms of2...Ease and <As,> would sullice to derive the aad hence predict source counts aid the redshift and Iuminositw distributions."," If GRBs were not subject to cosmological evolution, the input in terms of, and $<\alpha_{23}>$ would suffice to derive the and hence predict source counts and the redshift and luminosity distributions."
 Comparing these predictions with the observations. such as from Swift. would then constitute a check," Comparing these predictions with the observations, such as from , would then constitute a check"
We are interesting in case ~>1 and in all formulas that followed we use the expansion 37=1— 5.,We are interesting in case $\gamma \gg 1$ and in all formulas that followed we use the expansion $\beta =1-\frac 1{2\gamma ^2}$ .
" In the spherical coordinates dV=2:)7drdcosts 2807w(l cos064,)dzdrfz."," In the spherical coordinates $dV=2\pi r^2drd\cos \theta _2=$ $2\pi r^2\omega \left( 1-\cos
\theta _1\right) d\varepsilon dr/\varepsilon ^2$."
" When ¢Xw. then 65«1 and thus Now it is possible to rewrite (8)) as an integral over r where UN, is the number of photons going from the companion per second and No is the number of particles leaving the pulsar per second. nm;=Nofail. n Auer. πως ds determined from the condition 1."," When $\varepsilon \gg \omega $, then $\theta _2\ll 1$ and thus Now it is possible to rewrite \ref{def}) ) as an integral over $r$ where $\,N_\omega $ is the number of photons going from the companion per second and $N_{e^{\pm }}$ is the number of particles leaving the pulsar per second, $n_\omega =N_\omega \left/ 4\pi cl^2\right. $, $n_{e^{\pm
}}=N_{e^{\pm }}/4\pi cr^2$ , $r_{\max }$ is determined from the condition $%
\theta _2= 1."
 UW Lorentz factor of a particle is constant then from (11)). (12)) follows: where In the case A<<l1 in Thomson limit the Lorentz actor of the particles is practically constant in the effective scattering volume.," If Lorentz factor of a particle is constant then from \ref
{ct2}) ), \ref{ct1}) ) follows: where In the case $K\ll 1$ in Thomson limit the Lorentz factor of the particles is practically constant in the effective scattering volume."
 In the case wsmc1. this statement is also correct.," In the case $\omega
\gamma /mc^2>1$ this statement is also correct."
 In this case if the mean free path of a xwticle is much bigger then the binary separation: fy...=ONuusαπωαὁ> a (σι is a Cull IxIein-Nishina ¢ross- then the inlluence of the field. of soft. photons on he spectrum. ancl Εαν of the relativistic particles may. be neglected.," In this case if the mean free path of a particle is much bigger then the binary separation: $%
l_{free}=\sigma _{K_{full}}L_{*}/4\pi \omega a^2c\gg a $\sigma _{K_{full}}$ is a full Klein-Nishina cross-section) then the influence of the field of soft photons on the spectrum and flux of the relativistic particles may be neglected."
 In the highly relativistic case (as/me 1) Thainvmpmic;ως and we note that [rom Ac« T it follows hat /r4a.," In the highly relativistic case $\omega \gamma /mc^2\gg 1$ ) $%
\sigma _{K_{full}}\sim \sigma _Tmc^2/\omega \gamma and we note that from $%
K\ll 1 it follows that $l_{free}\gg a$."
 After a substitution of the expression (9)) in (13)) and he followingὃν integrationo under the assumption that > is constant we obtain: =nmBoya the function £F(i.2.5.c) is determined as follows: and is equal to zero when Sas.," After a substitution of the expression \ref{sigm}) ) in \ref{def1}) ) and the following integration under the assumption that $\gamma $ is constant we obtain: where the constant $\tau =\frac{3\sigma _T}{128\pi ^2ca}$ and in the range the function $F\left( \omega ,\varepsilon ,\gamma ,\psi \right) $ is determined as follows: and is equal to zero when $\varepsilon >\varepsilon _{\max }$."
 When the equation (17)) simplifies: In the Thomson limit Zi=257w(l|àcosc) and hence LASSe<PSS(11cose’)11.," When $\omega
\ll \varepsilon \ll \varepsilon _{\max }$ the equation \ref{Ff}) ) simplifies: In the Thomson limit $\varepsilon _{\max }=2\gamma ^2\omega (1+\cos \psi )$ and hence $1-\lambda =\frac \varepsilon {mc^2\gamma }<2\frac{\omega \gamma }{%
mc^2}(1+\cos \psi )\ll 1."
 Under these conditions equation (17)) becomes If we integrate. equation (15)) with £=fp over the energy of the scattered. photons 5 we receive the o- dependence of the total outgoing energy in according with the equation (6))., Under these conditions equation \ref{Ff}) ) becomes If we integrate equation \ref{Lg}) ) with $F=F_T$ over the energy of the scattered photons $\varepsilon $ we receive the $\psi $ -dependence of the total outgoing energy in according with the equation \ref{de1}) ).
 The relationship between functions &(o) and £u is: Under the real conditions the spectrum of the optical photons is not monochromatic., The relationship between functions $\Phi \left( \psi \right) $ and $F_T$ is: Under the real conditions the spectrum of the optical photons is not monochromatic.
 Lf the conpanion’s radiation is similar with a radiation from the black body with temperature Z. then the number of soft. photons in. the energy range [rom w low |ee is: lteplacing ANS in formula (15)) with αλ and. integrating over w we recelve: ligure 3 shows the dependence of £55 on 5 for dilferent values of c and 5.," If the companion's radiation is similar with a radiation from the black body with temperature $T$, then the number of soft photons in the energy range from $\omega $ to $\omega +d\omega $ is: Replacing $N_\omega $ in formula \ref{Lg}) ) with $dN_\omega $ and integrating over $\omega $ we receive: Figure 3 shows the dependence of $F_{bb}$ on $\varepsilon $ for different values of $\psi $ and $\gamma $ ."
 It can be seen that in the energy range where the dillerence of the cross-section from the Thomson one is significant the form of the spectrum changes considerably., It can be seen that in the energy range where the difference of the cross-section from the Thomson one is significant the form of the spectrum changes considerably.
 Figure 3 also illustrates the dependence. of the scattered: photons themaximumenergy.ofzu on ~ ( according to (16)) within the Vhomson limit ZuPas(1|cose) and in the opposite caseFall|cose)Se1 the maximum. energv ος=2mce52 ).," Figure 3 also illustrates the dependence of the maximum energy of the scattered photons $%
\varepsilon _{\max on $\gamma $ ( according to \ref{omax}) ) within the Thomson limit $\varepsilon _{\max }=2\omega \gamma ^2(1+\cos \psi )$ and in the opposite case$\frac{\omega \gamma }{mc^2}(1+\cos \psi )\gg 1$ the maximum energy $\varepsilon _{\max }=2mc^2\gamma $ )."
 EThe ce dependence of the energy. gone with the scattered. photons in a unit of time in a unit solid angle can be seen in 4., The $\psi -$ dependence of the energy gone with the scattered photons in a unit of time in a unit solid angle can be seen in Figure 4.
 The energy is normalized to Lint., The energy is normalized to 2$.
2Figure During the orbital motion c changes periodicallv., During the orbital motion $\psi $ changes periodically.
" In Figure 5 the dependence of £5, on 5 for dillerent. positions of the companions in the circular orbit. and. for.cülferent values of ~ is shown.", In Figure 5 the dependence of $F_{bb}$ on $\varepsilon $ for different positions of the companions in the circular orbit and fordifferent values of $\gamma $ is shown.
 In this figure / is an inclination angle and 4 is a true anomaly., In this figure $i$ is an inclination angle and $\varphi $ is a true anomaly.
 The relation. between cancl «5 is: cost’= sin/cosy.," The relation between $\psi $and $\varphi $ is: $\cos \psi =\sin
i\cos \varphi $ ."
 Figure 6 shows the orbital angle 4 dependence of the luminosity £-(47) going in the direction of the observer for the different values of inclination angles and Lorentz factors., Figure 6 shows the orbital angle $\varphi $ dependence of the luminosity $L_\gamma (\varphi )$ going in the direction of the observer for the different values of inclination angles and Lorentz factors.
 Formulas (14)) - (20)) were received. under the, Formulas \ref{rmax}) ) - \ref{fbb}) ) were received under the
"respectively, ave shown in the third and fourth coluuus of Tables 1 and 2..","respectively, are shown in the third and fourth columns of Tables \ref{tab:4Gtests} and \ref{tab:14Gtests}."
" For both test-case SSPs it can be seen that IL increases significantly svhen 7,,, is increased for the whole isochrone. as would be expected."," For both test-case SSPs it can be seen that $\beta$ increases significantly when $T_{eff}$ is increased for the whole isochrone, as would be expected."
" However when Toy, is increased for giauts ον, the Is value increases sleltly for the 11 Gyr SSP aud actually for the L Gyr SSP we attribute this to the fact that for eiant stars at these low temperatures (Thy,< 50001.) and gravities (logg <3). increasing ρε by 1001 iucreases the contiuumunu level around the I> line but does not sienificautly increase the streueth of the line itself. and so the measured ITJ value can decrease."," However when $T_{eff}$ is increased for giants only, the $\beta$ value increases slightly for the 14 Gyr SSP and actually for the 4 Gyr SSP – we attribute this to the fact that for giant stars at these low temperatures $T_{eff}<5000$ K) and gravities $g<$ 3), increasing $T_{eff}$ by 100K increases the continuum level around the $\beta$ line but does not significantly increase the strength of the line itself, and so the measured $\beta$ value can decrease."
" The higher order Baluer lues Uép aud σε behave in a qualitatively simular wav to IT? in that they increase significantly when T;,;; is increased for the whole isochrone. but increase only mareiually. or decrease slightly. for the elauts-only case."," The higher order Balmer lines $\delta_F$ and $\gamma_F$ behave in a qualitatively similar way to $\beta$ in that they increase significantly when $T_{eff}$ is increased for the whole isochrone, but increase only marginally, or decrease slightly, for the giants-only case."
" For all but oue of the metal indicators Guchiding Meb aud all the Fe lines) iudex streueths decrease when Toy, is increased for the whole isochrouc. je exception being Col668 for the 1 Gyr SSP. which increases slighth."," For all but one of the metal indicators (including $b$ and all the Fe lines) index strengths decrease when $T_{eff}$ is increased for the whole isochrone, the exception being $_2$ 4668 for the 4 Gyr SSP, which increases slightly."
 When {ερ is increased or elauts only. offsets are generally smaller. with ιο exception of Col668. but for mauv dudices offsets go iu the opposite scuse to those for the wholesochrone case. e.g. CNL. CN2 and most of 1ο Fe lines.," When $T_{eff}$ is increased for giants only, offsets are generally smaller, with the exception of $_2$ 4668, but for many indices offsets go in the opposite sense to those for the whole-isochrone case, e.g. CN1, CN2 and most of the Fe lines."
" Results of the decreased Lp, test (ising the runcated Lb Cyr SSP) are not Listed. as the neastredl values of offsets would not be applicable o the full SSP."," Results of the decreased $T_{eff}$ test (using the truncated 4 Gyr SSP) are not listed, as the measured values of offsets would not be applicable to the full SSP."
 Tlowever we note that. for all 23 neasured iudices. offsets for the decreased Τε runcated SSP were fouud to be svuunetrically opposite with respect to the increased τει runcated SSP case. uaplving that this would also e the case for the full SSP.," However we note that, for all 23 measured indices, offsets for the decreased $T_{eff}$ truncated SSP were found to be symmetrically opposite with respect to the increased $T_{eff}$ truncated SSP case, implying that this would also be the case for the full SSP."
 To demoustrate the effect of these offsets ou estimated ages and ictallicitics of SSPs. werave constructed representative index-index evids τοι the DaSTI models for a range of ages aud uetallicities and overplotted the test-case SSPs Figures 1.. 2 and 3. show eris of IL) vs. Meb. Febl0G and Co l668. respectively.," To demonstrate the effect of these offsets on estimated ages and metallicities of SSPs, wehave constructed representative index-index grids from the BaSTI models for a range of ages and metallicities and overplotted the test-case SSPs – Figures \ref{fig1}, \ref{fig2} and \ref{fig3} show grids of $\beta$ vs. $b$, Fe5406 and $_2$ 4668, respectively."
" In each feure the left haud panel shows the results of the ""holedsochrone tests and the right haud pancl shows the results of the eiauts-ouly tests.", In each figure the left hand panel shows the results of the whole-isochrone tests and the right hand panel shows the results of the giants-only tests.
" It cau be seeu that increasing T;,;,; bw 1000for the whole isochrone makes the appareut age of the 1 Car SSP approxiuatelv 0.75 Cir vouuger wlilst the li Gyr SSP looks (αν vounger. au apparent reduction in age of ~20% for both populations."," It can be seen that increasing $T_{eff}$ by 100Kfor the whole isochrone makes the apparent age of the 4 Gyr SSP approximately 0.75 Gyr younger whilst the 14 Gyr SSP looks $-$ 3 Gyr younger, an apparent reduction in age of $\sim$ for both populations."
 The effect ou apparent ages for the elauts-only test ds small. altering absolute ages by no more than ~: however since the offsets for the | Covr and Ll Gar populations eo in opposite scuses. the relative ages also change.," The effect on apparent ages for the giants-only test is small, altering absolute ages by no more than $\sim$, however since the offsets for the 4 Gyr and 14 Gyr populations go in opposite senses, the relative ages also change."
 Assessing the effect of temperaturechauges onu derived ietallicitics is more complicatect. since df depends on which iudex is used and whether metallicity is defined iu terms of [Fe/T]| or total metallicity. Z.," Assessing the effect of temperaturechanges on derived metallicities is more complicated, since it depends on which index is used and whether metallicity is defined in terms of [Fe/H] or total metallicity, $Z$."
 The three exids displayed in Figures 1.. 2 and 3. demonstrate the problem.," The three grids displayed in Figures \ref{fig1}, \ref{fig2} and \ref{fig3} demonstrate the problem."
" For a LOOW increase in Ty,. the inferred |Fo/1I] from the I> Meb plane is — 0.1 dex lower than the reference SSP. from the IL? Fe5106 plane it is virtually uuchanecd. and from the I> C» 1668 plane |Fe/TI| apparently increases by ~ 0.05 dex."," For a 100K increase in $T_{eff}$, the inferred [Fe/H] from the $\beta$ $b$ plane is $\sim$ 0.1 dex lower than the reference SSP, from the $\beta$ –Fe5406 plane it is virtually unchanged, and from the $\beta$ $_2$ 4668 plane [Fe/H] apparently increases by $\sim$ 0.05 dex."
" Offsets are similar for both the whole-isochrouc and “elauts-only cases,", Offsets are similar for both the `whole-isochrone' and `giants-only' cases.
 The IL) Cs L668 exid also demonstrates another niportant poiut or the Ll Gyr SSP. the offset for the increased 1001 case the absolute value of Co L668 slightly. however the inferred |Fo/TI] as this erid is not orthogonal (as is also the case for all other conumnonlv used age/nmetalliityv diagnostic erids).," The $\beta$ $_2$ 4668 grid also demonstrates another important point – for the 14 Gyr SSP, the offset for the increased 100K case the absolute value of $_2$ 4668 slightly, however the inferred [Fe/H] as this grid is not orthogonal (as is also the case for all other commonly used age/metallicity diagnostic grids)."
" Similarly constructed evids (not displaved hero) show that. with iucreasiug 7,,,. Cal227 behaves iu the same way as Meh. [MgFo] behaves as Fe5106 (but traces total Z rather than [Fe/II] sec ?)) and CNI and CN2 behave similarly to Ου σος,"," Similarly constructed grids (not displayed here) show that, with increasing $T_{eff}$, Ca4227 behaves in the same way as $b$, [MgFe] behaves as Fe5406 (but traces total $Z$ rather than [Fe/H] – see ) and CN1 and CN2 behave similarly to $_{2}$ 4668."
 The issue of the contradictory behavior of various diagnostic metal lines aud the inipact on iuferred chemical composition will be discussed. further iu Section 7.., The issue of the contradictory behavior of various diagnostic metal lines and the impact on inferred chemical composition will be discussed further in Section \ref{sec:discsum}.
 It was found that. for the logy tests. offsets in all iudices were found to be svnuunetrcal for the iucreased/decreased cases with respect to the reference case. heuce the results are listed iu Tables d and 2 as single values for cach pairing of tests 3a/3b and laflb. with the appropriate + or = sign.," It was found that, for the $g$ tests, offsets in all indices were found to be symmetrical for the increased/decreased cases with respect to the reference case, hence the results are listed in Tables \ref{tab:4Gtests} and \ref{tab:14Gtests} as single values for each pairing of tests 3a/3b and 4a/4b, with the appropriate $\pm$ or $\mp$ sign."
 For both test-case SSPs it can be seen that ITJ. Hog and τε increase sieuificautly when logy is decreased by 0.25 dex (aud decrease whenlogg is increased). affecting derived ages in a quantitatively similar wav to a 1001 increase in Toye.," For both test-case SSPs it can be seen that $\beta$ , $\delta_F$ and $\gamma_F$ increase significantly when $g$ is decreased by 0.25 dex (and decrease when$g$ is increased), affecting derived ages in a quantitatively similar way to a 100K increase in $T_{eff}$ ."
 Meb. Ca1227. Feb106. [MgFo] and C2 1665," $b$ , Ca4227, Fe5406, [MgFe] and $_2$ 4668"
In spite of the fact that nitrogen is one of the most abundant elements in the cosmos. only recently it has been possible to produce large samples of N abundances in metal poor stars. enabling us to quantitatively study its evolution in the Galaxy (see e.g. [sraelian et al.,"In spite of the fact that nitrogen is one of the most abundant elements in the cosmos, only recently it has been possible to produce large samples of N abundances in metal poor stars, enabling us to quantitatively study its evolution in the Galaxy (see e.g. Israelian et al."
 2004. and references therein).," 2004, and references therein)."
 Nitrogen tracks tron ({N/Fe]~ 0) down to a metallicity of ~—1.0., Nitrogen tracks iron $\sim 0$ ) down to a metallicity of $\sim -1.0$.
" The evolution at lower metallicity ts less clear, while many stars seem to indicate that N continues to track iron. a few field stars are found to be nitrogen-rich (the prototypes of the class being HD 74000 and HD 160617 Bessell&Nor-ris 1982)). while others seem to show [N/Fe] ratios below solar."," The evolution at lower metallicity is less clear; while many stars seem to indicate that N continues to track iron, a few field stars are found to be nitrogen-rich (the prototypes of the class being HD 74000 and HD 160617 \citealt{BN82}) ), while others seem to show [N/Fe] ratios below solar."
 In their investigation of the chemical composition i1 the lowest metallicity stars Spiteetal.(2005) suggestec that a possible interpretation of the data in the metallicity range -4.0x[Fe/H]€-2.5 is a bimodal distribution of nitrogen abundances. the majority of stars at [N/Fe]~+0.4 and the rest at [N/Fe]~—1.0.," In their investigation of the chemical composition in the lowest metallicity stars \citet{spite05}
 suggested that a possible interpretation of the data in the metallicity range $\rm -4.0 \le [Fe/H]\le -2.5$ is a bimodal distribution of nitrogen abundances, the majority of stars at $\sim +0.4$ and the rest at $\sim -1.0$."
 Given the complexity of its nucleosynthesis. iron is not the best reference element to study chemical evolution. as discussed in Cayreletal.(2004).. and ar c-chain element such as O or Mg is preferable.," Given the complexity of its nucleosynthesis, iron is not the best reference element to study chemical evolution, as discussed in \citet{cayrel04}, and an $\alpha$ -chain element such as O or Mg is preferable."
 The [N/a] ratio begins to decline as one moves from solar to lower metallicities (see e.g. figure 17 of Spiteetal. 2005)) and at low metallicities it reaches a constant value. a at [N/e]-—0.9. although it could be argued that there are in fact plateaus. the second found around [N/o]—1.4.," The $\alpha$ ] ratio begins to decline as one moves from solar to lower metallicities (see e.g. figure 17 of \citealt{spite05}) ) and at low metallicities it reaches a constant value, a at $\alpha$ $\sim -0.9$, although it could be argued that there are in fact plateaus, the second found around $\alpha$ $\sim -1.4$."
 From the nucleosynthetic point of view there is only one mode of nitrogen production: it is formed in the course of the CNO cycle. at the expenses of available C and O. Since this requires the pre-existence of nuclei of C and O. one may expect the production of N to increase with increasing abundance of C and O. re. with metallicity.," From the nucleosynthetic point of view there is only one mode of nitrogen production: it is formed in the course of the CNO cycle, at the expenses of available C and O. Since this requires the pre-existence of nuclei of C and O, one may expect the production of N to increase with increasing abundance of C and O, i.e. with metallicity."
 Nitrogen would then be à so-called purely element. to distinguish it from elements. like the a-chain elements. whose yield is independent of metallicity.," Nitrogen would then be a so-called purely element, to distinguish it from elements, like the $\alpha$ -chain elements, whose yield is independent of metallicity."
 The ratio of a purely element to a primary one should be a monotonically increasing function of metallicity., The ratio of a purely element to a primary one should be a monotonically increasing function of metallicity.
 The fact that [N/a] shows a plateau at low metallicity demonstrates that N cannot be a purely secondary element., The fact that $\alpha$ ] shows a plateau at low metallicity demonstrates that N cannot be a purely secondary element.
 We note here that although the concept of and elements is a useful simplification. for the majority of the elements the production history is complex enough that it is inappropriate to label it with either of these terms.," We note here that although the concept of and elements is a useful simplification, for the majority of the elements the production history is complex enough that it is inappropriate to label it with either of these terms."
 For nitrogen. although the nucleosynthetic process 1s unique. it is necessary to invoke different sites where the process occurs to explain. the two kinds of behaviour.," For nitrogen, although the nucleosynthetic process is unique, it is necessary to invoke different sites where the process occurs to explain the two kinds of behaviour."
 Intermediate mass stars in the AGB phase (IM-AGB stars. Ventura et al.," Intermediate mass stars in the AGB phase (IM-AGB stars, Ventura et al."
 2002). in which the H-burning shell operates through the CNO cycle. have always beer considered major producers of nitrogen.," 2002), in which the H-burning shell operates through the CNO cycle, have always been considered major producers of nitrogen."
 For these stars the necessary C and O nuclei are those already present at the time of the formation of the star and the production mode would besecondary., For these stars the necessary C and O nuclei are those already present at the time of the formation of the star and the production mode would be.
 Massive stars are other possible N producers., Massive stars are other possible N producers.
 In such stars the central H burning takes place through the CNO cycle. and again the mode of production would be secondary.," In such stars the central H burning takes place through the CNO cycle, and again the mode of production would be secondary."
 If. however. the star is rotating. as soon as the central He burning is ignited. fresh C and O produced in the core may be transported through rotational mixing to the H-burning shell.," If, however, the star is rotating, as soon as the central He burning is ignited, fresh C and O produced in the core may be transported through rotational mixing to the H-burning shell."
 This would then produce N in a primary fashion. since the necessary parent nuclei are created (see Maeder and Meynet (2005) for a complete discussion).," This would then produce N in a primary fashion, since the necessary parent nuclei are created (see Maeder and Meynet (2005) for a complete discussion)."
 In such a scenario N. production cannot be purely primary. since some secondary production 1s unavoidable.," In such a scenario N production cannot be purely primary, since some secondary production is unavoidable."
 A primary mode of production of N may occur also in IM-AGB stars. provided a way of mixing the products of the He burning shell to the H burning shell exists.," A primary mode of production of N may occur also in IM-AGB stars, provided a way of mixing the products of the He burning shell to the H burning shell exists."
 The two above-mentioned producers (IM-AGB stars and massive stars) release their nucleosynthetic products. which become available for the next generation of stars. on quite different. time-scales.," The two above-mentioned producers (IM-AGB stars and massive stars) release their nucleosynthetic products, which become available for the next generation of stars, on quite different time-scales."
 Massive stars will explode às core-collapse supernovae on time scales shorter that 10 Myr. while IM-AGB will release their products through winds or in the planetary nebula phase on a time scale between 300 Myr and IGyr. depending on their mass and metallicity.," Massive stars will explode as core-collapse supernovae on time scales shorter that 10 Myr, while IM-AGB will release their products through winds or in the planetary nebula phase on a time scale between 300 Myr and 1Gyr, depending on their mass and metallicity."
 The pattern of N abundances in Globular Clusters is somewhat more complex than in field stars., The pattern of N abundances in Globular Clusters is somewhat more complex than in field stars.
 In the first place. like other light elements.N displays sizeable star to star scatter," In the first place, like other light elements,N displays sizeable star to star scatter"
Iu this section. we will derive the relevant quantities of the PP aud ES for diphoton interaction. and of the Compton scattering for e.,"In this section, we will derive the relevant quantities of the PP and ES for diphoton interaction, and of the Compton scattering for $e^\pm$."
 Similar processes have been examined in the previous papers 2.3.1.0.0[.. however. there are several siguificautoO distinctions in our mauusceript The diphotou PP carries on via t and v-chanucls in the SAL. aud the amplitude is The diphoton PP can occur via exchanges of scalar aud tensor wnparticles ii s-clhamnel. which gives vise to the interference between uuparticle s-chaunel aud SM t+. e-chaunel amplitudes.," Similar processes have been examined in the previous papers \cite{EC1,EC2,EC3,2,Compton}, however, there are several significant distinctions in our manuscript The diphoton PP carries on via $t$ - and $u$ -channels in the SM, and the amplitude is The diphoton PP can occur via exchanges of scalar and tensor unparticles in $s$ -channel, which gives rise to the interference between unparticle $s$ -channel and SM $t$ -, $u$ -channel amplitudes."
" The scattering wuplitudes through scalar aud tensor uuparticle exchanges are where DyeP) is defined in (6)) aud A7""77(μι. for one photon εμίψμι). the other photon εν(ρὸ) and oue tensor poluuparticle with Lorentz iudices po is py)defined as The spiu-averaged amplitucle squared is given by where I stauds for the contribution from the SAL IE is that frou scalay unparticle ILI is that from teusor unparticle exchange. aud TV is the interference between the SAL ancl exchange.uuparticle auplitudes. respectively It is worth noticing that. due to the phase factor exptidix) related to the s-chauucl from the uuparticle sector. there exists interference terii IV. between the SM aud uuparticle amplitudes which is a clear signature of unparticle plivsics."," The scattering amplitudes through scalar and tensor unparticle exchanges are where $T_{\mu\nu,\rho\sigma}(P)$ is defined in \ref{note}) ) and $K^{\mu\nu\rho\sigma}(p_1^\mu,p_2^\nu)$ for one photon $\epsilon_\mu(p_1)$, the other photon $\epsilon_\nu(p_2)$ and one tensor unparticle with Lorentz indices $\rho\sigma$ is defined as The spin-averaged amplitude squared is given by where ${\bf I}$ stands for the contribution from the SM, ${\bf II}$ is that from scalar unparticle exchange, ${\bf III}$ is that from tensor unparticle exchange, and ${\bf IV}$ is the interference between the SM and unparticle amplitudes, respectively It is worth noticing that, due to the phase factor $\exp(-id_{\cal
U}\pi)$ related to the $s$ -channel from the unparticle sector, there exists interference term ${\bf IV}$ between the SM and unparticle amplitudes which is a clear signature of unparticle physics."
(200033) at GO. 90. 135 and 170 sau. After a few montlis of operation of the satellite it appeared that at 60 and 90 jan the on-off signal was strongly distorted by trausicuts in the responsitivity of the detectors.,") at 60, 90, 135 and 170 $\mu$ m. After a few months of operation of the satellite it appeared that at 60 and 90 $\mu$ m the on-off signal was strongly distorted by transients in the responsitivity of the detectors."
 Suuibuly. chopping appeared to be an inadequate observing mode at 135 iux at 170 jun because of confusion with structure in the ackeround roni iufrared cnrus.," Similarly, chopping appeared to be an inadequate observing mode at 135 and at 170 $\mu$ m because of confusion with structure in the background from infrared cirrus."
 We therefore switclie o the observing mode PITT22 aud made ας., We therefore switched to the observing mode PHT22 and made minimaps.
 ΑΠΛΗΣ consumed more observing time aud we therefore dropped the observatious at 90 and 135 µια. We trie ο reobserve in mniuimuap mode those targets tha hac already been observed in chopped ος ising extra iue allocated when ISO lived onger than expected) mt succeeded ouly partially: several arects had left the observing window., Minimaps consumed more observing time and we therefore dropped the observations at 90 and 135 $\mu$ m. We tried to reobserve in minimap mode those targets that had already been observed in chopped mode (using extra time allocated when ISO lived longer than expected) but succeeded only partially: several targets had left the observing window.
 Iu total we use 65 hrs of observations., In total we used 65 hrs of observations.
 lu this article we discuss oulv the stellar flux densities derived from the 60 and 17022. minimaps., In this article we discuss only the stellar flux densities derived from the 60 and $\mu$ m minimaps.
 Appoeudix A contains a ctailed description of our lueasurenieut o»ocedure., Appendix A contains a detailed description of our measurement procedure.
 lustrunienutal problems Guadulv detector luCniorv effects) made us postpoue the reduction of the chopped ueasureiieuts until a later date: lis applies also to the ΠΑΣ (all chopped) measurements atWu., Instrumental problems (mainly detector memory effects) made us postpone the reduction of the chopped measurements until a later date; this applies also to the many (all chopped) measurements at.
".. We added published (Ábraliinetal.1998) ISOPIIOT ucasurenieuts of five A-type stars (9 Ua. + Ua. à UMa € UMa anc SO λα),"," We added published \citep{abra:98} ISOPHOT measurements of five A-type stars $\beta$ UMa, $\gamma$ UMa, $\delta$ UMa, $\epsilon$ UMa and 80 UMa)."
 The measurements have been obtaiue in a different mode from our observations. but we treat all measurements equallv.," The measurements have been obtained in a different mode from our observations, but we treat all measurements equally."
 These five stars are all at about pc (Perrvinanctal.1997).. sufficieutlv -earby to allow detection of the photospheric flux.," These five stars are all at about pc \citep{perr:97}, sufficiently nearby to allow detection of the photospheric flux."
 These stars are spectroscopic doubles aud they do not fulfill all of our selection criteria: below we argue why we iucluded hem anvlow., These stars are spectroscopic doubles and they do not fulfill all of our selection criteria; below we argue why we included them anyhow.
" Abvahdmetal.(1998) present ISOPIIOT ueasurenmieuts of four more stars. which they assuue to be at the same distance because all nine stars are supposed o be mienibers of au equidistant eroup called the ""Ursa Major stream."," \citet{abra:98} present ISOPHOT measurements of four more stars, which they assume to be at the same distance because all nine stars are supposed to be members of an equidistant group called the “Ursa Major stream”."
 The Hipparcos measurements (Perirviuauetal. 1997).. however. show that four of the niue stars are at a distance of pc and thus too far away to be useful for our purposes.," The Hipparcos measurements \citep{perr:97}, however, show that four of the nine stars are at a distance of pc and thus too far away to be useful for our purposes."
 All our data have been reduced using standard calibration tables aud the processing steps of OLPG/PIAT., All our data have been reduced using standard calibration tables and the processing steps of OLP6/PIA7.
 These steps include the iustruinental corrections and photometric calibration of the data., These steps include the instrumental corrections and photometric calibration of the data.
 At the time when we reduced our data there did not vet exist a standard procedure to extract the flux., At the time when we reduced our data there did not yet exist a standard procedure to extract the flux.
 We therefore developed aud used our own moethlod- see appendix D. Later versions of he software which contain upgrades of the photometric calibration do not significautlv alter our photometric results and the conclusions of this paper remain unchanged., We therefore developed and used our own method- see appendix B. Later versions of the software which contain upgrades of the photometric calibration do not significantly alter our photometric results and the conclusions of this paper remain unchanged.
 For each filter the observing mode gave two internal calibration measurements which were closely tuned to the actual sky brightuess., For each filter the observing mode gave two internal calibration measurements which were closely tuned to the actual sky brightness.
 This males the absolute calibration insensitive to instrumental effects as filter-to-filter calibrations aud signal non-luearities which were among others the main photometric calibration iuproveieuts for the upgrades., This makes the absolute calibration insensitive to instrumental effects as filter-to-filter calibrations and signal non-linearities which were among others the main photometric calibration improvements for the upgrades.
 In addition. if is standard procedure to eusure that each uperade does not degrade the photometric calibration of the validated nodes of the previous processiue version.," In addition, it is standard procedure to ensure that each upgrade does not degrade the photometric calibration of the validated modes of the previous processing version."
 The flux deusitics at aare presented in Table 3.., The flux densities at are presented in Table \ref{tab:60um}.
 The couteut of cach column is as follows: (1): the WD nuuber: (2): the TDT πο as used in the ISO archive: (3) the flux deusity corrected for bandwidth effects (assunmüusg that the spectrun is characterized bv the Ravleigh-Jeaus equation) and for the fact that the stellar ux extended over more than 1 pixel: (1) the error estimate assigned by the ISOPIIOT software to the flux measurement in column (3): for the TRAS incasurements the error has been put at 12Jv: (5) the flux expected frou the stellar photosphere. £177 as derived from Eq.," The content of each column is as follows: (1): the HD number; (2): the TDT number as used in the ISO archive; (3) the flux density corrected for bandwidth effects (assuming that the spectrum is characterized by the Rayleigh-Jeans equation) and for the fact that the stellar flux extended over more than 1 pixel; (4) the error estimate assigned by the ISOPHOT software to the flux measurement in column (3); for the IRAS measurements the error has been put at mJy; (5) the flux expected from the stellar photosphere, $F^{\mathrm{pred}}_{\nu}$ as derived from Eq."
" 1 using the V and (2BV) values in Table 2:: (6): the difference between columus (3) and (5): we call it the ""excess flix”. FU: (7): the ratio of the excess flux compared to the measurement error οἴνοιι in column (1): when we conchiles hat the CXCCSS IS real and not a measurement οπως we recalculated the monochromatic flux density by assuming a flat spectra. within the ISOPIIOT 60 jan bandwidth: the result is in columm (8) aud is called ω"," \ref{eq:plets} using the $V$ and $(B-V)$ values in Table \ref{tab:stars}; (6): the difference between columns (3) and (5); we call it the “excess flux”, $F_{\nu}^{\mathrm{exc}}$; (7): the ratio of the excess flux compared to the measurement error given in column (4); when we concluded that the excess is real and not a measurement error we recalculated the monochromatic flux density by assuming a flat spectrum within the ISOPHOT 60 $\mu$ m bandwidth; the result is in column (8) and is called $F_{\nu}^{\mathrm{disk}}$."
" Column (9) shows au extimate of zi, the optical depth of the disk at visual wavelengths. but estimated from the flux density at 60piu: see below for a clefinition aud see appendix D for more details."," Column (9) shows an estimate of $\tau_{60}^{\mathrm{disk}}$, the optical depth of the disk at visual wavelengths, but estimated from the flux density at $60\,\mu$ m; see below for a definition and see appendix D for more details."
 Flux densities iu columms 3. L|. 6 aud 7 have been corrected for the poiut spread ποο. beiug larger than the pixel size of the detector and for Ravleigh-Jeaus colour-correction (cc): for ISO Huxes. the inbaud fix has been divided by 069 (the correction for the poiut spread function (—psf). sec Appendix B) aud by 1.06 (ος) and for IRAS fluxes. the IFSC or IPSC flux have been divided by 1.91 (cc).," Flux densities in columns 3, 4, 6 and 7 have been corrected for the point spread function being larger than the pixel size of the detector and for Rayleigh-Jeans colour-correction (cc); for ISO fluxes, the inband flux has been divided by 0.69 (the correction for the point spread function (=psf), see Appendix B) and by 1.06 (cc) and for IRAS fluxes, the IFSC or IPSC flux have been divided by 1.31 (cc)."
" The ""disk euussion iu colin ὃ is corrected” from column 6.", The “disk emission” in column 8 is ``de-colour-corrected'' from column 6.
 We have checked the quality of our results at in two wavs: (1i) bv comparing ISO with IRAS &ux densities: Gi) bv comparing fluxes measured by ISO with predictions based on the (B8T) photometric index., We have checked the quality of our results at in two ways: (i) by comparing ISO with IRAS flux densities; (ii) by comparing fluxes measured by ISO with predictions based on the $(B-V)$ photometric index.
 The second approach allows us to assess the quality of ISO flux densities below the IRAS sensitivity lint., The second approach allows us to assess the quality of ISO flux densities below the IRAS sensitivity limit.
via different methods.,via different methods.
 It appears that. a reasonable percentage of the values of {1 is higher than the data estimates. so one can not reject the hypothesis that the underlying model is the FARIMA(0.4.0) process.," It appears that a reasonable percentage of the values of $H$ is higher than the data estimates, so one can not reject the hypothesis that the underlying model is the $d$ ,0) process."
 Nevertheless. we now look for an enhanced model which would describe better the behavior of different estimators obtained [or the original ancl shulllec data. and would improve the fit in terms of the prediction error.," Nevertheless, we now look for an enhanced model which would describe better the behavior of different estimators obtained for the original and shuffled data, and would improve the fit in terms of the prediction error."
 Thus. we propose a slieht eeneralization of the model incorporating the short-dependence component. namely FARIMA (2.d.0) model.," Thus, we propose a slight generalization of the model incorporating the short-dependence component, namely FARIMA $d$ ,0) model."
 We estimated the AR(2) coellicients: ας (linear term) and es (quaclratic term) via the mean-square error (MSE) minimalization scheme (aking into account (ree statistics: FIRT. VR and RS. see Figure 4..," We estimated the AR(2) coefficients: $a_1$ (linear term) and $a_2$ (quadratic term) via the mean-square error (MSE) minimalization scheme taking into account three statistics: FIRT, VR and RS, see Figure \ref{fig_mse}."
" The estimated values are: à,=0.02 and ag=0.03.", The estimated values are: $a_1 = 0.02$ and $a_2 = 0.03$.
 The FARIALA processes were generated according to the algorithm presented by ((2004))., The FARIMA processes were generated according to the algorithm presented by ).
 We caleulate yyapy. Ly Rand digg estimators for the simulated FARIALA (2.4.0) processes (lop panel in Figs.," We calculate $H_{FIRT}$, $H_VR$ and $d_{RS}$ estimators for the simulated FARIMA $d$ ,0) processes (top panel in Figs."
 5 and 6)) and the corresponding shuflled data (bottom panel in Figs., \ref{fig_box_est} and \ref{fig_box_rs}) ) and the corresponding shuffled data (bottom panel in Figs.
 5 and 6))., \ref{fig_box_est} and \ref{fig_box_rs}) ).
" We generate 1000 trajectories of size 2!"" which is close to the length of the original solar data. ee. 1089 and present the results in form of the so-called box plots."," We generate $1000$ trajectories of size $2^{10}$ which is close to the length of the original solar data, e. 1089 and present the results in form of the so-called box plots."
 The box plot has lines at the lower quartile. median. and upper quartile values.," The box plot has lines at the lower quartile, median, and upper quartile values."
 The whiskers are lines extending from each end of the box to show the extent of the rest of the data., The whiskers are lines extending from each end of the box to show the extent of the rest of the data.
 Outliers are data with values bevond the ends of the whiskers., Outliers are data with values beyond the ends of the whiskers.
 If there is no data outside the whisker. a dot is placed at the bottom whisker.," If there is no data outside the whisker, a dot is placed at the bottom whisker."
 The results correspond to the analvsis ol the original solar data included in Table 1.., The results correspond to the analysis of the original solar data included in Table \ref{tab1}.
 Thus. our FARIMA simulations reconstruct well the structure of the original solar flares data.," Thus, our FARIMA simulations reconstruct well the structure of the original solar flares data."
We conclude that the proper model could be based on the FARIMA(.d.0) process with Pareto innovations with the parameters eq—0.02. ao=0.03. d—0.19. and a=1.25 which has the long-range dependence property since d>1—2/0 2008)).,"We conclude that the proper model could be based on the $d$ ,0) process with Pareto innovations with the parameters $a_1 = 0.02$, $a_2 = 0.03$, $d=0.19$, and $\alpha = 1.25$ which has the long-range dependence property since $d>1-2/\alpha$ )."
 The Pareto distribution is also convenient for the present analvsis because il gives a description of positive random variables whereas the Lévvy-stable one for 1<a«2 is related to both positive and negative random variables. but a (ime series of x-ray [lare energy is quite positive by definition.," The Pareto distribution is also convenient for the present analysis because it gives a description of positive random variables whereas the Lévvy-stable one for $1<\alpha<2$ is related to both positive and negative random variables, but a time series of x-ray flare energy is quite positive by definition."
 We also calculated the I-dav-ahead. prediction for the FARIAMA(2.d.0) time series (for the prediction discussion in (he infinite variance FARIALA case see 1995)).The results are depicted in Figure 7..," We also calculated the 1-day-ahead prediction for the $d$ ,0) time series (for the prediction discussion in the infinite variance FARIMA case see ).The results are depicted in Figure \ref{fig_pred}."
 In this paper we demonstrate how self-similar models driven by Lévvy stable noise, In this paper we demonstrate how self-similar models driven by Lévvy stable noise
The simplest explanation for this troud is simply that the radiative time is a steeper fiction of temperature than the advective time: advective efficiency is giveu roughly by the ratio of the radiative aud advective times ee:Cowan&Agol2010).,The simplest explanation for this trend is simply that the radiative time is a steeper function of temperature than the advective time: advective efficiency is given roughly by the ratio of the radiative and advective times \citep[eg:][]{Cowan_2010}.
. In the limit of Newtonian cota the radiative time scales as TaqXdoC.m ," In the limit of Newtonian cooling, the radiative time scales as $\tau_{\rm rad} \propto T^{-3}$."
"df one asstunes the wind speed to be of order the local sound speed. then the advective tine scales απ TuaeκT ""2. "," If one assumes the wind speed to be of order the local sound speed, then the advective time scales as $\tau_{\rm adv} \propto T^{-0.5}$ ."
One might therefore naivelv expect the advective efficioney to scale as Z7—77., One might therefore naively expect the advective efficiency to scale as $T^{-2.5}$.
 Such an explanation would not explain the apparcut sharp transition seen at 2100 Is. however.," Such an explanation would not explain the apparent sharp transition seen at 2400 K, however."
 The combination of optical observations of secondary eclipses and thermal observations of phase variations is the best wav to coustrain planetary albedo and circulation., The combination of optical observations of secondary eclipses and thermal observations of phase variations is the best way to constrain planetary albedo and circulation.
 The optical observations should be taken near the stars blackbody peak. both to maximize sigual-to-noise. aud to avoid contamination from the planets thermal euission. but this separation may not be possible for the hottest transiting planets.," The optical observations should be taken near the star's blackbody peak, both to maximize signal-to-noise, and to avoid contamination from the planet's thermal emission, but this separation may not be possible for the hottest transiting planets."
 The thermal observations. likewise. should be near the plauct’s blackbody peak to better coustrain its bolometric flux.," The thermal observations, likewise, should be near the planet's blackbody peak to better constrain its bolometric flux."
 Note that this waveleneth is shortward of the ideal contrast ratio. which typically falls on the planets Ravleigh-Jeaus tail.," Note that this wavelength is shortward of the ideal contrast ratio, which typically falls on the planet's Rayleigh-Jeans tail."
" Furthermore. the thermal phase observations should span a full planetary orbit: the light curve niünmuun is the most sensitive measure of 5, and should occur nearly half an orbit apart from the liebt curve maxiuimn. despite skewed diurnal heating patterus (Cowan&Agol2008.2010)."," Furthermore, the thermal phase observations should span a full planetary orbit: the light curve minimum is the most sensitive measure of $\varepsilon$, and should occur nearly half an orbit apart from the light curve maximum, despite skewed diurnal heating patterns \citep{Cowan_2008, Cowan_2010}."
 This means that observing campaigns that ouly cover a little more than half an orbit (transit » eclipse} are probably underestimating the real peak-trough phase amplitude., This means that observing campaigns that only cover a little more than half an orbit (transit $\to$ eclipse) are probably underestimating the real peak-trough phase amplitude.
 A possible improvement to this study would be to perforin a uniform data reduction for all the exoplauct observations of hot Jupiters., A possible improvement to this study would be to perform a uniform data reduction for all the exoplanet observations of hot Jupiters.
 These data make up the majority of the constraints presented iu our study and most are publicly available., These data make up the majority of the constraints presented in our study and most are publicly available.
 And while the published observations were analyzed im disparate wavs. a consensus approach to correcting detector svsteimaties is beeiuniung to cmerec.," And while the published observations were analyzed in disparate ways, a consensus approach to correcting detector systematics is beginning to emerge."
 N.D.C. acknowledges useful discussions of aspects of this work with T. Robinson. ALS. MLiulew. JJ. Fortuev. T.S. Barman and D.S. Spiegel.," N.B.C. acknowledges useful discussions of aspects of this work with T. Robinson, M.S. Marley, J.J. Fortney, T.S. Barman and D.S. Spiegel."
 Thanks to our referee BALS. [Hausen for iusightful feedback. and to E.D. Feigelson for suggestions about statistical methods.," Thanks to our referee B.M.S. Hansen for insightful feedback, and to E.D. Feigelson for suggestions about statistical methods."
 N.DB.C. was supported by the Natural Scieuces and Eneineering Research Council of Canada., N.B.C. was supported by the Natural Sciences and Engineering Research Council of Canada.
 E.À. i supported by a National Science Foundation Carecr Craut., E.A. is supported by a National Science Foundation Career Grant.
 Support for this work was provided bv NASA through an award issued by JPL/Caltech., Support for this work was provided by NASA through an award issued by JPL/Caltech.
 This research has made use of the Exoplanet Orbit Database and the Exoplanet Data Explorer at exoplaucts.ore., This research has made use of the Exoplanet Orbit Database and the Exoplanet Data Explorer at exoplanets.org.
From this state of the data reduction on. SNIFS usually reduces ils supernova observations with a data reduction pipeline carefully tuned. to the needs of the Nearby Supernova Factory project.,"From this state of the data reduction on, SNIFS usually reduces its supernova observations with a data reduction pipeline carefully tuned to the needs of the Nearby Supernova Factory project."
 While these pipeline data products were very useful for an initial inspection of the data Ilrom the Deep Impact observations. we decided for the final analysis of the data to bypass these photometric extraction functions of the SNIFS pipeline aud instead to extract the photometric information directly from the Euro3D data cubes using IRAF (Tock. 1986) scripts.," While these pipeline data products were very useful for an initial inspection of the data from the Deep Impact observations, we decided for the final analysis of the data to bypass these photometric extraction functions of the SNIFS pipeline and instead to extract the photometric information directly from the Euro3D data cubes using IRAF (Tody, 1986) scripts."
 The main reason was to achieve better control of atmospheric dispersion in the extraction of broad-band spectrophotometry [rom the data cubes. better elimination of residual cosmic ray events in (he data. and a better handling of the changing [ιν distribution over the course of (he observations due to the expansion of the debris cloud around the comet.," The main reason was to achieve better control of atmospheric dispersion in the extraction of broad-band spectrophotometry from the data cubes, better elimination of residual cosmic ray events in the data, and a better handling of the changing flux distribution over the course of the observations due to the expansion of the debris cloud around the comet."
 Initially. the 15x 15xA data cubes were processed into 225xA two-dimensional spectra.," Initially, the $\times$ $\times\lambda$ data cubes were processed into $\times\lambda$ two-dimensional spectra."
 In this format. the few residual cosmic rav events that had. escaped automatic removal could be easily identified and interactively eliminated using the IRAF image editing task (imedit).," In this format, the few residual cosmic ray events that had escaped automatic removal could be easily identified and interactively eliminated using the IRAF image editing task (imedit)."
 As an example. Fig.," As an example, Fig."
 1 shows the extracted spectra in the blue and red channel in {his two-dimensional format. which is essentially a set of 15 long (6) slit spectra stacked in vertical direction.," 1 shows the extracted spectra in the blue and red channel in this two-dimensional format, which is essentially a set of 15 long $\arcsec$ ) slit spectra stacked in vertical direction."
 Note that lor Fig., Note that for Fig.
 1. we show a spectrum after subtraction of an “empty skv frame taken a few minutes before the comet spectrum ab a position 5' [rom (he comet nucleus. outside of the coma.," 1, we show a spectrum after subtraction of an “empty sky” frame taken a few minutes before the comet spectrum at a position $\arcmin$ from the comet nucleus, outside of the coma."
 The sky subtraction was done to eliminate night skv emission lines [rom Fig., The sky subtraction was done to eliminate night sky emission lines from Fig.
 1 which would otherwise limit the visibility of Taint emission features in the comet's coma., 1 which would otherwise limit the visibility of faint emission features in the comet's coma.
 In this extracted. long-slit data format (Fig.," In this extracted, long-slit data format (Fig."
 1). the CN (0-0) emission band al 388 nm is seen prominently.," 1), the CN (0-0) emission band at 388 nm is seen prominently."
" The C, emission band at 405 nm is also faintly indicated in (he blue channel. but the signal-to-noise ratio is too low to analvze this emission in anv more detail."," The $_3$ emission band at 405 nm is also faintly indicated in the blue channel, but the signal-to-noise ratio is too low to analyze this emission in any more detail."
 The Cs Swan bands were not clearly detected., The $_2$ Swan bands were not clearly detected.
 The brightest Cs band at, The brightest $_2$ band at
QSO itself (Petitjean et al.,QSO itself (Petitjean et al.
 1996)., 1996).
 The lack of CO and thermal dust. emission from (his position argues for the latter. while the lack of N V and C IV emission lines argues for (he former (Fontana οἱ al.," The lack of CO and thermal dust emission from this position argues for the latter, while the lack of N V and C IV emission lines argues for the former (Fontana et al."
 1998: Ohta οἱ al., 1998; Ohta et al.
 2000)., 2000).
 QGuilloteau et al. (, Guilloteau et al. (
1997) have detected CO (54) line emission from BRI 13350417 at 2=44074. with a line FWIIM = 420 kms | and an integrated line intensity of 2.840. Jv kms |.,"1997) have detected CO (5–4) line emission from BRI 1335–0417 at $z = 4.4074$, with a line FWHM = $420$ km $^{-1}$ and an integrated line intensity of $\pm$ 0.3 Jy km $^{-1}$."
" The flux density of the dust continuum emission measured at 1.35nun with the Plateau de Bure Interferometer (PAB) at 2"" resolution is 5.61.1 mJy. while that measured al 1.25mm with the IRAM 30m telescope (10.6"" resolution) is 10.3+1.4 mJy."," The flux density of the dust continuum emission measured at 1.35mm with the Plateau de Bure Interferometer (PdBI) at $''$ resolution is $\pm$ 1.1 mJy, while that measured at 1.25mm with the IRAM 30m telescope $''$ resolution) is $10.3\pm1.4$ mJy."
 The implied [ar LR luminosity based on the 1.25 mm measurement is 3.1xLOM L.., The implied far IR luminosity based on the 1.25 mm measurement is $3.1 \times 10^{13}$ $_\odot$.
 The 135mm PdbI observations shows marginal evidence for an extended source. wilh a formal Gaussian size of L.02:0.47. with major axis oriented roughly north-south.," The 1.35mm PdBI observations shows marginal evidence for an extended source, with a formal Gaussian size of $1.0 \pm 0.4''$, with major axis oriented roughly north-south."
 No hieh resolution optical images have been published of 13350417., No high resolution optical images have been published of 1335–0417.
" The Dss shows an unresolved source with Mj = 27.3. although with a pixel scale of just 1"" the details of the source structure remain unknown."," The DSS shows an unresolved source with $_B$ = –27.3, although with a pixel scale of just $''$ the details of the source structure remain unknown."
 BRI 1335-0417 has been detected in the radio continuum. with an integrated flux density of 220+43 μ.]ν at 1.4 GIIz and 16—11 idv ad 4.9 Gllz (Carilli et al.," BRI $-$ 0417 has been detected in the radio continuum, with an integrated flux density of $220\pm
43$ $\mu$ Jy at 1.4 GHz and $76\pm 11$ $\mu$ Jy at 4.9 GHz (Carilli et al."
 1999)., 1999).
 Table 1 lists the observing parameters for BRI 12020725 and BRI 1335-0417., Table 1 lists the observing parameters for BRI 1202–0725 and BRI $-$ 0417.
 We also present a re-analvsis of the D configuration observations from. Crilli et al. (, We also present a re-analysis of the D configuration observations from Carilli et al. (
1999).,1999).
 Standard amplitude ancl phase calibration were applied. correcting for atmospheric," Standard amplitude and phase calibration were applied, correcting for atmospheric"
A complete journal of our observations is given in Table 1L..,A complete journal of our observations is given in Table \ref{tab:obs}.
 The number of radial velocity. observations which we List in this table are those which are presented in this paper. in acdition to the measurements previously presented in ?. and ?..," The number of radial velocity observations which we list in this table are those which are presented in this paper, in addition to the measurements previously presented in \citet{Maxted02} and \citet{MoralesRueda03}."
 The majority of the cata usec in this study were collected using the Intermediate. Dispersion Spectrograph (LDS) mounted on the 2.5m Isaac Newton Telescope., The majority of the data used in this study were collected using the Intermediate Dispersion Spectrograph (IDS) mounted on the 2.5m Isaac Newton Telescope.
 Our earliest observations in 2000/2001 used the 500mm camera with the 112001 grating centred on aand the PEW (Lk 1k) CCD eiving a dispersion of AA//pix and a resolution ofAA., Our earliest observations in 2000/2001 used the 500mm camera with the R1200R grating centred on and the TEK (1k $\times$ 1k) CCD giving a dispersion of /pix and a resolution of.
 Subsequent observations used the 235 mm camera with the I1200D erating and the thinned LEV10 (2k « 4k) CCD. giving a dispersion of AA//pix and a resolution ofAG.," Subsequent observations used the 235 mm camera with the R1200B grating and the thinned EEV10 (2k $\times$ 4k) CCD, giving a dispersion of /pix and a resolution of."
 The central wavelength used with this grating varied slightly between observing runs. but in all cases was chosen to cover aandσ," The central wavelength used with this grating varied slightly between observing runs, but in all cases was chosen to cover and."
ι In 2003 we expanded this project to cover southern targets with the first of a series of observing runs using the grating spectrograph plus the SITe CCD mounted on the 1.9mm telescope at the Sutherland site of the South African Astronomical Observatory (SAO)., In 2003 we expanded this project to cover southern targets with the first of a series of observing runs using the grating spectrograph plus the SITe CCD mounted on the m telescope at the Sutherland site of the South African Astronomical Observatory (SAAO).
 Crating 4. with 1200 grooves per mm. was used to obtain spectra covering aand wwith a dispersion of /pix ancl a resolution of better than aat.," Grating 4, with 1200 grooves per mm, was used to obtain spectra covering and with a dispersion of /pix and a resolution of better than at."
4600... We took two consecutive observations of cach object and bracketed them with Cur. frames to calibrate the spectra in wavelength., We took two consecutive observations of each object and bracketed them with CuAr frames to calibrate the spectra in wavelength.
 After debiasing and Latficlling the frames (using Tungsten. Latlicll frames for the INT115 observations and dome Lats for the SAAO observations). spectral extraction proceeded. according το the optimal algorithm of ?..," After debiasing and flatfielding the frames (using Tungsten flatfield frames for the INT/IDS observations and dome flats for the SAAO observations), spectral extraction proceeded according to the optimal algorithm of \citet{Marsh89}."
 The ares were extracted using the profile associated: with their corresponding target to avoid systematic errors caused by the spectra being tilted., The arcs were extracted using the profile associated with their corresponding target to avoid systematic errors caused by the spectra being tilted.
 Uncertainties on every point were propagated through every stage of the data reduction., Uncertainties on every point were propagated through every stage of the data reduction.
 To measure the radial velocities we use least squares fitting of a model line profile and follow the same procedure described. in ?.., To measure the radial velocities we use least squares fitting of a model line profile and follow the same procedure described in \citet{MoralesRueda03}.
 The model line profile consists of three CGaussians with cillerent widths and depths., The model line profile consists of three Gaussians with different widths and depths.
 For any. given sdD. the widths and depths of the Gaussians are optimised and then fixed while their. velocity. olfsets from the rest wavelengths of the lines in question are fitted separately for each spectrum: see 7. for further details of this procedure.," For any given sdB, the widths and depths of the Gaussians are optimised and then fixed while their velocity offsets from the rest wavelengths of the lines in question are fitted separately for each spectrum; see \citet{Maxted00b} for further details of this procedure."
 For the red spectra obtained in Apr 2000 - Alay 2001 we fit the Hine., For the red spectra obtained in Apr 2000 - May 2001 we fit the line.
 For all other spectra the fitting was performed simultaneously for aand, For all other spectra the fitting was performed simultaneously for and.
 In Tables 9 and Lo we list all of our racial velocity measurements., In Tables \ref{tab:rv1a} and \ref{tab:rv2a} we list all of our radial velocity measurements.
 Table 9— contains the measurements [or 28 svstems which we have found to be binary. and for which we have found the orbital period.," Table \ref{tab:rv1a} contains the measurements for 28 systems which we have found to be binary, and for which we have found the orbital period."
 The description of the orbital period. determination is given in Section 3.2.., The description of the orbital period determination is given in Section \ref{sec:periods}.
 ‘Table 1. contains measurements for all of the remaining sdDs covered by our project which we have not. previously published. a total of LOS objects.," Table \ref{tab:rv2a} contains measurements for all of the remaining sdBs covered by our project which we have not previously published, a total of 108 objects."
 Most of these objects show no evidence for orbital variations. and so are likely to be either single sdDs. or binary systems in which the mass function is too low for us to detect radial velocity variations (due to à low companion mass. a long binary period. a binary inclination close to zero or a combination of these factors).," Most of these objects show no evidence for orbital variations, and so are likely to be either single sdBs, or binary systems in which the mass function is too low for us to detect radial velocity variations (due to a low companion mass, a long binary period, a binary inclination close to zero or a combination of these factors)."
 For some objects the non-cetection of variation may be because the number of observations of that object is small., For some objects the non-detection of variation may be because the number of observations of that object is small.
 In some other cases. there are clear signs that the sdB is in à binary system. however there are still competing orbital aliases of comparable significance and so our data is insullicient to determine the true orbital period.," In some other cases, there are clear signs that the sdB is in a binary system, however there are still competing orbital aliases of comparable significance and so our data is insufficient to determine the true orbital period."
 These are strong candidates for future observation., These are strong candidates for future observation.
 In this paper we will concentrate on the 28 solved systems given in Table 9.., In this paper we will concentrate on the 28 solved systems given in Table \ref{tab:rv1a}.
 In Section 4.3. we discuss the binary fraction of the sdB population and in Section 44 we investigate the nature of the companion stars., In Section \ref{sec:binfrac} we discuss the binary fraction of the sdB population and in Section \ref{sec:composite} we investigate the nature of the companion stars.
 For these sections we consider the complete list of objects in our programme., For these sections we consider the complete list of objects in our programme.
 In racial velocity work. while one can very soon determine a star to be binary. it can take much longer to determine the orbital period.," In radial velocity work, while one can very soon determine a star to be binary, it can take much longer to determine the orbital period."
 With a small amount of data there is a danger of picking an incorrect alias., With a small amount of data there is a danger of picking an incorrect alias.
 This period can be very wrong. even when the quoted uncertainty is tiny. because the statistics are not Gaussian and so an error of LOO or even 1000 times the quoted uncertainty can happen.," This period can be very wrong, even when the quoted uncertainty is tiny, because the statistics are not Gaussian and so an error of 100 or even 1000 times the quoted uncertainty can happen."
 On the other hand. the process of collecting sulficient data to determine the true period bevond doubt can be very inellieient in terms of telescope time.," On the other hand, the process of collecting sufficient data to determine the true period beyond doubt can be very inefficient in terms of telescope time."
 It was therefore necessary for us to select a criterion bv which we consider a binary orbit to be solved., It was therefore necessary for us to select a criterion by which we consider a binary orbit to be solved.
 Our initial rule of thumb! was to consider the period in a system to be determined if the lowest minimum of the v function has à X7 at least 10 less than the next lowest minimum., Our initial `rule of thumb' was to consider the period in a system to be determined if the lowest minimum of the $\chi^2$ function has a $\chi^2$ at least 10 less than the next lowest minimum.
 From a Bayesian point of view. this z>10 criterion is equivalent to requiring that the second. best period is 2exp5=150 times less probable than the best.," From a Bayesian point of view, this $\Delta\chi^2>10$ criterion is equivalent to requiring that the second best period is $\ga \exp 5 = 150$ times less probable than the best."
 This argument. although true. is not. precise because while the peak of the second. alias may be 150 times less probable than the peak of the best. alias. there is no euarantee that the total probability. of other period is as low.," This argument, although true, is not precise because while the peak of the second alias may be $> 150$ times less probable than the peak of the best alias, there is no guarantee that the total probability of other period is as low."
 We therefore instead chose to determine. the probabilities of the true orbital period being within a certain range of our favoured value., We therefore instead chose to determine the probabilities of the true orbital period being within a certain range of our favoured value.
 The details of this Bayesian calculation are given in ? and ?.., The details of this Bayesian calculation are given in \citet{MoralesRueda03} and \citet{Marsh95}.
 For the purposes. of comparing the observed κά binary period distribution to heoretical evolutionary models. knowing the orbital perio o within £5 per cent is sullicient.," For the purposes of comparing the observed sdB binary period distribution to theoretical evolutionary models, knowing the orbital period to within $\pm 5$ per cent is sufficient."
 We therefore establishec he following criterion: we considered a system to be solvec if the probability that the true orbital period lies further han 5 per cent from our best alias is less than 0.1 per cent., We therefore established the following criterion: we considered a system to be solved if the probability that the true orbital period lies further than $5$ per cent from our best alias is less than $0.1$ per cent.
 For 28 of our targets we have enough racdia velocity measurements to satisfy this criterion., For 28 of our targets we have enough radial velocity measurements to satisfy this criterion.
— Me lve previously announced nine of these in conference sroceedings (DOGO0034|150. PC1230|052. P€H2441113. 611519|G40. PC1528|104. WPD2040|3955. 2:: EC'00404-4429. EC02200-2338. ο. IC12408-1427. 2)). but we present rere the more detailed: analysis.," We have previously announced nine of these in conference proceedings (PG0934+186, PG1230+052, PG1244+113, PG1519+640, PG1528+104, KPD2040+3955, \citealt{MoralesRueda03a}; EC00404-4429, EC02200-2338, \citealt{MoralesRueda05}, EC12408-1427 \citealt{MoralesRueda06}) ), but we present here the more detailed analysis."
 Phe orbit of one of our argets has been independently. determined. (PCL000|408. 7)) and so we present our results to corroborate this," The orbit of one of our targets has been independently determined (PG1000+408, \citealt{Shimanskii08}) ) and so we present our results to corroborate this"
abundance below the average derived from the rest of galaxies in our sample by a factor of e»3 and ~20. respectively.,"abundance below the average derived from the rest of galaxies in our sample by a factor of $\sim3$ and $\sim20$, respectively."
 We do not observe major differences in the derived abunclances among the AGNs and SBGs in our sample., We do not observe major differences in the derived abundances among the AGNs and SBGs in our sample.
 This homogeneity suggests that the tvpe of nuclear activity does nol plav a major role at large scales in the production/destruction of these species in the bulk of the molecular gas., This homogeneity suggests that the type of nuclear activity does not play a major role at large scales in the production/destruction of these species in the bulk of the molecular gas.
" This ellect is illustrated by the two positions observed in 11063. where the LINCO and ""CS abundances decrease by less than a factor of two from the nuclear AGN to the SB ring."," This effect is illustrated by the two positions observed in 1068, where the HNCO and $^{34}$ CS abundances decrease by less than a factor of two from the nuclear AGN to the SB ring."
 Of course we have to bear in mind that the two regions are not fully resolved by the beam of the IRAAI 30mm telescope at these frequencies., Of course we have to bear in mind that the two regions are not fully resolved by the beam of the IRAM m telescope at these frequencies.
 If. as suggested by the observations in theGalactic Center region (Martínetal.2008).. the abundance of these species is mostly driven by (he presence of shocks and strong UV radiation fields in the ISM. then the observed differences are likely linked to star formation aclivily in these galaxies.," If, as suggested by the observations in theGalactic Center region \citep{Martin08}, the abundance of these species is mostly driven by the presence of shocks and strong UV radiation fields in the ISM, then the observed differences are likely linked to star formation activity in these galaxies."
 Significant differences are observed within the galaxies 3342 and 22., Significant differences are observed within the galaxies 342 and 2.
 In 3342. the INCO/CA%O ratio increases towards the position in the minispiral and decreases in (he central region where ISAM is illuminated by the UV radiation from the central clusters.," In 342, the $^{18}$ O ratio increases towards the position in the minispiral and decreases in the central region where ISM is illuminated by the UV radiation from the central clusters."
" These chemical differences. smoothed to the 10""—20"" single-dish bean size. are much more contrasted in the high-resolution interferometric maps by 2005)."," These chemical differences, smoothed to the $10''-20''$ single-dish beam size, are much more contrasted in the high-resolution interferometric maps by \citep{Meier05}."
. Although not spatially. we resolve the nuclear region of 22 by the two velocity components.," Although not spatially, we resolve the nuclear region of 2 by the two velocity components."
 As seen from the high-resolution maps of CO (Meierοἱal.2008). the lower velocily component (—100 kms)) is associated to the nuclear and northeastern molecular arm. while the higher velocity component kms)) will be mostly associated to the nuclear and southwest molecular arm.," As seen from the high-resolution maps of $^{13}$ CO \citep{Meier08} the lower velocity component $-100$ ) is associated to the nuclear and northeastern molecular arm, while the higher velocity component ) will be mostly associated to the nuclear and southwest molecular arm."
 We clearly see a difference in the chemistry of both components with an increase of the WNCO abundance towards (he southwest molecular arm of a [actor of ~2 as well as a decrease of the Cs abundance of a factor of 3—4., We clearly see a difference in the chemistry of both components with an increase of the HNCO abundance towards the southwest molecular arm of a factor of $\sim2$ as well as a decrease of the CS abundance of a factor of $3-4$.
 Overall. the INCO/CS ratio increases bv almost an order of magnitude.," Overall, the HNCO/CS ratio increases by almost an order of magnitude."
 We can compare the ratios measured in galaxies with the abundances derived in the central region of the Milkv. Way. (Martínetal.2003)., We can compare the ratios measured in galaxies with the abundances derived in the central region of the Milky Way \citep{Martin08}.
. The vast majority of galaxies show IINCO abundances intermediate between (hose found in the shock dominated giant molecular clouds (GAICs) ancl the Galactic PDR. prototypes (see dashed lines in Fig 3))., The vast majority of galaxies show HNCO abundances intermediate between those found in the shock dominated giant molecular clouds (GMCs) and the Galactic PDR prototypes (see dashed lines in Fig \ref{fig.CSHNCO}) ).
 Two sources in our sample. namely 882 and 3342* clearly [all in the range of abundances found in GMCSs.," Two sources in our sample, namely 83 and $^*$ clearly fall in the range of abundances found in GMCs."
 On the contrary. ALT&82 represents the other extreme where LINCO emission is utterly vanished.," On the contrary, 82 represents the other extreme where HNCO emission is utterly vanished."
 The upper limit to the detection derived alter a deep 5.5 hours integration on the IINCO 5—4 is an order of magnitude lower (han (he previous limit to the detection in (he central position (Nguyen-Q-Tieuetal. 1991).. but still INCO remains undetected.," The upper limit to the detection derived after a deep 5.5 hours integration on the HNCO $5-4$ is an order of magnitude lower than the previous limit to the detection in the central position \citep{Nguyen91}, , but still HNCO remains undetected."
solar model with an accuracy of 1077 ou the solar radius and huninosity.,solar model with an accuracy of $10^{-5}$ on the solar radius and luminosity.
 The performance of our code in the rauge of masses (0.7. AL. to 3.5 AD.) ancl evolutionary stages studied iu this work was successfully tested by comparing our computed isochrones with those of Cürardictal.(2000).., The performance of our code in the range of masses (0.7 $_{\odot}$ to 3.5 $_{\odot}$ ) and evolutionary stages studied in this work was successfully tested by comparing our computed isochrones with those of \cite{art-Girardietal2000A&AS}.
 The stars computed in this work are embedded in a deuse halo of DAL, The stars computed in this work are embedded in a dense halo of DM.
 To account for the impact of the DAL particles on the stars we cousidered that some of the DM. particles that populate the halo are eravitationally captured by the stars aud accmuulate in their interior., To account for the impact of the DM particles on the stars we considered that some of the DM particles that populate the halo are gravitationally captured by the stars and accumulate in their interior.
 The nmuber of captured DM particles was computed using the integral expression of Could (1987).. as implemented in Coudoloetal.(2001).," The number of captured DM particles was computed using the integral expression of \cite{art-Gould1987}, as implemented in \cite{art-GondoloEdsjoDarkSusy2004}."
. Note that. for the capture process to be οποια. the DM particles are assiuned to have a non-negligible scattering cross section with barvous σε. which we chose to be sanaller than the present lanits frou direct detection experuuents σος=10 ‘You? (Alumedetal.2010) aud asp=1055 cur? (Behukeetal.2011) for a WIMP with a mass of LOO CeV. For these values of the spin-depeudent (SD) interactious with lydrogen σι.atoms always dominate over the spin-independeut (SI) ones with other stellar isotopes.," Note that, for the capture process to be efficient, the DM particles are assumed to have a non-negligible scattering cross section with baryons $\sigma_{\chi}$, which we chose to be smaller than the present limits from direct detection experiments: $\sigma_{\chi,SI}=10^{-44}\;$ $^2$ \citep{art-CDMS2010Sci} and $\sigma_{\chi,SD}=10^{-38}\;$ $^2$ \citep{art-COUPP2011PhRvL} for a WIMP with a mass of 100 GeV. For these values of $\sigma_{\chi}$, the spin-dependent (SD) interactions with hydrogen atoms always dominate over the spin-independent (SI) ones with other stellar isotopes."
" Iu the capture rate (C) calculation we assuned a stellar velocity e,=220 kni 1 aud a Maxwellian DM velocity, distribution with a dispersion e,=270 kin +.", In the capture rate $C_{\chi}$ ) calculation we assumed a stellar velocity $v_{\star}=220\;$ km $^{-1}$ and a Maxwellian DM velocity distribution with a dispersion $\bar{v_{\chi}}=270\;$ km $^{-1}$.
 These values apply for the solar case. but are certainly inaccurate for a nuclear cluster.," These values apply for the solar case, but are certainly inaccurate for a nuclear cluster."
 For iustauce. stars with velocities as high as 100 kins | are observed near the Calactic center (CC) (Luctal.2009).," For instance, stars with velocities as high as 400 km $^{-1}$ are observed near the Galactic center (GC) \citep{art-LuGhezetal2009}."
. Tn this case the capture rate would be reduced by a factor of 6 (for a more thorough analysis of how C. varics for different stellar and DAL characteristics sce Lopes.Casuiellas&Eneénio (2011)))., In this case the capture rate would be reduced by a factor of 6 (for a more thorough analysis of how $C_{\chi}$ varies for different stellar and DM characteristics see \cite{art-LopCasEug2011PhRevD}) ).
 At the same time. it is complex. to model the DAL velocity distribution iu the GC. as the motion of the DM particles is strongly influcuced by the eravitational potential of the stars and the central black hole.," At the same time, it is complex to model the DM velocity distribution in the GC, as the motion of the DM particles is strongly influenced by the gravitational potential of the stars and the central black hole."
" Tuterestinely, Scottetal.(2000). tested other DAL velocity distributions with the aim of grasping the possible varlatious on Cy."," Interestingly, \cite{art-Scottetal2009MNRAS} tested other DM velocity distributions with the aim of grasping the possible variations on $C_{\chi}$."
 When a non-CGaussiui distribution (desigued to fif a N-body simulation of a Milkv Way-size DAL halo) was iurplemieuted. thecapture rate was boosted by a factor of 3-5.," When a non-Gaussian distribution (designed to fit a N-body simulation of a Milky Way-size DM halo) was implemented, thecapture rate was boosted by a factor of 3-5."
 On the other laud. the same authors fouud that the truncation of the isothermal distribution at the local escape velocity reduces by a factor of 2.," On the other hand, the same authors found that the truncation of the isothermal distribution at the local escape velocity reduces $C_{\chi}$ by a factor of 2."
 Tho sane order of tucertainty ou Cy is expected in the cases preseuted in the present work., The same order of uncertainty on $C_{\chi}$ is expected in the cases presented in the present work.
 After some scatterings. the DM. particles sink to the core of the star and rapidly thermalize with stellar matter.," After some scatterings, the DM particles sink to the core of the star and rapidly thermalize with stellar matter."
 The umuber of DAL particles iu the stellar core increases until their scltamnuililation rate balances the capture rate., The number of DM particles in the stellar core increases until their self-annihilation rate balances the capture rate.
 This equilibziun is reached in a time scale below vr for all cases studied here., This equilibrium is reached in a time scale below $^4$ yr for all cases studied here.
 Thus. the aunihilation of DM. particles provides a new source of energv which contributes to the total luminosity of the star according to (Salat&Silk1989):: where mu ds tho mass of the DM particles. aud fi=2/3 to take iuto account that one third of the energy may escape the star in the form of neutrinos (loccoetal. 2008)..," Thus, the annihilation of DM particles provides a new source of energy which contributes to the total luminosity of the star according to \citep{art-SalatiSilk1989}: where $m_{\chi}$ is the mass of the DM particles, and $f_\chi=2/3$ to take into account that one third of the energy may escape the star in the form of neutrinos \citep{art-Ioccoetal2008}."
" This energy is injected to the stellar mocels following the thermal distribution of tle DAL particles. which characteristic radius is below aud of the stellar radius for in,=100 Ce aud 8 GoV respectively."," This energy is injected to the stellar models following the thermal distribution of the DM particles, which characteristic radius is below and of the stellar radius for $m_{\chi}=100\;$ GeV and 8 GeV respectively."
 The total iuput of energy from DM auniililation. aud thus also its imipact on stellar evolution. will depend mainly ou the product pio.," The total input of energy from DM annihilation, and thus also its impact on stellar evolution, will depend mainly on the product $\rho_{\chi}\sigma_{\chi}$."
 The lvdrostatic equilibriun (the balance between xessure aud eravitv) achieved by a star within a dense DAL halo differs from the one reached in the classical Xeture due to the new source of enerev added to he classical thermouuclear energev sources., The hydrostatic equilibrium (the balance between pressure and gravity) achieved by a star within a dense DM halo differs from the one reached in the classical picture due to the new source of energy added to the classical thermonuclear energy sources.
 This fact cads to three main consequeuces that will influence the characteristics of the whole cluster: l.speed:, This fact leads to three main consequences that will influence the characteristics of the whole cluster: 1.:
 The ceutral eniperature of stars that evolve within dense DAL halos is lower than that of classical stars due to their negative recat capacity., The central temperature of stars that evolve within dense DM halos is lower than that of classical stars due to their negative heat capacity.
 Another simple way to understand this is to inaegine a forming star iu the pre-main sequence., Another simple way to understand this is to imagine a forming star in the pre-main sequence.
 The cloud of gas that forms the proto-star shrinks. Increasing its central teniperature uutil the eravitational collapse is balanced by the thermonuclear reactions: if another source of energev helps to compensate eravity. the hydrostatie equilibrimm is reached. earlier. when the central temperature is lower.," The cloud of gas that forms the proto-star shrinks, increasing its central temperature until the gravitational collapse is balanced by the thermonuclear reactions; if another source of energy helps to compensate gravity, the hydrostatic equilibrium is reached earlier, when the central temperature is lower."
 Therefore. stars within dense DAL halos burn hydrogen at a lower rate. slowing down their evolution through later pliases.," Therefore, stars within dense DM halos burn hydrogen at a lower rate, slowing down their evolution through later phases."
" For example. a star of 1 AL. will spend more than 20 Car in the lain sequence (MS) if it evolves iu à DM halo of deusity py=2-10? GeV 7 (assuming «p=107 cu’. although other values of p, aud ayspeq, cau be cousidered. leading to the same effects as loug as the product pio is kept constant)."," For example, a star of 1 $_{\odot}$ will spend more than $20$ Gyr in the main sequence (MS) if it evolves in a DM halo of density $\rho_{\chi}=2\cdot10^9\;$ GeV $^{-3}$ (assuming $\sigma_{\chi,SD}=10^{-38}\;$ $^2$, although other values of $\rho_{\chi}$ and $\sigma_{\chi,SD}$ can be considered, leading to the same effects as long as the product $\rho_{\chi}\sigma_{\chi}$ is kept constant)."
 This is a significant difference from the classical picture. im which a star as the Sun is expected to exhaust its lvdrogen core in less than LO Cr.," This is a significant difference from the classical picture, in which a star as the Sun is expected to exhaust its hydrogen core in less than 10 Gyr."
 As shown in earlier works (Salati&Silk.1989).. the more massive the star ds. the less it is affected by WIAIP aunibilation.," As shown in earlier works \citep{art-SalatiSilk1989}, the more massive the star is, the less it is affected by WIMP annihilation."
 Considering the same DM halo of the previous example. a star of 3 AL. won't be affected.," Considering the same DM halo of the previous example, a star of 3 $_{\odot}$ won't be affected."
 2.diagram:, 2.:
 Since DM ming accounts for at least one third of the total enerev. the balance will be reached with a larger radius and a lower effective temperature than m the classical vnicture (Fairbairnetal.2008).," Since DM burning accounts for at least one third of the total energy, the balance will be reached with a larger radius and a lower effective temperature than in the classical picture \citep{art-Fairbairnetal2008}."
. Therefore. stars that evolve iu dense DM halos follow ποτ] differcut paths in the U-R diaeraimi.," Therefore, stars that evolve in dense DM halos follow slightly different paths in the H-R diagram."
 We found that. in addition to he different paths followed during the MS. which was already reported in previous works (Casanellas&Lopes 2009).. stars follow brighter tracks during the red eiaut xanch (ROB).," We found that, in addition to the different paths followed during the MS, which was already reported in previous works \citep{art-CasanellasLopes2009ApJ}, stars follow brighter tracks during the red giant branch (RGB)."
 This feature is illustrated in Figure 1.., This feature is illustrated in Figure \ref{fig-diffHRm1}.
 Even if the difference iu the paths is remarkable. its effect ou the cluster is 11all compared with the slowing of the evolutionary speed.," Even if the difference in the paths is remarkable, its effect on the cluster is small compared with the slowing of the evolutionary speed."
 3.states:, 3.:
 For extremely high DM densities. stars are powered only bv the enerey from DM annihilation.," For extremely high DM densities, stars are powered only by the energy from DM annihilation."
 Whether the star was formed in this enuvironnieut or arrived there a posteriori it will reach a state of equilibrimm in the Uvashi track. far from the MS where most stars are found (Casanellas&Lopes 2009)..," Whether the star was formed in this environment or arrived there a posteriori, it will reach a state of equilibrium in the Hyashi track, far from the MS where most stars are found \citep{art-CasanellasLopes2009ApJ}. ."
 Tn this case the star is fully convective aud remains in the sale position iu the II-R. diagram as long as there are DAL particles to be captured in the halo (an illustrative exaniple is shown in Figure 1))., In this case the star is fully convective and remains in the same position in the H-R diagram as long as there are DM particles to be captured in the halo (an illustrative example is shown in Figure \ref{fig-diffHRm1}) ).
"the redshift-space effect depends on f,,(<jgtz). and we rewrite Eq.(1)) as f == jestLf) LG where we have defined Om ae lem Sena I where ji(hr) is oMspherical Bessel function.","the redshift-space effect depends on $f_g(z)\sigma_{8m}(z)$, and we rewrite \ref{eq:P(k)b}) ) as ) = C_0(k) (z)+ ^2 ]^2 + where we have defined (z) ^2(k|z) 1cm, 0.5cm ^2(k|z) where $j_1(kr)$ is spherical Bessel function."
 Note that We have assumed linear bias for simplicity., Note that }(z)=b(z) We have assumed linear bias for simplicity.
 To study the expected impact of future galaxy redshift surveys. we use the Fisher matrix formalism.," To study the expected impact of future galaxy redshift surveys, we use the Fisher matrix formalism."
 In the limit where the length scale corresponding to the survey volume is much larger than the scale of any features in {1}. we can assume that the likelihood function for the band powers of a galaxy redshift survey is Gaussian (Feldman.Kaiser.&Peacock1994).," In the limit where the length scale corresponding to the survey volume is much larger than the scale of any features in $P(k)$, we can assume that the likelihood function for the band powers of a galaxy redshift survey is Gaussian \citep{FKP}."
. Then the Fisher matrix can be approximated as (Tegmark1997) where p; are the parameters to be estimated from data. and the derivatives are evaluated at parameter values of the fiducial model.," Then the Fisher matrix can be approximated as \citep{Tegmark97}
 }= where $p_i$ are the parameters to be estimated from data, and the derivatives are evaluated at parameter values of the fiducial model."
 The effective volume of the survey Pos where in the oe]second equation. the comoving number density n ds assumed to only depend on the redshift for simplicity.," The effective volume of the survey ) ^3 ]^2 ]^2, where in the second equation, the comoving number density $n$ is assumed to only depend on the redshift for simplicity."
 ote that the Fisher matrix. 75; is the inverse of the covariance matrix of the parameters p; if the p; are Gaussian distributed. Eq.(99) , Note that the Fisher matrix $F_{ij}$ is the inverse of the covariance matrix of the parameters $p_i$ if the $p_i$ are Gaussian distributed. \ref{eq:full Fisher}) )
"propagates the measurement error in 1nP(K) Gwhich is WOPorional to in""r(K)] L7, into measurement errors for the yarameters p;.", propagates the measurement error in $\ln P(\bfk)$ (which is proportional to $[V_{eff}(\bfk)]^{-1/2}$ ) into measurement errors for the parameters $p_i$.
 To minimize nonlinear effects. we restrict wavenumbers to he quasi-linear regime.," To minimize nonlinear effects, we restrict wavenumbers to the quasi-linear regime."
" We take A,,;,= 0. and ρω is given by requiring that the variance of matter fluctuations in a sphere of radius £2. 67(2)=0.25. forR=x /(2h,,,,)."," We take $k_{min}=0$ , and $k_{max}$ is given by requiring that the variance of matter fluctuations in a sphere of radius $R$, $\sigma^2(R)= 0.25$, for$R=\pi/(2k_{max})$ ."
 This gives husLIbMpe atz= 0.and luusc0:2 5Mpe. tats=1. wel within the quasi-linear regime.," This gives $k_{max} 
\simeq 0.1\,h$ $^{-1}$ at $z=0$, and $k_{max} \simeq 0.2\,h$ $^{-1}$ at $z=1$, well within the quasi-linear regime."
" In addition. we impose a uniform upper limit of Aue.0.2bMpe 5! (ie. A,=O2hMpe ta >I lk to ensure that we are only considering the conservative linear regime essentially unaffected by nonlinear effects."," In addition, we impose a uniform upper limit of $k_{max}\leq 0.2\,h$ $^{-1}$ (i.e. $k_{max}=0.2\,h$ $^{-1}$ at $z>1$ ), to ensure that we are only considering the conservative linear regime essentially unaffected by nonlinear effects."
" The observed galaxy power spectrum in a given redshift shel centered at redshift 2; can be described by a set of parameters. ΠΕ OQ. fubzi)mus(ue Oxy(tr) Boye eine ene ns hi] where wayXpfs=0) (matter density today). wh=Quh?xp(s0) (baryon density today). 1, is the power-law index of the primordial matter power spectrum. and { is the dimensionless Hubble constant."," The observed galaxy power spectrum in a given redshift shell centered at redshift $z_i$ can be described by a set of parameters, $H(z_i)$, $D_A(z_i)$, $f_g(z_i)\sigma_{8m}(z_i)$ , $\sigma_{8g}(z_i)$ , $P_{shot}^i$, $\omega_m$ , $\omega_b$, $n_s$, $h$, where $\omega_m=\Omega_m h^2 \propto \rho_m(z=0)$ (matter density today), $\omega_b=\Omega_b h^2 \propto \rho_b(z=0)$ (baryon density today), $n_s$ is the power-law index of the primordial matter power spectrum, and $h$ is the dimensionless Hubble constant."
" Note that 2,(&)xA""27(A). with the matter transfer function 7(&) only depending on w,,, and c, (Eisenstein&Hu1998)..7 if & were in units of L/Mpc. and if the dark energy dependence of 71a can be neglected."," Note that $P_m(k)\propto k^{n_s} T^2(k)$, with the matter transfer function $T(k)$ only depending on $\omega_m$ and $\omega_b$ \citep{EisenHu98}, if $k$ were in units of $1/$ Mpc, and if the dark energy dependence of $T(k)$ can be neglected."
" We marginalize over unnn) P1, in each redshift slice. and project (Htz;).Da(z;).-- Vosti).] Gus Whex06), Dau. P1] into a final set of parameters (Wang"," We marginalize over $\sigma_{8g}(z_i)$ , $P_{shot}^i$ in each redshift slice, and project $H(z_i)$, $D_A(z_i)$, $f_g(z_i)\sigma_{8m}(z_i)$ , $\omega_m$ , $\omega_b$ , $n_s$, $h$ into a final set of cosmological parameters \citep{Wang06}."
" We refer to this as the “full P(À) method. with growth information included"". in which the growth information is included assuming that general relativity is valid."," We refer to this as the “full $P(k)$ method, with growth information included”, in which the growth information is included assuming that general relativity is valid."
 For more conservative dark energy constraints. we do not assume general relativity. and marginalize over i)must(2;)] from each redshift slice (in sain to mshii.Pin bn and only project [Htz;).-. Di(zi). wine we s. into the final set of cosmological ," For more conservative dark energy constraints, we do not assume general relativity, and marginalize over $f_g(z_i)\sigma_{8m}(z_i)$ from each redshift slice (in addition to $\sigma_{8g}(z_i)$, $P_{shot}^i$ ), and only project $H(z_i)$, $D_A(z_i)$, $\omega_m$, $\omega_b$, $n_s$, $h$ into the final set of cosmological parameters."
"We refer to this as the ""full. (A) method. marginalized over growth ,U)information""."," We refer to this as the “full $P(k)$ method, marginalized over growth information”."
 The details of our implementation can be found in Wang(2006.2008a).," The details of our implementation can be found in \cite{Wang06,Wang08a}."
. For an ultra conservative approach. we can marginalize over the cosmological parameters that describe the shape of the power spectrum. Len wp. nau. At. and only project {Hiz;). cosmologicalD(2) or (Hir).MMDato). futsσκηση)} into the final set of ," For an ultra conservative approach, we can marginalize over the cosmological parameters that describe the shape of the power spectrum, $\omega_m$, $\omega_b$, $n_s$, $h$, and only project $H(z_i)$, $D_A(z_i)$ or $H(z_i)$, $D_A(z_i)$, $f_g(z_i)\sigma_{8m}(z_i)$ into the final set of cosmological parameters."
"We refer to this as the DP(k)-marginalized-over-shape"" method."," We refer to this as the $P(k)$ -marginalized-over-shape"" method."
" To change from one set of parameters to another. we use (Wang2006) where £),; is the Fisher matrix for a set of parameters p. and £18 the Fisher matrix for a set of equivalent parameters q."," To change from one set of parameters to another, we use \citep{Wang06}
 }= where $F_{\alpha \beta}$ is the Fisher matrix for a set of parameters $\bfp$, and $F_{ij}$ is the Fisher matrix for a set of equivalent parameters $\bfq$."
" Measurements of the growth rate. /,(2) and the BAO are correlated and need to be considered. simultaneously. (Ballinger.Peacock.&Heavens1996:SimpsonPeacock2010)."," Measurements of the growth rate $f_g(z)$ and the BAO are correlated and need to be considered simultaneously \citep{ballinger96,Simpson10}."
. Note that the BAO only approach from Seo&Eisenstein(2007) is similar to our “/?(A)-marginalize-over-shape™ approach. but we adopt a more een that includes correlations between LG)generalDaC) andthatprocedure[fofτισκι(i).," Note that the BAO only approach from \cite{SE07} is similar to our $P(k)$ -marginalize-over-shape” approach, but we adopt a more general procedure that includes correlations between $\{H(z_i),D_A(z_i)\}$ and $\{f_g(z_i)\sigma_{8m}(z_i)\}$."
 Similarly. our approach is more than of White. Song Percival (2009). who made predictions for RSD constraints in a way that does not take into account simultaneous BAO measurements.," Similarly, our approach is more general than that of White, Song Percival (2009), who made predictions for RSD constraints in a way that does not take into account simultaneous BAO measurements."
 We derive dark energy constraints with and without Planck priors., We derive dark energy constraints with and without Planck priors.
 The Planck priorsare included as discussed in Appendix B.., The Planck priorsare included as discussed in Appendix \ref{sec:Planck}. .
 We derivedark energy constraints with and without Planck priors. whose calculation is discussed in Appendix B. Given thatPlanck isalreadyoperating successfully. and the full Planck data will be available when aspace-based galaxy survey isconducted (estimatedto bearound 2017). results including Planck priors are the most interesting for cosmological constraints.," We derivedark energy constraints with and without Planck priors, whose calculation is discussed in Appendix B. Given thatPlanck isalreadyoperating successfully, and the full Planck data will be available when aspace-based galaxy survey isconducted (estimatedto bearound 2017), results including Planck priors are the most interesting for cosmological constraints."
 We have included results without Planck priors in order to show the level of the dependency on additional data. and to enable the reader to reproduce our results.," We have included results without Planck priors in order to show the level of the dependency on additional data, and to enable the reader to reproduce our results."
Are the Sculptor Ciroup cis really tiat different [roin tle Local Group dlIs?,Are the Sculptor Group dIs really that different from the Local Group dIs?
 The Local Group jas two galaxies with values of log (τρως) <9.cU aud the Sculptor (οἱ bas three galaxies with log (Tyas) Z 11.5. so. at least statistically these two distributions could be found to be different.," The Local Group has two galaxies with values of log $\tau_{gas}$ ) $\le$ 9, and the Sculptor Group has three galaxies with log $\tau_{gas}$ ) $\ge$ 11.5, so, at least statistically these two distributions could be found to be different."
 Using he Ixolmogorov-S1uiruov (hereafter Ix-8 test (specifically of Press et 11992) to determiue he significance of the dilfereuces betweeithe distributions iu Figure 3. we find that the probability hat the Sculptor Gvoup aud Local Cirotp galaxies are drawn (rom the same clistributions is.," Using the Kolmogorov-Smirnov (hereafter K-S) test (specifically of Press et 1992) to determine the significance of the differences between the distributions in Figure 3, we find that the probability that the Sculptor Group and Local Group galaxies are drawn from the same distributions is."
. This emphasizes the large degree of overap iu tlie two samples., This emphasizes the large degree of overlap in the two samples.
 Interestinely. tie. IWS test vields a srobability of ‘that the Seulpor Ciroup aid LSB samples are drawn [rom the same distribution.," Interestingly, the K-S test yields a probability of that the Sculptor Group and LSB samples are drawn from the same distribution."
 However. this reflects the fac ταat the W-S test has a ügher seusiivity to the meclian values of cistributious than to the spreads in clistributious (Press οtal. 11992): the standard deviaious of the two samples differ by a factor of two.," However, this reflects the fact that the K-S test has a higher sensitivity to the median values of distributions than to the spreads in distributions (Press et 1992); the standard deviations of the two samples differ by a factor of two."
 The F-test statistic Cron1 Press et 11992) indicates that the Sculptor Group auk Lue LSB dls have significantly «lifferent variances (siglUicanee of ))., The F-test statistic from Press et 1992) indicates that the Sculptor Group and the LSB dIs have significantly different variances (significance of ).
 All other Ix-S test eoiipa‘isous of the samples slOWLL lu Figure 3 vield probabilities of less thanLO%., All other K-S test comparisons of the samples shown in Figure 3 yield probabilities of less than.
. Note that since tl ecwrent star formation rates represent a suapshot in time. we have uo guarantee that these distriutlous might not have loosec very different just a few 10 years ago.," Note that since the current star formation rates represent a snapshot in time, we have no guarantee that these distributions might not have looked very different just a few $^8$ years ago."
 Another important caveat siould be added at this point: missing from Figure 3) are gas-rich £ealaxies with current star [ornaion rates of zero (uo Ha emission)., Another important caveat should be added at this point; missing from Figure 3 are gas-rich galaxies with current star formation rates of zero (no $\alpha$ emission).
 The Local Ciroup has five of these dIrr/dSph or “transition [n])alaxies (Mateo 1998) aud the Scu—Or Group lias tliree., The Local Group has five of these dIrr/dSph or “transition” galaxies (Mateo 1998) and the Sculptor Group has three.
 Although uot plotted in Figure 3. the consideratiou of these galaxies would iicrease the degree of overlap iu the two samples.," Although not plotted in Figure 3, the consideration of these galaxies would increase the degree of overlap in the two samples."
 The transition galaxies are discussed. ii mo'e cletail i Sl., The transition galaxies are discussed in more detail in 4.
 There are differeut physical reasous for a galaxy to have a very arge value of τοις., There are different physical reasons for a galaxy to have a very large value of $\tau_{gas}$.
 For example. van Zee et ((1997) and Ixeunicutt. Skillman (2001) fourd values iu excess of 200 €vr for the LSB dI DDO 151.," For example, van Zee et (1997) and Kennicutt Skillman (2001) found values in excess of 200 Gyr for the LSB dI DDO 154."
 This appears to be mainly due to the very large HI |alo surrouncdius this galaxy which: contains of the galaxys HL., This appears to be mainly due to the very large HI halo surrounding this galaxy which contains of the galaxy's HI.
 Du this case. most o ‘this gas is probably unavailable for star formation.," In this case, most of this gas is probably unavailable for star formation."
 Iu the case of episodie star formation. it is possible ο observe a galaxy with a current SFR that is well below its average SFR.," In the case of episodic star formation, it is possible to observe a galaxy with a current SFR that is well below its average SFR."
 This resuts in an oarilicially large valle OL Tyas, This results in an artificially large value of $\tau_{gas}$.
 With detailed recent star formation histories like those availe from iuagiug studies of the young stellar population (e.g.. Dolim-Paliier et 11997. 1998) 1 is possible ο get a better 1jeastre of the recent SER.," With detailed recent star formation histories like those available from imaging studies of the young stellar population (e.g., Dohm-Palmer et 1997, 1998) it is possible to get a better measure of the recent SFR."
 Often the Ha and total luminosities of galaxies are usec --- 10rder to iivestigate how tle current SEHR compares to some measure of the past average SFR., Often the $\alpha$ and total luminosities of galaxies are used in order to investigate how the current SFR compares to some measure of the past average SFR.
" ThiW. ratio has been referred to as both the “star lormation timescale” (Roberts 1963. Hodge 1993) which is the ratio of the mass ο ‘the stars present to the current rate of star formation aud the ""birthrate yaraineter” (Scalo 1986. ]venulcutt et 11991) which is the ratio of the past average SER to tle curreu SFR."," This ratio has been referred to as both the “star formation timescale” (Roberts 1963, Hodge 1993) which is the ratio of the mass of the stars present to the current rate of star formation and the “birthrate parameter” (Scalo 1986, Kennicutt et 1994) which is the ratio of the past average SFR to the current SFR."
 Deriving the past average SFR requires estimating the total mass of the stars [ort1ος ove ‘the lifetime of the galaxy., Deriving the past average SFR requires estimating the total mass of the stars formed over the lifetime of the galaxy.
 This can be done by couverting the color of the stellar population iuto a representative mass-to-light ratio (e.e.. Miller 1996).," This can be done by converting the color of the stellar population into a representative mass-to-light ratio (e.g., Miller 1996)."
 Alternatively. Ixennicutt et ((1901 convert the Ha equivalent widtl directly into a birthrate parameter.," Alternatively, Kennicutt et (1994) convert the $\alpha$ equivalent width directly into a birthrate parameter."
 Both methods depeud ou asstuuptious of stellar aud galaxy evolution. aud making assumptious about galaxy evolution iu he studs of galaxy evolutiou may," Both methods depend on assumptions of stellar and galaxy evolution, and making assumptions about galaxy evolution in the study of galaxy evolution may"
The constructive. helpful comments of the referee. Ixatrien Ixolenberg. are much appreciated.,"The constructive, helpful comments of the referee, Katrien Kolenberg are much appreciated."
 We thank Dr.. Ixatalin Oláhh for placing her photographic observations at our disposal., We thank Dr. Katalin Oláhh for placing her photographic observations at our disposal.
 The financial support of OTI. grant. Ik-068626 is acknowledged., The financial support of OTKA grant K-068626 is acknowledged.
 €. Clement thanks the Natural Sciences and Engineering Research Council of Canada for financial support., C. Clement thanks the Natural Sciences and Engineering Research Council of Canada for financial support.
 Zs., Zs.
 Il. thanks the Lendüllet program of the Llungarian Academy. of Sciences for supporting his work., H. thanks the `Lendüllet' program of the Hungarian Academy of Sciences for supporting his work.
(P varies as a function of £4.,$\chi^2$ varies as a function of $t_{\rm on}$.
 The well-defined τήτπτ is found at the biu closest to the correct value (32.6 Myr)., The well-defined minimum is found at the bin closest to the correct value (32.6 Myr).
 The other parameter associated with the time. fg cau also be fonuud in the same way.," The other parameter associated with the time, $t_{\rm off}$ can also be found in the same way."
" We uote that our test case quasars had ""top hat” Πο curves. but that in a more realistic observational case there müeht be a smoother variation with time of the quasar huninositv."," We note that our test case quasars had “top hat” light curves, but that in a more realistic observational case there might be a smoother variation with time of the quasar luminosity."
 With good enough data this would show up as several detected light echoes nest to cach other in the data. but with differing huninosities.," With good enough data this would show up as several detected light echoes next to each other in the data, but with differing luminosities."
 Iu Figure 7 we vary the Iuuinosity parameter m our search and show the 47 values., In Figure \ref{lumsearch} we vary the luminosity parameter in our search and show the $\chi^2$ values.
 We fiud. that the X? is approximately constaut for a wide rauge of hDuuinosities well below the correct one (5<1010 cre/s. with a change starting at about 2 orders of magnitude below it.)," We find that the $\chi^2$ is approximately constant for a wide range of luminosities well below the correct one $5\times 10^{46}$ erg/s, with a change starting at about 2 orders of magnitude below it.)"
 The wid is found at a value roughly au order of magnitude sanaller than the actual value. indicating that it is possible hat the huuinuositv of the quasar is more difficult to find accurately than the position.," The minimum is found at a value roughly an order of magnitude smaller than the actual value, indicating that it is possible that the luminosity of the quasar is more difficult to find accurately than the position."
 The Iumiuositv in this case was very high. though: and so it is possible that the ight echo saturated. making it more difficult to flud the iunmositv.," The luminosity in this case was very high, though and so it is possible that the light echo saturated, making it more difficult to find the luminosity."
 A iore detailed study of this awaits future work., A more detailed study of this awaits future work.
 There is then a rie d X7 to a plateau., There is then a rise in $\chi^2$ to a plateau.
 This dlateam has a lower 47 value that the case of uo light echo. indicating that in this case a template with regious conipletelv devoid of absorption is a worse fit.," This plateau has a lower $\chi^2$ value that the case of no light echo, indicating that in this case a template with regions completely devoid of absorption is a worse fit."
 Tu this paper we have presented a techuique for searching for quasar light echoes in fforest spectra., In this paper we have presented a technique for searching for quasar light echoes in forest spectra.
 Our conclusions cau be suunnnaarized as ((1) We &ud that our technique which involves passing a template through a realistic simulated dataset viclds a ⋅⋅ 9 . . ⋅ ⋯∐∐∐⋯⋯∖−↖↽⋜↧↕⋯∖↕∪↥⋅↑↕∐∖⋜∏∏∐⋅∪↻↥⋅↕⋜↧↑↸∖↕↕∶↴∙⊾∐↑↸∖↸⊳∐∪↻⋜∐⋅⋜∐⊔↸↳ ((2) The statistical significance of a detection is strongly dependent on the hnuninositv of the quasar aud the time since it shut off., Our conclusions can be summarized as (1) We find that our technique which involves passing a template through a realistic simulated dataset yields a minimum $\chi^{2}$ value for the appropriate light echo (2) The statistical significance of a detection is strongly dependent on the luminosity of the quasar and the time since it shut off.
 Quasars with D-baud luminosity 73«LOY erg/s are necessaryto make detections at 20 significance or greater when the quasar switched off ~30 Myr previously or ((3) The nuuber of spectra in a data sample is also a relatively important factor iu the significance of detection. but their resolution is uot.," Quasars with B-band luminosity $> 3\times10^{45}$ erg/s are necessaryto make detections at $2 \sigma$ significance or greater when the quasar switched off $\sim 30$ Myr previously or (3) The number of spectra in a data sample is also a relatively important factor in the significance of detection, but their resolution is not."
 An optimal dataset should be densely sampled with mauv sieltlines per square degree but the signal to noise or resolution of spectra can be low., An optimal dataset should be densely sampled with many sightlines per square degree but the signal to noise or resolution of spectra can be low.
 We lave only simulated the search for quasars light echoes for quasars with a lifetime of ~10* years., We have only simulated the search for quasars light echoes for quasars with a lifetime of $\sim 10^{7}$ years.
 We leave it up to future work to fiud whether longer lived quasars can be detected iiore casily (we suspect that they cau. as heir signature in the forest will extend much further)," We leave it up to future work to find whether longer lived quasars can be detected more easily (we suspect that they can, as their signature in the forest will extend much further)."
 We have seen that a restframie B-band luminosity of 3100 eres | is necessary o detect light echoes with this lifetime at the 20 level. with the significance increasing if the quasar switched off nore recently than 30 My previously.," We have seen that a restframe B-band luminosity of $3\times10^{45}$ erg $^{-1}$ is necessary to detect light echoes with this lifetime at the $2 \sigma$ level, with the significance increasing if the quasar switched off more recently than 30 Myr previously."
 Using observational data for the quasar hDunuinositv function we cau compute or a eiven lifetime how many quasar Lelt echocs might je within a particular dataset., Using observational data for the quasar luminosity function we can compute for a given lifetime how many quasar light echoes might lie within a particular dataset.
" For example. Wopkins (2007) fiud from a compilation of recent quasar bpuniuositv function data that the space density of quasars at redshift 2=3 with this resttrame D baud huninositv (corresponding to an AB magnitude  2[8) is vl«10""Mpe (we use he0.7 in this calculation)."," For example, Hopkins (2007) find from a compilation of recent quasar luminosity function data that the space density of quasars at redshift $z=3$ with this restframe $B$ band luminosity (corresponding to an AB magnitude $\sim -24.8$ ) is $8.1 \times 10^{-7} {\rm Mpc}^{-3}$ (we use h=0.7 in this calculation)."
 We have estimated this value using the software made available by Wopkins which includes data from such samples as Wolf (2003) aud Richards (2006).," We have estimated this value using the software made available by Hopkins, which includes data from such samples as Wolf (2003) and Richards (2006)."
 If we imagine a survey with a sky area of 2 square degrees. such as the COSMOS survey (Scoville 2006). then the survey volume between 2=2.5 and 2=3.5 1s 2.35<10! Mpc.," If we imagine a survey with a sky area of 2 square degrees, such as the COSMOS survey (Scoville 2006), then the survey volume between $z=2.5$ and $z=3.5$ is $2.3\times10^{7} {\rm Mpc}^{3}$ ."
 Tf we ignore space density evolution over this redshift rauge. then we would expect," If we ignore space density evolution over this redshift range, then we would expect"
(oe. the nuclear burning is smaller than the accretion huninosity). the mass accretion rate should be higher than MacemLOS10SAL because the cfficiency of the radiation effect must be sinaller than which is cousistent with our estimation of M44;=1.5«10'M. + derived from the envelope mass at the optical Wain.,"(e.g., the nuclear burning is smaller than the accretion luminosity), the mass accretion rate should be higher than $\dot M_{\rm acc} \gtrsim 1.0 \times 10^{-7} M_\odot$ $^{-1}$ because the efficiency of the irradiation effect must be smaller than, which is consistent with our estimation of $\dot M_{\rm acc}= 1.5 \times 10^{-7} M_\odot$ $^{-1}$ derived from the envelope mass at the optical maximum."
 These systems with relatively high mass accretion rates are exactly the same as those proposed by Iachisu ct al. (, These systems with relatively high mass accretion rates are exactly the same as those proposed by Hachisu et al. (
1999b) as a progenitor svsteii of SNe Ia (sce also Li&vandenIHeuvel 1997)).,1999b) as a progenitor system of SNe Ia (see also \cite{lih97}) ).
 Usine the same simplified evolutional model as described iu IHachisu et al. (, Using the same simplified evolutional model as described in Hachisu et al. (
"1999b). we have followed binary evolutions for various pairs with the initial sets of (AL,;. Ado). αι Ίνοι, for the iutial primary niasses of AL;= 15.6. T. and OAL... the initial secoudary lasses of Mo;—LT. 3.0AL. bx AM»;=(LA. step. aud the initial separations of 4;=SO GOOR. bv Aloga;=0.01 step.","1999b), we have followed binary evolutions for various pairs with the initial sets of $M_{1,i}$, $M_{2,i}$, $a_i$ ), i.e., for the initial primary masses of $M_{1,i}=4$, 5, 6, 7, and $9 M_\odot$, the initial secondary masses of $M_{2,i}=1.7$ $3.0 M_\odot$ by $\Delta M_{2,i}=0.1 M_\odot$ step, and the initial separations of $a_i=80$ $600 R_\odot$ by $\Delta \log a_i= 0.01$ step."
 Starting from the initial set (FAL... 2.047... Lb0R.). for exanuploe. we have obtained a binary svstem of νου= OOM... Missy=9.241... aud Py=1.375 davs. after the binary underwent the first common cuvelope evolution aud then the primary naked helium star evolved to a helm eiut and had transterred helium to the secondary. MS.," Starting from the initial set $7 M_\odot$, $2.0 M_\odot$, $150 R_\odot$ ), for example, we have obtained a binary system of $M_{\rm WD,0}= 0.9 M_\odot$ , $M_{\rm MS,0}= 2.2 M_\odot$, and $P_0= 1.375$ days, after the binary underwent the first common envelope evolution and then the primary naked helium star evolved to a helium giant and had transferred helium to the secondary MS."
 Then. the secondary MS has slightly evolved to expaud and filled its Roche lobe.," Then, the secondary MS has slightly evolved to expand and filled its Roche lobe."
 Mass transfer begius from the MS to the WD., Mass transfer begins from the MS to the WD.
 We have further followed evolution of the binary uutil the binary reaches AAyp=L37aAl.. aud 0.7577 days at the same tine. that is. we regard the binary as V391 CrA when both the conditions. Mp=121141. aud P=03577 davs. are satisfied at the same time.," We have further followed evolution of the binary until the binary reaches $M_{\rm WD}= 1.37 M_\odot$ and $P= 0.7577$ days at the same time, that is, we regard the binary as V394 CrA when both the conditions, $M_{\rm WD}= 1.37 M_\odot$ and $P= 0.7577$ days, are satisfied at the same time."
 Then. we obtain the preseut state of V391 CrA having the secondary mnass of Mj;=1.3997. aud the mass transfer rate of AL=1.6.10TAL. L ," Then, we obtain the present state of V394 CrA having the secondary mass of $M_{\rm MS}= 1.39 M_\odot$ and the mass transfer rate of $\dot M_2= 1.6 \times 10^{-7} M_\odot$ $^{-1}$ ."
"In our evolutionary model. this binary svsten will soon explode as an SN Ia when the WD mass reaches Mj,=δέν, "," In our evolutionary model, this binary system will soon explode as an SN Ia when the WD mass reaches $M_{\rm Ia}= 1.378 M_\odot$."
The mass transfer rate of our evolutionary model is consistent with M15<10TAL. Ὁ estimated from the light curve fitting., The mass transfer rate of our evolutionary model is consistent with $\sim 1.5 \times 10^{-7} M_\odot$ $^{-1}$ estimated from the light curve fitting.
 Finally. we may conclude that V391 CrA is the secoud strong candidate for Type Ia progenitors. next to U Sco (ITachisu et al.," Finally, we may conclude that V394 CrA is the second strong candidate for Type Ia progenitors, next to U Sco (Hachisu et al."
 2000a. 20005).," 2000a, 2000b)."
 This research. has been supported iu part by the Crant- for Scientific Research (09610325. 11610226) of theJapanese Miuistrv of Education. Science. Culture. aud Sports.," This research has been supported in part by the Grant-in-Aid for Scientific Research (09640325, 11640226) of theJapanese Ministry of Education, Science, Culture, and Sports."
oreerounds are therefore of ereat importance. aud. indeed. direct attacks on the raw data are very rarely justified.,"foregrounds are therefore of great importance, and, indeed, direct attacks on the raw data are very rarely justified."
 Practically any analysis must consider sky maps which have been processed in some way. either by explicit oreground corrections. or by introduce a sky cut.," Practically any analysis must consider sky maps which have been processed in some way, either by explicit foreground corrections, or by introducing a sky cut."
 The importance of foreground removal has been recognized iu the community for a long time BBauday Wolteudale 1991: Reacdhead Lawreuce 1992: Braudt ot 11991: Teemark Efstathiou 1996: Teeniark 1998: Teemark et 22000). as has he preferred) method for discriminating agaist such contanination. namely multi-frequency observations.," The importance of foreground removal has been recognized in the community for a long time Banday Wolfendale 1991; Readhead Lawrence 1992; Brandt et 1994; Tegmark Efstathiou 1996; Tegmark 1998; Tegmark et 2000), as has the preferred method for discriminating against such contamination, namely multi-frequency observations."
 While the CAMB itself contributes equally to all requencics (as nieasured in thermodvuamic temperature units) due to the black body nature of its spectrum. oreerounds are typically strongly frequeucy-depenudoeut.," While the CMB itself contributes equally to all frequencies (as measured in thermodynamic temperature units) due to the black body nature of its spectrum, foregrounds are typically strongly frequency-dependent."
 One may therefore distinguish between forcerounds aud eenumne CAB anisotropy by studyviug how the signal varies with frequency., One may therefore distinguish between foregrounds and genuine CMB anisotropy by studying how the signal varies with frequency.
 However. detailed subtraction of forcerounds lias ecuerally required one of two assuniptious to be made: either that all of the physical colmpoucuts of the foreground emission aud their spectral vchavior are known. or that accurate templates of all of he components are available and that the appropriate spectral behavior cau be determined by fitting the cluplates to the available iiulti-frequeucy data.," However, detailed subtraction of foregrounds has generally required one of two assumptions to be made: either that all of the physical components of the foreground emission and their spectral behavior are known, or that accurate templates of all of the components are available and that the appropriate spectral behavior can be determined by fitting the templates to the available multi-frequency data."
 Neither ucthod is casily adapted to accommodate real spatial varlatious in the spectral behavior of the foregrounds., Neither method is easily adapted to accommodate real spatial variations in the spectral behavior of the foregrounds.
 The project appears to lave systematic issues under control. whilst considerable effort has con expended on foreeround issues. and uncertainties nav still remain.," The project appears to have systematic issues under control, whilst considerable effort has been expended on foreground issues, and uncertainties may still remain."
 Receut cetectious of nou-Ciussiauitv (Cliaug et al., Recent detections of non-Gaussianity (Chiang et al.
 2003: Coles et al., 2003; Coles et al.
 2001: Eriksen oet al., 2004; Eriksen et al.
 20012: Naselskv et al., 2004a; Naselsky et al.
 2003: Park 2001: Vielva ct al., 2003; Park 2004; Vielva et al.
 2001). the continuing debate over the low amplitude of huge augular scale anisotropy (see e.g. Efstathiou 2001). aud a possible preferred direction or aligumnenut contained therein (de Oliveira-Costa et 232001: Exikseu et 220080) may vet be affected by improvements of our ability to remove non-cosmoloeical foregrounds.," 2004), the continuing debate over the low amplitude of large angular scale anisotropy (see e.g., Efstathiou 2004), and a possible preferred direction or alignment contained therein (de Oliveira-Costa et 2004; Eriksen et 2004b) may yet be affected by improvements of our ability to remove non-cosmological foregrounds."
 The satellite observes the sky at five frequencies (23. 33. [1. 61 and 91 QCGIIZ) and using at least in part this information the team have applied three different methods for removing. coustrainius or describing the foregrouuds (Bennett et 220035).," The satellite observes the sky at five frequencies (23, 33, 41, 61 and 94 GHz), and using at least in part this information the team have applied three different methods for removing, constraining or describing the foregrounds (Bennett et 2003b)."
 The first method is to produce a so-called. iuterual luear combination map (from now on denoted ILC). which asstuues nothing about the particular frequency dependencies or inorphologics of the foregrounds.," The first method is to produce a so-called internal linear combination map (from now on denoted ILC), which assumes nothing about the particular frequency dependencies or morphologies of the foregrounds."
 Tusteacd. a CAIB imap is reconstructed by co-adding the data at the five frequencics (ποπ: convolved to a Common aneular resolution of 1 degree) with a set of weights that nünimuizes the final variance of the map.," Instead, a CMB map is reconstructed by co-adding the data at the five frequencies (now convolved to a common angular resolution of 1 degree) with a set of weights that minimizes the final variance of the map."
 The details of the non-linear method adopted to derive these weights have not been described by the team., The details of the non-linear method adopted to derive these weights have not been described by the team.
 In order to acconunodate spectral variability of the foregrounds. the skv has been partitioned iuto 12 separate regions and the miniunun variance criterion applied to cach iu turn to determine the weights.," In order to accommodate spectral variability of the foregrounds, the sky has been partitioned into 12 separate regions and the minimum variance criterion applied to each in turn to determine the weights."
 Discoutinuitics between reeious have been minimized by πλουhing of the weights at the boundaries., Discontinuities between regions have been minimized by smoothing of the weights at the boundaries.
 The resultant CAIB map is visually satistactory but has complex noise properties. aud indeed the team warns against its use for cosinological analysis.," The resultant CMB map is visually satisfactory but has complex noise properties, and indeed the team warns against its use for cosmological analysis."
 Nevertheless. the map has beeu subjected to such studies iu the literature. and indeed he team do use the resultant map as an input o their second foreground removal techuique.," Nevertheless, the map has been subjected to such studies in the literature, and indeed the team do use the resultant map as an input to their second foreground removal technique."
 This involves the application of a Maxima Entropy Alethod. (MEMO) in order to construct a model of the oregrounds. component by component.," This involves the application of a Maximum Entropy Method (MEM) in order to construct a model of the foregrounds, component by component."
 The strength of this method in principle is its ability to recoustruct he svuchrotron. free-free and dust eimüssiou. and their detailed frequency dependence on a pixel-by-pixol basis.," The strength of this method in principle is its ability to reconstruct the synchrotron, free-free and dust emission and their detailed frequency dependence on a pixel-by-pixel basis."
 However. the initial stage of the analysis ust still utilize cluplates for these dominant foreground components. and also establishes priors on their spectral behavior wousing the data at the five frequencies after correction for a CAIB component as determined by he WILC inethod above.," However, the initial stage of the analysis must still utilize templates for these dominant foreground components, and also establishes priors on their spectral behavior by using the data at the five frequencies after correction for a CMB component as determined by the WILC method above."
" As we will see later. ILC-ike methods in general still allow some leakage of Oreground signal iuto the CMD reconstruction. aud whether this results im anv feedback is difficult. to determine,"," As we will see later, ILC-like methods in general still allow some leakage of foreground signal into the CMB reconstruction, and whether this results in any feedback is difficult to determine."
 Again as a consequence of complex noise xoperties. the resultant map has not been considered useful for cosmological purposes.," Again as a consequence of complex noise properties, the resultant map has not been considered useful for cosmological purposes."
 huistead. the eani has used the results to iuterpret the nature of theP oreground euission.," Instead, the team has used the results to interpret the nature of the foreground emission."
 Iu particular. they ideutifv a dust-correlated component in the lower frequency (23. 33 aud 11 GIIz) channels with a spectral iudex of j2.45 for a spectrum of the form 27.," In particular, they identify a dust-correlated component in the lower frequency (23, 33 and 41 GHz) channels with a spectral index of $\beta \sim -2.5$ for a spectrum of the form $\nu^{\beta}$."
 This is pluvsically interpreted as a flat spectrum svuchrotrou compoucut in reeijons of star formation near the Galactic plane. rather than to cluission from spinning dust. which had become the preferred solution to this anomalous. low frequency dust correlated cussion.," This is physically interpreted as a flat spectrum synchrotron component in regions of star formation near the Galactic plane, rather than to emission from spinning dust, which had become the preferred solution to this anomalous, low frequency dust correlated emission."
 The issue remains open. but recent reanalyses by Lagache (2003) auc Fiukbeiner (2001) &ud evidence for the latter interpretation.," The issue remains open, but recent reanalyses by Lagache (2003) and Finkbeiner (2004) find evidence for the latter interpretation."
 The origin of this controversy les simply with the fact that the component separation as implemented by is allowed oulv to produce a combined svuchrotron/spiuniug dust solution at cach frequency. with no attempt mace to separately discutaugle these two coupoucuts.," The origin of this controversy lies simply with the fact that the component separation as implemented by is allowed only to produce a combined synchrotron/spinning dust solution at each frequency, with no attempt made to separately disentangle these two components."
" The final method for foreground correction 19 perhaps the simplest of the που, aud it is also the preferred iethod for generating cleaned maps suitable for cosmological analysis."," The final method for foreground correction is perhaps the simplest of the three, and it is also the preferred method for generating cleaned maps suitable for cosmological analysis."
 Rather than inherently exploiting the frequency information contained in the data. one subtracts external templates of the various physical componucnts (ie. maps produced bv non-CNID observations made preferably at frequencies where/ a specific componucut dominates the cussion) with coupling cocticicuts deteriiuned by cross-correlation with the observed maps.," Rather than inherently exploiting the frequency information contained in the data, one subtracts external templates of the various physical components (i.e., maps produced by non-CMB observations made preferably at frequencies where a specific component dominates the emission) with coupling coefficients determined by cross-correlation with the observed maps."
 This avoids the noise aplication which occurs when oue co-adds the data alone. aud has the added benefit that the resultant maps have well-known noise properties. provided that the templates themselves do not contribute sieuificautlv to this.," This avoids the noise amplification which occurs when one co-adds the data alone, and has the added benefit that the resultant maps have well-known noise properties, provided that the templates themselves do not contribute significantly to this."
 Difiiculies associated with the method inchide uucertaiutiocs 1i the detailed morphologies of he templates as scaled to the waveleneths of mterest. aud the propagation of errors iu the coupling coefficients 1ito the error budges of specific analvses.," Difficulties associated with the method include uncertainties in the detailed morphologies of the templates as scaled to the wavelengths of interest, and the propagation of errors in the coupling coefficients into the error budgets of specific analyses."
 It should be roted that Teemark. de Oliveira-C'osta. Tailton (2003) have applied a generalization of he ILC inethod to he data.," It should be noted that Tegmark, de Oliveira-Costa, Hamilton (2003) have applied a generalization of the ILC method to the data."
 The basic idea is to allow the weielting of cach, The basic idea is to allow the weighting of each
seen at the upper right corner of the GIRAFFE CCD.,seen at the upper right corner of the GIRAFFE CCD.
 Therefore we reduced the data again using the Geneva pipeline (version 1.11)., Therefore we reduced the data again using the Geneva pipeline (version 1.11).
 To remove the bright spot from the data we obtained raw dark frames from the ESO for the time range covered by our data., To remove the bright spot from the data we obtained raw dark frames from the ESO for the time range covered by our data.
 First we created master bias frames by averaging the five bias frames that had been obtained for each observation., First we created master bias frames by averaging the five bias frames that had been obtained for each observation.
 Then the three dark frames observed on a given date were bias corrected and averaged with cosmic ray rejection., Then the three dark frames observed on a given date were bias corrected and averaged with cosmic ray rejection.
 To reduce the noise in the dark frames. we smoothed them with à 2x2 pixel box filter.," To reduce the noise in the dark frames, we smoothed them with a $\times$ 2 pixel box filter."
 We then used the three flat fields observed for each night to determine the positions of the spectra., We then used the three flat fields observed for each night to determine the positions of the spectra.
 We did not correct these data for dark current. as the bright spot showed no negative effect on the detection of the spectra.," We did not correct these data for dark current, as the bright spot showed no negative effect on the detection of the spectra."
 Afterwards we derived the full wavelength solution from the ThAr are frames observed for each date., Afterwards we derived the full wavelength solution from the ThAr arc frames observed for each date.
 Using this solution. we finally extracted and rebinned the science data.," Using this solution, we finally extracted and rebinned the science data."
 At this step we included the smoothed master dark frames., At this step we included the smoothed master dark frames.
 The wavelength calibration was adjusted using the simultaneously observed ThAr spectra., The wavelength calibration was adjusted using the simultaneously observed ThAr spectra.
 As the optimal extraction produced spectra of lower signal-to-noise than conventional averaging we used the average option to integrate the flux for each spectrum., As the optimal extraction produced spectra of lower signal-to-noise than conventional averaging we used the average option to integrate the flux for each spectrum.
 We also divided the spectra by extracted flat-field spectra to perform a first correction for the CCD sensitivity variations., We also divided the spectra by extracted flat-field spectra to perform a first correction for the CCD sensitivity variations.
 Unfortunately. no flux standard stars are observed by GIRAFFE in the MEDUSA mode. which would permit at least a relative flux calibration to be obtained.," Unfortunately, no flux standard stars are observed by GIRAFFE in the MEDUSA mode, which would permit at least a relative flux calibration to be obtained."
 For each exposure. we subtracted the median of the spectra from the sky fibers from the extracted spectra of the target stars.," For each exposure, we subtracted the median of the spectra from the sky fibers from the extracted spectra of the target stars."
 We corrected all spectra for barycentric motion., We corrected all spectra for barycentric motion.
 During this correction. we noted that we had made a sign error when performing this correction for the data published by Moehler et al. (2007)).," During this correction, we noted that we had made a sign error when performing this correction for the data published by Moehler et al. \cite{modr07}) )."
 This error led to a smearing of the line profiles in the averaged spectra., This error led to a smearing of the line profiles in the averaged spectra.
Tn the previous studies. blue E/SO galaxies were regarded as: 1) star-forming low-mass spheroids with inverse color eracdieuts (huctal.2001): 2) carly-type galaxies with a blue ceutral clump caused by prolonged sinele-collapse (Menauteauetal.2001) or caused bv the accretion of blue objects (Ehuegreenetal.2005a): 3) earb-tvpe ealaxies that suffered the secondary starburst (Ferrerasetal. 2005): or 14) ACON-host galaxies (Moeuauteauetal.2005).,"In the previous studies, blue E/S0 galaxies were regarded as: 1) star-forming low-mass spheroids with inverse color gradients \citep{im01}; 2) early-type galaxies with a blue central clump caused by prolonged single-collapse \citep{men01b} or caused by the accretion of blue objects \citep{elm05a}; 3) early-type galaxies that suffered the secondary starburst \citep{fer05}; ; or 4) AGN-host galaxies \citep{men05}."
. In this studs. we found the followines. focusing on the fine structures of the blue early-type galaxies.," In this study, we found the followings, focusing on the fine structures of the blue early-type galaxies."
 First. at least [754 (27/58) of the BECs have traces of tidal interactions (au elongated or off-centered core: au elougated or off-ceutered blue chuup: multiple blue chuups: tidally distorted outer structures: interacting companions: or a blue outer rie).," First, at least $47\%$ $27/58$ ) of the BEGs have traces of tidal interactions (an elongated or off-centered core; an elongated or off-centered blue clump; multiple blue clumps; tidally distorted outer structures; interacting companions; or a blue outer ring)."
 Secoud. at least 28% (16/58) of the DEC probably contain ACGNs in their ceuters (high N-vay hunuinositv or AGN spectra).," Second, at least $28\%$ $16/58$ ) of the BEGs probably contain AGNs in their centers (high X-ray luminosity or AGN spectra)."
 This result supports the idea that ACINs may be responsible for the blue color of some blue E/S0 galaxics (Menauteauetal.2005j., This result supports the idea that AGNs may be responsible for the blue color of some blue E/S0 galaxies \citep{men05}.
. It is interesting that lI (s/5s) of the Ώρας (CGNIIS30. CN5302. CNS292. (N21520. CN21667. CS23739. GS20018. and 6822653) show both the traces of tidal eveuts and the ACN features. because this partially supports the merecr origin of ACONs (Canalizo2001:Ἱνοσν]ον&Dopita2003:Sánchezetal. 2005).," It is interesting that $14\%$ $8/58$ ) of the BEGs (GN1430, GN5302, GN8292, GN24520, GN24667, GS3739, GS20048, and GS22653) show both the traces of tidal events and the AGN features, because this partially supports the merger origin of AGNs \citep{can01,kew03,san05}."
 Third. 19% (11/58) of the BECs have red inner rugs whose nature is still unclear.," Third, $19\%$ $11/58$ ) of the BEGs have red inner rings whose nature is still unclear."
 All of these red-ring BECs are aud 69( (31/16) of the BECs have red inner vines., All of these red-ring BEGs are and $69\%$ $11/16$ ) of the BEGs have red inner rings.
 This inclicates that the red imuer rings are possibly related to ACN activity or strong star-formation activity., This indicates that the red inner rings are possibly related to AGN activity or strong star-formation activity.
 However. it is not easy to explain how such rings form in detail.," However, it is not easy to explain how such rings form in detail."
" Fourth. as described in 8112. the BEGCs have sanaller sizes than the REGs ou the averaee,"," Fourth, as described in 4.2, the BEGs have smaller sizes than the REGs on the average."
" However. there are also several BECs with large SIZOS: (13/58) of the DEGs are larger than Ruz2-= kpc and two DECs (CNI1212 and (CS895) are even as large as A,cm10 kpc."," However, there are also several BEGs with large sizes: $22\%$ $13/58$ ) of the BEGs are larger than $R_{\rm hl,z}=2$ kpc and two BEGs (GN14242 and GS895) are even as large as $R_{\rm hl,z}\approx 10$ kpc."
 Considering such large sizes of those BECs. they are hardly “star-forming low-mass spheroids (huctal.2001) unless they are very unstable dynamically or them velocity dispersions are unusually siuall.," Considering such large sizes of those BEGs, they are hardly `star-forming low-mass spheroids \citep{im01}' ' unless they are very unstable dynamically or their velocity dispersions are unusually small."
 Almost a half of the BECs show evidence of idal events and at least a quarter of the BECGs wobably contain aun AGN in their ceuter., Almost a half of the BEGs show evidence of tidal events and at least a quarter of the BEGs probably contain an AGN in their center.
 A cast iu our sample. the role of extra-galactic disturbance appears to be more dominant than hat of a prolonged single collapse (Menauteatetal.2001b).," At least in our sample, the role of extra-galactic disturbance appears to be more dominant than that of a prolonged single collapse \citep{men01b}."
". Within the frame of the ""fwo- model’ (Ferrerasetal.2005).. the secondary ast” immav be interpreted as the result of a tida interaction or mereie with other objects."," Within the frame of the `two-burst model' \citep{fer05}, , the `secondary burst' may be interpreted as the result of a tidal interaction or merging with other objects."
 Recently. evidences are accunmlating. supporting a new scheme for early-type galaxy formation: the ealactic (Cowieetal.1996:Susa&Uineunura2000:etal.2005).," Recently, evidences are accumulating, supporting a new scheme for early-type galaxy formation: the galactic \citep{cow96,sus00,tre05}."
. According to the model. the more massive galaxies are. the older hey are.," According to the model, the more massive galaxies are, the older they are."
 Although this scenario started as a uodified version of the monolithic collapse model. it was sugeested more recently that the cau be also compatible with the hierarchical nereine scenario.," Although this scenario started as a modified version of the monolithic collapse model, it was suggested more recently that the can be also compatible with the hierarchical merging scenario."
 Duudyeal.(2005). found that he E/SO fraction decreases as redshift increases. while the stellar mass function for all galaxies shows little evolution.," \citet{bun05} found that the E/S0 fraction decreases as redshift increases, while the stellar mass function for all galaxies shows little evolution."
 At the same time. they also ound that the lowerauass limit of E/S0 galaxies is higher at earlier epoch. which indicates that ealaxy merging may be significantly related to the scheme.," At the same time, they also found that the lower-mass limit of E/S0 galaxies is higher at earlier epoch, which indicates that galaxy merging may be significantly related to the scheme."
 Faberetal.(2006). zupported lis idea with the evolution of the lhuuinositv πιοπο for blue and red galaxies to :~1 using he DEEP? (Davisetal.2003) aud the COMDO- (Wolfetal.2003) survey data., \citet{fab06} supported this idea with the evolution of the luminosity function for blue and red galaxies to $z\sim1$ using the DEEP2 \citep{dav03} and the COMBO-17 \citep{wol03} survey data.
" According to hem. the ""quenching nass from blue galaxies o red galaxies via gas-rich mergers. decreases as iue goes,"," According to them, the `quenching mass' from blue galaxies to red galaxies via gas-rich mergers, decreases as time goes."
 DeLuciaetal.(2006) showed that it is possible o reproduce the using he simulation based on the hierarchical mereiue scenario aud the ACDM model., \citet{del06} showed that it is possible to reproduce the using the simulation based on the hierarchical merging scenario and the $\Lambda$ CDM model.
 Tuterestinely. the BECs in our sample also show some relationship with the wodel in their size — redshift relation.as shown in Fie.," Interestingly, the BEGs in our sample also show some relationship with the model in their size – redshift relation,as shown in Fig."
 9 and Fie., 9 and Fig.
 10: the BECs at lower redshift are smaller than the BECs at higher redshift. while thesize," 10: the BEGs at lower redshift are smaller than the BEGs at higher redshift, while thesize"
fading with a rate α=(—0.167+0.015)96yr-! (Aller&Reynolds1985).,fading with a rate $\alpha=(-0.167\pm 0.015)\% \mathrm{yr}^{-1}$ \cite{Aller1985}.
". At higher frequency, above 104 GHz, the observation are also consistent with synchrotron emission with a power-law of spectral index —0.73 (Veron-Cetty&Woltjer1993)."," At higher frequency, above $^4$ GHz, the observation are also consistent with synchrotron emission with a power-law of spectral index $-0.73$ \cite{Veron1993}."
". The data from IRAS satellite (Marsden have been reanalyzed by (Strom&Greidanus1992) revealing a significant excess of emission over the low frequency synchrotron spectrum, well explained by a single dust component at a temperature T~46K, thus requiring a 0.02Mo dust mass."," The data from IRAS satellite \cite{Marsden1984} have been reanalyzed by \cite{Strom1992} revealing a significant excess of emission over the low frequency synchrotron spectrum, well explained by a single dust component at a temperature $\rm T \sim 46 \ K$, thus requiring a $\rm 0.02 \ M_\odot$ dust mass."
" Using MPifrR bolometer arrays at the IRAM 30 m telescope, (Bandieraetal.2002) gave the first evidence for a new component at milllimetre wavelengths."," Using MPifrR bolometer arrays at the IRAM 30 m telescope, \cite{Bandiera2002} gave the first evidence for a new component at milllimetre wavelengths."
 They have shown that this 1.3mm (230 GHz) excess flux cannot be interpreted as emission from a dust component whereas the data may be consistent with a low energy cutoff in the energy distribution of the emitting particles., They have shown that this $\rm 1.3 \ mm$ (230 GHz) excess flux cannot be interpreted as emission from a dust component whereas the data may be consistent with a low energy cutoff in the energy distribution of the emitting particles.
" However, at 847m (353 GHz) (Greenetal.2004) found a good agreement with the canonical radio synchrotron emission and hence no need for an extra component."," However, at $\mu$ m (353 GHz) \cite{Green2004} found a good agreement with the canonical radio synchrotron emission and hence no need for an extra component."
" As the Planck satellite mission will use the Crab nebula emission to perform polarisation cross-check in the range 30 to 353 GHz, a deep knowledge of the physical origin of this emission is needed."," As the Planck satellite mission will use the Crab nebula emission to perform polarisation cross-check in the range 30 to 353 GHz, a deep knowledge of the physical origin of this emission is needed."
" For this purpose, we use observations of the Crab Nebula by the WMAP satellite at 23, 33, 41, 61 and 94 GHz (Pageetal.2007) and by the Archeops balloon experiment at 148, 217, 345 and 545 GHz (Macías-Pérezetal.2007;Desertetal. 2008)."," For this purpose, we use observations of the Crab Nebula by the WMAP satellite at 23, 33, 41, 61 and 94 GHz \cite{Page2006} and by the Archeops balloon experiment at 143, 217, 345 and 545 GHz \cite{Macias-Perez2007,Desert2007}."
. These data in addition to already published data are used to study the Crab Nebula SED., These data in addition to already published data are used to study the Crab Nebula SED.
 The paper is organized as follows., The paper is organized as follows.
In Sect.,In Sect.
 2 we present the re-analysis of the Archeops and WMAP data., \ref{sec:archWMAP} we present the re-analysis of the Archeops and WMAP data.
 Sect., Sect.
 3 presents the SED of the Crab Nebula from 1 to 10° GHz and compares it to a model including the synchrotron and dust well known components., \ref{sec:EMspectrum} presents the SED of the Crab Nebula from 1 to $^6$ GHz and compares it to a model including the synchrotron and dust well known components.
 In Sect., In Sect.
 4 we perform a coherent analysis of the Crab SED over the full frequency range adding an extra component to the previous model to account for the possible millimetre excess., \ref{sec:modelEMspectrum} we perform a coherent analysis of the Crab SED over the full frequency range adding an extra component to the previous model to account for the possible millimetre excess.
 The implications of our results for the Planck polarisation calibration are discussed in Sect. 5.., The implications of our results for the Planck polarisation calibration are discussed in Sect. \ref{sec:planckpolarcalib}.
 Summary and conclusions are given in Sect. 6.., Summary and conclusions are given in Sect. \ref{sec:conclusion}. .
 In this paper we are mainly interested in the microwave and millimetre regime for which the best currently available data are the WMAP (Pageetal.2007) and Archeops (Desertetal.2008) data., In this paper we are mainly interested in the microwave and millimetre regime for which the best currently available data are the WMAP \cite{Page2006} and Archeops \cite{Desert2007} data.
 To avoid any bias on our analysis by a difference of treatment of these two data sets we have decided to re-estimate the Crab nebula unpolarised emission from 23 to 545 GHz using the publicly available WMAP (Pageetal.2007;Hinshawetal.2009) and Archeops (Macías-Pérez intensity maps.," To avoid any bias on our analysis by a difference of treatment of these two data sets we have decided to re-estimate the Crab nebula unpolarised emission from 23 to 545 GHz using the publicly available WMAP \cite{Page2006,Hinshaw2009}
 and Archeops \cite{Macias-Perez2007} intensity maps."
" Furthermore, to make the full data set as homogeneous as possible, the total intensity flux for each frequency was estimated using standard radial photometry techniques to mimic the flux integration performed for the high resolution data sets."," Furthermore, to make the full data set as homogeneous as possible, the total intensity flux for each frequency was estimated using standard radial photometry techniques to mimic the flux integration performed for the high resolution data sets."
" Notice that we integrate far from the radio extension of the Crab nebula, ~ 5 arcminutes radius, to recover the Crab nebula flux diluted over the entire beam pattern, Therefore, a key issue for this analysis is the subtraction of the background galactic emission which is mainly due to synchrotron and free-free emissions at low frequencies and dust at high frequencies."," Notice that we integrate far from the radio extension of the Crab nebula, $\sim$ 5 arcminutes radius, to recover the Crab nebula flux diluted over the entire beam pattern, Therefore, a key issue for this analysis is the subtraction of the background galactic emission which is mainly due to synchrotron and free–free emissions at low frequencies and dust at high frequencies."
 As the WMAP and Archeops maps are available on Healpix format we have develop our own radial photometrysoftware for this pixelization scheme., As the WMAP and Archeops maps are available on Healpix format we have develop our own radial photometrysoftware for this pixelization scheme.
 The results of this analysis are presented, The results of this analysis are presented
"Surveys with degree sized fields of view (FOV), covering total area of 1000 square degrees or larger, are likely to have enough exposures and stars to make accurate corrections lo PSF anisotropy.","Surveys with degree sized fields of view (FOV), covering total area of 1000 square degrees or larger, are likely to have enough exposures and stars to make accurate corrections to PSF anisotropy."
 Here we quantify the residual PSF systematics described above and discuss their possible impact on survey strategy., Here we quantify the residual PSF systematics described above and discuss their possible impact on survey strategy.
 The following parameters constitute our description of a lensing survey., The following parameters constitute our description of a lensing survey.
 The survey size is given by the number of poinungs as: Qs=pointOppo., The survey size is given by the number of pointings as: $\Omega_{\rm S} = N_{\rm point} \Omega_{\rm FOV}$.
" For PSF measurement NV, Includes only those stars that have well measured shapes, and which are robustly identified as stars."," For PSF measurement $N_*$ includes only those stars that have well measured shapes, and which are robustly identified as stars."
 Interloping small galaxies are an additional concern if one tries to push the stellar locus too close to the galaxy locus of the size-magnitude diagram., Interloping small galaxies are an additional concern if one tries to push the stellar locus too close to the galaxy locus of the size-magnitude diagram.
 The number of significant principal components used to, The number of significant principal components used to
subtraction. [latfiecld correction and. cosmic rav removal.,"subtraction, flatfield correction and cosmic ray removal."
 The fatfield: correction is performed by means of both a quartz-halogen lamp (taken with the Gemini. calibration unit. GCAL) and twilight spectra. which are. normalised ancl divided from all the spectra.," The flatfield correction is performed by means of both a quartz-halogen lamp (taken with the Gemini calibration unit, GCAL) and twilight spectra, which are normalised and divided from all the spectra."
 Εις process corrects for pixcl-to-pixel sensitivity. variations and for two-dimensional large-scale illumination patterns due to slit. vignetting., This process corrects for pixel-to-pixel sensitivity variations and for two-dimensional large-scale illumination patterns due to slit vignetting.
 Cosmic rays are identified. and. eliminated by interpolating each line of an image with a high order function., Cosmic rays are identified and eliminated by interpolating each line of an image with a high order function.
 The resicluals (cosmic rav hits) that are narrower than the expect instrumental line width are replaced with the fitted: value., The residuals (cosmic ray hits) that are narrower than the expect instrumental line width are replaced with the fitted value.
 The package GSWAVELENG is used. to. establish an accurate wavelength calibration TILfrom copper argon (Cur) are images taken just before and after each exposure., The package GSWAVELENGTH is used to establish an accurate wavelength calibration from $-$ argon (CuAr) arc images taken just before and after each exposure.
 Each spectrum is converted to a linear wavelength scale. using roughly 130 are lines fitted by 8rd-5th order. polynomials., Each spectrum is converted to a linear wavelength scale using roughly 130 arc lines fitted by 3rd-5th order polynomials.
 An accuracy of better than 0.1 is consistently achieved for the wavelength calibration solution along the entire spectral coverage., An accuracy of better than 0.1 is consistently achieved for the wavelength calibration solution along the entire spectral coverage.
 The spectra are also corrected. for. geometrical distortions along the spectral directions (S-distortion)., The spectra are also corrected for geometrical distortions along the spectral directions (S-distortion).
 The spectral resolution is derived. as the mean of the Gaussian JINILMSs measured for a number of unblended arc-lanip lines which are distributed: over the whole spectral range of a wavelength-calibrated spectrum., The spectral resolution is derived as the mean of the Gaussian FWHMs measured for a number of unblended arc-lamp lines which are distributed over the whole spectral range of a wavelength-calibrated spectrum.
 ‘To obtain reliable measurements of the kinematic parameters at. large galactocentric radii it ds critical to perform an accurate sky subtraction.," To obtain reliable measurements of the kinematic parameters at large galactocentric radii, it is critical to perform an accurate sky subtraction."
 Sky contamination becomes increasingly significant for the outer parts of the ealaxy., Sky contamination becomes increasingly significant for the outer parts of the galaxy.
 lo these regions. the level of galaxy. light. can be a few percent of the sky signal.," In these regions, the level of galaxy light can be a few percent of the sky signal."
 This effect represents one of the main sources of systematic errors in the analysis of the kinematic racial profiles., This effect represents one of the main sources of systematic errors in the analysis of the kinematic radial profiles.
 The generous field-of-view of GAIOS anc the intrinsic low Luminosity of our sample allow us to accurately estimate the sky continuum from the edges of the slit. where the galaxy light is negligible.," The generous field-of-view of GMOS and the intrinsic low luminosity of our sample allow us to accurately estimate the sky continuum from the edges of the slit, where the galaxy light is negligible."
 The sky level is estimated by interpolating along the outermost 10 20 aresee at the two edges of the slit ancl subtracted from the whole image., The sky level is estimated by interpolating along the outermost $-$ 20 arcsec at the two edges of the slit and subtracted from the whole image.
 A sky subtraction accurate to within one percent is achieved. guarantecing negligible: related systematics on the measured kinematic parameters.," A sky subtraction accurate to within one percent is achieved, guaranteeing negligible related systematics on the measured kinematic parameters."
 For each galaxy. we obtained four fully reduced galaxy frames of identical exposure time.," For each galaxy, we obtained four fully reduced galaxy frames of identical exposure time."
 These are then co-adcded to form a single frame., These are then co-added to form a single frame.
 Phe final galaxy [rane is a two-dimensional spectrum from which we extract 1D spectra along the slit., The final galaxy frame is a two-dimensional spectrum from which we extract 1D spectra along the slit.
" Phe spatial width (i.c. the number of CCD rows binned) for cach extracted 11) spectrum. increases with racius to achieve a minimum signal-to-noise (S/N) per of ~ 35 in the spectral region of Ale,.", The spatial width (i.e. the number of CCD rows binned) for each extracted 1D spectrum increases with radius to achieve a minimum signal-to-noise (S/N) per of $\sim$ 35 in the spectral region of $_{b}$.
 By adopting this criterion we have been able to obtain reliable measurements of kinematic parameters to galactocentric radii of  37., By adopting this criterion we have been able to obtain reliable measurements of kinematic parameters to galactocentric radii of $\sim$ $r_{e}$.
 The imaging observations were performed. with the Advanced Camera for Surveys (ACS) mounted. on the Llubble Space Telescope. (LIST)., The imaging observations were performed with the Advanced Camera for Surveys (ACS) mounted on the Hubble Space Telescope (HST).
 The cight Virgo cluster ealaxies were observed by the ACS Virgo Cluster Survey (ACSVCS: Cotéctal. 2004)). while the six galaxies in the Fornax cluster were observed. by the ACS Fornax Cluster Survey CACSECS: Jordánctal. 2007)).," The eight Virgo cluster galaxies were observed by the ACS Virgo Cluster Survey (ACSVCS; \citealt{cote04}) ), while the six galaxies in the Fornax cluster were observed by the ACS Fornax Cluster Survey (ACSFCS; \citealt{jordan07}) )."
 Each ealaxy was observed: using the ACS Wide Field Channel AVEC) mode (Fordοἱal.1998: Siriannietal. 2005))., Each galaxy was observed using the ACS Wide Field Channel (WFC) mode \citealt{ford98}; \citealt{sirianni05}) ).
" The WEC consists of two 4096 CCD detectors (WECT and WEC2) witha pixel size of 155 and a spatial scale of 0.049"" pixel+.", The WFC consists of two $\times$ 4096 CCD detectors (WFC1 and WFC2) with a pixel size of $\mu$ m and a spatial scale of $\arcsec$ $^{-1}$.
 The nucleus of cach galaxy was placed on the WECT detector at the reference. pixel. (2096. 200).," The nucleus of each galaxy was placed on the WFC1 detector at the reference pixel (2096, 200)."
 Fo avoid the gap between the WECI and WEC?2 detectors an additional οκο 5 lo, To avoid the gap between the WFC1 and WFC2 detectors an additional offset $\arcsec$ to
The eruption of the two OAs on the disk. namely. Events 15 and 3)) and 16 and 9)). were observed on 2006 May 1 ancl 2006 July 4. respectively.,"The eruption of the two OAs on the disk, namely, Events 15 and ) and 16 and ), were observed on 2006 May 1 and 2006 July 4, respectively."
 With the solar minimum impending. both events are much less energetic Chan the earlier ones.," With the solar minimum impending, both events are much less energetic than the earlier ones."
 Putting side bv side. they exhibit equite similar characteristics.," Putting side by side, they exhibit quite similar characteristics."
 Both arcades were overlving a cecavecl active region. namely. NOAA 10875 G((d)) ancl NOAA 10898 7((d)). respectively.," Both arcades were overlying a decayed active region, namely, NOAA 10875 (d)) and NOAA 10898 (d)), respectively."
 In both events. the gradual inflation of the OA is sustained for about 4.5 hrs. at a speed of <5I... and the subsequent eruption only resulted in a C-class flare S((d) and 9((0)).," In both events, the gradual inflation of the OA is sustained for about 4.5 hrs, at a speed of $<$, and the subsequent eruption only resulted in a C-class flare (d) and (d))."
 Both flares were associated with slow CAIEs. although in the former event a halo CME was produced 6((h)).," Both flares were associated with slow CMEs, although in the former event a halo CME was produced (h))."
 Like the PEA events. the morphology of the resultant CMEs bear similarities to the inflating arcacdes.," Like the PEA events, the morphology of the resultant CMEs bear similarities to the inflating arcades."
 On the other haad. in the 2006 May 1 event. (he arcade can be seen back to as early as 2006 April 28. 4 clays prior to its eruption. while in the 2006 July 4 event. the overlving loops of interest only became illuminated at 14:48 UT. and their growth and subsequent eruption were observed henceforth.," On the other hand, in the 2006 May 1 event, the arcade can be seen back to as early as 2006 April 28, 4 days prior to its eruption, while in the 2006 July 4 event, the overlying loops of interest only became illuminated at 14:48 UT, and their growth and subsequent eruption were observed henceforth."
 The loops on 2006 July 4 were located hieh in the corona from the beginning: the projected half length of the highest loop is about2., The loops on 2006 July 4 were located high in the corona from the beginning: the projected half length of the highest loop is about.
.. One max wonder how these bipolar. potential-like loops became eruplive and resulted in CMEs.," One may wonder how these bipolar, potential-like loops became eruptive and resulted in CMEs."
 κο X-ray images show hiehly complex loops underlving the inflating arcade in both active regions 6((eο) and 7((e&)). but there is no sign of twisted or sheared fields. such as the well-known soft X-ray sigmoids.," soft X-ray images show highly complex loops underlying the inflating arcade in both active regions (e–g) and (e–g)), but there is no sign of twisted or sheared fields, such as the well-known soft X-ray sigmoids."
 The quasi-static stage in both evenis were temporallv associated with helicity injection. except that the amount of helicity accumulation flattened after the flare in the 2006 May. 1 event S((0)). while the helicity injection rate displaved no obvious change throughout the flare in the 2006 July 4 event 9((c)).," The quasi-static stage in both events were temporally associated with helicity injection, except that the amount of helicity accumulation flattened after the flare in the 2006 May 1 event (c)), while the helicity injection rate displayed no obvious change throughout the flare in the 2006 July 4 event (c))."
 Zhangetal.(2001) showed a three-phase kinematic evolution lor three of four CALEs well observed by LASCO Cl. C2 and C3 coronagraphs from 1.1 toRz: the initiation phase. impulsive acceleration phase. ancl propagation phase.," \citet{zhang01} showed a three-phase kinematic evolution for three of four CMEs well observed by LASCO C1, C2 and C3 coronagraphs from 1.1 to: the initiation phase, impulsive acceleration phase, and propagation phase."
 The initiation phase is characterized by a slow ascension (< 1)) over tens of minutes. which always occurs before the onset of (he associated flare.," The initiation phase is characterized by a slow ascension $<$ ) over tens of minutes, which always occurs before the onset of the associated flare."
 The impulsive acceleration phase with a rapid acceleration of 100.500 nis 7 coincides very well with the flares's rise phase lasting [or a [ew to tens of minutes., The impulsive acceleration phase with a rapid acceleration of 100–500 m ${}^{-2}$ coincides very well with the flares's rise phase lasting for a few to tens of minutes.
 The, The
There are 131 MSX sources imnatehed with 2MASS extended sources across the LAIC.,There are 131 MSX sources matched with 2MASS extended sources across the LMC.
 Although: 241ASS detects extended sources over the entire 100. square-degree field. the central region of the galaxy will lave the preponderauce of matched extended sources. since most of the MSX structure is concentrated there.," Although 2MASS detects extended sources over the entire 100 square-degree field, the central region of the galaxy will have the preponderance of matched extended sources, since most of the MSX structure is concentrated there."
 matched sources are ecucrally faint., The matched sources are generally faint.
 The majority ( appear nebular. while a large uunuber (~31%)) appear point-like. and many of these are fairly red.," The majority $\sim$ ) appear nebular, while a large number $\sim$ ) appear point-like, and many of these are fairly red."
 About of the exteuded sources are galaxies behind the LMC., About of the extended sources are galaxies behind the LMC.
 Most of the extended sources fall iu the same color regine as the candidate PNe and. HII regious. aud a few are near the area on the diagram for the ACB stars.," Most of the extended sources fall in the same color regime as the candidate PNe and HII regions, and a few are near the area on the diagram for the AGB stars."
" The AISX footprint is large cuough that may sources iu the LMC that are extended. as seen by 2MÁASS. would be considered ""poiut-like aud would fall iuto the MSX PSC."," The MSX footprint is large enough that many sources in the LMC that are extended, as seen by 2MASS, would be considered “point-like” and would fall into the MSX PSC."
 All of the matched extended sources are associated with matched MSX/2MASS point sources., All of the matched extended sources are associated with matched MSX/2MASS point sources.
" This occurs. since many 2MASS extended sources contain associated poiut sources, such as ionizing stars within the nebulae. in the case of the candidate ΠΠ reeious or reflection nebulac. or central stars within candidate PNe."," This occurs, since many 2MASS extended sources contain associated point sources, such as ionizing stars within the nebulae, in the case of the candidate HII regions or reflection nebulae, or central stars within candidate PNe."
 Also. a few of the candidate ACD stars must show associated extended enission. e.g.. eircunistellar nebulae. as seeu by. 2MASS. which. together. are also bright iu the MSN A baud.," Also, a few of the candidate AGB stars must show associated extended emission, e.g., circumstellar nebulae, as seen by 2MASS, which, together, are also bright in the MSX A band."
 This work preseuts a classification scheme for resolved objects in external galaxies where uear-IR aud iuid-IR colors are available., This work presents a classification scheme for resolved objects in external galaxies where near-IR and mid-IR colors are available.
 In addition to future laree erounud-based IR observing programs. the next opportunity for large-scale application of this work would be to uudertake a survev ofthe LMC with the IR Array Camera (IRAC) iustrmment.," In addition to future large ground-based IR observing programs, the next opportunity for large-scale application of this work would be to undertake a survey of the LMC with the IR Array Camera (IRAC) instrument."
 Like MSN. TRAC has au ὃ jaa capability.," Like MSX, IRAC has an 8 $\mu $ m capability."
 Its projected scusitivity is approxiniatelv magnitude 12 iu this baud in the moderate iuteeration mode., Its projected sensitivity is approximately magnitude 12 in this band in the moderate integration mode.
 Combining an IRAC LAIC survey with the 2\[ASS point source catalog would vield a uch more comple5 population study of the LMC., Combining an IRAC LMC survey with the 2MASS point source catalog would yield a much more complete population study of the LMC.
 The linits of such a survey are shown on the maguitude-color diagrams iu Figures 5 ancl 7., The limits of such a survey are shown on the magnitude-color diagrams in Figures 5 and 7.
 The maenitude=7.5 Πιτ of the MSN A-baud survey was too high to detect the majority of ACB stars with low-to-mmoderate mass loss., The magnitude=7.5 limit of the MSX A-band survey was too high to detect the majority of AGB stars with low-to-moderate mass loss.
 We see from Figure 7 that a /2\TASS survev will detect the entire LAIC ACD population. as well as the fainter eud of the PN population. assuniue the Wainscoat et al. (," We see from Figure 7 that a /2MASS survey will detect the entire LMC AGB population, as well as the fainter end of the PN population, assuming the Wainscoat et al. ("
1992) absolute magnitudes are similar to what tle same stellar types would have under the metallicity conditions in the LAIC.,1992) absolute magnitudes are similar to what the same stellar types would have under the metallicity conditions in the LMC.
 Based on the IRAC performance provided by. Hora. Fazio. Willner (2000). an IRAC survey would detect Wsyou}=15.1 sources at —5 with the 30-s integration mode (A. Omont €. SNRCiülnore. priv.," Based on the IRAC performance provided by Hora, Fazio, Willner (2000), an IRAC survey would detect $m_{[8\ \mu {\rm m}]}=15.1$ sources at SNR=5 with the 30-s integration mode (A. Omont G. Gilmore, priv."
 comu.)., comm.).
 From the MSX 8.3 gan source couuts in Figure 1. coufusiou is estimated to set m at imsμμo16.," From the MSX 8.3 $\mu $ m source counts in Figure 1, confusion is estimated to set in at $m_{[8\ \mu {\rm m}]}\sim 16$."
 Indeed. confusion 15 relatively minimal iu the deusest regions of the LMC at a A. magnitude limit of 11.7.," Indeed, confusion is relatively minimal in the densest regions of the LMC at a $K_{s}$ magnitude limit of 14.7."
" For 221MÀSS. SNR&5 occurs at AN,=15.1."," For 2MASS, $\simeq $ 5 occurs at $K_{s}=15.1$."
 Thus. an IRAC [5 inu LAIC survey is well-natched to the 2\LASS survey. of the LMC.," Thus, an IRAC [8 $\mu $ m] LMC survey is well-matched to the 2MASS survey of the LMC."
 Aeain. from Figure 1. more than a million LMC sources would be matched with J. ΠΠ. Ay. aud the IRAC colors.," Again, from Figure 1, more than a million LMC sources would be matched with $J$, $H$, $K_{s}$, and the IRAC colors."
" 2\LTASS also conducted a ""GN survey (six thes the normal survey integration fine per scan} of the LMC. increasing the survey limit by about 11.5 magnitudes."," 2MASS also conducted a “6X” survey (six times the normal survey integration time per scan) of the LMC, increasing the survey limit by about 1–1.5 magnitudes."
 We recomuencl such a survey be doue as oue of the major Guest Observer projects., We recommend such a survey be done as one of the major Guest Observer projects.
 MSN iud 2MÁSS have both observed the LMC as part of their operations in the mid- aud ucar-IR. respectively.," MSX and 2MASS have both observed the LMC as part of their operations in the mid- and near-IR, respectively."
 We lave performed a cross-correlation of the source catalogs from the two IR surveys aud lave found 1661 »oiut source and 131 extended source matches., We have performed a cross-correlation of the source catalogs from the two IR surveys and have found 1664 point source and 131 extended source matches.
 It is clear hat the combination of ucar-IR and mid-IR colors cau iff some of the degeneracy in the IR color-color diagram or sources in the LAIC., It is clear that the combination of near-IR and mid-IR colors can lift some of the degeneracy in the IR color-color diagram for sources in the LMC.
 For instance. in the 2ALTASS-ouly color plauc. HII regious and PNe are mdistiuguishable Yon normal stars.," For instance, in the 2MASS-only color plane, HII regions and PNe are indistinguishable from normal stars."
 With the MSN A-baud as another axis o the color plane. such nebulae are niucli easier to ideutifv and distinguish.," With the MSX A-band as another axis to the color plane, such nebulae are much easier to identify and distinguish."
 Based ou the combined 2MIASS/MSX. colors. using he IR point source model by Waiuscoat et al. (," Based on the combined 2MASS/MSX colors, using the IR point source model by Wainscoat et al. ("
1992). and obtaining source names and spectral types from the SIMIBAD database. when available. we have ideutified 11 categories ofstellar populations aud red nebulae. iucliding nun sequence stars. giant stars. RSCGs C- aud O-rich ACD stars. PNe. IE IL regions. aud other dusty objects likely associated with carly-type stars.,"1992), and obtaining source names and spectral types from the SIMBAD database, when available, we have identified 11 categories of stellar populations and red nebulae, including main sequence stars, giant stars, RSGs, C- and O-rich AGB stars, PNe, H II regions, and other dusty objects likely associated with early-type stars."
 731 of these sources were previously unidentified., 731 of these sources were previously unidentified.
 Object detection Ην. the MSN 8.3 jan (band A) seusitivity., Object detection is limited by the MSX 8.3 $\mu$ m (band A) sensitivity.
 The spatial distribution of these IR sources can also add to our uuderstaudius of star formation aud stellar evolution iu the LMC. and of the evolutiou of the LMC as a ΠΟΙΟ of an interacting svstem with the Milkv. Wav. as well as a template for low-metallicity svstenis at higher redshift.," The spatial distribution of these IR sources can also add to our understanding of star formation and stellar evolution in the LMC, and of the evolution of the LMC as a member of an interacting system with the Milky Way, as well as a template for low-metallicity systems at higher redshift."
 Our compilation of MSN aud 2MASS photometry and astrometry for these sources will provide a sample of objects for future observations aud analysis by ground-based aud space-haseck observatories. such as aud SOFIIA.," Our compilation of MSX and 2MASS photometry and astrometry for these sources will provide a sample of objects for future observations and analysis by ground-based and space-based observatories, such as and IA."
 This research las made use of the SIMDAD database. operated at CDS. Strasboure. France. aud of NASA's Astrophysics Data Systei Abstract Service.," This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France, and of NASA's Astrophysics Data System Abstract Service."
 This publication mace use of data products from the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts aud the Dufrared Processing aud Analysis Ceuter/California Tustitute of Technology. funded by the National Aeronautics auc Space Achuinistration and the National Science Foundation.," This publication made use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
 We thank Martiu Coben for determining maguitudes in MSX bands for the 87. Wainscoat object types. and also thiuik Sean Carex. Don Mizuno aud Tom I&uchar. who assembled the MSX lhnage data.," We thank Martin Cohen for determining magnitudes in MSX bands for the 87 Wainscoat object types, and also thank Sean Carey, Don Mizuno and Tom Kuchar, who assembled the MSX Image data."
This relation could in principle be applied directly to the motion of the Local Group of galaxies (rough the Universe.,This relation could in principle be applied directly to the motion of the Local Group of galaxies through the Universe.
 Consequently. comparison of the peculiar velocity and acceleration of the LG serves as a lool to estimate (he 7 parameter.," Consequently, comparison of the peculiar velocity and acceleration of the LG serves as a tool to estimate the $\beta$ parameter."
" Independent. knowledge of biasing allows to estimate cosmological censitv O,,.", Independent knowledge of biasing allows to estimate cosmological density $\Omm$.
 In realitv. however. we do not observe continuous galaxy density field. but rather discrete objects. even if in a verv large number.," In reality, however, we do not observe continuous galaxy density field, but rather discrete objects, even if in a very large number."
" In the following derivation. which can be found iin ?.. we model galaxies as point sources: p,(r)=$5;Miop(r—r;). where dp is Dirac's delta: M; and m; are respectivelv (he mass and (he position of the -th galaxy."," In the following derivation, which can be found in \cite{ViSt}, we model galaxies as point sources: $\rho_\mrg(\bmr) = \sum_i M_i\, \delta_D(\bmr - \bmr_i)$, where $\delta_D$ is Dirac's delta; $M_i$ and $\bmr_i$ are respectively the mass and the position of the $i$ -th galaxy."
 Putting the coordinate svstem αἱ r=0 (harveenter of the LG) we obtain the acceleration of the LG as a sum of the force contributions from all sources in the Universe: This Newtonian is still not useful for caleulations based on observational data. as masses of individual galaxies are usually known very. poorly. if ever.," Putting the coordinate system at $\bmr=0$ (barycenter of the LG) we obtain the acceleration of the LG as a sum of the force contributions from all sources in the Universe: This Newtonian is still not useful for calculations based on observational data, as masses of individual galaxies are usually known very poorly, if ever."
 However. if the -th galaxy has an intrinsic luminosity L;. we can write where 5;=απ is the flux received from the /-th object.," However, if the $i$ -thgalaxy has an intrinsic luminosity $L_i$ , we can write where $S_i=L_i\slash 4\pi r_i^2$ is the flux received from the $i$ -th object."
 This relation means that if we know the behavior of the mass-to-light ratio in the band(s) of the survey. we can estimate the acceleration of the LG Irom a catalog. oone containing astro- and photometric data only (positions and fluxes).," This relation means that if we know the behavior of the mass-to-light ratio in the band(s) of the survey, we can estimate the acceleration of the LG from a catalog, one containing astro- and photometric data only (positions and fluxes)."
 Furthermore. if the mean mass-to-light ratio is a universal constant. GU/L).. welinallv," Furthermore, if the mean mass-to-light ratio is a universal constant, , wefinally"
"Using a pure luminosity evolution model, we obtain an evolution law given by L.(24um)ος(1+z)*5593 for the star-forming sample.","Using a pure luminosity evolution model, we obtain an evolution law given by $L_*(24~\micron) \propto (1+z)^{3.8 \pm 0.3}$ for the star-forming sample."
 This luminosity evolution is shown as a fit to the 1/Vmax LF in Figure 7.., This luminosity evolution is shown as a fit to the $V_{\rm max}$ LF in Figure \ref{LF_allbins}.
 One of the evolution descriptions of uses pure luminosity evolution with the Schechter form of LF., One of the evolution descriptions of uses pure luminosity evolution with the \citet{Schechter76} form of LF.
" They report luminosity evolution of (1+ for z—0.8, agreeing with our result."," They report luminosity evolution of $(1+z)^{3.6 \pm 0.3}$ for $z < 0.8$, agreeing with our result."
 The evolution of 2)*9503the LF was also quantified by determining the luminosity above which the cumulative number, The evolution of the LF was also quantified by determining the luminosity above which the cumulative number
effectively rules out these or aiy other stochastic processes.,effectively rules out these or any other stochastic processes.
 Coopereal.(2003) have suggested that the uniform red c‘olors of the cold classical IKBOs are caused by their existence iu au irradiation ninimiun enviroument. with iradiation increasing interior due to 1icreasiug fluxes [rom solar euergetic ious aud inereasing exerior by increase [lix of energetic ions from the termination shock. but this hypothesis lias no general explanation [or he remainiug colos. LOY wod it lead to the sizable observed poptlation neitrally colo'ed higher inclination objects in the sane envi‘OneLI.," \citet{2003EM&P...92..261C} have suggested that the uniform red colors of the cold classical KBOs are caused by their existence in an irradiation minimum environment, with irradiation increasing interior due to increasing fluxes from solar energetic ions and increasing exterior by increased flux of energetic ions from the termination shock, but this hypothesis has no general explanation for the remaining colors, nor would it lead to the sizable observed population neutrally colored higher inclination objects in the same environment."
 A new hvpothliesIs has suggested.OO tha the coloκ of IXBOs are se by early evaporation aud irraciation of volatiles followed by dyuamical mixi& (Brown.Schaller&Fraser2011a)., A new hypothesis has suggested that the colors of KBOs are set by early evaporation and irradiation of volatiles followed by dynamical mixing \citep{2011ApJ...739L..60B}.
. Ln this ypothesis. objects a‘e formecl outside o. Neptune in the ~15-30 AU reeion in the early solar W.vstem out of a diverse aucl variable mix «of inaterias.," In this hypothesis, objects are formed outside of Neptune in the $\sim$ 15-30 AU region in the early solar system out of a diverse and variable mix of materials."
 When the uebula clisperses and objects are irst exposed to suuligπο their surfaces hea «αμα volatiles are driven oll.," When the nebula disperses and objects are first exposed to sunlight, their surfaces heat, and volatiles are driven off."
 For wkely variable starting ‘olpositions. the surface comj)ositious of objects a the same cdistauce [rom he sun will become early kleutical.," For widely variable starting compositions, the surface compositions of objects at the same distance from the sun will become nearly identical."
 Figuο 9 shows a mocel f‘om Brow1.Schaller&Fraser(2011a) which shows the istance at which cilerent voatiles would be completely removed from the surface of an object as a function of size «X the object., Figure 9 shows a model from \citet{2011ApJ...739L..60B} which shows the distance at which different volatiles would be completely removed from the surface of an object as a function of size of the object.
 Regardless of the diversity ol initial compositious. tlie surfaces ol the objects will quickly have strong gradients in volatile composition.," Regardless of the diversity of initial compositions, the surfaces of the objects will quickly have strong gradients in volatile composition."
 These volatiles are then irradiated aud begin to develo2 the colors expected rom tleir particular nix of 'eumaliinug surlace volatiles., These volatiles are then irradiated and begin to develop the colors expected from their particular mix of remaining surface volatiles.
 Iu this mocel. metliauol is the most uuportan coloriug agent.," In this model, methanol is the most important coloring agent."
 Methanol is depleted on all Surfaces= inside of aboul 20 AU and present ou a| of those exterior.," Methanol is depleted on all surfaces inside of about 20 AU, and present on all of those exterior."
 While ivdrocarbous are generally expected to carbonize aid turn dark aud 1eutral upon srolonged irradiation. laboratory experi.rents have sueeestedMOD that irradiation of metianol to dosages expected o1 WBOs over the age of the solar systetn leads to brielit red surfaces.," While hydrocarbons are generally expected to carbonize and turn dark and neutral upon prolonged irradiation, laboratory experiments have suggested that irradiation of methanol to dosages expected on KBOs over the age of the solar system leads to bright red surfaces."
 In this hypotlesis. (lien. the objects interior to the lemethanol line bec‘ole the neural population. while those ouside become the red. populatio," In this hypothesis, then, the objects interior to the methanol line become the neutral population, while those outside become the red population."
 A nOar system scale 1ustability suc Las that envisioned by the Nice model tliei scatters the objects onto ile orbits where they currenty reside. mixing the neutra and red populatious.," A solar system scale instability such as that envisioned by the Nice model then scatters the objects onto the orbits where they currently reside, mixing the neutral and red populations."
 The cod ClassicalI KBOs fori in jxace (Batyvein.Brown&Betts2012).. auc their uuique colors form by the additioial preseice of amnmonia o heir su‘aces.," The cold classical KBOs form in place \citep{2012ApJ...744L...3B}, and their unique colors form by the additional presence of ammonia on their surfaces."
 As KWBOs scater inward and become ceMatis. heir surfaces remain uuchauged.," As KBOs scatter inward and become centaurs, their surfaces remain unchanged."
 Wtile the moclel of Brown.Sealer&EFser(2011a) is t eolly quautitative model to aemp o explain all of he ¢aurreut observa10is of colors of IXBOs. 1t is cle:uly speclative.," While the model of \citet{2011ApJ...739L..60B} is the only quantitative model to attempt to explain all of the current observations of colors of KBOs, it is clearly speculative."
 Much of t aboratory data require to trace tle speciic irradiation cheulstry las not been performed. aud he question of tre evolttion of the su‘faces of tlie ceutatus ‘elnaius open.," Much of the laboratory data required to trace the specific irradiation chemistry has not been performed, and the question of the evolution of the surfaces of the centaurs remains open."
 Extending this moc o the Hubble Space Teescope SULveY of Fuser&Brown(2011).. which sugeMODests mixing ou i surfaces of ΝΕΟΣ. we WOld conclude that the irradiated volatiles are the neutral aud red maer hat are mixed in with the hydrated silicate-like material.," Extending this model to the Hubble Space Telescope survey of \citet{hwtsoss}, which suggests mixing on the surfaces of KBOs, we would conclude that the irradiated volatiles are the neutral and red materials that are mixed in with the hydrated silicate-like material."
 No obvious explanation exists for t different amounts of1nixlug between the irradiated. volatiles aud the silicates., No obvious explanation exists for the different amounts of mixing between the irradiated volatiles and the silicates.
Using WD cooling models from Foutaineetal.(2001).. we can derive a relationship between surface eravity aud effective temperature for a even mass (see Fieure 5)).,"Using WD cooling models from \cite{fontaine01}, we can derive a relationship between surface gravity and effective temperature for a given mass (see Figure \ref{fig:massconts}) )."
" Additionally. for any given mass and οσο), we can calculate the radius from Equation 3.."," Additionally, for any given mass and log(g), we can calculate the radius from Equation \ref{eqn:radius}."
 Putting the mass dependent Tie - logs(g) relation ogether with Equation 3 niuis that. of the our parameters. Toy. log(g) mass. and radius. the combination of Tag aud any of the other three allows oue o calculate the remaining two parameters.," Putting the mass dependent $_{\rm{eff}}$ - log(g) relation together with Equation \ref{eqn:radius} means that, of the four parameters, $_{\rm{eff}}$, log(g), mass, and radius, the combination of $_{\rm{eff}}$ and any of the other three allows one to calculate the remaining two parameters."
 Of these four. we have chosen to parametrize our erid over Toy aud nass. since previous work allows us to put a reasonable xior on the mass of our objects.," Of these four, we have chosen to parametrize our grid over $_{\rm{eff}}$ and mass, since previous work allows us to put a reasonable prior on the mass of our objects."
 This reduces the six xuwaneters that control cach WD spectrum (effective cluperature. mass. distance. reddening. surface gravity. and the radius of the WD) to four parameters that weed to be fit in our procedure.," This reduces the six parameters that control each WD spectrum (effective temperature, mass, distance, reddening, surface gravity, and the radius of the WD) to four parameters that need to be fit in our procedure."
 We are therefore able o generate a model SED in these four filters eiven a nass. effective temperature. reddening. aud a clistance.," We are therefore able to generate a model SED in these four filters given a mass, effective temperature, reddening, and a distance."
 The first three paraicters control the shape of the SED. while the distance controls the overall scaling.," The first three parameters control the shape of the SED, while the distance controls the overall scaling."
 In order o determine temperatures for our WDs and to display th iuodel and WD SEDs. we have first held mass aud reddening as constants m our fitting procedure.," In order to determine temperatures for our WDs and to display both model and WD SEDs, we have first held mass and reddening as constants in our fitting procedure."
 We lave assuned a mass of 0.53 ML..., We have assumed a mass of $0.53$ $_{\odot}$.
 This follows from the the prediction of Reuzini&FusiPecci(1958). and Beuziuietal.(1996). that the WDs forming today in globular clusters should have a mass of 0.53+0.02 AL... Further support of our mass sclection comes from recent direct spectroscopic nieasureiments frou Ialiraietal.(2009). of 0292001. AL. for the mass of WDs in ADL and 0.5340.03 AL. for the mass of WDs in NGC 6752 (Alochleret 1)., This follows from the the prediction of \cite{renzini88} and \cite{renzini96} that the WDs forming today in globular clusters should have a mass of $0.53 \pm 0.02$ $_{\odot}$ Further support of our mass selection comes from recent direct spectroscopic measurements from \cite{kalirai09} of $0.53 \pm 0.01$ $_{\odot}$ for the mass of WDs in M4 and $0.53 \pm 0.03$ $_{\odot}$ for the mass of WDs in NGC 6752 \citep{moehler04}.
. We also asstune a reddemime of 0.01 and then fit the distance modulus aud the effective temperature to each individual object., We also assume a reddening of $0.04$ and then fit the distance modulus and the effective temperature to each individual object.
 We ignore the mdividual distance moduli here aud include the fitted effective teniperatures iu Table 2. for our 59 WDs., We ignore the individual distance moduli here and include the fitted effective temperatures in Table \ref{tab:wds} for our 59 WDs.
 Figure 6. shows the SEDs for our selected. WD sample., Figure \ref{fig:sed_prob} shows the SEDs for our selected WD sample.
 Iu Section 5.. we determine the distauce to 17 Tuc using Equation 2 to fit to our WD phlotomoetry. however we do so using a 1miaxinimun likelihood fiction aud mareinalise over various parameters.," In Section \ref{sec:distance}, we determine the distance to 47 Tuc using Equation \ref{eqn:mag2} to fit to our WD photometry, however we do so using a maximum likelihood function and marginalise over various parameters."
 To determine tie distance to 17 Tuc. we will use white dwarf specral anodels iu order to predict au expected absolute magnitude for our suuple WDs in our I filters. similarly to our techuique in Section L.ILlowever iu this section. we fit the distance modulus to LF Tuc by comparing the expected absolute magnitude," To determine the distance to 47 Tuc, we will use white dwarf spectral models in order to predict an expected absolute magnitude for our sample WDs in our 4 filters, similarly to our technique in Section \ref{sec:seds}.However in this section, we fit the distance modulus to 47 Tuc by comparing the expected absolute magnitude"
"near the Si IV, CIIV, Al IJ, and MglII lines.","near the Si IV, IV, Al III, and II lines."
" For analysis of individual lines, we exclude quasars with a BAL associated with that wavelength."," For analysis of individual lines, we exclude quasars with a BAL associated with that wavelength."
" Because the presence of a BAL can also affect the spectral shape in the ultraviolet regime, we exclude every BAL quasar from our analysis of spectral slopes in Section 3.4;; we also exclude BALs from our composite spectra in Section 5.1.."," Because the presence of a BAL can also affect the spectral shape in the ultraviolet regime, we exclude every BAL quasar from our analysis of spectral slopes in Section \ref{subsec:color excess}; we also exclude BALs from our composite spectra in Section \ref{subsec:composites}."
 Radio morphology is crucial for categorizing quasars and understanding their underlying physical properties., Radio morphology is crucial for categorizing quasars and understanding their underlying physical properties.
" The apparent morphology of a radio source is determined by intrinsic parameters such as age, size, and jet power, for example, but also by orientation."," The apparent morphology of a radio source is determined by intrinsic parameters such as age, size, and jet power, for example, but also by orientation."
" In this section, we present the classification of radio quasars by radio morphology and define the morphology classes used throughout the rest of this paper."," In this section, we present the classification of radio quasars by radio morphology and define the morphology classes used throughout the rest of this paper."
 We then discuss the fraction of highly-reddened quasars as a function of morphology class., We then discuss the fraction of highly-reddened quasars as a function of morphology class.
" We begin by giving a brief overview of work by ?,hereafterKIO8 to classify radio sources by morphology and to describe the optical and radio properties of the different morphology classes.", We begin by giving a brief overview of work by \citet[][hereafter KI08]{ki08} to classify radio sources by morphology and to describe the optical and radio properties of the different morphology classes.
 Their work motivates the more detailed morphology investigations we now pursue., Their work motivates the more detailed morphology investigations we now pursue.
" KI08 compiled a large, homogeneous catalog of radio and optical sources for the purpose of exploring the multi-wavelength properties of AGNs."," KI08 compiled a large, homogeneous catalog of radio and optical sources for the purpose of exploring the multi-wavelength properties of AGNs."
 A simple method was used for automatic morphology classification based on flux densities obtained at varying spatial resolution from different radio sky surveys., A simple method was used for automatic morphology classification based on flux densities obtained at varying spatial resolution from different radio sky surveys.
" They used this method to define three morphology classes: *compact? (unresolved at “resolved” at bbut fairly concentrated),5’’)), and (resolved“complex” (having significant extended emission beyond 5'))."," They used this method to define three morphology classes: “compact"" (unresolved at ), “resolved"" (resolved at but fairly concentrated), and “complex"" (having significant extended emission beyond )."
" KI08 found that these three classes, defined entirely by their radio emission, correspond to different types of optical sources (with significant overlap among the classes)."," KI08 found that these three classes, defined entirely by their radio emission, correspond to different types of optical sources (with significant overlap among the classes)."
" For example, the complex and resolved classes consist of optically-identified galaxies with steep radio spectra, while flat-spectrum quasars typically fall into the compact class."," For example, the complex and resolved classes consist of optically-identified galaxies with steep radio spectra, while flat-spectrum quasars typically fall into the compact class."
" These results are consistent with the basic unification theory of radio-loud objects, which suggests that radio AGNs viewed along the jet axis appear as compact and flat-spectrum in the radio (because core-boosting dominates the radio emission) and as quasars in the optical (because the central object is "," These results are consistent with the basic unification theory of radio-loud objects, which suggests that radio AGNs viewed along the jet axis appear as compact and flat-spectrum in the radio (because core-boosting dominates the radio emission) and as quasars in the optical (because the central object is unobscured)."
"Conversely, AGNs viewed perpendicular to the jet axis unobscured).tend to be dominated by steep-spectrum extended radio emission, and appear as galaxies in the optical regime."," Conversely, AGNs viewed perpendicular to the jet axis tend to be dominated by steep-spectrum extended radio emission, and appear as galaxies in the optical regime."
" With the new quasar classifications presented here, we extend the analysis of KI08 by examining the dependence of quasar spectral properties on radio morphology."," With the new quasar classifications presented here, we extend the analysis of KI08 by examining the dependence of quasar spectral properties on radio morphology."
" Additionally, our classifications separate the “complex” class of KIO8 into finer morphology categories."," Additionally, our classifications separate the “complex"" class of KI08 into finer morphology categories."
" We defined our initial quasar sample using 77,258 quasars from S07 (the full catalog of 77,429 sources from the SDSS DR5 catalog, less the 171 sources that were dropped from the updated DR7 catalog; see refsubsec:sdss))."," We defined our initial quasar sample using 77,258 quasars from S07 (the full catalog of 77,429 sources from the SDSS DR5 catalog, less the 171 sources that were dropped from the updated DR7 catalog; see \\ref{subsec:sdss}) )."
 We divided the S07 sample into a main sample of “radio quasars” and a secondary sample of “radio-quiet quasars”.," We divided the S07 sample into a main sample of “radio quasars"" and a secondary sample of “radio-quiet quasars""."
" The radio sample consists of 4714 quasars with a nearby FIRST source whose integrated flux density is $99>2 mJy, in order to ensure a sufficiently bright sample for reliable morphology classification."," The radio sample consists of 4714 quasars with a nearby FIRST source whose integrated flux density is $S_{20}\geq2$ mJy, in order to ensure a sufficiently bright sample for reliable morphology classification."
" The radio-quiet sample comprises 65,253 quasars that lie in the area of FIRST coverage but do not have a FIRST We exclude from our analysis an additional 5525 quasars that lie outside the FIRST survey area."," The radio-quiet sample comprises 65,253 quasars that lie in the area of FIRST coverage but do not have a FIRST We exclude from our analysis an additional 5525 quasars that lie outside the FIRST survey area."
" We also exclude 1449 ""faint radio? quasars with a faint (Soo<2 mJy) FIRST counterpart within2"".", We also exclude 1449 “faint radio” quasars with a faint $S_{20}<2$ mJy) FIRST counterpart within.
. 'The radio sample may contain a small number of false matches., The radio sample may contain a small number of false matches.
 The distribution of FIRST—SDSS positional offsets is a Rayleigh distribution with o=0.2 and a small tail out to, The distribution of FIRST—SDSS positional offsets is a Rayleigh distribution with $\sigma=0.2$ and a small tail out to
This classification has been applied. to the 10. σ-άνορ galaxies and the 10 control sample galaxies for which we have images.,This classification has been applied to the 10 $\sigma$ -drop galaxies and the 10 control sample galaxies for which we have images.
 Results for the c-drop sample are listed in Table 2.. and for the control sample in Table 3.. and are summarised as histograms in Fig. 4..," Results for the $\sigma$ -drop sample are listed in Table \ref{sigmaproperties}, and for the control sample in Table \ref{controlproperties}, and are summarised as histograms in Fig. \ref{halfa}."
 There are no significant differences at all in the very central eemission morphology., There are no significant differences at all in the very central emission morphology.
 For the inner kkpe region. however. we notice that star-forming activity is much higher in c-drop galaxies than in control-sample galaxies.," For the inner kpc region, however, we notice that star-forming activity is much higher in $\sigma$ -drop galaxies than in control-sample galaxies."
 Patehes and filaments in c-drop host galaxies are more extended (covering a larger surface area) and more obvious in the images than in the control sample galaxies., Patches and filaments in $\sigma$ -drop host galaxies are more extended (covering a larger surface area) and more obvious in the images than in the control sample galaxies.
 Most of the control galaxies are devoid of formation or have a weak diffuse emission., Most of the control galaxies are devoid of formation or have a weak diffuse emission.
 Galaxies in the σ- sample often show strong starbursts with bright rings and patches., Galaxies in the $\sigma$ -drop sample often show strong starbursts with bright rings and patches.
 The differences between these two subsamples could be due to the fact that they are not as well matched as the two samples of 20 galaxies., The differences between these two subsamples could be due to the fact that they are not as well matched as the two samples of 20 galaxies.
 The mean axis ratio for the c-drop galaxies is 0.56 while it is 0.62 for the control galaxies., The mean axis ratio for the $\sigma$ -drop galaxies is 0.56 while it is 0.62 for the control galaxies.
 Four o-drop galaxies have axis ratios smaller than 0.4., Four $\sigma$ -drop galaxies have axis ratios smaller than 0.4.
 Two of these galaxies (GC 1808 and NGC 3593) show clear evidence of uclear star-burst rings. which means that we are seeing the aner region without strong extinction.," Two of these galaxies (NGC 1808 and NGC 3593) show clear evidence of nuclear star-burst rings, which means that we are seeing the inner region without strong extinction."
 In the other two cases (NGC 3623 and NGC 3627) there is no emission. possibly due to dust obscuration.," In the other two cases (NGC 3623 and NGC 3627) there is no emission, possibly due to dust obscuration."
 In the control sample dust seems to hide the inner regions in NGC 3169 (d/D=0.62 and diffuse circumnuclear emission) and in NGC 3718 (4/0=0.49 with o circumnuclear emission)., In the control sample dust seems to hide the inner regions in NGC 3169 $d/D=0.62$ and diffuse circumnuclear emission) and in NGC 3718 $d/D=0.49$ with no circumnuclear emission).
 It is difficult to explain using only obscuration the virtual absence of circumnuclear emission in the two other galaxies. NGC 3031 which shows a tiny nuclear spiral and the nearly face-on galaxy NGC 6340.," It is difficult to explain using only obscuration the virtual absence of circumnuclear emission in the two other galaxies, NGC 3031 which shows a tiny nuclear spiral and the nearly face-on galaxy NGC 6340."
 The other difference between the two subsamples ts that the mean morphological type of the control sample is earlier than that of the c-drop sample (mean Hubble morphological type 1.8 and 2.6. respectively).," The other difference between the two subsamples is that the mean morphological type of the control sample is earlier than that of the $\sigma$ -drop sample (mean Hubble morphological type 1.8 and 2.6, respectively)."
 Both subsamples have galaxies with types ranging between zero and six., Both subsamples have galaxies with types ranging between zero and six.
 In Knapen et al. (, In Knapen et al. (
2006) galaxies with no circumnuclear emission tend to be early type and patchy emission tends to be in later-type galaxies which is coherent with what we have found in the two subsamples.,2006) galaxies with no circumnuclear emission tend to be early type and patchy emission tends to be in later-type galaxies which is coherent with what we have found in the two subsamples.
 Diffuse emission is also more prevalent in the control sample., Diffuse emission is also more prevalent in the control sample.
 As the diffuse emission is an intermediate case between no emission and patchy emission. the overabundance of this type in the control sample is probably due to the fact that it contains earlier-type galaxies.," As the diffuse emission is an intermediate case between no emission and patchy emission, the overabundance of this type in the control sample is probably due to the fact that it contains earlier-type galaxies."
 But in Knapen et al. (, But in Knapen et al. (
2006) the nuclear starburst ring distribution peaks around T=|.case.,"2006) the nuclear starburst ring distribution peaks around $T=1$,."
 According to the statistics of Knapen et al. (, According to the statistics of Knapen et al. (
2006). corrected for the nuclear ring fractions reported there for each morphological type of the host galaxy. four of our o-drop galaxies and six control galaxies should have a circumnuclear ring—the actually observed numbers are five and zero.,"2006), corrected for the nuclear ring fractions reported there for each morphological type of the host galaxy, four of our $\sigma$ -drop galaxies and six control galaxies should have a circumnuclear ring—the actually observed numbers are five and zero."
 Therefore it is reasonable to think that the higher proportion of circumnuclear star-forming rings in o galaxies is a true effect correlated with the c-drop., Therefore it is reasonable to think that the higher proportion of circumnuclear star-forming rings in $\sigma$ -drop galaxies is a true effect correlated with the $\sigma$ -drop.
shot noise.,shot noise.
" LILO derived that weight function which optimizes a,..", H10 derived that weight function which optimizes $\sigma_w^2$.
 We derive the analogous optimal weight for ἐς., We derive the analogous optimal weight for $E$.
 Suppose we are given a list of halo masses ancl positions., Suppose we are given a list of halo masses and positions.
 Suppose that this list is complete down to some limiting mass my., Suppose that this list is complete down to some limiting mass $m_d$.
" We can assign the remainder of the mass to a ""dust bin” which contains the fraction i,=1SZay of mass that is not in the halos."," We can assign the remainder of the mass to a “dust bin,” which contains the fraction $\eta_d = 1 - \sum \eta_i$ of mass that is not in the halos."
" If 6, is the relative clensity Ductuation field of the mass in the dust bin. we can define the power of the dust field as and a bias vector οι between halos and dust by )ecause the bias and optimal weighting for the halo bins against the mass are: In this formulation it is apparent that as our halo catalog extends to lower masses and. teQO. the optimal weight 20,44n. mass weighting. and. £70."," If $\delta_d$ is the relative density fluctuation field of the mass in the dust bin, we can define the power of the dust field as and a bias vector $\vc_d$ between halos and dust by Because the bias and optimal weighting for the halo bins against the mass are: In this formulation it is apparent that as our halo catalog extends to lower masses and $\eta_d \rightarrow 0$, the optimal weight $\vw_{\rm opt}\rightarrow \veta$, mass weighting, and $E^2\rightarrow 0$."
 The further question of interest for finite apy is: How well do the known halo Uuetuations 4; predict the dust bin density 947, The further question of interest for finite $\eta_d$ is: How well do the known halo fluctuations $\delta_i$ predict the dust bin density $\delta_d$?
 So far the analysis has been completely. general as to the generation of the mass field ancl the designation of halos within it., So far the analysis has been completely general as to the generation of the mass field and the designation of halos within it.
 Now we examine mocdels for the relation between halos ancl mass., Now we examine models for the relation between halos and mass.
 The most common assumption about the distribution of halos (or galaxies) is that they are drawn from the mass distribution in a biased Poisson process., The most common assumption about the distribution of halos (or galaxies) is that they are drawn from the mass distribution in a biased Poisson process.
 Lf we are examining the Fourier coellicients of the mass and halo distributions. the biased. Poisson model can be broken down into three assumptions: Asstunption (iii) of the biased-Poisson model has been noted to be inconsistent with the assumption that the jio catalog can be extended to comprise all of the mass.," If we are examining the Fourier coefficients of the mass and halo distributions, the biased Poisson model can be broken down into three assumptions: Assumption (iii) of the biased-Poisson model has been noted to be inconsistent with the assumption that the halo catalog can be extended to comprise all of the mass."
 We will therefore consider models in which assumption (i) olds without (iii). such that the halos sample a field tha is distinct. from the mass distribution.," We will therefore consider models in which assumption (i) holds without (iii), such that the halos sample a field that is distinct from the mass distribution."
 We will take care herefore to distinguish οι. the bias of the halos with respec o the halo field 85. from5.. which we abwavs celine via the covariance with mass as per (τὸ).," We will take care therefore to distinguish , the bias of the halos with respect to the halo field $\delta_h$, from, which we always define via the covariance with mass as per \ref{bcov}) )."
 The assumptions that halos are linearly biased. and tha he halo generating process has no spatial correlations are idealizations: in fact. halos in simulations do not overlap (almost by definition). and their bias is non-linear.," The assumptions that halos are linearly biased, and that the halo generating process has no spatial correlations are idealizations: in fact, halos in simulations do not overlap (almost by definition), and their bias is non-linear."
 We wil return to the limits of these assumptions later., We will return to the limits of these assumptions later.
 When takes the form. (26)). two things are of note: first. this description is stable under re-binning of the halos in the imit of narrow bins.," When takes the form \ref{vvt}) ), two things are of note: first, this description is stable under re-binning of the halos in the limit of narrow bins."
 More specifically. i 0; ancl f; are varving functions of the mass m; of halos in bin 7. then the v; and f; do not change if two adjacent bins are. merged.," More specifically, if $v_i$ and $f_i$ are slowly-varying functions of the mass $m_i$ of halos in bin $i$ , then the $v_i$ and $f_i$ do not change if two adjacent bins are merged."
 In other words we can write functions e(m) ancl f(m). and all of our linear-alecbra formulations can be carried over into integrals over halo mass m.," In other words we can write functions $v(m)$ and $f(m)$, and all of our linear-algebra formulations can be carried over into integrals over halo mass $m$."
 The second useful. [act about (26)) is that it can be inverted analytically using the Sherman-Morrison formula: When all three conditions of the biased-Poisson mocel are met. the optimal weight function and stochasticity are found simply from (15)) and (28)): This recovers the result. from 7) that. the. optimal linear mass estimator in a Poisson model weightsD cach halo by its bias 5; (times a mass-indepencent factor) and the stochasticity. of the estimator is determined. by Mud.," The second useful fact about \ref{vvt}) ) is that it can be inverted analytically using the Sherman-Morrison formula: When all three conditions of the biased-Poisson model are met, the optimal weight function and stochasticity are found simply from \ref{wopt}) ) and \ref{vvtinv}) ): This recovers the result from \citet{Percival04} that the optimal linear mass estimator in a Poisson model weights each halo by its bias $b_i$ (times a mass-independent factor), and the stochasticity of the estimator is determined by $\sum n_i b_i^2 P$."
 Conveniently the weights scale with the bias. independent of the range of halo masses included in the estimator.," Conveniently the weights scale with the bias, independent of the range of halo masses included in the estimator."
" This property does hold for more general forms of aandδι, it fails when vz6 and condition (ii) is violated."," This property does hold for more general forms of and, it fails when $\vv \ne \vb$ and condition (iii) is violated."
 In this paper we will assume that the halos occupy the matter distribution via a biased Poisson process., In this paper we will assume that the halos occupy the matter distribution via a biased Poisson process.
 We will examine the muniatrix for halos in numerical simulations. examine what if any of the three biased-Poisson conditions actually holds. and then use the general formulae (15)) and (16))to find how the optimal stochasticity clillers fromthe Poisson predictions.," We will examine the matrix for halos in numerical simulations, examine what if any of the three biased-Poisson conditions actually holds, and then use the general formulae \ref{wopt}) ) and \ref{eopt}) )to find how the optimal stochasticity differs fromthe Poisson predictions."
 We now examine the case where all halos. and. the dust.," We now examine the case where all halos, and the dust,"
kink mode in filament threads.,kink mode in filament threads.
 For coronal loops the kink speed is outside the slow continuum., For coronal loops the kink speed is outside the slow continuum.
" A coronal loop is presumably hotter and denser than its surrounding corona, so the ordering of sound, Altvénn, and kink speeds is cy.€€;va;<c;va."," A coronal loop is presumably hotter and denser than its surrounding corona, so the ordering of sound, Alfvénn, and kink speeds is $\csco < \csp < \vap < c_k < \vac$."
" Therefore, there is no slow resonance for the kink mode in coronal loops and the present study does not apply to coronal loops."," Therefore, there is no slow resonance for the kink mode in coronal loops and the present study does not apply to coronal loops."
 This Letter is organized as follows: 2 contains a description of the model configuration and the mathematical method., This Letter is organized as follows: \ref{sec:math} contains a description of the model configuration and the mathematical method.
 The results are presented and discussed in 3.., The results are presented and discussed in \ref{sec:results}.
" Finally, our conclusions are ejiven in 4."," Finally, our conclusions are given in \ref{sec:conclusion}."
 The model for the equilibrium configuration considered here is equivalent to that of (2008).., The model for the equilibrium configuration considered here is equivalent to that of \citet{arregui}.
" We model a filament thread as a straight cylinder with prominence-like conditions embedded in an unbounded corona, with a transverse inhomogencous layer between both media (see Fig. 1)."," We model a filament thread as a straight cylinder with prominence-like conditions embedded in an unbounded corona, with a transverse inhomogeneous layer between both media (see Fig. \ref{fig:model}))."
" We use cylindrical coordinates, namely r. v. and z for the radial, azimuthal, and longitudinal coordinates."," We use cylindrical coordinates, namely $r$, $\varphi$, and $z$ for the radial, azimuthal, and longitudinal coordinates."
" We adopt the density profile by Ruderman&Roberts(2002) which only depends on the radial direction, with In these expressions, p; 1s the internal density, p, is the external (coronal) density, @ is the tube mean radius, and / is the transitional layer width."," We adopt the density profile by \citet{rudermanroberts} which only depends on the radial direction, with In these expressions, $\rho_i$ is the internal density, $\rho_e$ is the external (coronal) density, $a$ is the tube mean radius, and $l$ is the transitional layer width."
" The limits //e=0 and //a=2correspond to a homogeneous tube and a fully inhomogencous tube, respectively."," The limits $l/a=0$ and $l/a=2$correspond to a homogeneous tube and a fully inhomogeneous tube, respectively."
" We use the following densities: p=5x107!kgm? and p,=2.5x10UKgm.", We use the following densities: $\rho_i=5\times10^{-11}~{\rm kg}~{\rm m}^{-3}$ and $\rho_e=2.5\times 10^{-13}~{\rm kg}~{\rm m}^{-3}$.
" Therefore. the density contrast between the internal and external plasma is p;/p,=200."," Therefore, the density contrast between the internal and external plasma is $\rho_i/\rho_e=200$."
 The plasma temperature is related to the density through the usual ideal gas equation., The plasma temperature is related to the density through the usual ideal gas equation.
" We consider T;=8000 K and T,=10° K for the internal and external temperatures, respectively."," We consider $T_i=8000$ K and $T_e=10^{6}$ K for the internal and external temperatures, respectively."
" The magnetic field is taken homogeneous and orientated along the z-direction, By= Boyz. with By=5 G everywhere."," The magnetic field is taken homogeneous and orientated along the $z$ -direction, ${\mathbf B}_0=B_0\hat{\mathbf z}$ , with $B_0=5$ G everywhere."
" With these conditions, 6« 0.04."," With these conditions, $\beta\approx0.04$ ."
" We consider the linear, ideal MHD equations for the 5z0 case."," We consider the linear, ideal MHD equations for the $\beta \neq 0$ case."
 Since y and z are ignorable, Since $\varphi$ and $z$ are ignorable
as given by the DoPLLIOT program.,as given by the DoPHOT program.
 Stars where variations of at least de occur in two or more of the passbancls are considered to be variables., Stars where variations of at least $\sigma$ occur in two or more of the passbands are considered to be variables.
 Three of the obscured stars are included. together with five carbon stars: two from the ALAV ist. two in common between ALW and DD. and one in the DB list only.," Three of the obscured stars are included, together with five carbon stars: two from the ALW list, two in common between ALW and DB, and one in the DB list only."
 Phere are hints that some of the other carbon stars are also variable., There are hints that some of the other carbon stars are also variable.
 Note that the obscured star C must rave à very laree amplitude and probably a period ofat least l vear., Note that the obscured star C must have a very large amplitude and probably a period of at least 1 year.
 This is consistent with the fact that stars with Ην colours similar to these stars are generally Mira. variables (sce Whitclock 2000. fis.," This is consistent with the fact that stars with JHK colours similar to these stars are generally Mira variables (see Whitelock 2000, fig."
 1)., 1).
 A full discussion of the nature of the obscured stars must obviously be deferred. till further variability studies have been mace., A full discussion of the nature of the obscured stars must obviously be deferred till further variability studies have been made.
 Neverthless it is of interest to craw some preliminary The results we have obtained are strikingly similar to those of Nishida et al. (, Neverthless it is of interest to draw some preliminary The results we have obtained are strikingly similar to those of Nishida et al. (
2000) (see also Tanabóé et al.,2000) (see also Tanabé et al.
 1997. 1999) who found an obscurecl carbon Mira in cach of the intermediate age clusters. NGC 419. NGC 1783 and NGC 1978. in the Magellanic Clouds.," 1997, 1999) who found an obscured carbon Mira in each of the intermediate age clusters, NGC 419, NGC 1783 and NGC 1978, in the Magellanic Clouds."
 These variables have values OL (JWW) (SAAO svstem) between 3.75 and 4.76., These variables have values of $(J-K)$ (SAAO system) between 3.75 and 4.76.
 Thus. in colour. the obscured Leo L stars lic between the Sagittarius dwarf spheroidal Miras.. mentioned above. and the LMC cluster. Miras.," Thus, in colour, the obscured Leo I stars lie between the Sagittarius dwarf spheroidal Miras, mentioned above, and the LMC cluster Miras."
 In the Magellanic Cloud: clusters the tip of the unobseured ACD is at AM—S4 and AM~53 (using data from Frogel ct al., In the Magellanic Cloud clusters the tip of the unobscured AGB is at $M_{K} \sim -8.4$ and $M_{bol} \sim -5.3$ (using data from Frogel et al.
 1990 ancl distance. moduli of 18.6 ancl 19.0 for the LMC ancl SAIC)., 1990 and distance moduli of 18.6 and 19.0 for the LMC and SMC).
" In. Leo | the corresponding values are 8.3 and 5.1 which however remain uncertain pending full variability In the Magellanic Cloud. clusters the Aliras have a mean Al, of SNO and Af,=5.1 or slightly brighter.", In Leo I the corresponding values are –8.3 and –5.1 which however remain uncertain pending full variability In the Magellanic Cloud clusters the Miras have a mean $M_{K}$ of –8.0 and $M_{bol} = -5.1$ or slightly brighter.
 The corresponding figures for the obseurecl stars in. Leo ] are quite uncertain both because three. ancl possibly all. are variable ancl also because the estimation of bolometric corrections for such stars [rom ο){1{ν colours is rather uncertain.," The corresponding figures for the obscured stars in Leo I are quite uncertain both because three, and possibly all, are variable and also because the estimation of bolometric corrections for such stars from $JHK$ colours is rather uncertain."
 The three obscured stars found to be variable have Aly=SQ and ως4.9 whilst all five obscured stars viele) Aly=7.6 and Adj—46., The three obscured stars found to be variable have $M_{K} = -8.0$ and $M_{bol} \sim -4.9$ whilst all five obscured stars yield $M_{K} = -7.6$ and $M_{bol} \sim -4.6$.
 To derive these results for Leo Lowe have used a relation of (4ÁN) to bolometric correction for carbon Miras derived by Whitelock (to be published) which includes the use of LSO data., To derive these results for Leo I we have used a relation of $(H-K)$ to bolometric correction for carbon Miras derived by Whitelock (to be published) which includes the use of ISO data.
" Thus the tip of the unobscured. AGB in AJ, and M4 is very similar in Leo | to that in the three Magellanic Cloud clusters.", Thus the tip of the unobscured AGB in $M_{K}$ and $M_{bol}$ is very similar in Leo I to that in the three Magellanic Cloud clusters.
 Phe obseured AGB may be fainter in As; than the cluster Miras but that is not lt is of some interest to note that if the intrinsic (underlving) colours of the five. obscured stars were (JIvy)~2.0. they would all move to a position near the top of the AGB in Avs. when dereddened: using the reddening law (AA~ACS A) found for the circumstellar envelope of the galactic carbon Mira. Ro kor (Feast. et al.," The obscured AGB may be fainter in $M_{bol}$ than the cluster Miras but that is not It is of some interest to note that if the intrinsic (underlying) colours of the five obscured stars were $(J-K_{s}) \sim 2.0$, they would all move to a position near the top of the AGB in $K_{s}$, when dereddened using the reddening law $\Delta K \sim \Delta (J-K)$ ) found for the circumstellar envelope of the galactic carbon Mira R For (Feast et al."
 The Alagellanie Cloud clusters are estimated to have ages in the range 1.6 to 2.0 Gyr. metallicities. Fe/ll 0.6. and turm-olf masses of ~1.5M..," The Magellanic Cloud clusters are estimated to have ages in the range 1.6 to 2.0 Gyr, metallicities, [Fe/H] $\sim -0.6$, and turn-off masses of $\sim 1.5 M_{\odot}$."
 In view of the above discussion it seems likely that the Leo 1. obscure stars are in the same age and mass range., In view of the above discussion it seems likely that the Leo I obscured stars are in the same age and mass range.
 Lt is therefore interesting to note that Gallart ct al (1999) suggest. tha major star formation in Leo I stopped about 2 Gyr ago and that a metallicity of that population as high as Vefl] ~0.6 is possible., It is therefore interesting to note that Gallart et al (1999) suggest that major star formation in Leo I stopped about 2 Gyr ago and that a metallicity of that population as high as [Fe/H] $\sim -0.6$ is possible.
 The results thus suggest tha the obscured stars. and the most. luminous unobscurec AGB stars. at least. belong to this voungest major stellar component of Leo The very rec star (star E) in [field Bois especially interesting.," The results thus suggest that the obscured stars, and the most luminous unobscured AGB stars, at least, belong to this youngest major stellar component of Leo The very red star (star E) in field B is especially interesting."
 Ht is 8 arcmin from the centre of Leo L much further out than any of the other stars we discuss.," It is $\sim 8$ arcmin from the centre of Leo I, much further out than any of the other stars we discuss."
 Even the C stars outside field A. mentioned in section 2.. are within ~4.5 aremin of the centre.," Even the C stars outside field A, mentioned in section 2.2, are within $\sim 4.5$ arcmin of the centre."
 Phere is HE in à wide area around Leo IL and apparently associated. with it (Blitz Robishaw 2000)., There is HI in a wide area around Leo I and apparently associated with it (Blitz Robishaw 2000).
 Star counts (Irwin IHatzidimitriou 1995) show that Leo LE has a tidal radius of 13 arcmin. which extends well bevond the distance of star LE from the centre.," Star counts (Irwin Hatzidimitriou 1995) show that Leo I has a tidal radius of 13 arcmin, which extends well beyond the distance of star E from the centre."
 However the stellar density at these distances is very low compared with that in the main body of the galaxy., However the stellar density at these distances is very low compared with that in the main body of the galaxy.
 Furthermore. in cwarf spheroidals the vounger populations are eenerally more concentrated to the centre than the older ones (e.g. Llarbeck et al.," Furthermore, in dwarf spheroidals the younger populations are generally more concentrated to the centre than the older ones (e.g. Harbeck et al."
 2001)., 2001).
 Thus if we have correctly interpreted this star as similar to the obscured ACD stars in field. A. ancl hence relatively voung. its position at such a large distance from the centre is remarkable.," Thus if we have correctly interpreted this star as similar to the obscured AGB stars in field A, and hence relatively young, its position at such a large distance from the centre is remarkable."
" ""This is particularly so when it is recalled that such objects are short lived and. are therefore tracers of much larger populations.", This is particularly so when it is recalled that such objects are short lived and are therefore tracers of much larger populations.
Pulling these results into context with other high spatial resolution studies of FU Ori and other FUORs we see that:,Putting these results into context with other high spatial resolution studies of FU Ori and other FUORs we see that:
"(blue bars), leads to a much more accurate and resolution independent description of recombinations (within ~25% of Ncccorrect 0n all grid resolutions).","(blue bars), leads to a much more accurate and resolution independent description of recombinations (within $\sim 25\%$ of $\dot{N}_\mr{rec,correct}$ on all grid resolutions)."
" Note that the assumption of a fully ionized simulation box is a special case, chosen for illustrative purposes."," Note that the assumption of a fully ionized simulation box is a special case, chosen for illustrative purposes."
 A more relevant case for the study of reionization is a partially ionized cosmological density field., A more relevant case for the study of reionization is a partially ionized cosmological density field.
" The most important consequence of using Ciocai instead of Cgiopai is a much lower average recombination rate, for whichever overdensity regions are ionized at any time."," The most important consequence of using ${\cal C}_\mr{local}$ instead of ${\cal C}_\mr{global}$ is a much lower average recombination rate, for whichever overdensity regions are ionized at any time."
 The convergence of recombination rates obtained in Fig., The convergence of recombination rates obtained in Fig.
" 3 with the use of Cijocai also leads to the convergence of I- speeds during reionization as we will show in?,, where we also discuss the ionized fraction as a function of overdensity during reionization."," \ref{fig3} with the use of ${\cal C}_\mr{local}$ also leads to the convergence of I-front speeds during reionization as we will show in, where we also discuss the ionized fraction as a function of overdensity during reionization."
" Also, Ciocal is computed assuming a full ionization and does not provide a perfectly accurate recombination rate estimate for partially ionized sub-volumes."," Also, ${\cal C}_\mr{local}$ is computed assuming a full ionization and does not provide a perfectly accurate recombination rate estimate for partially ionized sub-volumes."
" Therefore, the use of a higher resolution RT computational mesh is always preferable to employing any kind of clumping."," Therefore, the use of a higher resolution RT computational mesh is always preferable to employing any kind of clumping."
" Clumping factors are often used in simulations of reionization to represent the unresolved matter inhomogeneity, below the computational mesh resolution, which may significantly contribute to the recombination ratesmcquinn07."," Clumping factors are often used in simulations of reionization to represent the unresolved matter inhomogeneity, below the computational mesh resolution, which may significantly contribute to the recombination rates."
"e.g.) Here we have shown that, because the density field during reionization is so inhomogeneous, using a single clumping factor in general leads to a significant over estimate of the recombination rate (factors of several), which may in fact get worse as the grid resolution is improved."," Here we have shown that, because the density field during reionization is so inhomogeneous, using a single clumping factor in general leads to a significant over estimate of the recombination rate (factors of several), which may in fact get worse as the grid resolution is improved."
" As a consequence the speed of ionization fronts is artificially depressed, reionization delayed, and the outcome of the radiative transfer simulations significantly resolution dependent."," As a consequence the speed of ionization fronts is artificially depressed, reionization delayed, and the outcome of the radiative transfer simulations significantly resolution dependent."
" improved on this issue by analytically deriving the clumping factor as a function of overdensity, C=C(A), and also took into account some density dependence by splitting their simulation volumes into 8 equal sub-volumes."," improved on this issue by analytically deriving the clumping factor as a function of overdensity, ${\cal
 C}={\cal C}(\Delta)$, and also took into account some density dependence by splitting their simulation volumes into 8 equal sub-volumes."
" However, both ignored the fact that clumping still depends on the volume over which it is computed the size ofthe sub-volumes), and that dependence is not negligible as shown in Fig. 2.."," However, both ignored the fact that clumping still depends on the volume over which it is computed the size ofthe sub-volumes), and that dependence is not negligible as shown in Fig. \ref{fig2}."
 Not taking the volume into account will always lead to the RT results depending on the computational mesh resolution., Not taking the volume into account will always lead to the RT results depending on the computational mesh resolution.
 Fortunately it is not computationally intensive to compute a local clumping factor on a reasonably fine grid., Fortunately it is not computationally intensive to compute a local clumping factor on a reasonably fine grid.
" The recombination rate obtained using this clumping factor is very close to the vvalue inferred directly form the particles themselves, and is not very sensitive to the resolution of the mesh used (see Fig. 3))."," The recombination rate obtained using this clumping factor is very close to the value inferred directly form the particles themselves, and is not very sensitive to the resolution of the mesh used (see Fig. \ref{fig3}) )."
" In our dark matter only simulations the cell size should be such that most of the structures are resolved, and should ideally be close to the Jeans mass of the photo-ionised IGM(?)."," In our dark matter only simulations the cell size should be such that most of the structures are resolved, and should ideally be close to the Jeans mass of the photo-ionised IGM."
. We therefore argue that any future fits of clumping factors aimed for use in RT simulations of reionization must include both density and volume dependence., We therefore argue that any future fits of clumping factors aimed for use in RT simulations of reionization must include both density and volume dependence.
" The alternative of computing the recombination rate per particle directly (?),, is far more computationally expensive."," The alternative of computing the recombination rate per particle directly , is far more computationally expensive."
We will consider the evolution of short-wavelength compressive sliwaves in which only the initial velocity is perturbed. both for simplicity aud lor consistency with our calculation of the iucompressive slaves.,"We will consider the evolution of short-wavelength compressive shwaves in which only the initial velocity is perturbed, both for simplicity and for consistency with our calculation of the short-wavelength incompressive shwaves."
 As before. we will assume that the initial kinetic energy is cistributed isotropically.," As before, we will assume that the initial kinetic energy is distributed isotropically."
" We use the WEB solutious to equation (??)) with Hh,>>].] With the initial density perturbation set to zero (cousisteut with our assumption of ouly initial velocity perturbatious). the uniformly-valid asvimptotic solution to leading order in (HÀ)1 is given by =Pvyen cos((W.—Wy). =Fe where the WIXB eikoual is given by W fokdl = euro ο. “right))) |. with Wi beiug the value of W at /=ott Using equation (?2)). the euergy integral for the compressive shiwaves in the short-wavelenetl lint is = L? (Phy = L? ΜΑΣ"," We use the WKB solutions to equation \ref{SOUND}) ) with $H k_y \gg 
1$ With the initial density perturbation set to zero (consistent with our assumption of only initial velocity perturbations), the uniformly-valid asymptotic solution to leading order in $(H k_y)^{-1}$ is given by = (W-W_0), =, _c = where the WKB eikonal is given by W c_s k dt = = + + ) ), with $W_0$ being the value of $W$ at $t=0$ Using equation \ref{VXWKB}) ), the energy integral for the compressive shwaves in the short-wavelength limit is = L^2 d^2k_0 = L^2 d^2k_0 _0 ."
"Earlier, we already reported the chemical and evolutionary distinction of the nucleus of NGC 7743 with our integral-field spectroscopic results for the central part of the galaxy.","Earlier, we already reported the chemical and evolutionary distinction of the nucleus of NGC 7743 with our integral-field spectroscopic results for the central part of the galaxy."
" By obtaining the data with the Multi-Pupil Fiber Spectrograph (MPFS) of the Russian 6-m telescope, Sil’chenko has measured the Lick indices and found that the (2006)nucleus was 5 times more metal-rich, and by more than 2 Gyr younger with respect to the neighboring bulge."," By obtaining the data with the Multi-Pupil Fiber Spectrograph (MPFS) of the Russian 6-m telescope, \citet{Silchenko2006} has measured the Lick indices and found that the nucleus was 5 times more metal-rich, and by more than 2 Gyr younger with respect to the neighboring bulge."
" However, the absolute values of the ages and metallicities obtained here differ fromthose, obtained by Sil’chenko there +0.7 was found for the nucleus and(2006):: 0.0 for the[Z/H]— bulge, and the mean age of the bulge was determined as 4 Gyr old."," However, the absolute values of the ages and metallicities obtained here differ fromthose, obtained by \citet{Silchenko2006}: there $=+0.7$ was found for the nucleus and $0.0$ for the bulge, and the mean age of the bulge was determined as 4 Gyr old."
" The emission lines were very strong in the nuclear and circumnuclear spectra, and the excitation mechanism of those was unclear, so perhaps the correction of the Lick index Hf for the emission by Sil’chenko(2006) was not completely correct; hence, the stellar population parameters for NGC 7743 in Sil’chenko (2006) may be biased."," The emission lines were very strong in the nuclear and circumnuclear spectra, and the excitation mechanism of those was unclear, so perhaps the correction of the Lick index $\beta$ for the emission by \citet{Silchenko2006} was not completely correct; hence, the stellar population parameters for NGC 7743 in Sil'chenko (2006) may be biased."
" Later, Si’chenko&Chilingarian(2011) have analysed the SAURON data (which are analysed here as well), by the method of direct spectrum fitting in the pixel area, that allows to exclude all the spectral regions, affected by emission."," Later, \citet{SilchenkoChilingarian2011} have analysed the SAURON data (which are analysed here as well), by the method of direct spectrum fitting in the pixel area, that allows to exclude all the spectral regions, affected by emission."
" The results of the SAURON data analysis by Si’chenko&Chilingarian(2011) agree perfectly with the present results, derived by fitting the long-slit SCORPIO spectra: they have obtained +0.2 and the age of 1.2 Gyr for the nucleus, and —0.33 and the age of 2.2 Gyr for the bulge (Sil'chenkoChilingarian2011) while Fig."," The results of the SAURON data analysis by \citet{SilchenkoChilingarian2011}
 agree perfectly with the present results, derived by fitting the long-slit SCORPIO spectra: they have obtained $=+0.2$ and the age of 1.2 Gyr for the nucleus, and $=-0.33$ and the age of 2.2 Gyr for the bulge \citep{SilchenkoChilingarian2011}
 while Fig."
" 2 implies [Z/H]=+0.2 and T=1 Gyr in the nucleus and [Z/H]——0.3+0.1, and the age of 2.5 Gyr for the bulge."," \ref{fig_longslit_results} implies $=+0.2$ and $T=1$ Gyr in the nucleus and $=-0.3\pm 0.1$, and the age of 2.5 Gyr for the bulge."
" Therefore, the data of several spectrographs give the evidence for the young age and enhanced metallicity in the stellar nucleus of NGC 7743; we conclude that there is a signature of a recent nuclear star formation burst."," Therefore, the data of several spectrographs give the evidence for the young age and enhanced metallicity in the stellar nucleus of NGC 7743; we conclude that there is a signature of a recent nuclear star formation burst."
 In Fig., In Fig.
" 2 the stellar population parameter profile in PA= 268°, almost along the bar, is quite outstanding among the others."," \ref{fig_longslit_results} the stellar population parameter profile in $PA=268\degr$ , almost along the bar, is quite outstanding among the others."
" The semiaxis of the bar is (2.93.4 kpc), according to the estimates by"," The semiaxis of the bar is $31\arcsec -37\arcsec$ (2.9–3.4 kpc), according to the estimates by"
Recently. the ILIZS.5. collaboration carried out a search [or Very High Energy (VIIE. E> 100 Gev) 5-rayv emission from Galactic voung stellar clusters and found positive detections in the direction of Westerlund 1 (Well) ancl Westerlund 2 (Wd2) (Ohm et al.,"Recently, the H.E.S.S. collaboration carried out a search for Very High Energy (VHE, $>$ 100 Gev) $\gamma$ -ray emission from Galactic young stellar clusters and found positive detections in the direction of Westerlund 1 (Wd1) and Westerlund 2 (Wd2) (Ohm et al."
 2009: Ohm09 henceforth)., 2009; Ohm09 henceforth).
 The motivation behind these searches comes from the fact that the stellar clusters are potential acceleration sites of VILE particles. since thev host a," The motivation behind these searches comes from the fact that the stellar clusters are potential acceleration sites of VHE particles, since they host a"
and re-radition.,and re-radition.
 But in any case. the molecular absorption is such pronounced. especially for the coldest very low-mass T-dwarfs. that it is easily detectable even in low signal-to-noise spectra CCuby et 11999).," But in any case, the molecular absorption is such pronounced, especially for the coldest very low-mass T-dwarfs, that it is easily detectable even in low signal-to-noise spectra Cuby et 1999)."
 It further seems that these strong absorption bands are also present in young (< IMvr). low gravity. low-mass objects.," It further seems that these strong absorption bands are also present in young $\leq 1$ Myr), low gravity, low-mass objects."
 For instance. distinct water vapor absorption bands have been observed in near-infrared spectra of brown dwarfs and planetary mass objects in the very young p Oph and Orion Trapezium cluster (Wilking. Greene Meyer 1999; Lucas et al.," For instance, distinct water vapor absorption bands have been observed in near-infrared spectra of brown dwarfs and planetary mass objects in the very young $\rho$ Oph and Orion Trapezium cluster (Wilking, Greene Meyer 1999; Lucas et al."
 2001: Lucas et 22006)., 2001; Lucas et 2006).
" None of the features described above are present in our spectrum of TMR-1C. In Figure 5. we show for comparison the synthetic spectrum of a young MMyj,, object having Tyy=1900 KK and log 223.5.", None of the features described above are present in our spectrum of TMR-1C. In Figure \ref{fittmr1c} we show for comparison the synthetic spectrum of a young $_{\rm Jup}$ object having $_{\rm eff}=1900$ K and log g=3.5.
 This model spectrum. generated by the Lyon group based on the Ames-DUSTY stellar evolutionary model (Chabrier et 22000. Allard et 22001). is clearly different from the spectrum of TMR-IC. Furthermore. as implied by Figure 6.. any young planetary mass object should be brighter than our measured upper detection limits at mid-infrared wavelengths.," This model spectrum, generated by the Lyon group based on the Ames-DUSTY stellar evolutionary model (Chabrier et 2000, Allard et 2001), is clearly different from the spectrum of TMR-1C. Furthermore, as implied by Figure \ref{sed}, any young planetary mass object should be brighter than our measured upper detection limits at mid-infrared wavelengths."
 An alternative solution may be that TMR-1C is physically associated with TMR-1AB. but instead of being an extremely low-mass object. it may be a slightly higher mass object having its light significantly dimmed by a surrounding circumstellar disk seen edge-on.," An alternative solution may be that TMR-1C is physically associated with TMR-1AB, but instead of being an extremely low-mass object, it may be a slightly higher mass object having its light significantly dimmed by a surrounding circumstellar disk seen edge-on."
 In this case. the emergent flux in the is dominated by scattered light from the disk surface. TMR-IC.," In this case, the emergent flux in the near-IR is dominated by scattered light from the disk surface. TMR-1C,"
 if occulted by a disk. would appear subluminous for its spectral type and could exhibit an observed near-IR spectral energy distribution different from the object’s intrinsic colours.," if occulted by a disk, would appear subluminous for its spectral type and could exhibit an observed near-IR spectral energy distribution different from the object's intrinsic colours."
 This hypothesis for the nature of TMR-1C was actually already mentioned in an earlier paper by Hartmann et (1999)., This hypothesis for the nature of TMR-1C was actually already mentioned in an earlier paper by Hartmann et (1999).
 Direct and indirect evidence for the presence of disks around young substellar objects have been found by several authors NNatta et 2002. Luhman et 22007).," Direct and indirect evidence for the presence of disks around young substellar objects have been found by several authors Natta et 2002, Luhman et 2007)."
 Such studies show that disks around young brown dwarfs have properties which are. in general. similar to those derived for circumstellar disks of low-mass stars.," Such studies show that disks around young brown dwarfs have properties which are, in general, similar to those derived for circumstellar disks of low-mass stars."
 Also. low-resolution near-IR spectra of objects with edge-on disks often appear featureless BBrandner et 22000. Scholz et 22008). as it is the case for the TMR-IC spectrum.," Also, low-resolution near-IR spectra of objects with edge-on disks often appear featureless Brandner et 2000, Scholz et 2008), as it is the case for the TMR-1C spectrum."
 To explore this possibility. we used the SED model fitting by Robitaille et (2007) that has been developed to analyze the SEDs of voung stellar objects.," To explore this possibility, we used the SED model fitting by Robitaille et (2007) that has been developed to analyze the SEDs of young stellar objects."
 A particular advantage of this tool is that it allows the inclusion of upper flux limits., A particular advantage of this tool is that it allows the inclusion of upper flux limits.
 In order to limit the solutions to models that are physically relevant for the case of TMR-1C. we postulate that the sources distance must be within 117-157pc. the range in distance considered for the Taurus molecular cloud (Torres et 22007). and that its age must be x1 MMyr.," In order to limit the solutions to models that are physically relevant for the case of TMR-1C, we postulate that the source's distance must be within 117-157pc, the range in distance considered for the Taurus molecular cloud (Torres et 2007), and that its age must be $\lesssim$ Myr."
 The model that then agrees best with the observed SED of TMR-IC is shown in Figure 7.., The model that then agrees best with the observed SED of TMR-1C is shown in Figure \ref{robi_model}.
" The model SED shows a steep decline beyond 2.,m. and a strong excess at (|.»IOmic."," The model SED shows a steep decline beyond $\mu$ m, and a strong excess at $\lambda >$ 10mic."
" The underlying model is that of a central source of ~0.13M. at a distance of ppc. and with the following parameters: age of the source ~4x IOyr. A,=I8mmag. T.4;23000K. inclination angle of the circumstellar disk. /=87°. inner disk radius. 7;=1.5 AAU. and disk outer radius. 7,=94 AAU."," The underlying model is that of a central source of $\sim 0.13{\rm M}_{\odot}$ at a distance of pc, and with the following parameters: age of the source $\sim 4\times10^5$ yr, $_{\rm v}=18$ mag, $_{\rm eff}$ =3000K, inclination angle of the circumstellar disk, $i=87^{\circ}$, inner disk radius, $r_i=1.8$ AU, and disk outer radius, $r_o=94$ AU."
 The predicted size of the disk seems large. given that TMR- appears unresolved in all the HST and ISAAC images.," The predicted size of the disk seems large, given that TMR-1C appears unresolved in all the HST and ISAAC images."
" An edge-on disk as large as o,294 AAU and its associated bipolar scattered light structures should have appeared resolved in the high spatial resolution HST images (FWHM~0715—072) and possibly also in the best ISAAC images.", An edge-on disk as large as $r_o=94$ AU and its associated bipolar scattered light structures should have appeared resolved in the high spatial resolution HST images $\sim 0\farcs15 - 0\farcs2$ ) and possibly also in the best ISAAC images.
 The fact that TMR-IC appears as a point-source implies that the actual disk must be smaller than ~20 AU., The fact that TMR-1C appears as a point-source implies that the actual disk must be smaller than $\sim$ 20 AU.
 A decreasing outer disk size will mostly affect TMR-1C's spectral energy distribution longward of IOjm. in a regime that is unconstrained due to the lack of observations.," A decreasing outer disk size will mostly affect TMR-1C's spectral energy distribution longward of $\mu$ m, in a regime that is unconstrained due to the lack of observations."
 In a picture where TMR-1C had been ejected in adynamical interaction within the small proto-stellar cluster of TMR-1. it would be expected that any primordial disk of IC has been truncated SSec. 4)).," In a picture where TMR-1C had been ejected in a dynamical interaction within the small proto-stellar cluster of TMR-1, it would be expected that any primordial disk of TMR-1C has been truncated Sec. \ref{formation_theo}) )."
 We note that if one allows models for sources at slightly larger distances (up to ~170pe) or slightly older objects (up to ~1.2x 10°yr). one finds that also objects with ~0.2—0.4M.« having disk sizes as small as AAU provide reasonable fits to the SED of TMR-1C. But since the SED is unconstrained beyond Sem. it is currently not possible to distinguish any further between the model solutions.," We note that if one allows models for sources at slightly larger distances (up to $\sim$ 170pc) or slightly older objects (up to $\sim1.2\times10^6$ yr), one finds that also objects with $\sim0.2-0.4{\rm M}_{\odot}$ having disk sizes as small as AU provide reasonable fits to the SED of TMR-1C. But since the SED is unconstrained beyond $\mu$ m, it is currently not possible to distinguish any further between the model solutions."
 SED models for objects with central source masses <0.1M. are not available., SED models for objects with central source masses $<0.1{\rm M}_{\odot}$ are not available.
 However. an important point we wish to emphasize. is that all models that provide a good fit imply a circumstellar disk with high inclination angle ofi>80°.," However, an important point we wish to emphasize, is that all models that provide a good fit imply a circumstellar disk with high inclination angle of $i>80^{\circ}$."
 Observational evidence for an edge-on disk surrounding TMR-IC would be an infrared excess emission att> IOum. Unfortunately. the spatial resolution of available observations at mid to far-infrared wavelengths are too low to disentangle a potential emission from TMR-IC from the emission of TMR-," Observational evidence for an edge-on disk surrounding TMR-1C would be an infrared excess emission at $\lambda>10\mu$ m. Unfortunately, the spatial resolution of available observations at mid to far-infrared wavelengths are too low to disentangle a potential emission from TMR-1C from the emission of TMR-1AB."
 Other evidence. as discussed above. could be provided by near-infrared images with high sensitivity and a spatial," Other evidence, as discussed above, could be provided by near-infrared images with high sensitivity and a spatial"
 the fact that magnetic fields 1n. galaxy clusters are too Despiteweak to be important. they can play a fundamental role in the mechanicallydynamical of the dilute channelling the of heat. stabilit,"Despite the fact that magnetic fields in galaxy clusters are too weak to be mechanically important, they can play a fundamental role in the dynamical stability of the dilute gas by channelling the transport of heat."
yThecolliston-gasless characterby of the hot intra-clustertransport medium (ICM). which is characterized stable gradients to genericallySchwarzschild’s criterionby (Piffarettientropyetal.2005:accordingCavagnoloetal. 2009).. enables the action of magnetic instabilities that are sensitive to (Balbus2000. 2004)..," Thecollision-less character of the hot intra-cluster medium (ICM), which is generically characterized by stable entropy gradients according to Schwarzschild's criterion \citep{2005A&A...433..101P,
2009ApJS..182...12C}, enables the action of magnetic instabilities that are sensitive to temperature gradients \citep{2000ApJ...534..420B, 2004ApJ...616..857B}."
 In particular. the temperaturemagneto-thermal gradientsinstability (MTI) when magnetic field lines are to a operatesgradient to the gravitational field orthogonal(Balbus 2001). temperaturewhereas the parallelheat-flux-driven instability (HBI) acts when magnetic field lines are buoyancyparallel to a to the gravitational fielc temperature(Quataert20068)... gradien," In particular, the magneto-thermal instability (MTI) operates when magnetic field lines are orthogonal to a temperature gradient parallel to the gravitational field \citep{2001ApJ...562..909B}, whereas the heat-flux-driven buoyancy instability (HBI) acts when magnetic field lines are parallel to a temperature gradient anti-parallel to the gravitational field \citep{2008ApJ...673..758Q}."
tNumerical anti-parallelsimulations suggest that the non-linear evolution of these instabilities leads to magnetic field that tend to suppress the of the to into configurationsthe free the abilitybackgrounc plasma tap (Parrishenergyetal.supplied2009;byBogdanoviétemperature2009:Ruszkowsktgradi," Numerical simulations suggest that the non-linear evolution of these instabilities leads to magnetic field configurations that tend to suppress the ability of the plasma to tap into the free energy supplied by the background temperature gradient \citep{2009ApJ...703...96P, 2009ApJ...704..211B,
2010ApJ...713.1332R}. ."
ent&Oh 2010). These fielc can nevertheless support overstable magneticgravity modes configurations(Balbus&Reynolds2010)., These magnetic field configurations can nevertheless support overstable gravity modes \citep{2010ApJ...720L..97B}.
. While the of thermal instabilities that render homogeneous. dilute landscapeplasmas unstable has been well explorec (Kunz2011).. and even extended to account for the effects of cosmie (Chandran&Dennis2006).. little is know1 about the rayseffects that verycan have on the of the dilute ICM.," While the landscape of thermal instabilities that render homogeneous, dilute plasmas unstable has been well explored \citep{2011MNRAS.417..602K}, and even extended to account for the effects of cosmic rays \citep{2006ApJ...642..140C}, very little is known about the effects that composition gradients can have on the stability of the dilute ICM."
 If composition gradientsfields do not prevent the stabilityefficient diffusion of tons magnetic(Narayan&Medvedev2001:Chuzhoy&Nusser2003:Loeb2004) then the gradients in mean molecular can be as as the gradients in temperature (see weightSection + and Qinimportant&Wu2000;etPengal. 2011)) Nagaiand another Shtykovskiysource of free energy to feed instabilities.," If magnetic fields do not prevent the efficient diffusion of ions \citep{2001ApJ...562L.129N, 2003MNRAS.342L...5C,
2004MNRAS.349L..13C} then the gradients in mean molecular weight can be as important as the gradients in temperature (see Section \ref{sec:discussion} and \citealt{2000ApJ...529L...1Q,
2009ApJ...693..839P, 2010MNRAS.401.1360S, 2011A&A...533A...6B}) ) and provide another source of free energy to feed instabilities."
 In order to provideobtain a more of the of the ICM. it 15 thus completeimportant to picturerelax the stabilityassumption of a propertieshomogeneous plasma.," In order to obtain a more complete picture of the stability properties of the ICM, it is thus important to relax the assumption of a homogeneous plasma."
 As a first towards the role. of composition gradientsstep in dilute plasmas. understandingwhere magnetic fields a key role by the conduction of heat and the playdiffusion of tons. we channelling the of modes for which heatconduction ts slow investigatecompared to the stabilitydynamical timescale involved.," As a first step towards understanding the role of composition gradients in dilute plasmas, where magnetic fields play a key role by channelling the conduction of heat and the diffusion of ions, we investigate the stability of modes for which heatconduction is slow compared to the dynamical timescale involved."
 These modes contain. among others. the overstable g-modes studied by Balbus&Reynolds(2010).," These modes contain, among others, the overstable $g$ -modes studied by \citet{2010ApJ...720L..97B}."
 In order to highlight the phenomena that emerge when composition physical are accounted for. we focus our attention on a dilutegradients binary mixture (e.g.. hydrogen and in a fixed gravitational fielddescribed bedot((pe))—0.. dt [ρι right))+g.. bbB)) .. ΕΠ 5o» .," In order to highlight the physical phenomena that emerge when composition gradients are accounted for, we focus our attention on a dilute binary mixture (e.g., hydrogen and in a fixed gravitational fielddescribed )=0 , + - ) + , ) , )= - ,"
 In order to highlight the phenomena that emerge when composition physical are accounted for. we focus our attention on a dilutegradients binary mixture (e.g.. hydrogen and in a fixed gravitational fielddescribed bedot((pe))—0.. dt [ρι right))+g.. bbB)) .. ΕΠ 5o» ..," In order to highlight the physical phenomena that emerge when composition gradients are accounted for, we focus our attention on a dilute binary mixture (e.g., hydrogen and in a fixed gravitational fielddescribed )=0 , + - ) + , ) , )= - ,"
rom the clusters.,from the clusters.
 Here we modeled the kinematic evolution of XItBs from the star clusters., Here we modeled the kinematic evolution of XRBs from the star clusters.
 The calculated. results are oresented. below., The calculated results are presented below.
 As stated before we constructed S models to investigate row the final results are. influenced by the. adopted xwameters., As stated before we constructed 8 models to investigate how the final results are influenced by the adopted parameters.
 Specifically the input parameters in our control moclel (1.0... model MI) are SEII—20 Myr. maa=190 uus toa=O. ace=1.0. and the Ixroupa IME.," Specifically the input parameters in our control model (i.e., model M1) are $=20$ Myr, $\sigma_{\rm kick}=190$ $^{-1}$, $\alpha=0$, $\alpha_{\rm CE}=1.0$, and the Kroupa IMF."
 In other models we change only one. parameter cach time. and. the model parameters are listed in Table 1.," In other models we change only one parameter each time, and the model parameters are listed in Table 1."
 Figures 1 and 2 show he simulated cumulative distribution of the X-ray source displacements (top: Total: middle: NS-NRBs: bottom: ARBs) at the age of 20 Myr for models MI-MA and M5- respectively.," Figures 1 and 2 show the simulated cumulative distribution of the X-ray source displacements (top: Total; middle: NS-XRBs; bottom: BH-XRBs) at the age of 20 Myr for models M1-M4 and M5-M8, respectively."
" Note that we only select sources in the luminosity range 10""«Ly<107 in order to compare with Ixaaretetal.(2004).", Note that we only select sources in the luminosity range $10^{36}<L_X<10^{38}$ $^{-1}$ in order to compare with \citet{kaaret04}.
.. We also normalize each histogram by the total number of NltBs within 1 kpe ofa star cluster for each galaxy., We also normalize each histogram by the total number of XRBs within 1 kpc of a star cluster for each galaxy.
 We find that the results are in general agreement with Ixaaretetal.(2004.Fig.2 ).., We find that the results are in general agreement with \citet[][Fig.~2 ]{kaaret04}. .
 Figures 3 and 4 show the moceled. distributions of the X-rav luminosities (Lx) at different displacement (2?) from. the star cluster for models MI-MA and M5-MS. respectively.," Figures 3 and 4 show the modeled distributions of the X-ray luminosities $L_{\rm X}$ ) at different displacement $R$ ) from the star cluster for models M1-M4 and M5-M8, respectively."
 The top. middle. and bottom panels are for all XSRBs. NS-ARBs. and BU-NRBs. respectively.," The top, middle, and bottom panels are for all XRBs, NS-XRBs, and BH-XRBs, respectively."
 The color bar represents the normalized. number ratio of NRBs in the £&Lx plane., The color bar represents the normalized number ratio of XRBs in the $R-L_{\rm X}$ plane.
 Note that NS-XItDs have relatively lower maximum Iuminosities than BIENBs in all of the mocdels. due to a lower Eddington accretion rate limit.," Note that NS-XRBs have relatively lower maximum luminosities than BH-XRBs in all of the models, due to a lower Eddington accretion rate limit."
 The predicted. Ls vs. AC relations in all the models are roughly: compatible with the observations. but ciffercnecs also exist.," The predicted $L_{\rm X}$ vs. $R$ relations in all the models are roughly compatible with the observations, but differences also exist."
" For models ALL-ALL. there is a scarcity. of bigh-luminosity (Lx.>107 +) sources in the 30100 pe region. while there is an overabundance of very Luminous sources (Lx710?"" 1j in the same region for models M5-M8."," For models M1-M4, there is a scarcity of high-luminosity $L_X>10^{38}$ $^{-1}$ ) sources in the $30-100$ pc region, while there is an overabundance of very luminous sources $L_X>10^{39}$ $^{-1}$ ) in the same region for models M5-M8."
 In addition. the correlations originate from cdilferent. stellar populations in dilferent. models.," In addition, the correlations originate from different stellar populations in different models."
 On one hand. models MI-MA have similar Lx vs. R correlations which are constructed by both BII-ARBs and NS-XRBs.," On one hand, models M1-M4 have similar $L_{\rm X}$ vs. $R$ correlations which are constructed by both BH-XRBs and NS-XRBs."
 Notethat BUENDs dominate at the hieh-luminosity (Ly71077 1). smalbollset. (10.<Ho300 pe) region. (named as region A). while NS-XItDs dominate at the low-Iuminositv (Lx«1077 1). small-olfset (10«£?300 pe) region (named as regions D) and large-olfset (300.«21000. pc) region (named as region ©).," Notethat BH-XRBs dominate at the high-luminosity $L_X>10^{38}$ $^{-1}$ ), small-offset $10<R<300$ pc) region (named as region A), while NS-XRBs dominate at the low-luminosity $L_X< 10^{38}$ $^{-1}$ ), small-offset $10<R<300$ pc) region (named as regions B) and large-offset $300<R<1000$ pc) region (named as region C)."
 On the other hand. for models M5-M8S. the correlations mainly result. from BII-XIUDs instead.," On the other hand, for models M5-M8, the correlations mainly result from BH-XRBs instead."
 Lt is interesting to note that there is a lat-roofed edge in the 10100 pc region. which is not caused by the Exddington luminosity limit for BILEXRBs.," It is interesting to note that there is a flat-roofed edge in the $10-100$ pc region, which is not caused by the Eddington luminosity limit for BH-XRBs."
" In order to explore the nature of the NIUDs in dilferent regions. we need to examine their observational properties (Le. current mass Als and spectral type o ""the donor star. orbital period {μι and system. velocity ¢istribution)."," In order to explore the nature of the XRBs in different regions, we need to examine their observational properties (i.e., current mass $M_2$ and spectral type of the donor star, orbital period $P_{\rm orb}$, and system velocity distribution)."
 We present the LyMo (left). LxPus (nicdle). and aAM» (right) distributions of NiBs in the 1)«HR300 pc region for models MI-M4A and M5-MS in Figures 5 and. 6. respectively.," We present the $L_{\rm X}-M_2$ (left), $L_{\rm X}-P_{\rm orb}$ (middle), and $P_{\rm orb}-M_2$ (right) distributions of XRBs in the $10<R<300$ pc region for models M1-M4 and M5-M8 in Figures 5 and 6, respectively."
 Figures 7 and S show the cisributions of the same parameters in the region of 300«£i1000 pc., Figures 7 and 8 show the distributions of the same parameters in the region of $300<R<1000$ pc.
 The detailed: source tvpes in regions A. B au C for different niocdels are also listed in Tables 2-4. respecively.," The detailed source types in regions A, B and C for different models are also listed in Tables 2-4, respectively."
 Figures 5r and 6 show that. the NiBs in region A are in short orbital periods ( 1 hr). with relatively low-mass ( 1-3 AL.) donors for all models.," Figures 5 and 6 show that, the XRBs in region A are in short orbital periods $\sim$ 1 hr), with relatively low-mass $\sim$ 1-3 $M_{\odot}$ ) donors for all models."
" They are mainly RLOF DII-NItDs with Helium main-sequence (HeMS) companions (sce ""able 2).", They are mainly RLOF BH-XRBs with Helium main-sequence (HeMS) companions (see Table 2).
 Phe binary velocities are ~1020 1 for models MI-\LE and ~3000100 ! for models M5-M8.," The binary velocities are $\sim 10-20$ $^{-1}$ for models M1-M4, and $\sim 30-100$ $^{-1}$ for models M5-M8."
 Considering that they have similar evolutionary timescales (sce Fig., Considering that they have similar evolutionary timescales (see Fig.
 9 below). the dillerence in the velocities explains the cillerent maximum olfsets of these high-Iuminosity. sources in cillerent mocels.," 9 below), the difference in the velocities explains the different maximum offsets of these high-luminosity sources in different models."
 Phe XRBs in region D have larger orbital »iods and companion masses., The XRBs in region B have larger orbital periods and companion masses.
 For example. the orbital »eriods are ~120 hr for models AIL ane M3-MS.. and can reach ~10 hr for model M2: the companion masses are ~1]—£A. lor models MI anc M4. ~1.20M. for models M5-Ms8. and can reach ~GOAL. for models M2 and ALS.," For example, the orbital periods are $\sim 1-20$ hr for models M1 and M3-M8, and can reach $\sim 10^3$ hr for model M2; the companion masses are $\sim1-4 M_{\odot}$ for models M1 and M4, $\sim 1-20 M_{\odot}$ for models M5-M8, and can reach $\sim 60
M_{\odot}$ for models M2 and M3."
 Table 3 indicates that the majority (> SO) of these sources are NS-N1Bs., Table 3 indicates that the majority $>$ $\%$ ) of these sources are NS-XRBs.
 They mainly have HeMS companions ( 50% in model MA to ~ 85% in model MIT). with mass ransfer either through ΟΙ or by wind capture.," They mainly have HeMS companions $\sim$ $\%$ in model M4 to $\sim$ $\%$ in model M7), with mass transfer either through RLOF or by wind capture."
 Figures 7 and S show that. the NRBs in region € all have relatively ow-nmass ( 1-4 AL.) companions.," Figures 7 and 8 show that, the XRBs in region C all have relatively low-mass $\sim$ 1-4 $M_{\odot}$ ) companions."
 Their orbital periods are around. several hours for all models., Their orbital periods are around several hours for all models.
 The typical velocities are ~150300 kinsο larger than those of high-unminosity sotrces.," The typical velocities are $\sim 150-300$ $^{-1}$, much larger than those of high-luminosity sources."
 We tentatively conclude that the dillerence. in. the uminosities for the sources in regions A.B and € is caused w dillerent types of the compact objects (BLENRBs vs. NRBs). of the donors (HeMS star vs. MS star). and of the accretion modes (RLOEF vs. wind-capture).," We tentatively conclude that the difference in the luminosities for the sources in regions A, B and C is caused by different types of the compact objects (BH-XRBs vs. NS-XRBs), of the donors (HeMS star vs. MS star), and of the accretion modes (RLOF vs. wind-capture)."
 Sources in region A are mainly. DII-XIDs with IRLOFE. mass ransfer from an Ηλ companion. hence reaching very high Iuminosities.," Sources in region A are mainly BH-XRBs with RLOF mass transfer from an HeMS companion, hence reaching very high luminosities."
 Sources in region B have similar properics as those in region A. but are mainly NS-NRBs. with relatively lower luminosities.," Sources in region B have similar properties as those in region A, but are mainly NS-XRBs, with relatively lower luminosities."
 Sources in region. C are either NS-XRBs (in models MI-MA). or contain a portion of BU-NRBs with MS companions. so they also cannot reach luminosities as high as those in region A. The displacement of the SRBs from the star clusters depends on their velocities at. the moment of the SN explosion ancl the delay time from the SN to the beginning of RLOEF.," Sources in region C are either NS-XRBs (in models M1-M4), or contain a portion of BH-XRBs with MS companions, so they also cannot reach luminosities as high as those in region A. The displacement of the XRBs from the star clusters depends on their velocities at the moment of the SN explosion and the delay time from the SN to the beginning of RLOF."
 In Fig., In Fig.
 9 we present the distribution of the delay imes lor sources in regions X (clash-clot-clotted line). B (solid line) ancl € (dotted line). respectivelv.," 9 we present the distribution of the delay times for sources in regions A (dash-dot-dotted line), B (solid line) and C (dotted line), respectively."
 We normalize cach the histograms by the total number of X-ray sources in each region., We normalize each the histograms by the total number of X-ray sources in each region.
 They. generally range from 1 to 8 Myr. but the imes are a bit shorter in mocdoels M1I-MA than in M5-MS.," They generally range from 1 to 8 Myr, but the times are a bit shorter in models M1-M4 than in M5-M8."
 Ln he latter they peak at ~5 Myr., In the latter they peak at $\sim 5$ Myr.
 Note that the Ly vs. # relation has clillerent origin tween models MI-MA (ace= 1.0) and models. M5-MS (ace= 3.0)., Note that the $L_{\rm X}$ vs. $R$ relation has different origin between models M1-M4 $\alpha_{\rm CE}=1.0$ ) and models M5-M8 $\alpha_{\rm CE}=3.0$ ).
 In the ace=LO cases. it is constructed. by noth high-luminositv (Ly>107 +) BU-XRBs and ow-luminosity (Ly<107 +) NS-NRBs. while in the Ace=3.0 cases. it results from BLI-XNRBs instead. so we need to discuss them separately.," In the $\alpha_{\rm CE}=1.0$ cases, it is constructed by both high-luminosity $L_X>10^{38}$ $^{-1}$ ) BH-XRBs and low-luminosity $L_X<10^{38}$ $^{-1}$ ) NS-XRBs, while in the $\alpha_{\rm CE}=3.0$ cases, it results from BH-XRBs instead, so we need to discuss them separately."
 In models MI-MA. luminous DII-NItDs are constrained within 300 pe because of their ow velocities. while NS-X1ItDs. which are relatively ciümamoerand faster. can move to region. C within ~10 Myr. (see ‘Table 4).," In models M1-M4, luminous BH-XRBs are constrained within $\sim 300$ pc because of their low velocities, while NS-XRBs, which are relatively dimmerand faster, can move to region C within $\sim 10$ Myr (see Table 4)."
 In models M5-MS8 . à number of RLOE. DII-XIUDs with MS companions are produced. significantly: different [rom mocdels MI-MA.," In models M5-M8 , a number of RLOF BH-XRBs with MS companions are produced, significantly different from models M1-M4."
 These NRBs. occupying region C. have relatively lower Iuminosities and higher velocities compared to the BIENDs with LIeMS donors in region A. 1n order to understand. the dillerence in the binary velocities in dillerent models. we examine the distribution of the companion masses Moy. and orbital periods Pis ," These XRBs, occupying region C, have relatively lower luminosities and higher velocities compared to the BH-XRBs with HeMS donors in region A. In order to understand the difference in the binary velocities in different models, we examine the distribution of the companion masses $M_{\rm 2, SNe}$ , and orbital periods $P_{\rm orb, SNe}$ "
number of absorbers seen along sight lines to quasars in Iux-limited catalogues.,number of absorbers seen along sight lines to quasars in flux-limited catalogues.
 In general. most authors agree that the differences between ttoward quasar and CRB sight lines cannot be due to a singleLl cllect. but. clusty absorbers have not been thought to be a significant factor.," In general, most authors agree that the differences between toward quasar and GRB sight lines cannot be due to a single effect, but dusty absorbers have not been thought to be a significant factor."
 1n light of our findings on the existence of a significant population of dusty aabsorbers. we can provide a more definitive determination of the effect of dust on the GItB/quasar sight line discrepancy.," In light of our findings on the existence of a significant population of dusty absorbers, we can provide a more definitive determination of the effect of dust on the GRB/quasar sight line discrepancy."
 Using a sample of CRB sight lines with intervening absorbers we can assess how this sample would sulfer. from. incompleteness if the sigh lines were instead. illuminated by quasars., Using a sample of GRB sight lines with intervening absorbers we can assess how this sample would suffer from incompleteness if the sight lines were instead illuminated by quasars.
 We assume tha the GRB sample is unbiased and. probes all intervening aabsorbers up to a moderate absorber £(2VW) of 0.2 mag., We assume that the GRB sample is unbiased and probes all intervening absorbers up to a moderate absorber $E(B-V)$ of 0.2 mag.
 Given an EW ane redshift’ distribution of aabsorbers in GRBs similar to our quasar sample (Pig. 17)).lt ," Given an EW and redshift distribution of absorbers in GRBs similar to our quasar sample (Fig. \ref{cap:cordists}) ),"
we would expect the GIUD sight lines to show the missing pper cent of absorbers (Section ?27))., we would expect the GRB sight lines to show the missing per cent of absorbers (Section \ref{mgstatssec}) ).
 However. in genera GRB spectra contain absorbers with cillerent redshift anc EW distribution to the absorbers in quasar spectra.," However, in general GRB spectra contain absorbers with different redshift and EW distribution to the absorbers in quasar spectra."
 The observed discrepancy in number density due to dust will be sensitive to such dilferences., The observed discrepancy in number density due to dust will be sensitive to such differences.
 Two GIU absorber samples in the literature are the ? sample and the ? sample.," Two GRB absorber samples in the literature are the \citet{2009A&A...503..771V} sample and the \citet{2009ApJS..185..526F}
 sample."
 Phe ? sample contains 22 strong systems (EN1 AA) identified. using high-resolution UVES sample and a series of high- and. low-resolution GIU alterglow spectra from the literature., The \citeauthor{2009A&A...503..771V} sample contains 22 strong systems $>1$ ) identified using high-resolution UVES sample and a series of high- and low-resolution GRB afterglow spectra from the literature.
 Phe ?7/— sample. contains 15 strong systems present in the follow-up. spectra of 77 optical afterglows of detected CRBs., The \citeauthor{2009ApJS..185..526F} sample contains 15 strong systems present in the follow-up spectra of 77 optical afterglows of detected GRBs.
 Phe redshift distributions of the two samples are shown in Fig. I8s-- (feff))., The redshift distributions of the two samples are shown in Fig. \ref{cap:grbzewdists}- ).
 Phe redshift distributions for the two GRB samples both exhibit a larger fraction of absorbers at. hieh recshilt compared with the quasar absorber redshift. clistribution., The redshift distributions for the two GRB samples both exhibit a larger fraction of absorbers at high redshift compared with the quasar absorber redshift distribution.
 Fie. [s--(righ!)), Fig. \ref{cap:grbzewdists}- )
 shows the degree to which both samples suller from EW incompleteness bv comparing the EW distributions with an exponential (2) of the form where WV is the exponential scale. factor., shows the degree to which both samples suffer from EW incompleteness by comparing the EW distributions with an exponential \citep{2005ApJ...628..637N} of the form where $W^{*}$ is the exponential scale factor.
 A fit to the Corrected. quasar absorber EW distribution (Fig. 17)), A fit to the corrected quasar absorber EW distribution (Fig. \ref{cap:cordists}) )
 gives a value of M —0.746-0.008., gives a value of $W^{*}$ $\pm$ 0.008.
 Scaling the corrected: quasar IN distribution using the redshift paths appropriate to the two CltD-derived: samples results in the predicted number versus EW distributions lor the ? and 2. GRB sunples shown in Fig. I8--(righ! ))., Scaling the corrected quasar EW distribution using the redshift paths appropriate to the two GRB-derived samples results in the predicted number versus EW distributions for the \citeauthor{2009A&A...503..771V} and \citeauthor{2009ApJS..185..526F} GRB samples shown in Fig. \ref{cap:grbzewdists}- ).
 Both samples are expected to be incomplete at low LEW but the scaled distribution functions accurately describe the ?. saniple with Z2.0AA. and the ?) sample is consistent over the interval 22.AA.," Both samples are expected to be incomplete at low EW but the scaled distribution functions accurately describe the \citeauthor{2009A&A...503..771V} sample with $\ga$, and the \citeauthor{2009ApJS..185..526F} sample is consistent over the interval $\simeq$."
 Our Uux-limited quasar-derived aabsorber distribution has no information regarding hieh-EW systems with 427755.0 Gvhere the associated large LOBV) essentially removesA nearly all such svstems from the SDSS quasar sample)., Our flux-limited quasar-derived absorber distribution has no information regarding high-EW systems with $W_{0}^{\lambda2796}$$>$ (where the associated large $E(B-V)$ essentially removes nearly all such systems from the SDSS quasar sample).
 Lhe extension of the distribution of aabsorbers bevoud ΛΣ out to SLOA is clear anc stronely suggests that the absorber dust-induced differences between: quasar-derivecl and CGltB-derived. samples. wil be even larger than calculated using our ΑΝ] parametrization., The extension of the distribution of absorbers beyond $W_{0}^{\lambda2796}$ =5.0 out to $\simeq$ is clear and strongly suggests that the absorber dust-induced differences between quasar-derived and GRB-derived samples will be even larger than calculated using our $N(W_{0}^{\lambda2796})$ parametrization.
 We can calculate. the fraction of absorbers tha would) be missed. in. corresponding Lux-limitecd quasar sample bv estimating an £02V) according to equation 7 and assigning a weighting according to the cus obscuration bias/signal-to-noise ratio ellects (equation 6))., We can calculate the fraction of absorbers that would be missed in corresponding flux-limited quasar sample by estimating an $E(B-V)$ according to equation \ref{eq:fullebvew} and assigning a weighting according to the dust obscuration bias/signal-to-noise ratio effects (equation \ref{eq:weightebv_old}) ).
" The predicted:ciscrepaney is obtained. by calculating the [ractraction of [alabsorbers| that are missed Mfrom a quasar-derivejul sample with an yy,here distribution corresponding to Equation LL (with W= 0.746) and redshift distributions ofthe ? ancl?)", The predicteddiscrepancy is obtained by calculating the fraction of absorbers that are missed from a quasar-derived sample with an $W_{0}^{\lambda2796}$ distribution corresponding to Equation \ref{nestorexp} (with $W^{*}=0.746$ ) and redshift distributions of the \citeauthor{2009A&A...503..771V} and \citet{2009ApJS..185..526F} .
"the zeros of Mj propagate as zero minors of Af, ancl (he origin of neutrino mixing is solely from My.",the zeros of $M_R$ propagate as zero minors of $M_\nu$ and the origin of neutrino mixing is solely from $M_R$.
 We presented (the svinmetry realization of these patlerus using a cvelic group Za., We presented the symmetry realization of these patterns using a cyclic group $Z_3$.
 It was found that class D; ancl D; predict a near maximal atmospheric mixing anele in the Init of large M..., It was found that class $B_5$ and $B_6$ predict a near maximal atmospheric mixing angle in the limit of large $M_{ee}$.
 Furthermore. (his prediction is independent of the values of the reactor and the solar mixing angles.," Furthermore, this prediction is independent of the values of the reactor and the solar mixing angles."
" The assumption of large M, is testable in the ongoing experiments for NDB decay since the rate of NDB decay is proportional to A.", The assumption of large $M_{ee}$ is testable in the ongoing experiments for NDB decay since the rate of NDB decay is proportional to $M_{ee}$.
" These experiments will either confirm or rule out a large value of M. in the next few νους,", These experiments will either confirm or rule out a large value of $M_{ee}$ in the next few years.
 The atmospheric mixing angle approaches 7/4 with the increasing value of Af..., The atmospheric mixing angle approaches $\pi/4$ with the increasing value of $M_{ee}$.
 A reactor mixing angle equal to zero is nol allowed in these textures. thus. naturally accommodating a non-zero 844 as suggested bv the recent results of the T2Ix experiment.," A reactor mixing angle equal to zero is not allowed in these textures, thus, naturally accommodating a non-zero $\theta_{13}$ as suggested by the recent results of the T2K experiment."
 Due to the relatively large value of 644 the Dirac-tvpe CP violating phase is fixed near 7/2 or 37/2 predicting almost maximal CP violation for these lexiures., Due to the relatively large value of $\theta_{13}$ the Dirac-type CP violating phase is fixed near $\pi/2$ or $3\pi/2$ predicting almost maximal CP violation for these textures.
 The research. work of ο. D. and L. S. is supported bx the University Grants Commission. Government of India. Grant No.," The research work of S. D. and L. S. is supported by the University Grants Commission, Government of India Grant No."
 3432/2008 (SR)., 34-32/2008 (SR).
 RoR. G. acknowledge the financial support provided by the Council for Scientific and Industrial Research (CSIR). Government of India.," R. R. G. acknowledge the financial support provided by the Council for Scientific and Industrial Research (CSIR), Government of India."
At first. let us consider the blazar-like(b/) emission of FR I active nuclei. which accordingly to the wnilication scheme (e.g..Urryand.Padovani1995) belong to the low-Iuminous (BL Lac) blazar subclass.,"At first, let us consider the blazar-like$bl$ ) emission of FR I active nuclei, which -- accordingly to the unification scheme \citep[e.g.,][]{urr95} – belong to the low-luminous (BL Lac) blazar subclass."
 Due to relativistic effects. this radiation is strongly beamed into the cone with an opening angle 85~l/Dy. where Ly is the nuclear jet bulk Lorentz factor.," Due to relativistic effects, this radiation is strongly beamed into the cone with an opening angle $\theta_{bl}
 \sim 1 / \Gamma_{bl}$, where $\Gamma_{bl}$ is the nuclear jet bulk Lorentz factor."
 In order to evaluate enerev density. of this emission in the kpe-scale jet rest frame. we assume that it ‘enters’ into the considered. emitting region directly [rom the jet base.," In order to evaluate energy density of this emission in the kpc-scale jet rest frame, we assume that it `enters' into the considered emitting region directly from the jet base."
 In other words. we neglect any possible misalienment. between Che pc-scale and the kpc-scale jets.," In other words, we neglect any possible misalignment between the pc-scale and the kpc-scale jets."
 The respective energy clensity of the blazar photons is then where Lj(0) is the isotropic luminosity of the blazar emission evaluated by the stationary observer located at 9=0 (Appendix D)., The respective energy density of the blazar photons is then where $L_{bl}(0)$ is the isotropic luminosity of the blazar emission evaluated by the stationary observer located at $\theta = 0$ (Appendix B).
 With the spatial scales of the kpe-seale jet emittingregion much larger than the characteristic nuclear scales (22.729 ry). the discussed source of the seed photons for (he inverse-Compton scattering should be considered as asfatzonary one.," With the spatial scales of the kpc-scale jet emittingregion much larger than the characteristic nuclear scales $R, \, r \gg r_0$ ), the discussed source of the seed photons for the inverse-Compton scattering should be considered as a one."
 A typical short time scale of blazar variability justilies this statement., A typical short time scale of blazar variability justifies this statement.
 Therefore. where 9j=[Py(1—ηcos8)| ‘and Ly is the isotropic luminosity of the considered stationary blazar source measured by the observer located αἱ the angle 8 to the jet axis.," Therefore, where $\delta_{bl} = [\Gamma_{bl} \, (1 - \beta_{bl} \, \cos \theta)]^{-1}$ and $L_{bl}$ is the isotropic luminosity of the considered stationary blazar source measured by the observer located at the angle $\theta$ to the jet axis."
" IIence. where we put. Ly,i»=Ly/10Peres|! (οἱ.equation1inCelottietal.2001)."," Hence, where we put $L_{bl, \, 42} \equiv L_{bl} / 10^{42} \, {\rm erg \, s^{-1}}$ \citep[cf. equation 1 in][]{cel01}."
. Typically. for the high energy peaked DL Lacs (IIBLs) the observed svnchrotron luminosity is 10!—10eres !. while lor the low energv peaked BL Lacs (LDLs) it is in the range 10—10eres! (Fossatietal.1998).," Typically, for the high energy peaked BL Lacs (HBLs) the observed synchrotron luminosity is $10^{44} - 10^{46} \, {\rm erg \, s^{-1}}$ while for the low energy peaked BL Lacs (LBLs) it is in the range $10^{45} - 10^{47} \, {\rm erg
 \, s^{-1}}$ \citep{fos98}."
. An another subclass of blazar sources — the flat spectrum radio quasars (FSRQs) have even higher observed. luminosities., An another subclass of blazar sources – the flat spectrum radio quasars (FSRQs) – have even higher observed luminosities.
 However. in a framework of the unification scheme FSRQs correspond to the FR 1 sources.," However, in a framework of the unification scheme FSRQs correspond to the FR II sources."
" The critical observed svuchrotron frequency of blazaremission. which anticorrelates with the blazar svnchrotron luminosity. is usually 10P—10"" Hz for UBLs and 107—10"" Hz for LBLs 1993).."," The critical observed synchrotron frequency of blazaremission, which anticorrelates with the blazar synchrotron luminosity, is usually $10^{15} -
 10^{17}$ Hz for HBLs and $10^{13} - 10^{15}$ Hz for LBLs \citep[again][]{fos98}. ."
" Below. we take Ly~[vL, Ji... where the blazar svnchrotron break frequencyin the jet comoving frame is Lj4,7Vyn(0)/ T.nd iyn(0)=vyo (OTy/0u)."," Below, we take $L_{bl} \sim [\nu L_{\nu}]_{bl, \, br}$ , where the blazar synchrotron break frequencyin the jet comoving frame is $\nu'_{bl, \, br} \sim \nu_{bl, \, br}(0) / \Gamma$ ,and $\nu_{bl, \, br}(0) =
 \nu_{bl, \, br} \, (2 \Gamma_{bl} / \delta_{bl})$ ."
 We, We
slit spectra of the stars.,slit spectra of the stars.
 The grating was .112 centered on (BLMRD mode) and the detector was a Tektronix CCD TK1034 with 10247 pixels of size 24 jum. The covered. wavelength range was and the dispersion mmm. giving resolutions of 2.70+0.10 pixels or 248x0.13 for a 17.0 slit.," The grating was 12 centered on (BLMRD mode) and the detector was a Tektronix CCD TK1034 with $^{2}$ pixels of size 24 $\mu$ m. The covered wavelength range was and the dispersion $^{-1}$, giving resolutions of $2.70\pm0.10$ pixels or $2.48\pm0.13$ for a .0 slit."
 At each position we first took a short 5 min exposure followed by one or two longer 15 min exposures., At each position we first took a short 5 min exposure followed by one or two longer 15 min exposures.
 The instrument response was derived from observations ofthe, The instrument response was derived from observations ofthe
Table | contains a list of the 22 observation runs performed between June 9. 2007 and November 11. 2008.,"Table \ref{tab:obs} contains a list of the 22 observation runs performed between June 9, 2007 and November 11, 2008."
 The number of raw triggers per hour is much higher than usual for the runs of June 9. 2007 and June 7/8. 2008.," The number of raw triggers per hour is much higher than usual for the runs of June 9, 2007 and June 7/8, 2008."
 The top panel in Figure shows the number of triggers per data block for one hour of one of these runs., The top panel in Figure \ref{june} shows the number of triggers per data block for one hour of one of these runs.
 The number of triggers is only plotted when it exceeds 20., The number of triggers is only plotted when it exceeds 20.
 The maximum number of triggers per block is equal to the number of traces: 200., The maximum number of triggers per block is equal to the number of traces: 200.
 This maximum is often, This maximum is often
flexion and found similar structure using weak lensing shear and flexion measurements and a Fourier mass reconstruction.,flexion and found similar structure using weak lensing shear and flexion measurements and a Fourier mass reconstruction.
 In this section we apply the AIM method of flexion measurement to substructure detection in A1689 as a real-world validation of our method., In this section we apply the AIM method of flexion measurement to substructure detection in A1689 as a real-world validation of our method.
" It is important to note that in both previous flexion analyses of A1689, the authors mistreat the mass-sheet degeneracy in some way."," It is important to note that in both previous flexion analyses of A1689, the authors mistreat the mass-sheet degeneracy in some way."
 As shown by (2008). this is not a valid approximation for the interior region of the cluster where κ&1.," As shown by , this is not a valid approximation for the interior region of the cluster where $\kappa\not\ll1$."
" This means that direct, quantitative comparisons between this paper and previous work are not feasible."," This means that direct, quantitative comparisons between this paper and previous work are not feasible."
" Instead, we look for qualitative morphological similarities between the structures detected, with the intent of performing a direct comparison of flexion measurement methods in future Work."," Instead, we look for qualitative morphological similarities between the structures detected, with the intent of performing a direct comparison of flexion measurement methods in future work."
 The data were obtained from the Hubble Legacy (HLA) and reduced using the standard HLA pipeline., The data were obtained from the Hubble Legacy (HLA) and reduced using the standard HLA pipeline.
 These data were originally obtained for proposal HST-GTO/ACS9289 (PI Ford)., These data were originally obtained for proposal HST-GTO/ACS9289 (PI Ford).
" The data include deep observations in F475W (9.6 ks), F625W (9.6 ks), F775W (11.8 ks), and F850LP (16.6 ks)."," The data include deep observations in F475W (9.6 ks), F625W (9.6 ks), F775W (11.8 ks), and F850LP (16.6 ks)."
" The observed field is 3/4x9.4, and oon the sky corresponds to 3.1 kpc at the redshift z=0.187 of the cluster, assuming a standard cosmology."," The observed field is $3\farcm4\times3\farcm4$, and on the sky corresponds to 3.1 kpc at the redshift $z=0.187$ of the cluster, assuming a standard cosmology."
Iu real space. we find that perurbative approaches based ou he fluid description can exend the wavenumber up to which systematic analytical predicloli can be obtained from fo=0.0330 1 up to O.23h to at 2= 0. and from ~07h 3 up to ~22h b at: =doas οnua with linear theory (if we liec απ accuracy Q| 1€),"In real space, we find that perturbative approaches based on the fluid description can extend the wavenumber up to which systematic analytical predictions can be obtained from $k\simeq 0.033 h$ $^{-1}$ up to $\simeq 0.23 h$ $^{-1}$ at $z=0$ , and from $\simeq 0.07 h$ $^{-1}$ up to $\simeq 2.2 h$ $^{-1}$ at $z=3$, as compared with linear theory (if we need an accuracy of $1\%$ )."
 We also give cletailed POSIts foY Oher redshifts aud ‘OL 10X aud 5O% accracy levex., We also give detailed results for other redshifts and for $10\%$ and $50\%$ accuracy levels.
 €4(oueMο to higherMAN & requiresM ταιςMH shel-crossing effects into account., Going to higher $k$ requires taking shell-crossing effects into account.
 Neverheless. these results show hat such erturbative approaches. based ou lvdrodvuamical ος1atious of motion. have a significant potential. especially at Heh aud moderate redshifts.," Nevertheless, these results show that such perturbative approaches, based on hydrodynamical equations of motion, have a significant potential, especially at high and moderate redshifts."
 Since it is unlikely hat i wil be possible to explicitly compute high-order perturbaive ternis up to the last order nj. which is above the ronperturbative correction (i.e.=t) at τO and νο=66 at : 3). this provicCR a strong incentive to clevelop resununuation schemes.," Since it is unlikely that it will be possible to explicitly compute high-order perturbative terms up to the last order $n_{\rm s.c.}$ which is above the nonperturbative correction $n_{\rm s.c.}=9$ at $z=0$ and $n_{\rm s.c.}=66$ at $z=3$ ), this provides a strong incentive to develop resummation schemes."
 Then. the hope is that suc1 ijethods can efficiently resi nost of the contributions associated with higher order ternis aud achieve a aster convergence.," Then, the hope is that such methods can efficiently resum most of the contributions associated with higher order terms and achieve a faster convergence."
 This analysis is also useful iuo the coutest of cosmological reconstructiou. where oue tries to recover he ανασα. hüstorv of a given region of the skv roni its present density field (and next to derive frou lis reconstruction the present velocity field).," This analysis is also useful in the context of cosmological reconstruction, where one tries to recover the dynamical history of a given region of the sky from its present density field (and next to derive from this reconstruction the present velocity field)."
 Iudeed. hese methods usually neglect shelbcrossimg effects.," Indeed, these methods usually neglect shell-crossing effects."
 Iu utieulu. we find that the Mouge-Amiperre-Rautorovicli ucthod is close to otinal at 2=0. as it fares well «OW1L ο the scale where s1Cl-crossiug contributions douinate.," In particular, we find that the Monge-Ampèrre-Kantorovich method is close to optimal at $z=0$, as it fares well down to the scale where shell-crossing contributions dominate."
" With respect to the rec'oustiuction of the barvon acoustic yeas or of the acoustic oscillations of the power spectrin. rere seclus to be on.V nodest roo for improvement over ""urreut methods based ou the linear displacement fiek at 2=O. in aerccmeut with only a few orders of perturbation wory being rclevan (ive= 9)."," With respect to the reconstruction of the baryon acoustic peak or of the acoustic oscillations of the power spectrum, there seems to be only modest room for improvement over current methods based on the linear displacement field at $z=0$, in agreement with only a few orders of perturbation theory being relevant $n_{\rm s.c.}=9$ )."
 At higher : one could 11i xinciple do much better by includiug higher order terllis. mt since barvon acoustic oscillations of the linear power PAxectrma have a snall ;uuplitude at high A. this may iot sjenificautle improve the signal-to-noise ratio.," At higher $z$ one could in principle do much better by including higher order terms, but since baryon acoustic oscillations of the linear power spectrum have a small amplitude at high $k$, this may not significantly improve the signal-to-noise ratio."
" Ποπονα, Us remains useful for other observational probes tlat we sensitive to the shape and amplitude of the power spectrum on weakly noulirear scales. such as weak-leusine SULPVCVS"," However, this remains useful for other observational probes that are sensitive to the shape and amplitude of the power spectrum on weakly nonlinear scales, such as weak-lensing surveys."
", We have also poiute Olt that the behavior of the syste after shell crossing pavs a kev role in the shape of the densiY power specrun on nildlv noulirear scales, κASQ)«100."," We have also pointed out that the behavior of the system after shell crossing plays a key role in the shape of the density power spectrum on mildly nonlinear scales, $1< \Delta^2(k) < 100$."
 Thus. arguneuts based on the spherical collapse dvuauics are not sificient to explai1 the steep rise of the ionlinear power iu this intermediate regine.," Thus, arguments based on the spherical collapse dynamics are not sufficient to explain the steep rise of the nonlinear power in this intermediate regime."
 ludeed. whie for the Zeldovich αναος suc ra erowth is almost entirely. wiped out wo the escape of articles to lufinity (as he raudon Luear displacements erase siuall-scale features and lead to a decay of A*(h) at lueh 4). for the “sticzv model there is a regular growt1 up to the high-k asviuptote. ος)—KR7.," Indeed, while for the Zeldovich dynamics such a growth is almost entirely wiped out by the escape of particles to infinity (as the random linear displacements erase small-scale features and lead to a decay of $\Delta^2(k)$ at high $k$ ), for the “sticky model” there is a regular growth up to the $k$ asymptote, $\Psticky(k) \sim k^{-2}$."
 Therefore. it is not sufficieit to consider the spherical collapse of a “typical” OVOTCchsitya. to explain the shape of P(h) on these scales. as one mst take the Gaussian average ovya whole rauge of initial density fluctuations. which show a siguificaut auoul of shell crossing that has a strong impact on the resulting power spectrum.," Therefore, it is not sufficient to consider the spherical collapse of a “typical” overdensity to explain the shape of $P(k)$ on these scales, as one must take the Gaussian average over a whole range of initial density fluctuations, which show a significant amount of shell crossing that has a strong impact on the resulting power spectrum."
 Finally. we have obtained simular results for the redshift“Space power spoectrun.," Finally, we have obtained similar results for the redshift-space power spectrum."
 Again. the scope of )orturtive approaches (based. on the fluid description) is greater at higjr Ll.," Again, the scope of perturbative approaches (based on the fluid description) is greater at higher $z$."
 As compared with the real-space )oWer spectruni. the characteristic wavenunibers where üeher order terus of the perturbative expansions and ΠΟΙΟurbative contributions come iuto play are süfted o lower values oA.," As compared with the real-space power spectrum, the characteristic wavenumbers where higher order terms of the perturbative expansions and nonperturbative contributions come into play are shifted to lower values of $k$."
 TIudeed. because of the amplification of perturbations youn the uniform: Iubble backerouid by he additional contribution from peculiar velocities. the 1onlinear regime is reached on larger scales than in real space.," Indeed, because of the amplification of perturbations from the uniform Hubble background by the additional contribution from peculiar velocities, the nonlinear regime is reached on larger scales than in real space."
 The “sticky inodel preseuted in this article. aud nore eenerallv the expression (7)) of the| density power spectitun in terms of the Lagraungku displacement. could serve as a basis for models for the noulinear o»wer spectrum.," The “sticky model” presented in this article, and more generally the expression \ref{Pkxq}) ) of the density power spectrum in terms of the Lagrangian displacement, could serve as a basis for models for the nonlinear power spectrum."
 In particular. the spiit of this model could be incorporated iuto the Lagraugiai-based models. associated with the parametric svsteni (T8) -CEO)). or into the “halo model”. to make the widege between the perturbaIve reelne and the lughly nonlinear. nouperturbative regine (?2)..," In particular, the spirit of this model could be incorporated into the Lagrangian-based models, associated with the parametric system \ref{kL}) \ref{f-def}) ), or into the “halo model”, to make the bridge between the perturbative regime and the highly nonlinear, nonperturbative regime \citep{Valageas2010a}."
 This could aIso be extended to redshiftspace bv inchiding a model or the velocitv dispersion within virialized objects., This could also be extended to redshiftspace by including a model for the velocity dispersion within virialized objects.
 Ilowever. this Is bevond the scope of the present article so we leave it to future works.," However, this is beyond the scope of the present article so we leave it to future works."
to classify galaxies into early- and late-type galaxies.,to classify galaxies into early- and late-type galaxies.
 T'he reliability and completeness of this automated method is about9096., The reliability and completeness of this automated method is about.
. Morphology information given by this scheme has been used in many studies on relation between galaxy properties and the environment (Choiettoraetal. 2010).," Morphology information given by this scheme has been used in many studies on relation between galaxy properties and the environment \citep{choi+07,choi+09,park07,park08,pnc09,pnh08,lee+08,lee+10a,lee+10b,lee+10c,han+10,cervantes10,cervantes11,
hwang+10,hwang+11,hwangnpark10,lara+10,tortora+10}."
. To improve the accuracy of morphology classification we perform additional visual check., To improve the accuracy of morphology classification we perform additional visual check.
" The morphology of 2,427 galaxies in the S1 sample, which is about7.3%,, is changed by the additional visual check."," The morphology of 2,427 galaxies in the S1 sample, which is about, is changed by the additional visual check."
" As a result, the S1 sample includes 13,867 early types, 19,431 late types, and 93 unclassified galaxies, while the 92 sample consists of 5,334 early types, 11,393 late types, and 1 unclassified galaxy, as listed in Table 1.."," As a result, the S1 sample includes 13,867 early types, 19,431 late types, and 93 unclassified galaxies, while the S2 sample consists of 5,334 early types, 11,393 late types, and 1 unclassified galaxy, as listed in Table \ref{table_morph}."
 We exclude unclassified galaxies in the following analysis., We exclude unclassified galaxies in the following analysis.
" We compare our classification with two previous works: Nair&Abraham(2010a,hereafter and Huertas-Companyetal. (2011).", We compare our classification with two previous works: \citet[][hereafter NA10]{nair10a} and \citet{huertas+11}.
. We find a good NA10)agreement among three morphology classifications., We find a good agreement among three morphology classifications.
 Figure 2 shows the results., Figure \ref{compmc} shows the results.
" NAI10 presented a visual classification for ~14,000 galaxies in the SDSS DR4 with g«16 mag at 0.01« 0.1."," NA10 presented a visual classification for $\sim14,000$ galaxies in the SDSS DR4 with $g<16$ mag at $0.01<z<0.1$ ."
" They assigned T-Types to all galaxies in their sample: T-Type=0 for «0 for E and 90, 1—9 for Sa-Sm, 10—11 for Im, S0/a,and 99 for unknown galaxies (peculiar galaxies, or unclassified ones)."," They assigned T-Types to all galaxies in their sample: $=0$ for S0/a, $<0$ for E and S0, $1-9$ for Sa-Sm, $10-11$ for Im, and 99 for unknown galaxies (peculiar galaxies, or unclassified ones)."
" For 5,334 galaxies included in both our and NA10’s sample, we classify these galaxies into 2,257 early-type galaxies and 3,077 late-type galaxies."," For 5,334 galaxies included in both our and NA10's sample, we classify these galaxies into 2,257 early-type galaxies and 3,077 late-type galaxies."
" The majority (88.3%)) of these early-type galaxies are assigned to T-Type<0 (E or S0), while of them have T-Types later than $0/a and of them are classified as unknown galaxies."," The majority ) of these early-type galaxies are assigned to $\leq0$ (E or S0), while of them have T-Types later than S0/a and of them are classified as unknown galaxies."
" On the other hand, of the late-type galaxies have T-Types for $0/a-Sm, and only 105 galaxies (3.4%)) have T-Types earlier than SO/a, and 79 galaxies (2.6%)) are classified as unknown galaxies."," On the other hand, of the late-type galaxies have T-Types for S0/a-Sm, and only 105 galaxies ) have T-Types earlier than S0/a, and 79 galaxies ) are classified as unknown galaxies."
" Huertas-Companyetal.(2011) presented a Bayesian automated classification for ~700,000 galaxies from the SDSS DR7 spectroscopic sample, and they quantified probabilities to each galaxy of being in four types (E, $0, Sab, Scd)."," \citet{huertas+11} presented a Bayesian automated classification for $\sim700,000$ galaxies from the SDSS DR7 spectroscopic sample, and they quantified probabilities to each galaxy of being in four types (E, S0, Sab, Scd)."
" Probabilities of being in early-type galaxy, P(Early), are only used in this comparison."," Probabilities of being in early-type galaxy, $P({\rm Early})$, are only used in this comparison."
" For 45,781 galaxies common in our and their sample, we divide them into 17,343 early-type and 28,438 late-type galaxies."," For 45,781 galaxies common in our and their sample, we divide them into 17,343 early-type and 28,438 late-type galaxies."
" We find that median values of P(Early) are 0.84 and 0.07 for the early-type and the late-type galaxies, respectively."," We find that median values of $P({\rm Early})$ are 0.84 and 0.07 for the early-type and the late-type galaxies, respectively."
" In addition, of early-type galaxies have P(Early) higher than 0.5, while of late-type galaxies have P(Early) lower than 0.5."," In addition, of early-type galaxies have $P({\rm Early})$ higher than 0.5, while of late-type galaxies have $P({\rm Early})$ lower than 0.5."
" In conclusion, our morphology classification shows a good agreement with two previous works."," In conclusion, our morphology classification shows a good agreement with two previous works."
" We use the physical parameters of galaxies to study the dependence of fa, on the properties of galaxies: absolute Petrosianmagnitude 9!M., 9:!(u—r) color, Petrosian radius (Ep«) in i-band, color gradient in 9.1(g—i), inverse concentration index (cn), central velocity dispersion (c), and equivalent width of the Ha line."," We use the physical parameters of galaxies to study the dependence of $\bfr$ on the properties of galaxies: absolute Petrosianmagnitude ${}^{0.1}M_r$, ${}^{0.1}(u-r)$ color, Petrosian radius $R_{Pet}$ ) in $i$ -band, color gradient in ${}^{0.1}(g-i)$, inverse concentration index $c_{in}$ ), central velocity dispersion $\sigma$ ), and equivalent width of the $\ha$ line."
 These parameters reflect most of major physical properties of galaxies from morphology and mass to kinematics and star formation activity (Park&Choi 2009).., These parameters reflect most of major physical properties of galaxies from morphology and mass to kinematics and star formation activity \citep{pnc09}. .
" The rest-frame absolute magnitudes of individual galaxies are computed in fixed bandpass, shifted to z=0.1, using the Galactic reddening correction of Schlegeletal.(1998) and K-corrections as described by Blantonetal.(2003)."," The rest-frame absolute magnitudes of individual galaxies are computed in fixed bandpass, shifted to $z=0.1$, using the Galactic reddening correction of \citet{schle1998} and K-corrections as described by \citet{blan03}."
". The mean evolution correction given by Tegmarketal. (2004), E(z)=1.6(z— 0.1), is also applied."," The mean evolution correction given by \citet{teg04}, , $E(z)=1.6(z-0.1)$ , is also applied."
" The superscript 0.1 indicates that the absolute magnitude is the rest-frame magnitude when the galaxy is at z= 0.1, after K-correction and luminosity evolution correction."," The superscript 0.1 indicates that the absolute magnitude is the rest-frame magnitude when the galaxy is at $z=0.1$ , after K-correction and luminosity evolution correction."
" To compute colors, we use extinction and K-corrected model magnitudes."," To compute colors, we use extinction and K-corrected model magnitudes."
" 9-14—r) and ""1(g—r) colors also follow this convention.", ${}^{0.1}(u-r)$ and ${}^{0.1}(g-r)$ colors also follow this convention.
" Hereafter, the superscript 0.1 will be omitted."," Hereafter, the superscript 0.1 will be omitted."
" The g—i color gradient is the difference between the color with aperture radius R«0.5Rpe and thecolor with the annulus 0.5Rpe,« Rpez.", The $g-i$ color gradient is the difference between the color with aperture radius $R<0.5R_{Pet}$ and thecolor with the annulus $0.5R_{Pet}<R<R_{Pet}$ .
" When the color gradient has a negative value, it indicates that the galaxy has a redder inner part and a bluer outer part."," When the color gradient has a negative value, it indicates that the galaxy has a redder inner part and a bluer outer part."
" The inverse concentration index (cin) is defined by H5o/Rgo where Hs5o and {ορ are semi-major axis lengths of ellipses enclosing 5096 and 9096 of Petrosian flux in the i-band image, respectively."," The inverse concentration index $c_{in}$ ) is defined by $R_{50}/R_{90}$ where $R_{50}$ and $R_{90}$ are semi-major axis lengths of ellipses enclosing $50\%$ and $90\%$ of Petrosian flux in the $i$ -band image, respectively."
 The central velocity dispersion value is adopted from NYU-VAGC (Blantonetal. 2005).., The central velocity dispersion value is adopted from NYU-VAGC \citep{blan05}. .
" The value ofHa equivalent widthis taken from MPA/JHU- (Tremontietal. whichwas computed using the straight integration 2004),over the fixed bandpass from the continuum-subtracted emission line with the modelof Bruzual&Charlot (2003).."," The value of$\ha$ equivalent widthis taken from MPA/JHU-VAGC \citep{trem04}, , whichwas computed using the straight integration over the fixed bandpass from the continuum-subtracted emission line with the modelof \citet{bruz03}. ."
 We adopt a flat, We adopt a flat
obviously invoke a «question on the comparison between the correlations of RC-ICN versus RC-CO: is the RC-IICN correlation signilicantly better than the RC-CO correlation?,obviously invoke a question on the comparison between the correlations of RC-HCN versus RC-CO: is the RC-HCN correlation significantly better than the RC-CO correlation?
 llow do the FIR and RC as SFR indicators relate with the star formation materials. the molecular gas (CO). and particularly the dense molecular gas (ICN)?," How do the FIR and RC as SFR indicators relate with the star formation materials, the molecular gas (CO), and particularly the dense molecular gas (HCN)?"
" Stars are born mainlv in GAICs,", Stars are born mainly in GMCs.
 The total mass of molecular gas of GMCs can be determined trom the CO luminosity., The total mass of molecular gas of GMCs can be determined from the CO luminosity.
 Llowever. the excesses of the CO luminosity in GAIC cores where active star formation occurs are not specilic enough (o reveal the star formation potential of the dense cores.," However, the excesses of the CO luminosity in GMC cores where active high-mass star formation occurs are not specific enough to reveal the star formation potential of the dense cores."
 The physical conditions of stu-Iorming GAIC cores (Evans 1999) are better revealed by emission from high dipole-moment molecules. like WON. whose emission (races the dense molecular gas (Πο)x 10'em7) associated with the cores of the GMCsS (GS04hb).," The physical conditions of star-forming GMC cores (Evans 1999) are better revealed by emission from high dipole-moment molecules, like HCN, whose emission traces the dense molecular gas $n(H_2) \approxgt 3 \times 10^4$ $^{-3}$ ) associated with the cores of the star-forming GMCs (GS04b)."
 FIB-IICN is shown to be linearly correlated (GS04a)., FIR-HCN is shown to be linearly correlated (GS04a).
 Does HCN also strongly correlate with RC?, Does HCN also strongly correlate with RC?
 What is the role of dense molecular gas in the FIR-RC relationship?, What is the role of dense molecular gas in the FIR-RC relationship?
 llere. we utilize (his HCN sample to analyze the various relationships among the global IICN. FIR. CO. and RC luminosities to answer some questions raised here. and to further demonstrate (he possibility of using the RC huninosity instead of FIR luminosity as an inclicator of the SER.," Here, we utilize this HCN sample to analyze the various relationships among the global HCN, FIR, CO, and RC luminosities to answer some questions raised here, and to further demonstrate the possibility of using the RC luminosity instead of FIR luminosity as an indicator of the SFR."
 The HCN sample and the total RC Iuminositv of the sample are presented in Section 2., The HCN sample and the total RC luminosity of the sample are presented in Section 2.
 Section 3 presents (he results and analysis of (he various correlations ol these global luminosities and their ratios and shows that the global correlation between RC and HCN appears to be a combinational result of (wo Geghtest correlations between and between FIB-IICN even though. RC-HCN correlation is significantly better (han the RC-CO correlation., Section 3 presents the results and analysis of the various correlations of these global luminosities and their ratios and shows that the global correlation between RC and HCN appears to be a combinational result of two tightest correlations between FIR-RC and between FIR-HCN even though RC-HCN correlation is significantly better than the RC-CO correlation.
 Discussion on (he possible physical relationship between (he ICN (dense molecular gas tracer) and the RC emission (the indicator of the rate of high-mass star formation) in galaxies is presented in Section 4. followed by the main points of this study in Seclion 5.," Discussion on the possible physical relationship between the HCN (dense molecular gas tracer) and the RC emission (the indicator of the rate of high-mass star formation) in galaxies is presented in Section 4, followed by the main points of this study in Section 5."
 Our RC sample is drawn entirely from (he IICN sample of Gao Solomon (GS04a.b).," Our RC sample is drawn entirely from the HCN sample of Gao Solomon (GS04a,b)."
 The sample includes almost all galaxies with strong CO and FIR emission in the northern sky (53 in total. GS04b) and 12 additional galaxies (mostly. LIRGs/ULIRGs) from literatures. mainly [rom Solomon. Downes Itadford (1992).," The sample includes almost all galaxies with strong CO and FIR emission in the northern sky (53 in total, GS04b) and 12 additional galaxies (mostly LIRGs/ULIRGs) from literatures, mainly from Solomon, Downes Radford (1992)."
 Details about the sample properties and ihe measurements of the total IICN emission in galaxies can be found in G904b., Details about the sample properties and the measurements of the total HCN emission in galaxies can be found in GS04b.
 The RC flux is mainly obtained from the NRAO VLA Sky Survey (NVSS: Condon οἱ al., The RC flux is mainly obtained from the NRAO VLA Sky Survey (NVSS; Condon et al.
 1993). which covers the sky. north of J2000.0 decl.2—40* at 1.4 Giz (Condon et al.," 1998), which covers the sky north of J2000.0 $-40\arcdeg$ at 1.4 GHz (Condon et al."
 1993).," 1998),"
suppression depends only on the so-called radiation parameter Aj}.a which is (up to a constant factor) the ratio of the accretion bhuninositv of the disk GALAL/R to the DIuuinositv Ly that a star would have had in the abseuce of irradiation.,"suppression depends only on the so-called irradiation parameter ), which is (up to a constant factor) the ratio of the accretion luminosity of the disk $GM\dot M/R$ to the luminosity $L_0$ that a star would have had in the absence of irradiation."
 Suppression factor (ΑΛ)>» Las A<1 while \(A)X lwhen A291., Suppression factor $\chi(\Lambda)\to 1$ as $\Lambda\ll 1$ while $\chi(\Lambda)\lesssim 1$ when $\Lambda\gg 1$.
 Tn a simple case considered by Rafikov (2007) the dependence of ν on the opacity behavior comes only through the where V4 is the adiabatic temperature eradicut near he stellar surface., In a simple case considered by Rafikov (2007) the dependence of $\chi$ on the opacity behavior comes only through the where $\nabla_{ad}$ is the adiabatic temperature gradient near the stellar surface.
 Iu the stronglv irradiated case the temperature of the stellar surface varies as a function of latitude: equatorial olt is strouglv heated by the hot immer parts of tho disk while the polar regious of the star are virtually unaffected wo irraciation and preserve their temperature at the evel of Ty=(LyfAcR?o)!/1., In the strongly irradiated case the temperature of the stellar surface varies as a function of latitude: equatorial belt is strongly heated by the hot inner parts of the disk while the polar regions of the star are virtually unaffected by irradiation and preserve their temperature at the level of $T_0=(L_0/4\pi R^2\sigma)^{1/4}$.
 In this situation one may wouder whether it is reasonable to assume a fixed opacity aw (as was done in Rafikov 2007) for the calculation of Iunünositv suppression given that the behavior of & can be different between the polar aud the equatorial regions of the star., In this situation one may wonder whether it is reasonable to assume a fixed opacity law (as was done in Rafikov 2007) for the calculation of luminosity suppression given that the behavior of $\kappa$ can be different between the polar and the equatorial regions of the star.
 Iowever. is was shown iu Rafikov (2007) that for the opacity behavior typical for stellar plotospleres in the teniperatuxe interval from ~2.5«10? K to 10! K the cooling of irradiated stars occurs mainly through their polar (even though the opacity scaling with temperature and pressure changes quite drastically within this femiperature interval at around 5000 I&).," However, is was shown in Rafikov (2007) that for the opacity behavior typical for stellar photospheres in the temperature interval from $\sim 2.5\times 10^3$ K to $10^4$ K the cooling of irradiated stars occurs mainly through their polar (even though the opacity scaling with temperature and pressure changes quite drastically within this temperature interval at around $5000$ K)."
 As a result. no matter how hot the equatorial parts of the star become and how complicated the opacity behavior is in this portion of the stellar surface. the integrated stellar luminosity L docs not depeud on these details very stronglv but is rather determined by the properties of the polar regions of the star: the size of the cool polar caps in which the photospherie temperature is preserved at the level of Ty. aud the opacity behavior in the adjacent parts of the stellar surface.," As a result, no matter how hot the equatorial parts of the star become and how complicated the opacity behavior is in this portion of the stellar surface, the integrated stellar luminosity $L$ does not depend on these details very strongly but is rather determined by the properties of the polar regions of the star: the size of the cool polar caps in which the photospheric temperature is preserved at the level of $T_0$, and the opacity behavior in the adjacent parts of the stellar surface."
 This provides motivation for using the stellar huninositv prescription represented by equation (9))., This provides motivation for using the stellar luminosity prescription represented by equation \ref{eq:Lambda}) ).
 Iu this paper we adopt & characteristic for the teiiperatine interval 2.5&10? IK <7T<5<lo?’ W (Bell Lin 1991):, In this paper we adopt $\kappa$ characteristic for the temperature interval $2.5\times 10^3$ K $\lesssim T\lesssim 5\times 10^3$ K (Bell Lin 1994):.
 This scaling should be reasonable in the polar regious of iradiated stars where the photospleric temperature Jy is not strongly affected by radiation and is close to the photospheric temperature that an isolated non-accreting star would have possessed in the Hayashi phase., This scaling should be reasonable in the polar regions of irradiated stars where the photospheric temperature $T_0$ is not strongly affected by irradiation and is close to the photospheric temperature that an isolated non-accreting star would have possessed in the Hayashi phase.
 The opacity lav (13)) corresponds to ©=6.5. assuniug Vad=2/5 as appropriate for the fully ionized stellar interior (5=5/3).," The opacity law \ref{eq:opacity}) ) corresponds to $\xi=6.5$ , assuming $\nabla_{ad}=2/5$ as appropriate for the fully ionized stellar interior $\gamma=5/3$ )."
 Having specified the opacity law we have thus fully determined the behavior of VCÀ) (which we take from Rafikov [2007]. sce the curve corresponding to €=6.5 in Fig.," Having specified the opacity law we have thus fully determined the behavior of $\chi(\Lambda)$ (which we take from Rafikov [2007], see the curve corresponding to $\xi=6.5$ in Fig."
 tof that paper) necessary to compute L, 4 of that paper) necessary to compute $L$.
 One remaining ineredicut of the calculation is the choice of Ly(Al.R) he luminosity of a nou-irracliated star.," One remaining ingredient of the calculation is the choice of $L_0(M,R)$ – the luminosity of a non-irradiated star."
 Hartinaun (1997) have adopted the following fit to the stellar evolution. tracks of D'AÀntona Mazztelli (1991): LyfAL.Ry=-1L.," Hartmann (1997) have adopted the following fit to the stellar evolution tracks of D'Antona Mazzitelli (1994): L_0(M,R)=1."
{| Evolution tracks in D'Autoua Mazzitelli (1991) have been caleulated using the equation of state from Maenui Mazzitelli (1979) that has been superceded by the more refined treatineuts (Sammon 1995)., Evolution tracks in D'Antona Mazzitelli (1994) have been calculated using the equation of state from Magni Mazzitelli (1979) that has been superceded by the more refined treatments (Saumon 1995).
 Also. D'Antona Mazzitelli's treatment of convection is based ou Canuto Alazzitellii (1991) which has previously raised. some concerns (Demarque 1999: Nordluud Stein 1999).," Also, D'Antona Mazzitelli's treatment of convection is based on Canuto Mazzitelli (1991) which has previously raised some concerns (Demarque 1999; Nordlund Stein 1999)."
 Despite these deficiencies. we lave choseu to adopt the prescription (L1)) i our work because of its aud also to allow direct comparison with the results of ITartinaun (1997).," Despite these deficiencies, we have chosen to adopt the prescription \ref{eq:L_0}) ) in our work because of its and also to allow direct comparison with the results of Hartmann (1997)."
 Equatious (5)). 0663). (29). (999. (LL) supplemented with dA//dt=M(f) and the depeudeuce \(A) from Rafikov (2007) fully cletermine the evolution of an accreting protostar ivradiated by its own disk.," Equations \ref{eq:master}) ), \ref{eq:L_D}) ), \ref{eq:f}) ), \ref{eq:Lambda}) ), \ref{eq:L_0}) ) supplemented with $dM/dt=\dot M(t)$ and the dependence $\chi(\Lambda)$ from Rafikov (2007) fully determine the evolution of an accreting protostar irradiated by its own disk."
 This system of equations is then evolved ummerically assuiiug that the prescription for A(f) is given., This system of equations is then evolved numerically assuming that the prescription for $\dot M(t)$ is given.
 Tere we present the results of our calculations., Here we present the results of our calculations.
 As initial conditions we choose Af=0.1 AL. and fp1., As initial conditions we choose $M=0.1$ $_\odot$ and $f_D=1$.
 We vary AR aud the prescription for AJ(f) to see their effect on the evolution of stellar properties., We vary $R$ and the prescription for $\dot M(t)$ to see their effect on the evolution of stellar properties.
 Tn Figure l we display protostellar evolution for initial R-125RBR.audauufoii Af=ls10° Mv |! with and without the effects of disk ivracdiation included., In Figure \ref{fig:f1} we display protostellar evolution for initial $R=1.5$ $_\odot$ and a uniform $\dot M=4\times 10^{-6}$ $_\odot$ $^{-1}$ with and without the effects of disk irradiation included.
 One can see that in both irradiated aud nou-imradiated cases evolution is pretty mich the same: star initially5 contracts until its central density aud temperaturebecome high enough for the deuterium to iguite., One can see that in both irradiated and non-irradiated cases evolution is pretty much the same: star initially contracts until its central density and temperaturebecome high enough for the deuterium to ignite.
 Right after that D burning strouely dominates over the stella huninositv and D abundance fp starts going down., Right after that D burning strongly dominates over the stellar luminosity and D abundance $f_D$ starts going down.
 Resulting energv, Resulting energy
technique. and to note that the results presented here are consistent (if marginally) with some other analyses which use only low-redshift data.,"technique, and to note that the results presented here are consistent (if marginally) with some other analyses which use only low-redshift data."
 If the tightness of the constraints seems surprising. consider firstly that in the absence of uncertainty about galaxy bias. the constraints would be extremely tight as the amplitude of the correlation function is very well determined.," If the tightness of the constraints seems surprising, consider firstly that in the absence of uncertainty about galaxy bias, the constraints would be extremely tight as the amplitude of the correlation function is very well determined."
 Secondly. by comparing model galaxies to real galaxies of the same abundance. we factor out most of the dependence on the details of the semi- model.," Secondly, by comparing model galaxies to real galaxies of the same abundance, we factor out most of the dependence on the details of the semi-analytic model."
 Any change that makes galaxies monotonically brighter or fainter will have no effect on our comparison between models and data., Any change that makes galaxies monotonically brighter or fainter will have no effect on our comparison between models and data.
 Thirdly. the number of galaxies assigned to each halo in the semi-analytio model is principally determined by the merger history of that halo. and this is well described by the extended Press-Schechter theory.," Thirdly, the number of galaxies assigned to each halo in the semi-analytic model is principally determined by the merger history of that halo, and this is well described by the extended Press-Schechter theory."
 Hence this is not a major uncertainty in the semi-analytic models: unlike purely statistical descriptions as provided by the HOD or CLF. our semi-analytic models do not have the freedom to define arbitrary HODs.," Hence this is not a major uncertainty in the semi-analytic models: unlike purely statistical descriptions as provided by the HOD or CLF, our semi-analytic models do not have the freedom to define arbitrary HODs."
 These are instead largely determined by the merger trees., These are instead largely determined by the merger trees.
 We ulso remark that other constraints using galaxy data alone. which at first sight seem inconsistent with ours. use different techniques or different data or both.," We also remark that other constraints using galaxy data alone, which at first sight seem inconsistent with ours, use different techniques or different data or both."
 These inconsistencies arise even though galaxy formation models can reproduce. statistics which implicitly average over a range of halo mass — such as the unconditional galaxy luminosity function — reasonably well., These inconsistencies arise even though galaxy formation models can reproduce statistics which implicitly average over a range of halo mass – such as the unconditional galaxy luminosity function – reasonably well.
 This illustrates that matching the luminosity- and colour-dependence of clustering. at small and large scales. is an important and very stringent requirement on galaxy formation models.," This illustrates that matching the luminosity- and colour-dependence of clustering, at small and large scales, is an important and very stringent requirement on galaxy formation models."
" We have compared the SDSS projected two-point correlation function at a galaxy space density n,=0.00308fh?Alpe7 to a suite of populated simulations generated using the ;N-body code and the semi-analytic codeGALFORM."," We have compared the SDSS projected two-point correlation function at a galaxy space density $\bar{n}_\mathrm{g}=0.00308\ h^3\
\mathrm{Mpc}^{-3}$ to a suite of populated simulations generated using the $N$ -body code and the semi-analytic code."
. Because we require ;/V -body data in a great numberof different cosmologies. we have relabelled and rescaled some simulation outputs using the techniques of ? to avoid the need to run a full simulation for each cosmology in our grid.," Because we require $N$ -body data in a great number of different cosmologies, we have relabelled and rescaled some simulation outputs using the techniques of \citet{ZHE02} to avoid the need to run a full simulation for each cosmology in our grid."
 The galaxy catalogues are self-consistent. being run afresh for each cosmology we study.," The galaxy catalogues are self-consistent, being run afresh for each cosmology we study."
 We have attempted to estimate the systematic error in our value of ox due to the particular choice of semi-analytic model by running three different variants in each cosmology., We have attempted to estimate the systematic error in our value of $\sigma_8$ due to the particular choice of semi-analytic model by running three different variants in each cosmology.
 For each of these variants we generate a catalogue in. which he parameters are adjusted to match the SDSS [trband luminosity function and a catalogue in which the parameters ake the same values as they take in the (aij.0.)=(0.3.0.8) cosmology.," For each of these variants we generate a catalogue in which the parameters are adjusted to match the SDSS $^{0.1}r$ -band luminosity function and a catalogue in which the parameters take the same values as they take in the $(\Omega_\mathrm{M},\sigma_8)=(0.3,0.8)$ cosmology."
 We obtain the result a=0.97E O.O4CGtatistical) +0.04 (systematic)., We obtain the result $\sigma_8=0.97\pm 0.04$ (statistical) $\pm 0.04$ (systematic).
 This constraint is impressively tight. given we lave attempted to narrow the range of assumptions we require to oroduee an estimate of σε by using only one well understood. ow redshift dataset.," This constraint is impressively tight, given we have attempted to narrow the range of assumptions we require to produce an estimate of $\sigma_8$ by using only one well understood, low redshift dataset."
" By choosing grids of cosmologies which lie on cluster-normalized curves. σκΟν]=const. we have shown hat the degeneracies inherent in our approach are different to hose inherent in cosmic shear measurements. which provide an important low redshift constraint in O3, and ox."," By choosing grids of cosmologies which lie on cluster-normalized curves, $\sigma_8\Omega_\mathrm{M}^{0.5}=\mathrm{const.}$ we have shown that the degeneracies inherent in our approach are different to those inherent in cosmic shear measurements, which provide an important low redshift constraint in $\Omega_\mathrm{M}$ and $\sigma_8$."
 In fact our method gives an almost pure constraint on ox., In fact our method gives an almost pure constraint on $\sigma_8$.
 We have shown that in our model. halo assembly bias does not severely affect our constraint on ox. though this may not be universally the case for other analytic codes.," We have shown that in our model, halo assembly bias does not severely affect our constraint on $\sigma_8$, though this may not be universally the case for other semi-analytic codes."
 Π it were not the case. we would expect it to bias our estimate of σε high. since failing to account for halo assembly bias tends to lower the amplitude of model correlation functions. requiring an increased as in the model to compensate.," If it were not the case, we would expect it to bias our estimate of $\sigma_8$ high, since failing to account for halo assembly bias tends to lower the amplitude of model correlation functions, requiring an increased $\sigma_8$ in the model to compensate."
 We obtain similar values for ox if we use samples with a lower or higher galaxy space density. but the error analysis is less secure.," We obtain similar values for $\sigma_8$ if we use samples with a lower or higher galaxy space density, but the error analysis is less secure."
 We also obtain similar values using a principal component analysis of 2dFGRS data. for à sample of similar space density to that used for our primary constraint.," We also obtain similar values using a principal component analysis of 2dFGRS data, for a sample of similar space density to that used for our primary constraint."
 We note. though. that the clustering of galaxies selected in bluer wavebands appears to be more dependent. as one might expect.," We note, though, that the clustering of galaxies selected in bluer wavebands appears to be more model-dependent, as one might expect."
 Our estimate of c looks high compared to the values obtained by many recent measurements. in particular those from he WMAP experiment.," Our estimate of $\sigma_8$ looks high compared to the values obtained by many recent measurements, in particular those from the WMAP experiment."
 While we note that this tension has some interesting consequences if it persists. we have also pointed out YW apparent inconsistencies between our results and other low redshift constraints may arise.," While we note that this tension has some interesting consequences if it persists, we have also pointed out how apparent inconsistencies between our results and other low redshift constraints may arise."
 Small and intermediate scales in he correlation functions contribute strongly to 47 and hence to our constraint on as. and vet are not as well understood as he large scales.," Small and intermediate scales in the correlation functions contribute strongly to $\chi^2$ and hence to our constraint on $\sigma_8$, and yet are not as well understood as the large scales."
 This is clearly an area where further modelling effort is required., This is clearly an area where further modelling effort is required.
" Moreover. we will not be completely assured hat semi-analytic models capture the phenomenology of the galaxy population sufficiently well for high precision cosmological constraints until they are able to match the observed colour- and uminosity-dependent clustering o""galaxies."," Moreover, we will not be completely assured that semi-analytic models capture the phenomenology of the galaxy population sufficiently well for high precision cosmological constraints until they are able to match the observed colour- and luminosity-dependent clustering of galaxies."
" The models need to be able to reproduce the properties o""the observed galaxy population on a halo-by-halo basis. not just he properties averaged spatially or over luminosity."," The models need to be able to reproduce the properties of the observed galaxy population on a halo-by-halo basis, not just the properties averaged spatially or over luminosity."
 This implies that if cosmological parameters can be tightly constrained by other techniques. measurements of galaxy clustering will continue to orovide stringent tests for models of galaxy formation.," This implies that if cosmological parameters can be tightly constrained by other techniques, measurements of galaxy clustering will continue to provide stringent tests for models of galaxy formation."
 Most of the work for this paper was undertaken while GH was in receipt of a PhD studentship from the Particle Physics and Astronomy Research Council., Most of the work for this paper was undertaken while GH was in receipt of a PhD studentship from the Particle Physics and Astronomy Research Council.
 Thanks to Peder Norberg for providing us with his principal component analysis of the 2dFGRS data. and his help with the use thereof.," Thanks to Peder Norberg for providing us with his principal component analysis of the 2dFGRS data, and his help with the use thereof."
 David Weinberg helped to initiate and plan this project. and provided the SDSS data on which our main analysis is based.," David Weinberg helped to initiate and plan this project, and provided the SDSS data on which our main analysis is based."
We note that no calculation to date shows any such peak. so it already seems very likely that any these instabilities are simply overwhelmed by the turbulent Hall cascade.,"We note that no calculation to date shows any such peak, so it already seems very likely that any these instabilities are simply overwhelmed by the turbulent Hall cascade."
 Nevertheless. we test this idea here in greater detail. by carefully selecting the initial conditions to favour the development of a Hall instability. and inhibit the development of a Hall cascade (at least initially).," Nevertheless, we test this idea here in greater detail, by carefully selecting the initial conditions to favour the development of a Hall instability, and inhibit the development of a Hall cascade (at least initially)."
 However. we find that even under these optimised circumstances. there is no evidence that Hall instabilities play any significant role in Eq. (1)).," However, we find that even under these optimised circumstances, there is no evidence that Hall instabilities play any significant role in Eq. \ref{eq:A}) )."
 R&GG begin by considering a large-scale background field Bo. which has no Hall term. that ts. it must satisfy which ts of course already a rather restrictive assumption.," G begin by considering a large-scale background field ${\bf B}_0$, which has no Hall term, that is, it must satisfy which is of course already a rather restrictive assumption."
 They next linearise Eq. (1)), They next linearise Eq. \ref{eq:A}) )
 about this background field to obtain describing the behaviour of small perturbationsb., about this background field to obtain describing the behaviour of small perturbations.
 If one finally ignores the very gradual Ohmic decay of Bo. and instead treats it as being constant in time. then Eq. (3))," If one finally ignores the very gradual Ohmic decay of ${\bf B}_0$, and instead treats it as being constant in time, then Eq. \ref{eq:B}) )"
 becomes a standard linear eigenvalue problem. with solutions that either decay or grow exponentially.," becomes a standard linear eigenvalue problem, with solutions that either decay or grow exponentially."
 Furthermore. whereas in Eq. (1))," Furthermore, whereas in Eq. \ref{eq:A}) )"
 the Hall term conserves magnetic energy /[B|-4V. in Eq. (3) ," the Hall term conserves magnetic energy $\int|{\bf B}|^2dV$, in Eq. \ref{eq:B}) )"
the now two linearised Hall terms do not conserve [IbPav., the now two linearised Hall terms do not conserve $\int|{\bf b}|^2dV$.
 It is indeed possible therefore to obtain exponentially growing solutions. corresponding to a transfer of energy from the background field to the perturbation.," It is indeed possible therefore to obtain exponentially growing solutions, corresponding to a transfer of energy from the background field to the perturbation."
 What we wish to consider in this work ts the subsequent nonlinear evolution of these perturbations. including also the no longer background field.," What we wish to consider in this work is the subsequent nonlinear evolution of these perturbations, including also the no longer constant-in-time background field."
 R&GG consider a plane layer geometry. periodic in: x and y and bounded in z. with either vacuum or perfectly conducting boundaries at the top and bottom.," G consider a plane layer geometry, periodic in $x$ and $y$ and bounded in $z$, with either vacuum or perfectly conducting boundaries at the top and bottom."
 For their background field. they specify By=f(<)ey.," For their background field, they specify ${\bf B}_0 = f(z)\,{\bf e}_x$."
 This satisfies Eq. (2)), This satisfies Eq. \ref{ZeroHall}) )
" for any choice of f(<). and also has the additional advantage of decoupling the horizontal wavenumbers κι &, for b (because Bo is independent of x and v)."," for any choice of $f(z)$, and also has the additional advantage of decoupling the horizontal wavenumbers $k_x$, $k_y$ for $\bf b$ (because ${\bf B}_0$ is independent of $x$ and $y$ )."
 Equation (3)) has therefore been reduced to a linear. one-dimensional eigenvalue problem. in which only the z structure still has to be solved.," Equation \ref{eq:B}) ) has therefore been reduced to a linear, one-dimensional eigenvalue problem, in which only the $z$ structure still has to be solved."
" For suitable choices of f(z) in particular with at least quadratic curvature. so that the second derivative £""z0. they then find that one can indeed obtain exponentially growing modes b. that ts. instabilities of the large-scale field Bo."," For suitable choices of $f(z)$ in particular with at least quadratic curvature, so that the second derivative $f''\neq0$, they then find that one can indeed obtain exponentially growing modes $\bf b$, that is, instabilities of the large-scale field ${\bf B}_0$."
" However. these instabilities only occur for horizontal wavenumbers &, and &, up to around 5 or so. which is nowhere nearly large enough to be considered truly small-scale."," However, these instabilities only occur for horizontal wavenumbers $k_x$ and $k_y$ up to around 5 or so, which is nowhere nearly large enough to be considered truly small-scale."
 The z structure is also not really small-scale. except for a narrow boundary layer that forms at large Bo.," The $z$ structure is also not really small-scale, except for a narrow boundary layer that forms at large $B_0$."
 On the scale of the whole neutron star. the crust 1s only a thin layer. so the large scales of R&GG could already be considered to be relatively small.," On the scale of the whole neutron star, the crust is only a thin layer, so the large scales of G could already be considered to be relatively small."
 But in the context of Hall cascades versus instabilities. which is of interest here. R&GG's own," But in the context of Hall cascades versus instabilities, which is of interest here, G's own"
specified by both the part of fu)(yp.ers) where ClaseC)sing and by the part where [ensery)sun? (Moerrifield Ixuijken 1994).,"specified by both the part of $F_{maj}(r_p, v_{los})$ where $|v_{los}|
> v_c(r_p)\sin i$ and by the part where $|v_{los}| < v_c(r_p)\sin i$ (Merrifield Kuijken 1994)."
 Using just one part of the LOSVD allows us to simplify the quadrature in equation (4))., Using just one part of the LOSVD allows us to simplify the quadrature in equation \ref{abelinv}) ).
" For the high-velocity side of he LOSVD in a realistic disk. f,(M...5) is a monotonically decreasing function of Wo, at a given. value of the angular momentum £."," For the high-velocity side of the LOSVD in a realistic disk, $f_\phi(W_\phi,L)$ is a monotonically decreasing function of $W_\phi$ at a given value of the angular momentum $L$."
" One can show that any positive distribution 'unction such that oy vields a tangential velocity distribution. f£,(M.L) with the above property."," One can show that any positive distribution function such that < 0 yields a tangential velocity distribution $f_\phi(W_\phi,L)$ with the above property."
 In stellar dynamics one usually considers systems with distribution. functions that obey equation (7))., In stellar dynamics one usually considers systems with distribution functions that obey equation \ref{dfeneg}) ).
 Definite criteria for the stability of spherical and axisvmmoetric stellar svstems have been established for distribution functions which are decreasing functions of the οποιον (CXntonov. 1962)., Definite criteria for the stability of spherical and axisymmetric stellar systems have been established for distribution functions which are decreasing functions of the energy (Antonov 1962).
 Moreover it is found that most distribution functions used. to. produce realistic models. of ealaxies do actually obey that condition., Moreover it is found that most distribution functions used to produce realistic models of galaxies do actually obey that condition.
 On the other hand. very little work. either numerical or theoretical. has been done using clistribution Functions which do not satisfy such a condition and. very little is known about the stability of these svstems (Llénnon 1973. Perez Aly 1996).," On the other hand, very little work, either numerical or theoretical, has been done using distribution functions which do not satisfy such a condition and very little is known about the stability of these systems (Hénnon 1973, Perez Aly 1996)."
 Therefore. in our paper. we follow the generalv accepted. assumption given by equation (7)).," Therefore, in our paper, we follow the generally accepted assumption given by equation \ref{dfeneg}) )."
" In this case WW), can be written as à monotonic function of £F(M.L) at given L."," In this case $W_\phi$ can be written as a monotonic function of $f_\phi(W_\phi,L)$ at given $L$."
 This property allows us to transform the above relation into (E.L) ," This property allows us to transform the above relation into f(E,L) =."
Provided that the function. {ΕΕονL) can be numerically inverted to eive WCG.L). this equation means that the DE can be calculated with a simple one-dimensional integration. bypassing the computationally-unstable estimation of the derivative of f.0V..£) in equation (4)).," Provided that the function $f_\phi(W_\phi,L)$ can be numerically inverted to give $W_\phi(f_\phi,L)$, this equation means that the DF can be calculated with a simple one-dimensional integration, bypassing the computationally-unstable estimation of the derivative of $f_\phi(W_\phi,L)$ in equation \ref{abelinv}) )."
 As mentioned in the Introduction. an inclined thin disk is rather simpler than most realistic galaxy models. since any line of sight will only intersect it at a single point.," As mentioned in the Introduction, an inclined thin disk is rather simpler than most realistic galaxy models, since any line of sight will only intersect it at a single point."
 Further. for an inclined svstem. one cannot tell from the photometry whether the assumption that the disk is thin is a valid one itis only when the svstem is viewed edge-on that its razor-thinness will be unambiguously apparent.," Further, for an inclined system, one cannot tell from the photometry whether the assumption that the disk is thin is a valid one – it is only when the system is viewed edge-on that its razor-thinness will be unambiguously apparent."
 We therefore now turn to consider the case of such a system. viewed. edge-on., We therefore now turn to consider the case of such a system viewed edge-on.
 In this case. the LOSVD is not just a rescalecl version of [O8LE) as in equation r (5)). but instead. is generated by an integration along the line of sight through the entire disk.," In this case, the LOSVD is not just a rescaled version of $f_\phi(W_\phi,L)$ as in equation \ref{Fmaj}) ), but instead is generated by an integration along the line of sight through the entire disk."
" ὃν choosing the z-axis as the direction of the line of sight. one can write the LOSVD at some distance +, from the centre of the galaxy as Clos) = dz de, .fGZ.L) (O)where (Css01) are the components of the velocity parallel ancl perpendicular to the direction. of the line of sight respectively."," By choosing the $z$ -axis as the direction of the line of sight, one can write the LOSVD at some distance $r_p$ from the centre of the galaxy as ) = dz d f(E,L) where $(v_{los},\vpe)$ are the components of the velocity parallel and perpendicular to the direction of the line of sight respectively."
" “Phe two components of the velocity (6i.6s) can be written in terms of (ej,01) as: yih = sind= ror The integrals> of motion are then The extra integration along the line of sight clearly complicates the relationship between the observed kinematics and the intrinsic ονnamics of the galaxy."," The two components of the velocity $(v_r,v_\phi)$ can be written in terms of $(v_{los},\vpe)$ as: with = = r. The integrals of motion are then The extra integration along the line of sight clearly complicates the relationship between the observed kinematics and the intrinsic dynamics of the galaxy."
 llowever. the properties that we iive derived. above. for inclined. disks suggest à wav to proceed.," However, the properties that we have derived above for inclined disks suggest a way to proceed."
" In. particular. the [act that the DE is completely specified by the high-velocity tail of f£, can be used to our advantage."," In particular, the fact that the DF is completely specified by the high-velocity tail of $\tilde f_\phi$ can be used to our advantage."
 In general. the nmean-streaming circular motions of the stars in a realistic clisk will be significantly greater than their random motions. so one can think of stars at each radius moving with a mean streaming velocity of close to the local circular speed. with a smaller amount of randoni moion superin.»osed.," In general, the mean-streaming circular motions of the stars in a realistic disk will be significantly greater than their random motions, so one can think of stars at each radius moving with a mean streaming velocity of close to the local circular speed, with a smaller amount of random motion superimposed."
" “Thus. the LOSVD at a projected racius ry in an edge-on disk wil be peaked at à velocity close to vb,(πμ)."," Thus, the LOSVD at a projected radius $r_p$ in an edge-on disk will be peaked at a velocity close to $v_c(r_p)$."
 At lower line-of-sigh velocities. there will be contributions both from stars tha have intrinsically low velocities and from stars that lic a large raclii in the galaxy. so that twir large circular motions are oriented mostly transverse to the line of sight. resulting in a small line-of-sight component.," At lower line-of-sight velocities, there will be contributions both from stars that have intrinsically low velocities and from stars that lie at large radii in the galaxy, so that their large circular motions are oriented mostly transverse to the line of sight, resulting in a small line-of-sight component."
 However. at line-of-igh velocities greater than e(ri). the majority of the stars contributing to the LOSVD will be those where the circular motion is oriented along the linc-o[-sight. which occurs only [or stars at radii rry.," However, at line-of-sight velocities greater than $v_c(r_p)$, the majority of the stars contributing to the LOSVD will be those where the circular motion is oriented along the line-of-sight, which occurs only for stars at radii $r \sim r_p$."
 Further. for the stars at radii close to ry. the linc-of-sight component of their random motions measures their velocities in the oO clireetion.," Further, for the stars at radii close to $r_p$, the line-of-sight component of their random motions measures their velocities in the $\phi$ direction."
 Thus. to a simple approximation. [or (js2te. Viri ens) LAr Up = Clos]AL. (15)where A is the length of path through the galaxy that lies sullictently close tor—ry for there to be a significant contribution to f(r.ciu).," Thus, to a simple approximation, for $v_{los} >
v_c$, F(r_p, ) (r = r_p, = ), where $\Lambda$ is the length of path through the galaxy that lies sufficiently close to $r = r_p$ for there to be a significant contribution to $F(r_p, v_{los})$."
 In terms used in the study of the Milky. Was. this path-leneth covers the region of the tangent point. over which n will not change rapidly (since r is not changing significantly). so the integral along the line of sight can be replaced. by the simple product of equation (15)).," In terms used in the study of the Milky Way, this path-length covers the region of the tangent point, over which $\tilde f_\phi$ will not change rapidly (since $r$ is not changing significantly), so the integral along the line of sight can be replaced by the simple product of equation \ref{Lambdadef}) )."
 Thus. equation (15)) implies that by observing the part of the LOSVD at high velocities. we obtain an estimate for the shape of the high- part of the tangential velocity. distribution. which we need to estimate the DE.," Thus, equation\ref{Lambdadef}) ) implies that by observing the part of the LOSVD at high velocities, we obtain an estimate for the shape of the high-velocity part of the tangential velocity distribution, which we need to estimate the DF."
a given object. unlike the DNOs no clear period-Iuminosity relationship has vet been deduced.,"a given object, unlike the DNOs no clear period-luminosity relationship has yet been deduced."
 Our new observations of VW Livi provide the first evidence for such a relationship., Our new observations of VW Hyi provide the first evidence for such a relationship.
 QPOs are rarely. visible at high energies: four distinct possibilities are of. amplitude at 585 s in U Gen (Cordova Mason. LOS4). of very low amplitude at S3 s in ds Cve (Mauche 1997). of 2240 s in OY Car just after an outburst (Ramsay et al.," QPOs are rarely visible at high energies: four distinct possibilities are of amplitude at 585 s in U Gem (Cordova Mason 1984), of very low amplitude at 83 s in SS Cyg (Mauche 1997), of 2240 s in OY Car just after an outburst (Ramsay et al."
 2001) and a modulation increasing in period. from 63 s to 68 s and in amplitude from lo in VW Livi (van der Woerd et al., 2001) and a modulation increasing in period from 63 s to 68 s and in amplitude from to in VW Hyi (van der Woerd et al.
 LOST)., 1987).
 Wheatley οἱ al. (, Wheatley et al. (
1996) observed flux variations of large amplitude and time scale 7500 s in Ginga (2-10 keV) X-ray observations of VW List made at the end of outburst.,1996) observed flux variations of large amplitude and time scale $\sim$ 500 s in Ginga (2-10 keV) X-ray observations of VW Hyi made at the end of outburst.
" As this is the place in the light curve where we see optical QPOs in VW νι, it is extremely probable that the N-rav observations are a dillerent manifestation of the same phenomenon."," As this is the place in the light curve where we see optical QPOs in VW Hyi, it is extremely probable that the X-ray observations are a different manifestation of the same phenomenon."
 We return to this later., We return to this later.
 Other possible X-ray. QPOs are. 290 s in AB Dra on the rise to outburst. 121 and 135 s in U Gem at quiescence and 254 s in the nova-like RW Sex (Cordova Mason 1984).," Other possible X-ray QPOs are 290 s in AB Dra on the rise to outburst, 121 and 135 s in U Gem at quiescence and 254 s in the nova-like RW Sex (Cordova Mason 1984)."
 lt is possible that there is more than one cause for the QPOs., It is possible that there is more than one cause for the QPOs.
 For example. it has been noted that the longest QPOs have periods close to the expected rotation periods at the outer edges of the accretion disces in CVs (Warner 1995a). and Lasota. Ixuulkers Charles (1999) have suggested a mocel for one of the DNOs in WZ See which uses a plasma blob at the disc rim.," For example, it has been noted that the longest QPOs have periods close to the expected rotation periods at the outer edges of the accretion discs in CVs (Warner 1995a), and Lasota, Kuulkers Charles (1999) have suggested a model for one of the DNOs in WZ Sge which uses a plasma blob at the disc rim."
 Nevertheless. it is also possible that all QPOs are caused by oscillations in the inner accretion disc.," Nevertheless, it is also possible that all QPOs are caused by oscillations in the inner accretion disc."
 Perturbation analyses of such dises by Carroll et al. (, Perturbation analyses of such discs by Carroll et al. (
1985). Lubow Prinele (1993). and Collins. Helfer van Horn (2000) show the possibility of a wide spectrum of non-raclial oscillations. analogous to p-mocde and g-mocde oscillations in stars.,"1985), Lubow Pringle (1993), and Collins, Helfer van Horn (2000) show the possibility of a wide spectrum of non-radial oscillations, analogous to p-mode and g-mode oscillations in stars."
 The brightness oscillations are usually ascribed to intrinsic Luminosity variations in the cise itself. but we point out in Paper LL that they may also be caused by variation ofinfercepfed radiation from the central high luminosity region of the disc. if there are vertical oscillations in thickness of some cise aunuli.," The brightness oscillations are usually ascribed to intrinsic luminosity variations in the disc itself, but we point out in Paper II that they may also be caused by variation of radiation from the central high luminosity region of the disc, if there are vertical oscillations in thickness of some disc annuli."
 Vhe cdwarl nova VW. Livi is a rich. source of short. period luminosity modulations., The dwarf nova VW Hyi is a rich source of short period luminosity modulations.
 It is the CV. in which DNOs were first. directly. observable in the light curve (Warner Llarwoocl 1973) (rather than only in the Fourier transform)., It is the CV in which DNOs were first directly observable in the light curve (Warner Harwood 1973) (rather than only in the Fourier transform).
 Unlike most dwarf novae. the optical DNOs in VW. νι appear most conspicuously in the final stage of decline from an outburst (they could be present at a similar brightness at the start of an outburst. but no high time resolution photometry has been made that early) anc follow a period-luminosity correlation over the range 2036 s (Warner Brickhill 1974. 1978: Robinson Warner 1984: Haefner. Schoembs Voet 1979: Schoembs Vogt. 1980).," Unlike most dwarf novae, the optical DNOs in VW Hyi appear most conspicuously in the final stage of decline from an outburst (they could be present at a similar brightness at the start of an outburst, but no high time resolution photometry has been made that early) and follow a period-luminosity correlation over the range 20–36 s (Warner Brickhill 1974, 1978; Robinson Warner 1984; Haefner, Schoembs Vogt 1979; Schoembs Vogt 1980)."
PF9-00065 awarded by the Chaudra XN-rav Center. which is operated by the Sinithsonian Astroplivsical Observatory for NASA uuder contract NASS-03060.,"PF9-00065 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060."
 E.B. acknowledges support for this work from the National Science Foundation through αμ AST-, E.B. acknowledges support for this work from the National Science Foundation through Grant AST-1107973.
neighbours (zz0.5 mae bluer in VI than the bluest field stars},neighbours $\approx 0.5$ mag bluer in $V-R$ than the bluest field stars).
 We do not believe our photometry sullers from this oblenm since the coellicients of the “colour-ainmass” terms in the transformations (equation 1) are all consistent with zero., We do not believe our photometry suffers from this problem since the coefficients of the “colour-airmass” terms in the transformations (equation 1) are all consistent with zero.
 Several of the standard stars we observed. are bluer han IXPD 042215421 RAV 149 with BY=0.129 and (2B=/—0.779. Landolt 1992) ancl were observed to zürlv high airmass (7 2). so we would be able to measure colour-airmass terms that were significantly. dillerent. [rom Zero.," Several of the standard stars we observed are bluer than KPD 0422+5421 RU 149 with $B-V=-0.129$ and $U-B=-0.779$, Landolt 1992) and were observed to fairly high airmass $\approx 2$ ), so we would be able to measure colour-airmass terms that were significantly different from zero."
" As an experiment. we considered. only the IKPD M22,|5421 observations obtained at airmass [ess than 1.4 (1977 images) and analvzed the residual light curve tthe time series ata with the ellipsoidal componen removed)."," As an experiment, we considered only the KPD 0422+5421 observations obtained at airmass less than 1.4 (1977 images) and analyzed the residual light curve the time series data with the ellipsoidal component removed)."
 “Phe resicuals still showed a periodicity a 7.8 hours and the residuals folded. on 7.8 hours. closely resembled. the folded. curve shown in Figure 3.., The residuals still showed a periodicity at 7.8 hours and the residuals folded on 7.8 hours closely resembled the folded curve shown in Figure \ref{fig3}.
. Thus the TS hour modulation does not seem to be an artifact. of colour-dependent extinction., Thus the 7.8 hour modulation does not seem to be an artifact of colour-dependent extinction.
 None of the other bright field stars show any apparen coherent periodicity., None of the other bright field stars show any apparent coherent periodicity.
 We show the O statistic for a comparison star in Ligure 2., We show the $\Theta$ statistic for a comparison star in Figure 2.
 The minimum © value is about 0.89., The minimum $\Theta$ value is about 0.89.
 Furthermore. the O values are greater than 0.95 at trial periods near 7.8 hours.," Furthermore, the $\Theta$ values are greater than 0.95 at trial periods near 7.8 hours."
 Thus the [act thatony IXPD 0422/5421 shows an apparent coherent perioclicity is an indication that the cause is not instrumental. and would have to be related to the observing conditions (Lor example the potential svsteniatic error associated with the Xue colour of IKPPD 0422|5421 discussed above).," Thus the fact that KPD 0422+5421 shows an apparent coherent periodicity is an indication that the cause is not instrumental, and would have to be related to the observing conditions (for example the potential systematic error associated with the blue colour of KPD 0422+5421 discussed above)."
 Llowever. if he extra signal in the KWPD 0422|5421 light curve is indeed an artifact of the observing conditions and/or procedure. hen it is not obvious how this signal could have a periodicity ab 7.8 hours.," However, if the extra signal in the KPD 0422+5421 light curve is indeed an artifact of the observing conditions and/or procedure, then it is not obvious how this signal could have a periodicity at 7.8 hours."
 One might be able to invent a periodicity at 7.98 hours (one-third of a sidereal day) if the signal were due o the rapid change of airmass as the source sets., One might be able to invent a periodicity at 7.98 hours (one-third of a sidereal day) if the signal were due to the rapid change of airmass as the source sets.
 “Phe fact hat the periocicity of the signal is not obviously. related o the rotation of the Earth is further evidence that it is intrinsic to WPL) 0422|5421 and not an artifact., The fact that the periodicity of the signal is not obviously related to the rotation of the Earth is further evidence that it is intrinsic to KPD 0422+5421 and not an artifact.
 With the 7.8 hour modulation removed from the time series data. it was possible to properly analyze the ellipsoidal light curve shown in Figure 4..," With the 7.8 hour modulation removed from the time series data, it was possible to properly analyze the ellipsoidal light curve shown in Figure \ref{fig4}."
" We mocdelled the folded light curve using ""mode 27 of the W-D code. which is the usual mode used For detached binaries."," We modelled the folded light curve using “mode 2” of the W-D code, which is the usual mode used for detached binaries."
 We assumed a circular orbit and synchronous rotation., We assumed a circular orbit and synchronous rotation.
 The polar temperature of the sd star was fixed at the spectroscopically determined. value of, The polar temperature of the sdB star was fixed at the spectroscopically determined value of
"OPO were LOEOO.7 and 3,900.6 respectively.","QPO were 0.7 and 0.6, respectively."
 Iu the euergy range 822L1 keV the lower frequeney QPO was not detectable (amplitude upper limit of 10 rns) the raus amplitude of the higher frequeucy OPO was 5.1-11., In the energy range 8.2–24.1 keV the lower frequency QPO was not detectable (amplitude upper limit of 10 rms); the rms amplitude of the higher frequency QPO was 1.1.
1 The VLFN during the Anenst 1 observation had a rums amplitude of 5.9400.6 and a power law index of L.61400.07., The VLFN during the August 1 observation had a rms amplitude of 0.6 and a power law index of 0.07.
 The peaked ΠΕΝ had a amplitude of L6t00.1 rius. a power lew index of 0.82500.6. and a cut-off frequency of 224113 Iz.," The peaked HFN had a amplitude of 0.4 rms, a power law index of 0.6, and a cut-off frequency of 13 Hz."
 The VLEN divingof the August 31 observation had au rims amplitude ο) and a power law mdex of 1.51-00.06.," The VLFN during the August 31 observation had an rms amplitude of 0.6, and a power law index of 0.06."
 Although no ΠΕΝ couldbe detected in that observation a OPO at the frequeney 26.9+22.3 Iz is seen du the power spectrum (Fig. 1..," Although no HFN could be detected in that observation a QPO at the frequency 2.3 Hz is seen in the power spectrum (Fig. \ref{powerspectrum},"
 right)., right).
 The FWIIM aud amplitude of this OPO is 11.14£55.2 Iz aud LE00.6 (5.72L1 keV)., The FWHM and amplitude of this QPO is 5.2 Hz and 0.6 (5.7–24.1 keV).
 We detected for the first time two simultaneous kilohertz QPOs in KS 260. at frequencies near SOS Tz aud 1159 Uz.," We detected for the first time two simultaneous kilohertz QPOs in KS $-$ 260, at frequencies near 898 Hz and 1159 Hz."
 The amplitudes of these QPOs increase with increasing photon enerev., The amplitudes of these QPOs increase with increasing photon energy.
 Both QPOs are Visible at the lowest count rates., Both QPOs are visible at the lowest count rates.
 At higher coun rates the lower frequency QPO disappears. the amplitude of the higher frequency QPO decreases. iux u-s frequency increases.," At higher count rates the lower frequency QPO disappears, the amplitude of the higher frequency QPO decreases, and its frequency increases."
 At the highest count rates the higher frequency QPO is only marginally detected at 120722111 Tz., At the highest count rates the higher frequency QPO is only marginally detected at 11 Hz.
 If that ΟΡΟ is real then its frequency is the lighes so far observed in an N-ray binary., If that QPO is real then its frequency is the highest so far observed in an X-ray binary.
 When the lower frequency ΟΡΟ is detected its frequency is cousisteu with cine constant., When the lower frequency QPO is detected its frequency is consistent with being constant.
 The munor differences in its frequency between slightly ifercut data sets are within the statistical errors., The minor differences in its frequency between slightly different data sets are within the statistical errors.
 When we combine all power spectra from the August 1 observation. the peak separation of the two QPOs (276.1433.8the Tz) is not consistent with being equal to half frequency. ofthe 523.92+00.05 Tz oscill:ious found by Sunith et al. (," When we combine all power spectra from the August 1 observation, the peak separation of the two QPOs 3.8 Hz) is not consistent with being equal to half the frequency of the 0.05 Hz oscillations found by Smith et al. ("
1997) in a type I burst.,1997) in a type I burst.
 However. this incousistencev is artifical.," However, this inconsistency is artifical."
 Iu order to obtain a better signal-to-noise ratio we sunuuced the power spectra of all three data seeimoeuts obtained during this observation., In order to obtain a better signal-to-noise ratio we summed the power spectra of all three data segments obtained during this observation.
 During the first scement we see only the higher: frequency. QPOApe near --.1176 Iz. during: the second seginent we see both QPOs↽↽ uear δ0δ aud 1159 Tz. aud durius the third scement ouly the higher frequency peak can be sieuificantly etected near 1183 Tz.," During the first segment we see only the higher frequency QPO near 1176 Hz, during the second segment we see both QPOs near 898 and 1159 Hz, and during the third segment only the higher frequency peak can be significantly detected near 1183 Hz."
 Combination of these three ata segments elves rise to an artifical imerease of the peak separation. due to the different focus in time of the two QPOs.," Combination of these three data segments gives rise to an artifical increase of the peak separation, due to the different focus in time of the two QPOs."
 When both Xilohertz QPOs ie present (scgment 2) the peals separation is 260.32299.6 Uz. which is aliuost exactly halfthe oscillation frequency (261.96) during the type IE burst.," When both kilohertz QPOs are present (segment 2) the peak separation is 9.6 Hz, which is almost exactly half the oscillation frequency (261.96) during the type I burst."
 This result stronely supports models in which the neutron spin frequency is 261.96 Iz {3.8 ic higher frequeney QPO a preferred. orbital frequencymus). around the neutrou star. and the lower frequency OPO the beat between these two.," This result strongly supports models in which the neutron spin frequency is 261.96 Hz (3.8 ms), the higher frequency QPO a preferred orbital frequency around the neutron star, and the lower frequency QPO the beat between these two."
 Such a model. with the preferred. orbital radius the magnetospheric radius. was already suggested for the three QPOs found in IU 31 (Strolunaver et al.," Such a model, with the preferred orbital radius the magnetospheric radius, was already suggested for the three QPOs found in 4U $-$ 34 (Strohmayer et al."
 1996)., 1996).
 Miller. Laaah Psaltis (1997) proposed a model where the preferred radius is the sonic radius.," Miller, Lamb Psaltis (1997) proposed a model where the preferred radius is the sonic radius."
 A problem for these models ix the case of Sco N-1. in which the frequency difference is not constant (van der Klis et al.," A problem for these models is the case of Sco X-1, in which the frequency difference is not constant (van der Klis et al."
 1997) as would be expected., 1997) as would be expected.
The models which have been used. to explain the nature of RX J1914|24 and RA JOSOG|15 fall into two gencral categories: those involving accretion and those which do not.,The models which have been used to explain the nature of RX J1914+24 and RX J0806+15 fall into two general categories: those involving accretion and those which do not.
 At [ace value. the results which show both IX J1914|24 and UN. JOSOG|15 spinning up appear to rule out most of the accretion driven models: the double degenerate polar mocel (Cropper et al 1998) and the direct impact. model (Marsh Steeghs 2002. Ramsay et al 2002).," At face value, the results which show both RX J1914+24 and RX J0806+15 spinning up appear to rule out most of the accretion driven models: the double degenerate polar model (Cropper et al 1998) and the direct impact model (Marsh Steeghs 2002, Ramsay et al 2002)."
 Unlike these moclels. in which the observed. intensity. variation is attributed. to he binary orbital period. Norton. Llaswell Wynn (2004) attribute it to the spin period of an accreting white dwarf.," Unlike these models, in which the observed intensity variation is attributed to the binary orbital period, Norton, Haswell Wynn (2004) attribute it to the spin period of an accreting white dwarf."
 o second. period is detected since in this model the binary inclination is close to face-on., No second period is detected since in this model the binary inclination is close to face-on.
 The non-accreting model is the Unipola-Inductor (UL) model proposed by Wu ct al (2002)., The non-accreting model is the Unipolar-Inductor (UI) model proposed by Wu et al (2002).
 In this mocel. a non-magnetic white cwarf traverses the magnetic field. of the primary magnetic white dwarf. which causes large currents to be driven.," In this model, a non-magnetic white dwarf traverses the magnetic field of the primary magnetic white dwarf, which causes large currents to be driven."
 Resistive dissipation occurs at the foot-points of the primary white dwarf and. N-ravs are released., Resistive dissipation occurs at the foot-points of the primary white dwarf and X-rays are released.
 This nmiocdel also predicts that the svstem is an electron-cvclotron maser source in the UL phase (Willes. Wu Ixuncic 2004).," This model also predicts that the system is an electron-cyclotron maser source in the UI phase (Willes, Wu Kuncic 2004)."
 Lt is important to note that these svstems have been observed. for less than LO vears and therefore the period decrease may simply be due to a slow variation rather than a long term trend., It is important to note that these systems have been observed for less than 10 years and therefore the period decrease may simply be due to a slow variation rather than a long term trend.
 In spite of the complicated. interaction between the stellar spin rates and the binary orbital period. the change in the binary orbital period is determined only by the energy and angular momentum losses and redistribution in the svsteor.," In spite of the complicated interaction between the stellar spin rates and the binary orbital period, the change in the binary orbital period is determined only by the energy and angular momentum losses and redistribution in the system."
 We illustrate this using the Ul mocdel given in Wu et al (2002)., We illustrate this using the UI model given in Wu et al (2002).
" For a binary with a secondary star in svnehronous rotation with the orbit. it can be shown that the change in the orbital angular velocity a, is given by where Loy is the energy. loss due to gravitational radiation. fea. is the energv loss due to dissipation. and aw, is the angular velocity of the primary (sec Wu et al (2002) for the explicit expression of g(w.,))."," For a binary with a secondary star in synchronous rotation with the orbit, it can be shown that the change in the orbital angular velocity $\omega_{\rm o}$ is given by where ${\dot E}_{\rm gr}$ is the energy loss due to gravitational radiation, ${\dot E}_{\rm diss}$ is the energy loss due to dissipation, and $\omega_{1}$ is the angular velocity of the primary (see Wu et al (2002) for the explicit expression of $g(\omega_{\rm o})$ )."
 In the Ul model. the energy. dissipation is caused by the resistive heating at the foot-points of the magnetic field lines on the magnetic white cbwarf.," In the UI model, the energy dissipation is caused by the resistive heating at the foot-points of the magnetic field lines on the magnetic white dwarf."
 In terms of the asynchronicity parameter of the spin of the primary. a(=wife). the energy dissipation ]t can be shown that As (Gofw)=(wife)|(Aga). the rate of change in the orbital angular velocity may he expressed as ‘This implies that the secular evolution of the orbital angular velocity is also determined. by the coupling. between the spin of the primary ancl the orbital rotation in spite of the ultimate driver being the encrey and angular momentuni losses due to gravitational radiation.," In terms of the asynchronicity parameter of the spin of the primary, $\alpha~(\equiv \omega_1/ \omega_{\rm o})$, the energy dissipation It can be shown that As $({\dot \omega}_1/\omega_1) = ({\dot \omega}_{\rm
o}/\omega_{\rm o}) + ({\dot \alpha}/\alpha)$, the rate of change in the orbital angular velocity may be expressed as This implies that the secular evolution of the orbital angular velocity is also determined by the coupling between the spin of the primary and the orbital rotation in spite of the ultimate driver being the energy and angular momentum losses due to gravitational radiation."
 More detailed work is needed to determine the range of a which could give rise to the observed spin-up rate., More detailed work is needed to determine the range of $\alpha$ which could give rise to the observed spin-up rate.
 The lower panel of Figure 3. shows that the observed count rate of RN J1914|24 varies by an order of magnitude., The lower panel of Figure \ref{pdot} shows that the observed count rate of RX J1914+24 varies by an order of magnitude.
 At the epoch. the unabsorbed. bolometric [ux determined using the model fits is shown in Table 2..," At the epoch, the unabsorbed, bolometric flux determined using the model fits is shown in Table \ref{fits}."
 Due to projection elfects we correct the observed count rate by the ratio of the peak to mean count rate (a factor of 3.4) which gives a luminosity of ~L1107 (for Ίκρο. Steeghs et al 2004).," Due to projection effects we correct the observed count rate by the ratio of the peak to mean count rate (a factor of 3.4) which gives a luminosity of $\sim1\times10^{35}$ (for 1kpc, Steeghs et al 2004)."
 For the same distance this is ~2 orders of magnitude lower than the X-ray luminosity estimate in Cropper et al (1998): this due to the higher absorption originally derived using theROSAL PSPC spectrum., For the same distance this is $\sim$ 2 orders of magnitude lower than the X-ray luminosity estimate in Cropper et al (1998): this due to the higher absorption originally derived using the PSPC spectrum.
 Figure 3. implies that it has reached ~4.107 lin the past., Figure \ref{pdot} implies that it has reached $\sim4\times10^{35}$ in the past.
 If the N-ray [lux is driven by aceretion and all of the, If the X-ray flux is driven by accretion and all of the
The possibility of. observiugoo eravitomaeuetic (beusce-Thirring) precession of accretion disks around compact objects was considered carly in the study of accretion outo black holes and ucutron stars.,The possibility of observing gravitomagnetic (Lense-Thirring) precession of accretion disks around compact objects was considered early in the study of accretion onto black holes and neutron stars.
 The conchisiou reached then (Dardeeu Petterson 11975: soe also Petterson 1197172. 1319715. 11978: Hatchett. Degelinau. Saraziu 11981) was that the combined action of eravitomagnctic: :id internal viscous torques forces disk flow. iuto the. rotation equator⋅ of the∩⊾↸∖∪⋯↸∖⊓⋅↕↸⋈↧∐↖⇁↑↕∐∐∙↕↘⊽↸∖⊓∖∏⋪⊔∩≼∐↴∖↴↽↴∖↴∙≼↙≽⋯⋅↖↖⇁∪↥⋅⊪∩⊾≺⋈∖↴∖↴ compact D object at —100 eravitational radii and holds it there. if the outer disk is πήραΠοιος or the disk is divenu by time-independenti; torques.," The conclusion reached then (Bardeen Petterson \markcite{Pett75}1 1975; see also Petterson \markcite{Pett77a}1 1977a, \markcite{Pett77b}1 1977b, \markcite{Pett78}1 1978; Hatchett, Begelman, Sarazin \markcite{HBS81}1 1981) was that the combined action of gravitomagnetic and internal viscous torques forces the disk flow into the rotation equator of the compact object at $\sim$ 100 gravitational radii and holds it there, if the outer disk is misaligned or the disk is driven by time-independent torques."
 As a result. it: has bec widely thought that eravitomaeuetic precession o “the inner disk is unpronmisiug as an explanaion for X“Tan⋅⇁ ni'evious," As a result, it has been widely thought that gravitomagnetic precession of the inner disk is unpromising as an explanation for X-ray oscillations."
" Recently Tpser (1996) has investigated “trapped ιο, localized: see Nato Wouma 1991) pressure perturbations that warp the imnermnost parts of accretion inks. found that the frame drageimgao near a herr black hole causes precession. of. such perturbations.. aud iscussed their possible relation to the low-frequeney quasi-periodic.⋅⋅ oscillations⋅⋅ (QPOs)""pe observed in. the X-lrav brightness. of. black hole candidates. (see vau der Ilis-Mo dar1995)."," Recently Ipser (1996) has investigated `trapped' (i.e., localized; see Kato Honma 1991) pressure perturbations that warp the innermost parts of accretion disks, found that the frame dragging near a Kerr black hole causes precession of such perturbations, and discussed their possible relation to the low-frequency quasi-periodic oscillations (QPOs) observed in the X-ray brightness of black hole candidates (see van der Klis 1995)."
 The existence⋅ of. these modes depends ou the presence of specific (but not implausible} non-Keplerian angular velocity profiles., The existence of these modes depends on the presence of specific (but not implausible) non-Keplerian angular velocity profiles.
 [pser’s geucral relativistic analvsis did not take iuto account the viscosityb of the aneas in the diskTUM., Ipser's general relativistic analysis did not take into account the viscosity of the gas in the disk.
 More recently. Stella Vietri (1998) have suggested that C»eravitomagueticOo precession of the iuner disk may be responsible for the broad bumps that are observed between 20 aud 60 Tz in the power spectra of some of the accreting neutrou stars iu low-nass binary systcis called atoll sources aud for the 565 Iz horizoutal branch quasi-periodic oscillation (ΟΡΟ)↽↽ observed iu∙ the X-ray- brightness: of: the so- ↸⊳⋜↧∐↸∖≼↧∑↴∖↴⋯∐⋅↸⊳↸∖↴∖↴⋖↴∖↴↸∖↸∖↖↽⋜⋃≺∐∖↥⋅↕↘∐∖↴↓∩∩⋅↱⊐⋟∙↕↑∐⋜↧," More recently, Stella Vietri \markcite{SV98}( (1998) have suggested that gravitomagnetic precession of the inner disk may be responsible for the broad bumps that are observed between 20 and 60 Hz in the power spectra of some of the accreting neutron stars in low-mass binary systems called atoll sources and for the $\sim$ 15--65 Hz horizontal branch quasi-periodic oscillation (QPO) observed in the X-ray brightness of the so-called Z sources (see van der Klis \markcite{vdK95}1 1995)."
↴∖↴↽⋅↽↽ been conjectured (e.g... Stella 119972. 119975) that the tilt of the iuner disk required for eravitomaguetic : : ↻↥⋅↸∖↸⊳↸∖↴∖∷∖↴↕∪∐⋯∏∐↴⋝↸∖↸⊳⋜⋯↴∖↴↸∖≼↧↸∖↕↑∐↸∖↥⋅↴⋝∙↖↽∐↸∖↥⋅⋜∥∐⋜↧⊓∪∐ ⋅⋅ torque thought to produce radiation warping of accretion disks: (see Iping∙↜ Petterson 11990: Prinele 11996: Maloney. Degeliuan. Pringle 11996: Maloney Begchnan 11997) or by forcing of the iuner disk by the neutron stars magnetic field.," It has been conjectured (e.g., Stella \markcite{Stella97a}1 1997a, \markcite{Stella97b}1 1997b) that the tilt of the inner disk required for gravitomagnetic precession could be caused either by the radiation torque thought to produce radiation warping of accretion disks (see Iping Petterson \markcite{IP90}1 1990; Pringle \markcite{Pringle96}1 1996; Maloney, Begelman, Pringle \markcite{MBP96}1 1996; Maloney Begelman \markcite{MB97}1 1997) or by forcing of the inner disk by the neutron star's magnetic field."
 Cui. Zhang. Chen ((1998) subsequeutlv suggested that the ~6300 Iz QPOs observed iu the X-ray brightuess of the ealactic black hole candidates (Miyiunoto et markeiteMvaanoto9111991: Norean. Remillard. Cremer 11997: Remillard 11997) max be produced by eravitomaguetic precession.," Cui, Zhang, Chen \markcite{CZC97}( (1998) subsequently suggested that the $\sim\,$ 6–300 Hz QPOs observed in the X-ray brightness of the galactic black hole candidates (Miyamoto et \\markcite{Miyamoto91}1 1991; Morgan, Remillard, Greiner \markcite{MRG97}1 1997; Remillard \markcite{Rem97}1 1997) may be produced by gravitomagnetic precession."
 These sugeestiousoo have renewed interest in the possibility of observing eravitomagnetic precessiou near compact objects., These suggestions have renewed interest in the possibility of observing gravitomagnetic precession near compact objects.
 Tu this. paper wereport calculations. of. the properties. of eravitomaguetic and radiation warping modes of ithe] ‘ἱ ," In this paper we report calculations of the properties of gravitomagnetic and radiation warping modes of geometrically thin, Keplerian $\alpha$ disks."
to bevond previous work iu several wavs., Our work goes beyond previous work in several ways.
" First. we lave computed the radiation warping mode functions and theireir frequencifrequencies audand ο...erowth ratrates up fo YveryTS highhie mode muubers,"," First, we have computed the radiation warping mode functions and their frequencies and growth rates up to very high mode numbers."
 Our results for the tilt functions and precession frequencies of the loworder radiation agreeresults∙ well] with previous- results(Maloues (Maloney∙ ctoscillations.↕∖↕↕ warpingal. markcitcMIBP9611996)., Our results for the tilt functions and precession frequencies of the low-order radiation warping modes agree well with previous results (Maloney et \\markcite{MBP96}1 1996).
 we have explored the time-dependent eravitomacuetic modes of the disk. including both uudiven (free) nodes aud modes that are being driven at the inner boundary.," Second, we have explored the time-dependent gravitomagnetic modes of the disk, including both undriven (free) modes and modes that are being driven at the inner boundary."
 Third. we investigated the.⋅↜ combined effects⋅ of⋅ gravit↜⋅ omaguetichave and radiation⋅⋅ torques.," Third, we have investigated the combined effects of gravitomagnetic and radiation torques."
 Finally.⋅ we have considered⋅ not ouly isothermal. disks. but also disks. with. a varicty. of. radial. temperature structures. aud inner. aud outer boundary conuditious appropriate to a variety of astroplivsical situations.," Finally, we have considered not only isothermal disks but also disks with a variety of radial temperature structures, and inner and outer boundary conditions appropriate to a variety of astrophysical situations."
" We. have ford. a fuiimE of. k""mBeqnenes. Bravitomagnetie. (ΤΕΛΙ modes with precession frequencies that rauge from the lowest frequency allowed by the size of the disk up to a certain critical frequency weir. Which is ~1 Tz for a compact object of solar mass."," We have found a family of low-frequency gravitomagnetic (LFGM) modes with precession frequencies that range from the lowest frequency allowed by the size of the disk up to a certain critical frequency $\omega_{\rm crit}$, which is $\sim$ 1 Hz for a compact object of solar mass."
 The lowest-order (fundamental) LFCAL mode is the frequency iode found earlier by Bardeen Petterson (1975)., The lowest-order (fundamental) LFGM mode is the zero-frequency mode found earlier by Bardeen Petterson (1975).
mn The other low-order LECUM) modes are loug-waveleueth. precessing∙ corrugations: of. the outeruter disk disksimilarsimular to the previously known low-order ow ↥⋅⋪∥∐⋪↧↑↕∪∐↖↖↽⋪∐⋅≻↕∐∩⊾⊔∪≼∐∖↴∖↴⋪⋯≼↧↕⊔↖↽↸∖≼↧⋪⋯∩↕∐∩↥⋅⋪↧↑↸∖↴∖↴i5 t i anus m =>LO times their precession frequencies.," The other low-order LFGM modes are long-wavelength, precessing corrugations of the outer disk similar to the previously known low-order radiation warping modes and have damping rates $\gta10$ times their precession frequencies."
 The highest-requeucy LEGAL modes are very tightly wound spiral ↸⊳∪↥⋅↥⋅∏∩⊾⋪↧↑↕∪∐↴∖↴∪≯↑∐↸∖↕∐∐∖↥⋅≼∐∖↴↽↑∐↕↑↸∖⊼↑↸∖∐≼↧∪∐↕↖⇁↑∪↕∩ eee ~ ines the MEEiuner radius: of the disk aud have damping ‘ 108 o ↥⋅⋪↧↑↸∖↴∖↴↽∩⊽≱∐⊔↸∖↴∖↴↑∐∖∐⋅∐⋅↸∖↸⊳↸∖↴∖↴↴∖↴↕∪∐↥⋅↸∖≺⋯∖∐↸⊳↕↸∖↴∖↴∙⊀lici |., The highest-frequency LFGM modes are very tightly wound spiral corrugations of the inner disk that extend only to $\sim$ 10 times the inner radius of the disk and have damping rates $\gta10^3$ times their precession frequencies.
 É. A radiation warping torque makes all the low-order nodes m this family time-dependent aud can cause a few of the lowest-frequency LEGAM modes to grow with time. but even a strong radiation warping torque las essentially no effect ou the modes with frequeucies >LO1 TIz.," A radiation warping torque makes all the low-order modes in this family time-dependent and can cause a few of the lowest-frequency LFGM modes to grow with time, but even a strong radiation warping torque has essentially no effect on the modes with frequencies $\gta10^{-4}$ Hz."
interesting feature of beiug variable. by altering the value of e.,"interesting feature of being variable, by altering the value of $\epsilon$."
 A hieh value of €c0.1 will keep AV switched off in normal thermal motions superimposed ou Ixepleriau shear. but will only switch on for particle / if ucighbourine particles ; have uou-WNeplerian approach velocities ο 4gercater than 105€ of ο," A high value of $\epsilon\simeq 0.1$ will keep AV switched off in normal thermal motions superimposed on Keplerian shear, but will only switch on for particle $i$ if neighbouring particles $j$ have non-Keplerian approach velocities $v_{ij}$ greater than $10\%$ of $v_i$."
 Ife 20.01. AV will be applied to ualler discrepancies from Ikepleriau motion. which may iuclude some thermal motions within the disk.," If $\epsilon \simeq 0.01$, AV will be applied to smaller discrepancies from Keplerian motion, which may include some thermal motions within the disk."
 Note that NK. does iof identify convergence. merolv the absence of Iepleriaun flow.," Note that $\cal{NK}$ does not identify convergence, merely the absence of Keplerian flow."
 However. as it is then used in conjunction with AV. which does detect convergence. +1ο two methods in combination are effective in applying a decelerating viscous force only when SPII particles are oe1 truly convergent flow within a Ixeplerian disk.," However, as it is then used in conjunction with AV, which does detect convergence, the two methods in combination are effective in applying a decelerating viscous force only when SPH particles are in truly convergent flow within a Keplerian disk."
 Finally. it should be noted that the flaw iu the calculation of |V«v also applies in the calculation of |V.a|.," Finally, it should be noted that the flaw in the calculation of ${\mid\nabla \cdot \bf v\mid}$ also applies in the calculation of ${\mid\nabla \cdot \bf a\mid}$."
 Lavee divergence in the acceleration or velocity fields can apparently be detected when iu fact the variation is due to noise in the particle positioning., Large divergence in the acceleration or velocity fields can apparently be detected when in fact the variation is due to noise in the particle positioning.
 Both quantities should be used with caution. particularly as diagnostic indicators in the creation of sink particles.," Both quantities should be used with caution, particularly as diagnostic indicators in the creation of sink particles."
 Both B (Balsara 1989) aud the TDV approach (Morris aud \Louaghan 1997) to controlling AV in circular shear flow fail for the sale reasons SPII estimates of Vv appear to detec convergence in steady shear flow., Both $B$ (Balsara 1989) and the TDV approach (Morris and Monaghan 1997) to controlling AV in circular shear flow fail for the same reasons: SPH estimates of $\nabla \cdot {\bf v}$ appear to detect convergence in steady shear flow.
 AV controlled by these methods is therefore applied iu all areas of a simulated I&epleriau accretion disk. with results which will be uuprepreseutative of real viscous evolution. and compromise the results (Clarke. 2009).," AV controlled by these methods is therefore applied in all areas of a simulated Keplerian accretion disk, with results which will be unprepresentative of real viscous evolution, and compromise the results (Clarke, 2009)."
 AV results iu the selective formation of prograde alieuments of SPI particles. which can be measured using a correlation coefficient method.," AV results in the selective formation of prograde alignments of SPH particles, which can be measured using a correlation coefficient method."
 We find that a disk populated with randomly placed SPI particles initially has 185€ particles in proerade aliguimeuts aud IS retroerade., We find that a disk populated with randomly placed SPH particles initially has $\%$ particles in prograde alignments and $\%$ retrograde.
 After being allowed to evolve usine xessure forces and AV. these figures change to 28% aud 12%," After being allowed to evolve using pressure forces and AV, these figures change to $\%$ and $\%$."
 Clearly the use of AV changes the arrauecient of particles iu the experiment. even when very low levels of AV are applied.," Clearly the use of AV changes the arrangement of particles in the experiment, even when very low levels of AV are applied."
" An alternative method of controlling ΑΝ. based ou ideutifvinge Keplerian velocity characteristics. shows promise, being more effective both at switching on iu reeions of convergence and off in pure I&epleriau flow. aud siguificautly reducing the accumulation of excess prograce aligninents of particles."," An alternative method of controlling AV, based on identifying Keplerian velocity characteristics, shows promise, being more effective both at switching on in regions of convergence and off in pure Keplerian flow, and significantly reducing the accumulation of excess prograde alignments of particles."
 AAnnabel Cartwright is the holder of a Roval Society Dorothy Todekin Fellowship., Annabel Cartwright is the holder of a Royal Society Dorothy Hodgkin Fellowship.
 Some of this work was uudertaken as part of a PPARC funded studentship., Some of this work was undertaken as part of a PPARC funded studentship.
 D.Stamatcllos acknowledges post-doctoral support from STFC under the auspices of the Cardiff Astronomy Rolling Grant., D.Stamatellos acknowledges post-doctoral support from STFC under the auspices of the Cardiff Astronomy Rolling Grant.
 Grateful thanks to Authouy Whitworth for his euidaunce and insight., Grateful thanks to Anthony Whitworth for his guidance and insight.
"In addition to the stars from the and surveys, we selected 18 stars from the literature using the following criteria: They have announced radial-velocity companions with minimum mass of 13-80 M;, were published after 2002, and are not flagged as binary stars in the new Hipparcos reduction.","In addition to the stars from the and surveys, we selected 18 stars from the literature using the following criteria: They have announced radial-velocity companions with minimum mass of 13-80 $M_J$, were published after 2002, and are not flagged as binary stars in the new Hipparcos reduction."
 The selected stars are listed in Table 5 and details are given below., The selected stars are listed in Table \ref{tab:littargets} and details are given below.
"Although we have so far only considered the redshift range 0.1 to 1.6, SuMIRe PFS as modelled here is capable of observing redshifts up to about 4.9.","Although we have so far only considered the redshift range 0.1 to 1.6, SuMIRe PFS as modelled here is capable of observing redshifts up to about 4.9."
" At redshifts greater than 2 the spectrograph can measure the Lyman-alpha spectral features, however for 1.6«z<2 there exists an effective blind spot in which no spectral features are observable by this survey."," At redshifts greater than 2 the spectrograph can measure the Lyman-alpha spectral features, however for $1.6 < z < 2$ there exists an effective blind spot in which no spectral features are observable by this survey."
 This optimization therefore considers two independent redshift mega-bins; the low-redshift mega-bin covering 0.1«z(low)1.6 and the high-redshift mega-bin covering 2«z(high)4.9., This optimization therefore considers two independent redshift mega-bins; the low-redshift mega-bin covering $0.1 < z(\rm low) < 1.6$ and the high-redshift mega-bin covering $2 < z(\rm high) < 4.9$.
 Figure 9 depicts the mega-bin modelling and Table 5 details the survey variables., Figure \ref{fig:bins2} depicts the mega-bin modelling and Table \ref{tbl:vary2} details the survey variables.
 The modelling details of the high redshift mega-bin can be found in P10., The modelling details of the high redshift mega-bin can be found in P10.
 Two types of optimization are performed., Two types of optimization are performed.
" Firstly a full MCMC optimization using the Savage-Dickey InB(—1,0) FoM is executed using 3 different optimization settings: From optimization (i) we establish the optimum time split between the low and high mega-bins; the optimum redshift and exposure times for each are then found from (ii) and (iii)."," Firstly a full MCMC optimization using the Savage–Dickey $\ln B(-1,0)$ FoM is executed using 3 different optimization settings: From optimization (i) we establish the optimum time split between the low and high mega-bins; the optimum redshift and exposure times for each are then found from (ii) and (iii)."
 Discrete optimizations can then be performed with the Savage-Dickey area’ FoM. To do this we fix the redshift limits and exposure time according to the InB optimization results., Discrete optimizations can then be performed with the Savage–Dickey $^{-1}$ FoM. To do this we fix the redshift limits and exposure time according to the $\ln B$ optimization results.
 By manually varying the time (and therefore area) we can examine this FoM's performance with total time allocation and its split across the redshift mega-bins., By manually varying the time (and therefore area) we can examine this FoM's performance with total time allocation and its split across the redshift mega-bins.
" The SuMIRe project has moved on a great deal from the version modelled here, for example the current design has no redshift blind spot and can in principle observe as deep as z— if targets are available (Murayama2011)."," The SuMIRe project has moved on a great deal from the version modelled here, for example the current design has no redshift blind spot and can in principle observe as deep as $z=10$ if targets are available \cite{mur2011}."
 As such direct comparison cannot be drawn; we instead use the observational technique outlined in the 2010 SuMIRe PFS white paper (Takada&Silverman2010) as our reference., As such direct comparison cannot be drawn; we instead use the observational technique outlined in the 2010 SuMIRe PFS white paper \cite{tak2010a} as our reference.
" In this set-up the survey only covers an area of 2000 sq.deg.,"," In this set-up the survey only covers an area of 2000 sq.deg.,"
" limited by the projected area that the Hyper Suprime-Cam (HSC), needed for pre-selecting PFS's target galaxies, will cover during its operational lifetime."," limited by the projected area that the Hyper Suprime-Cam (HSC), needed for pre-selecting PFS's target galaxies, will cover during its operational lifetime."
" Its redshift range is 0.6 to 1.6, exposure time is 15 minutes, and therefore total time is roughly 500 hours."," Its redshift range is 0.6 to 1.6, exposure time is 15 minutes, and therefore total time is roughly 500 hours."
" The full MCMC optimization using InB(—1,0) found that an improvement in our confidence in ACDM (if it is the underlying model) will be gained for a wide range of time allocation to the low redshift mega-bin, but it is clearly preferable to focus all time in this mega-bin."," The full MCMC optimization using $\ln B(-1,0)$ found that an improvement in our confidence in $\Lambda\rm CDM$ (if it is the underlying model) will be gained for a wide range of time allocation to the low redshift mega-bin, but it is clearly preferable to focus all time in this mega-bin."
" This is summarised in Figure 10,, plotting Tow against InB(—1,0)."," This is summarised in Figure \ref{fig:timelow}, plotting $\tau_{\rm low}$ against $\ln B(-1,0)$."
 Figure 11 shows how the value of zmin(low) affects the survey's ability to prefer ACDM., Figure \ref{fig:zminlow} shows how the value of $z_{\rm min}(\rm low)$ affects the survey's ability to prefer $\Lambda\rm CDM$.
" We find that it is best to make use of the full redshift range in the low mega-bin, j1.6."," We find that it is best to make use of the full redshift range in the low mega-bin, $0.1 <z <1.6$ ."
" However, aslongasthelowerlimitisnotgreaterthanzzl.Othereisnogreatl "," However, as long as the lower limit is not greater than $z=1.0$ there is no great loss of performance."
iggleZand pointsmeasuredbyWBOS Scovertheredshi ftrangebetween0., This has a great deal to do with the fact that the data-points measured by WiggleZ and BOSS cover the redshift range between 0.1 and 1.0.
land is reasonable.," It also indicates that the reference survey's choice of $z_{\rm
 min} =0.6$ is reasonable."
" As with the DETF optimization there is a preference for maximising the survey volume, and we therefore find that the optimal survey minimises the exposure time and maximises survey area."," As with the DETF optimization there is a preference for maximising the survey volume, and we therefore find that the optimal survey minimises the exposure time and maximises survey area."
that both the accretion flow and the jets themselves are radiatively inefficient (di Matteo et al.,that both the accretion flow and the jets themselves are radiatively inefficient (di Matteo et al.
 2003)., 2003).
 This indicates that the flow must be advection dominated. 1.e.. the gravitational energy released in the flow must be carried into the centre rather than radiated locally (see Narayan. Mahadevan Quataert 1998: Kato. Fukue Mineshige 2008: Narayan McClintock 2008: for reviews of advection-dominated accretion flows. or ADAFs)}.," This indicates that the flow must be advection dominated, i.e., the gravitational energy released in the flow must be carried into the centre rather than radiated locally (see Narayan, Mahadevan Quataert 1998; Kato, Fukue Mineshige 2008; Narayan McClintock 2008; for reviews of advection-dominated accretion flows, or ADAFs)."
 A fraction of the energy must then be efficiently transferred to the jets once the accreting gas reaches the centre., A fraction of the energy must then be efficiently transferred to the jets once the accreting gas reaches the centre.
 The above conclusion is supported by a study of 8 other massive nearby elliptical galaxies where the gas properties close o the Bondi radius can be observed or reasonably extrapolated (Allen et al., The above conclusion is supported by a study of 8 other massive nearby elliptical galaxies where the gas properties close to the Bondi radius can be observed or reasonably extrapolated (Allen et al.
 2006)., 2006).
 In all these cases. the Bondi mass accretion rate Af determined at the Bondi radius rp correlates well with he power of the jets Pi. where the latter is measured from the bubbles inflated by the jets in the surrounding gas.," In all these cases, the Bondi mass accretion rate $\dot{M}_{\rm B}$ determined at the Bondi radius $r_{\rm B}$ correlates well with the power of the jets $P_{\rm j}$, where the latter is measured from the bubbles inflated by the jets in the surrounding gas."
 Writing the jet power as P|—nMpc. the jet production efficiency factor rj is ‘ound to be about 2 per cent.," Writing the jet power as $P_{\rm j}=\eta_{\rm j}\dot M_{\rm B}c^2$, the jet production efficiency factor $\eta_{\rm j}$ is found to be about 2 per cent."
 This is a rather large efficiency and underscores the need for mass at approximately the Bondi accretion rate reaching the gravitational radius of the black hole re., This is a rather large efficiency and underscores the need for mass at approximately the Bondi accretion rate reaching the gravitational radius of the black hole $r_{\rm g}$.
 There is ittle room for any inefficiency in the transport of mass to the centre. e.g.. through mass loss in outflows along the way.," There is little room for any inefficiency in the transport of mass to the centre, e.g., through mass loss in outflows along the way."
 We ure concerned here whether an ADAF can be established in galactic nuclei and whether the mass accretion rate is comparable to the Bondi rate., We are concerned here whether an ADAF can be established in galactic nuclei and whether the mass accretion rate is comparable to the Bondi rate.
 The range of jet power in the systems discussed above is between 10710%ergs /.sothe Eddington ratio (power emitted to Eddington limit) is 10*7 fora black hole of mass 1094.. and ten times less for 10/247...," The range of jet power in the systems discussed above is between $10^{43} - 10^{45}~{\rm erg\,s^{-1}}$, so the Eddington ratio (power emitted to Eddington limit) is $10^{-4} - 10^{-2}$ for a black hole of mass $10^9 M_\odot$, and ten times less for $10^{10}M_\odot$."
 This is very much in the regime where an ADAF is expected (Narayan Yi 1995b: Narayan McClintock 2008)., This is very much in the regime where an ADAF is expected (Narayan Yi 1995b; Narayan McClintock 2008).
 Moreover. as noted by Narayan Yi (1994. 199543) and Fabian Rees (1995). and confirmed in more recent investigations (Narayan McClintock 2008). the large thermal pressure of an ADAF may be especially good for the production and collimation of jets.," Moreover, as noted by Narayan Yi (1994, 1995a) and Fabian Rees (1995), and confirmed in more recent investigations (Narayan McClintock 2008), the large thermal pressure of an ADAF may be especially good for the production and collimation of jets."
 Thus. it is natural to consider an accretion model for systems with powerful jets.," Thus, it is natural to consider an ADAF-like accretion model for systems with powerful jets."
 ADAFs have been well studied since the work of Narayan Yi (1994. 1995ab) and Abramowiez et al. (," ADAFs have been well studied since the work of Narayan Yi (1994, 1995ab) and Abramowicz et al. ("
1995).,1995).
 However. in much of the previous work. the outer edge of the solution was generally taken to be either of a self-similar form (e.g.. Chen. Abramowiez Lasota 1997: Popham Gammie 1998) or a geometrically thin disk that evaporates to form the ADAF (e.g. Narayan. Kato Honma 1997: Manmoto et al.," However, in much of the previous work, the outer edge of the solution was generally taken to be either of a self-similar form (e.g., Chen, Abramowicz Lasota 1997; Popham Gammie 1998) or a geometrically thin disk that evaporates to form the ADAF (e.g., Narayan, Kato Honma 1997; Manmoto et al."
 2000)., 2000).
 Neither of these boundary conditions is relevant for understanding accretion from an external medium., Neither of these boundary conditions is relevant for understanding accretion from an external medium.
 Since an ADAF is essentially space filling. we expect the accretion flow to match smoothly on to the external medium without any shocks or other kinds of discontinuities.," Since an ADAF is essentially space filling, we expect the accretion flow to match smoothly on to the external medium without any shocks or other kinds of discontinuities."
 We investigate in this paper exactly how this matching occurs when we have a slowly rotating external medium (eg., We investigate in this paper exactly how this matching occurs when we have a slowly rotating external medium (eq.
 23. gives a quantitative measure of what we mean by slow rotation)., \ref{rotpar} gives a quantitative measure of what we mean by slow rotation).
 Previous. studies of Bondi-like accretion with angular momentum have generally considered inviscid flows., Previous studies of Bondi-like accretion with angular momentum have generally considered inviscid flows.
 Proga Begelman (2003) carried out two-dimensional axisymmetric simulations and showed that an equatorial torus forms because of the angular momentum barrier and that this torus constrains the amount of polar accretion., Proga Begelman (2003) carried out two-dimensional axisymmetric simulations and showed that an equatorial torus forms because of the angular momentum barrier and that this torus constrains the amount of polar accretion.
 Krumholz. McKee Klein (2005) extended their work and developed approximate formulae for the mass accretion rate as a function of the vorticity of the external gas. and Cuadra et al. (," Krumholz, McKee Klein (2005) extended their work and developed approximate formulae for the mass accretion rate as a function of the vorticity of the external gas, and Cuadra et al. ("
2006) carried out detailed simulations of inviscid accretion on to Sagittarius at the Galactic Centre.,2006) carried out detailed simulations of inviscid accretion on to Sagittarius $^*$ at the Galactic Centre.
 Recently. Inogamov Sunyaev (2010) proposed an accretion model for M87.," Recently, Inogamov Sunyaev (2010) proposed an accretion model for M87."
 As in the other studies cited here. the centrifugal barrier causes the inviscid accreting gas to form a torus well inside rp.," As in the other studies cited here, the centrifugal barrier causes the inviscid accreting gas to form a torus well inside $r_{\rm B}$."
 Inogamov Sunyaev assume that viscosity then turns on at smaller radii and suggest that the torus will thus feed a standard. thin accretion dise on the inside. which might evaporate into an ADAF at vet smaller radii.," Inogamov Sunyaev assume that viscosity then turns on at smaller radii and suggest that the torus will thus feed a standard thin accretion disc on the inside, which might evaporate into an ADAF at yet smaller radii."
 The presence of the thin disk segment in their model causes the total inflow time of the gas from the Bondi radius ry to the black hole gravitational radius ry to be far longer than or a Bondi flow or (as we shall see) an ADAF., The presence of the thin disk segment in their model causes the total inflow time of the gas from the Bondi radius $r_{\rm B}$ to the black hole gravitational radius $r_g$ to be far longer than for a Bondi flow or (as we shall see) an ADAF.
 Self-adjustment of he feedback. in which the jet power responds to conditions (e.g. cooling time) beyond rpg. then becomes very difficult. with large iysteresis expected.," Self-adjustment of the feedback, in which the jet power responds to conditions (e.g. cooling time) beyond $r_{\rm B}$, then becomes very difficult, with large hysteresis expected."
 In contrast to the above studies. we are interested in viscous accretion.," In contrast to the above studies, we are interested in viscous accretion."
 The closest paper to our work is Park (2009)., The closest paper to our work is Park (2009).
 For echnical reasons. that work focused on extremely hot external media CF710? KK) for which the Bondi radius is much closer o the black hole than in real systems.," For technical reasons, that work focused on extremely hot external media $T_{\rm ext} > 10^{9}$ K) for which the Bondi radius is much closer to the black hole than in real systems."
 We consider more realistic external conditions (Z4-109.7 KK»., We consider more realistic external conditions $T_{\rm ext} \sim 10^{6-7}$ K).
 We also study in more detail he transition from a Bondi flow to an ADAF as the external rotation is varied., We also study in more detail the transition from a Bondi flow to an ADAF as the external rotation is varied.
 As in Park (2009). we require the flow to be continuous out through rp and beyond.," As in Park (2009), we require the flow to be continuous out through $r_{\rm B}$ and beyond."
 Such a model ensures that the üccretion. power is as well coupled with the conditions in the outer gas as possible. thereby allowing for the most efficient feedback.," Such a model ensures that the accretion power is as well coupled with the conditions in the outer gas as possible, thereby allowing for the most efficient feedback."
 We moreover require that outflows. and significant radial exchanges of energy within the ADAF. are suppressed.," We moreover require that outflows, and significant radial exchanges of energy within the ADAF, are suppressed."
 We postulate that relativistic jets are created and mechanically powered very efficiently (but very radiatively inefficiently) by the accreting gas close to the black hole. but how this occurs is beyond the scope of the present work.," We postulate that relativistic jets are created and mechanically powered very efficiently (but very radiatively inefficiently) by the accreting gas close to the black hole, but how this occurs is beyond the scope of the present work."
 We limit ourselves to a more basic question: Can an idealised ADAF transfer a high enough mass accretion rate from beyond rg down to ry?, We limit ourselves to a more basic question: Can an idealised ADAF transfer a high enough mass accretion rate from beyond $r_{\rm B}$ down to $r_{\rm g}$?
 Since we are primarily interested in slowly-rotating. steady. viscous accretion flows. we assume that the density and pressure of the gas are distributed spherically at each radius.," Since we are primarily interested in slowly-rotating, steady, viscous accretion flows, we assume that the density and pressure of the gas are distributed spherically at each radius."
 We also assume that all quantities are independent of time (steady state assumption)., We also assume that all quantities are independent of time (steady state assumption).
 We thus focus only on radial variations., We thus focus only on radial variations.
 Under these assumptions. the mass accretion rate AY at radius r is given by where p(r) is the density and v(r) is the radial velocity: the latter is taken to be negative when gas flows inwards.," Under these assumptions, the mass accretion rate $\dot{M}$ at radius $r$ is given by where $\rho(r)$ is the density and $v(r)$ is the radial velocity; the latter is taken to be negative when gas flows inwards."
 When considering accretion flows in which rotational support is important. e.g.. geometrically thin disks. or ADAFs with more rotation than we consider here. the factor «πι in the above relation is replaced by (2z7z7)(2H). where H(r) is the “vertical” scale height of the gus at radius r.," When considering accretion flows in which rotational support is important, e.g., geometrically thin disks, or ADAFs with more rotation than we consider here, the factor $4\pi r^2$ in the above relation is replaced by $(2\pi r)(2H)$, where $H(r)$ is the “vertical” scale height of the gas at radius $r$."
 In the simpler approximation considered here. we effectively set H=r. which could be interpreted as a geometrically very thick disk.," In the simpler approximation considered here, we effectively set $H=r$, which could be interpreted as a geometrically very thick disk."
 Except for this difference. the equations we consider are identical to those described in Narayan et al. (," Except for this difference, the equations we consider are identical to those described in Narayan et al. ("
1997).,1997).
 To mock up relativistic gravity in our Newtonian model. we assume a gravitational potential (Paczynsski Wiita 1980)," To mock up relativistic gravity in our Newtonian model, we assume a gravitational potential (Paczyńsski Wiita 1980)"
The xeciselv measured frequencies of solar oscillations provide us with a powerful tool to probe the solar interior with sutiicicnt accuracy.,The precisely measured frequencies of solar oscillations provide us with a powerful tool to probe the solar interior with sufficient accuracy.
 These frequencies are primarily deteruined by the iiecliiuical quantities like sound speed. densi voor the adiahatic index of the solar material.," These frequencies are primarily determined by the mechanical quantities like sound speed, density or the adiabatic index of the solar material."
 The primary inversions of the observed frequencies vield oulv the sound speed and deusitv profiles inside the Sun., The primary inversions of the observed frequencies yield only the sound speed and density profiles inside the Sun.
 On t1ο other hand. in order to infer the temperature aud cremical composition profiles. additional assunnptious regarding the input phivsies such as opacities. equation of state and unclear energy generation rates are required.," On the other hand, in order to infer the temperature and chemical composition profiles, additional assumptions regarding the input physics such as opacities, equation of state and nuclear energy generation rates are required."
 Cougτὰς Isosovichey (1988)) aud I&osovichev (19963) lave eniploved t1e equations of thermal equilibrium to express the changes in primary variables (p. T4) in terms ofthose in secondary variables QZ) aud thus obtained equations connecting the frequency differences to variations in abundance profiles.," Gough Kosovichev \cite{dog88}) ) and Kosovichev \cite{kos96}) ) have employed the equations of thermal equilibrium to express the changes in primary variables $\rho,\Gamma_1$ ) in terms of those in secondary variables $Y,Z$ ) and thus obtained equations connecting the frequency differences to variations in abundance profiles."
 Shibahashi Takata (1996)). Takata Shibahlhashi (1998)) and Shibahashi. Wircmath Taxata (1998)) adopt the equatious of thermal equilibriuu. standard opacitics and nuclear reaction rates to deduce the temperature aud bydrogen abundance profiles wili the use of oilv. the inverted sound speed profile.," Shibahashi Takata \cite{st96}) ), Takata Shibahashi \cite{tak98}) ) and Shibahashi, Hiremath Takata \cite{shi98}) ) adopt the equations of thermal equilibrium, standard opacities and nuclear reaction rates to deduce the temperature and hydrogen abundance profiles with the use of only the inverted sound speed profile."
 Autia Chitre (1995... 1998)) followed a simular approach. but they used t1ο inverted density profile. ln acleliticn to the sound speed profile. or calculating the temperature and hydrogen abundance profiles. for a prescribe heavy clement abundance CZ) profile.," Antia Chitre \cite{ac95}, \cite{ac98}) ) followed a similar approach, but they used the inverted density profile, in addition to the sound speed profile, for calculating the temperature and hydrogen abundance profiles, for a prescribed heavy element abundance $Z$ ) profile."
" Iu general. the compted luminosity du a seiswically coniputed soar model is not expected to be in agreement with the observed solu ποπ]τν,"," In general, the computed luminosity in a seismically computed solar model is not expected to be in agreement with the observed solar luminosity."
 Dv appli18o the observed luminosity consradut it is possible to estimate he cross-section of protoproton (pp) nuclear reaction., By applying the observed luminosity constraint it is possible to estimate the cross-section of proton-proton (pp) nuclear reaction.
 Antia Chitre (1998)) esinated this cross-section to be 5j=(L15d:0.25)«10?5 ΛΙΟΝ barns., Antia Chitre \cite{ac98}) ) estimated this cross-section to be $S_{11}= (4.15\pm0.25)\times10^{-25}$ MeV barns.
 Similar values ie been obtained bv cOnpannues the computed solar nodels with bhelioseisiuic ¢ata (DeelTunoceuti. Fioreutini Ricci 1998: Schlattl. Beauno Paterno 1998)).," Similar values have been obtained by comparing the computed solar models with helioseismic data (Degl'Innocenti, Fiorentini Ricci \cite{inn98}; Schlattl, Bonanno Paterno \cite{bon99}) )."
 The hain source of error in these estinates is the nnucertaiuties in Z profiles., The main source of error in these estimates is the uncertainties in $Z$ profiles.
 In this work we try to find the reedon iu he Z $44 plane that is consistent wih the constraints inposed by the helioseisimic data., In this work we try to find the region in the $Z$ $S_{11}$ plane that is consistent with the constraints imposed by the helioseismic data.
 It iiw even be argued that one c‘an decrude the oressure. in addition to the sound s)ved ancl deusity. roni priniarv iuversions using the equaion of livdrostatic equilibrimu.," It may even be argued that one can determine the pressure, in addition to the sound speed and density, from primary inversions using the equation of hydrostatic equilibrium."
 This profile can then be ux (Las an adiional constraint for deteriuuing the heavy οClueit albnudaice xofile., This profile can then be used as an additional constraint for determining the heavy element abundance profile.
 Iu this work we explore the possibiliv of determiuiug the Z profile in acdilon to he .X profile usine his additional iuput., In this work we explore the possibility of determining the $Z$ profile in addition to the $X$ profile using this additional input.
 AlternateVe We ¢“ALL CLOcrue he Z profile (or opacities) instead of the A profie (TriratLiv 1998)).," Alternately, we can determine the $Z$ profile (or opacities) instead of the $X$ profile (Tripathy \cite{tri98}) )."
 Rox»reh (1996) las also examined No profiles which are obaimed by suitadv scaling the hydrogeu abundance profiles from a staidarel solar model in order to eeucrate the observed Ipnuiuosity., Roxburgh \cite{rox96}) ) has also examined $X$ profiles which are obtained by suitably scaling the hydrogen abundance profiles from a standard solar model in order to generate the observed luminosity.
 The motivation of this study was to expore the possibiitv of reducing the neutrino fiuxes viclded bw the seinuic models by. allowing for variations iu both the composition profiles as well as selected nuclear reaction rates., The motivation of this study was to explore the possibility of reducing the neutrino fluxes yielded by the seismic models by allowing for variations in both the composition profiles as well as selected nuclear reaction rates.
 The sound speed aud density profiles inside the Sun aro inferred frou the observed frequencies usiug a Reeularized Least Squares technique (Αα 1996))., The sound speed and density profiles inside the Sun are inferred from the observed frequencies using a Regularized Least Squares technique (Antia \cite{a96}) ).
 The, The
systematic. errors that push them well upward in the CMD.,systematic errors that push them well upward in the CMD.
 However. the region {«28 preserves much useful structure representing the metallicity range. losing only the reddest part of the range.," However, the region $I < 28$ preserves much useful structure representing the metallicity range, losing only the reddest part of the range."
 The tests of the photometry described above lead us to restrict the MDF analysis to the brightest part of the RGB at 28.0> 27.0., The tests of the photometry described above lead us to restrict the MDF analysis to the brightest part of the RGB at $28.0 > I > 27.0$ .
 In this range. even the random z0.40—mag scatter in (V—7) i5 much smaller than the 1.5-mag intrinsic spread in the actual CMD. enough to give us workable leverage on the underlying metallicity distribution.," In this range, even the random $\pm 0.40-$ mag scatter in $(V-I)$ is much smaller than the 1.5-mag intrinsic spread in the actual CMD, enough to give us workable leverage on the underlying metallicity distribution."
 In addition. the = 0.2—mag systematic error in the colors is far smaller than the intrinsic color range of the red-giant tracks as they fan out toward the top of the RGB (Figure 4)). creating a shift at worst of 0.1-0.2 dex in mean metallicity.," In addition, the $\simeq 0.2-$ mag systematic error in the colors is far smaller than the intrinsic color range of the red-giant tracks as they fan out toward the top of the RGB (Figure \ref{cmd_fiducials}) ), creating a shift at worst of 0.1-0.2 dex in mean metallicity."
 To estimate the combined effects of the quality of photometry on the deduced MDF. we used the NGC 5128 data (Figure 10)) again as a testbed.," To estimate the combined effects of the quality of photometry on the deduced MDF, we used the NGC 5128 data (Figure \ref{cmd5128}) ) again as a testbed."
 With the procedures described in detail in previous papers (Harris et al. 2002.. 2007b.," With the procedures described in detail in previous papers (Harris et al. \cite{har02}, \cite{har07},"
. Rejkuba et al. 2005)).," Rejkuba et al. \cite{rej05}) ),"
 interpolation within the grid of RGB tracks yields the MDF in the form of number of stars per unit interval in [m/H] = log(Z/Z..)., interpolation within the grid of RGB tracks yields the MDF in the form of number of stars per unit interval in [m/H] = $(Z/Z_{\odot})$.
 The inner halo stars within 8 kpe in NGC 5128 are now well established to have a broad. predominantly metal-rich MDF peaking near [m/H] =—0.3. and other giant ellipticals are expected to be similar (see the extensive discussions in Harris et al. 2002..," The inner halo stars within 8 kpc in NGC 5128 are now well established to have a broad, predominantly metal-rich MDF peaking near [m/H] $\simeq -0.3$, and other giant ellipticals are expected to be similar (see the extensive discussions in Harris et al. \cite{har02},"
 Rejkuba et al., Rejkuba et al.
 2005 among others)., \cite{rej05} among others).
 The RGB-grid interpolation procedure applied to the NGC 5128 stars is shown in three different versions in Figure ΤΙ., The RGB-grid interpolation procedure applied to the NGC 5128 stars is shown in three different versions in Figure \ref{simhisto}.
 In the top panel. we take the 507 stars in Figure 10((a) that lie in our adopted range 27«J28 and derive the MDF in histogram form as shown there.," In the top panel, we take the 507 stars in Figure \ref{cmd5128}( (a) that lie in our adopted range $27 < I
< 28$ and derive the MDF in histogram form as shown there."
 This can be thought of as the intrinsic. input distribution free of photometric selection or scatter.," This can be thought of as the intrinsic, input distribution free of photometric selection or scatter."
 In the second panel. the MDF derived through. the interpolation grid is shown for only the 167 NGC 5128 stars that were in the simulation tests (that Is. those in Figure 10((b)).," In the second panel, the MDF derived through the interpolation grid is shown for only the 167 NGC 5128 stars that were in the simulation tests (that is, those in Figure \ref{cmd5128}( (b))."
 The (1.V—7) values entered into the RGB grid are their true (input) values.," The $(I, V-I)$ values entered into the RGB grid are their true (input) values."
 The visible change in the MDF compared with the intrinsic. version is the direct effect of the V-band cutoff. which ts responsible for the loss of the highest-metallicity stars and thus also decreases the metallicity of the histogram peak.," The visible change in the MDF compared with the intrinsic version is the direct effect of the $V-$ band cutoff, which is responsible for the loss of the highest-metallicity stars and thus also decreases the metallicity of the histogram peak."
" In the third panel. the MDF ts now shown for the recovered stars at δω=115""—155” or 9.3—12.5 kpe between ]=27—28 derived from their actually measured photometry (Figure 10((0))."," In the third panel, the MDF is now shown for the recovered stars at $R_{gc} =
115''-155''$ or $9.3 - 12.5$ kpc between $I=27 - 28$ derived from their actually measured photometry (Figure \ref{cmd5128}( (c))."
 The overall histogram range and peak remain similar to the second panel. but relatively more stars appear at low metallicity because of the systematic bias of the colors toward the blue as mentioned above.," The overall histogram range and peak remain similar to the second panel, but relatively more stars appear at low metallicity because of the systematic bias of the colors toward the blue as mentioned above."
 Lastly. the bottom panel of Figure 11. shows the ratio of the first and third histograms (input versus measured). normalized to an overall average of 1.00.," Lastly, the bottom panel of Figure \ref{simhisto} shows the ratio of the first and third histograms (input versus measured), normalized to an overall average of 1.00."
 We use this ratio to make a reasonable correction of the MDF for both the incompleteness (which affects mainly the high-metallicity end) and the color bias (which affects mainly the low-metallicity end)., We use this ratio to make a reasonable correction of the MDF for both the incompleteness (which affects mainly the high-metallicity end) and the color bias (which affects mainly the low-metallicity end).
 The final results for our estimated MDF are shown in Figure 12.., The final results for our estimated MDF are shown in Figure \ref{fehhisto}.
 The 13972 M87 stars from Figure 4 1n the range J= as described above are put through the interpolation erid of RGB tracks to construct the raw MDF shown in the bottom panel (shaded histogram)., The 13972 M87 stars from Figure \ref{cmd_fiducials} in the range $I=27.0-28.0$ as described above are put through the interpolation grid of RGB tracks to construct the raw MDF shown in the bottom panel (shaded histogram).
 This distribution is then multiplied by the correction factor shown in the bottom panel of Figure 11. to produce the shown as the heavy unshaded line., This distribution is then multiplied by the correction factor shown in the bottom panel of Figure \ref{simhisto} to produce the shown as the heavy unshaded line.
 This corrected distribution represents, This corrected distribution represents
Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs}) are isolated neutron stars with periods of several seconds (Po~ 2-12 sj rapid spin down Qon lotess ty bright C107 -107* ty) and highly variable X-ray. AXPs and SGRs are commonly interpreted in terms of the model.,"Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are isolated neutron stars with periods of several seconds $P\sim2$ –12 s), rapid spin down $\dot{P}\sim10^{-11}$ s $^{-1}$ ), bright $\sim$$10^{34}$ $10^{35}$ ) and highly variable X-ray AXPs and SGRs are commonly interpreted in terms of the model."
 Magnetars are. ultra-magnetized neutron stars with magnetic fields largely in excess of the quantum critical field Bory=TLdt1Y G (2222).," Magnetars are ultra-magnetized neutron stars with magnetic fields largely in excess of the quantum critical field $B_{\rm{QED}}=\frac{m_e^2c^3}{\hbar e}\simeq 4.4$$\times$$10^{13}$ G \citep{paczynski92,duncan92,thompson95,thompson96}."
 Contrary to what happens in ordinary radio pulsars. the X-ray luminosity is larger than their rotational energy loss.," Contrary to what happens in ordinary radio pulsars, the X-ray luminosity is larger than their rotational energy loss."
 Since no stellar companions have been detected thus far. aso accretion is unlikely to be responsible for the emission of AXPs and SGRs.," Since no stellar companions have been detected thus far, also accretion is unlikely to be responsible for the emission of AXPs and SGRs."
 Their persistent X-ray luminosity. as well as tje. bursts and flares typical of these sources. are instead believed to be powered by the decay of their ultra-strong magnetic field (see ? and ? for recent AXPs were first recognisec as a class of persistent X-ray pulsars. with the peculiarity that the X-ray luminosity exceeds," Their persistent X-ray luminosity, as well as the bursts and flares typical of these sources, are instead believed to be powered by the decay of their ultra-strong magnetic field (see \citealt{woods06} and \citealt{mereghetti08} for recent AXPs were first recognised as a class of persistent X-ray pulsars, with the peculiarity that the X-ray luminosity exceeds"
"can then calculate the ratio of the gyrofrequency to the momentum exchange rate, for both the ions(fsi) and the electrons (f.).","can then calculate the ratio of the gyrofrequency to the momentum exchange rate, for both the ions$\beta_{Si}$ ) and the electrons $\beta_e$ )."
" We find that 6g°*ez7.9x10', while Bm2.7x 107."," We find that $\beta_\mathrm{Si}^\mathrm{max} \approx 7.9
 \times 10^{-7}$, while $\beta_e^\mathrm{max} = 2.7 \times 10^{-4}$ ."
" In general, we find over the entire domain for r in the equatorial plane, that Bs;(r)<fe(r)«1."," In general, we find over the entire domain for $r$ in the equatorial plane, that $\beta_{Si}(r) \ll \beta_e(r) \ll 1$."
" As mentioned earlier, one the central assumptions in the model is that the Lorentz force vanishes in the fluid, i.e., we have force-free MHD, therefore it is to be expected that the ratios 6 of the gyrofrequencies to the momentum transfer rate, would also accordingly be small."," As mentioned earlier, one the central assumptions in the model is that the Lorentz force vanishes in the fluid, i.e., we have force-free MHD, therefore it is to be expected that the ratios $\beta$ of the gyrofrequencies to the momentum transfer rate, would also accordingly be small."
" This then implies, that for the radial range considered in this study, ie. 0€r80Ro, the Hall and ambipolar diffusion terms arising from electron-ion drift and ion-neutral drift respectively, are negligible."," This then implies, that for the radial range considered in this study, i.e. $0 \leq r \leq 80R_0$, the Hall and ambipolar diffusion terms arising from electron-ion drift and ion-neutral drift respectively, are negligible."
" Additionally, the conductivity of the plasma can be estimated using the usual relation, c=10'T?K-9?Q-'com-! (eg.?), where Te is the electron temperature, such that T.(r)>Tzas(r)."," Additionally, the conductivity of the plasma can be estimated using the usual relation, $\sigma = 10^7 T_e^{3/2} \mathrm{K}^{-3/2} \Omega^{-1} \mathrm{cm}^{-1}$ \citep[e.g.][]{Vishniac2003}, where $T_e$ is the electron temperature, such that $T_e(r) > T_\mathrm{gas}(r)$."
" Figure 1 shows that the minimum temperature for the gas far away from the photosphere is about Tz""~640 K. We can therefore find an estimate for the lower limit of the conductivity using this temperature and estimate a value of Omin71.6x10%Q7!cm*!, yielding a magnetic Reynolds number (e.g.7) of Rem=ULpoomin~1073, so Ohmic diffusion is unimportant (U~10 km/s and L£zRo are typical velocity and length scales for Betelgeuse and Lo is the permeability of the vacuum, 4vx 107’Qs/m)."," Figure \ref{fig:figure1} shows that the minimum temperature for the gas far away from the photosphere is about $T_\mathrm{gas}^\mathrm{min} \approx 640~$ K. We can therefore find an estimate for the lower limit of the conductivity using this temperature and estimate a value of $\sigma_\mathrm{min} \approx 1.6 \times 10^{11} ~ \Omega^{-1}~\textrm{cm}^{-1}$, yielding a magnetic Reynolds number \citep[e.g.][]{Davidson} of $\mathrm{Re}_m = U L \mu_0 \sigma_{min} \sim 10^{23}$, so Ohmic diffusion is unimportant $U \sim 10~$ km/s and $L \approx R_0$ are typical velocity and length scales for Betelgeuse and $\mu_0$ is the permeability of the vacuum, $4\pi \times 10^{-7} \Omega \mathrm{s}/\mathrm{m}$ )."
" The assumption of effectively infinite plasma conductance, is then a reasonable assumption."," The assumption of effectively infinite plasma conductance, is then a reasonable assumption."
" Thus, we see that the central assumption of ideal-MHD in the equatorial plane of Betelgeuse is a reasonable first approximation to make, for the purpose of conveying a simple picture."," Thus, we see that the central assumption of ideal-MHD in the equatorial plane of Betelgeuse is a reasonable first approximation to make, for the purpose of conveying a simple picture."
" Finally, the Mach numbers in the inner region of the circumstellar envelope are small and actually, do not exceed unity before the sonic point, located at around 5.27Ro."," Finally, the Mach numbers in the inner region of the circumstellar envelope are small and actually, do not exceed unity before the sonic point, located at around $5.27R_0$."
 Therefore our model does not suffer from the same drawbacks as many Alfvénn wave models that have the artifact of having large Mach numbers close to the surface of the star., Therefore our model does not suffer from the same drawbacks as many Alfvénn wave models that have the artifact of having large Mach numbers close to the surface of the star.
" Indeed, in our model the wind only becomes super-Alfvénnic beyond the Alfvénn point at around 25Ro and even then, only mildly so."," Indeed, in our model the wind only becomes super-Alfvénnic beyond the Alfvénn point at around $25R_0$ and even then, only mildly so."
" In Figure 2 we have shown the azimuthal velocity, ug(r) of the gas (red solid line)."," In Figure \ref{fig:figure2} we have shown the azimuthal velocity, $u_{\phi}(r)$ of the gas (red solid line)."
" The dust is assumed to co-rotate with the gas, thus the azimuthal velocity profile shown in the figure is also true for the dust (for r> ra), since there is no drag between the dust and the gas in the azimuthal direction."," The dust is assumed to co-rotate with the gas, thus the azimuthal velocity profile shown in the figure is also true for the dust (for $r \geq r_d$ ), since there is no drag between the dust and the gas in the azimuthal direction."
 The azimuthal velocity profile is obtained once the radial velocity profile is known (see?).., The azimuthal velocity profile is obtained once the radial velocity profile is known \citep[see][]{Thirumalai2010}.
 The velocity profile obtained is typical for a magneto-rotational wind., The velocity profile obtained is typical for a magneto-rotational wind.
 The dotted lines in the figure and in its inset represent the rotation of the star (see?).., The dotted lines in the figure and in its inset represent the rotation of the star \citep[see][]{Uitenbroek1998}.
 It is readily seen that the azimuthal wind velocity closely traces the observed rotation of the star and only begins to depart markedly after rzz4.5Ro., It is readily seen that the azimuthal wind velocity closely traces the observed rotation of the star and only begins to depart markedly after $r \approx 4.5 R_0$.
" This decoupling from stellar rotation occurs in the inner wind region of the magneto-rotational wind, at a physical distance corresponding to ~2.03x 10'*cm, for Betelgeuse."," This decoupling from stellar rotation occurs in the inner wind region of the magneto-rotational wind, at a physical distance corresponding to $ \approx 2.03 \times 10^{14}$ cm, for Betelgeuse."
 It is also interesting to note that at large distances the azimuthal velocity of the gas is strongly de-coupled from the stellar rotation rate., It is also interesting to note that at large distances the azimuthal velocity of the gas is strongly de-coupled from the stellar rotation rate.
" Therefore, the chromospheric component in the atmosphere at a few stellar radii may also appear to be de-coupled from stellar rotation as was observed by (?).."," Therefore, the chromospheric component in the atmosphere at a few stellar radii may also appear to be de-coupled from stellar rotation as was observed by \cite[][]{Harper2006}."
" However, it is to be remembered that the hybrid-MHD-dust-driven theory presented in this study, concerns itself with the equatorial plane alone, while the observations of ? are spatially resolved in the cross-dispersion direction; this may have an averaging effect along the dispersion direction."," However, it is to be remembered that the hybrid-MHD-dust-driven theory presented in this study, concerns itself with the equatorial plane alone, while the observations of \citet[][]{Harper2006} are spatially resolved in the cross-dispersion direction; this may have an averaging effect along the dispersion direction."
" Thus, the theoretically obtained rotational decoupling that occurs in the equatorial plane, in the current study, may well be quite a bit different from the observations of ?.."," Thus, the theoretically obtained rotational decoupling that occurs in the equatorial plane, in the current study, may well be quite a bit different from the observations of \citet[][]{Harper2006}."
" Therefore any inferences drawn in this regard must be tempered by the observation that the current model is a rather simple picture, with concomitant limitations."," Therefore any inferences drawn in this regard must be tempered by the observation that the current model is a rather simple picture, with concomitant limitations."
 We now turn our attention to the question of spots on the surface of Betelgeuse and the related question of temperature inhomogeneities., We now turn our attention to the question of spots on the surface of Betelgeuse and the related question of temperature inhomogeneities.
" Recent observations (see7???) suggest that there may be a few (on the order of 2 or 3) large spots on the surface of Betelgeuse, indicating that there are temperature inhomogeneities on the surface of Betelgeuse."," Recent observations \citep[see][]{Wilson1992,Wilson1997,Buscher1990,Haubois2009} suggest that there may be a few (on the order of 2 or 3) large spots on the surface of Betelgeuse, indicating that there are temperature inhomogeneities on the surface of Betelgeuse."
 We therefore chose to model Betelguese with two spots on its equator (see?).., We therefore chose to model Betelguese with two spots on its equator \citep[see][]{Soker1999}.
" In the current study however, these spots were taken to be colder than the effective temperature so that the temperature profile close to the photosphere 1.0<r3.5Ro may be investigated so as place constraints on the alumina formation region, if at all, close to the star."," In the current study however, these spots were taken to be colder than the effective temperature so that the temperature profile close to the photosphere $1.0 \leq r \leq 3.5 R_0$ may be investigated so as place constraints on the alumina formation region, if at all, close to the star."
" The spots were considered to have temperatures of 2600 K and 3000 K respectively, well within fluctuations of about 1000 K or so, about the effective temperature (see ?).."," The spots were considered to have temperatures of $2600$ K and $3000$ K respectively, well within fluctuations of about $1000$ K or so, about the effective temperature \citep[see][]{Schwarzschild1975}. ."
 The magnetic field was left unchanged so that the effect of changing a single parameter could be patently established., The magnetic field was left unchanged so that the effect of changing a single parameter could be patently established.
" The solitary requisite constraint that was placed however, was that regardless of the photospheric spot temperature, the temperature at a distance of about 30Ro should be"," The solitary requisite constraint that was placed however, was that regardless of the photospheric spot temperature, the temperature at a distance of about $30R_0$ should be"
Detailed measurements of the temperature distribution of the plasma of solar and stellar coronae are of critical importance to understand how they are heated to multimillion degree temperatures. so it is no surprise that the thermal structure of the solar corona has been studied ever since the publication of the seminal paper of Edlen (1942). who discovered its high temperature for the first time.,"Detailed measurements of the temperature distribution of the plasma of solar and stellar coronae are of critical importance to understand how they are heated to multimillion degree temperatures, so it is no surprise that the thermal structure of the solar corona has been studied ever since the publication of the seminal paper of Edlen (1942), who discovered its high temperature for the first time."
 These studies have known a revival d recent years. thanks to the availability of a large body of observations obtained with the new instrumentation carried on board the SOHO. TRACE. STEREO. Hinode and SDO satellites.," These studies have known a revival in recent years, thanks to the availability of a large body of observations obtained with the new instrumentation carried on board the SOHO, TRACE, STEREO, Hinode and SDO satellites."
 Plasma loops are a fundamental component of the solar corona. since they are ubiquitous both in quiet and active solar plasmas.," Plasma loops are a fundamental component of the solar corona, since they are ubiquitous both in quiet and active solar plasmas."
 Steady-state models of loops predict that the temperature of these structures varies along the loop (Rosner 1978). and that loops with different length can have very different temperatures: nanoflare-based models predict finely-structured loops filled with multithermal plasma all along their length (Patsourakos Klimehuk 2005).," Steady-state models of loops predict that the temperature of these structures varies along the loop (Rosner 1978), and that loops with different length can have very different temperatures; nanoflare-based models predict finely-structured loops filled with multithermal plasma all along their length (Patsourakos Klimchuk 2005)."
 A consequence of these models is that quiet coronal plasmas observed should be multithermal because they are the sum of a multitude of loops with different lengths and temperatures along the line of sight., A consequence of these models is that quiet coronal plasmas observed should be multithermal because they are the sum of a multitude of loops with different lengths and temperatures along the line of sight.
 The question of clearly distinguishing between multithermal and isothermal plasma along the line of sight is thus very important in the framework of coronal heating models (e.g. Klimchuk 2006. Reale 2010).," The question of clearly distinguishing between multithermal and isothermal plasma along the line of sight is thus very important in the framework of coronal heating models (e.g. Klimchuk 2006, Reale 2010)."
 However. recent measurements have provided increasing evidence of a nearly isothermal nature of both coronal hole and quiet Sun plasmas when they are observed above the solar limb. contrary to model predictions.," However, recent measurements have provided increasing evidence of a nearly isothermal nature of both coronal hole and quiet Sun plasmas when they are observed above the solar limb, contrary to model predictions."
" Several authors (e.g. Feldman 1998 and 1999, Warren 1999, Doschek 2001. Landi 2008) reported single-temperature determination which are remarkably similar even if taken with different instruments and at different times during the solar cycle."," Several authors (e.g. Feldman 1998 and 1999, Warren 1999, Doschek 2001, Landi 2008) reported single-temperature determination which are remarkably similar even if taken with different instruments and at different times during the solar cycle."
 A constant temperature was also found to characterize an entire streamer by Landi (2006)., A constant temperature was also found to characterize an entire streamer by Landi (2006).
 Even a complex active region. when observed at the solar west limb. seemed to be composed by three near isothermal plasmas (Landi Feldman 2008).," Even a complex active region, when observed at the solar west limb, seemed to be composed by three near isothermal plasmas (Landi Feldman 2008)."
 Further studies carried out. with Hinode/EIS observations. on the contrary. provided evidence that the quiescent solar corona ts not isothermal. but it is characterized by a rather narrow temperature structure with a tail that extends to higher temperatures (Warren Brooks 2009. Brooks 2009. Hahn 2011).," Further studies carried out with Hinode/EIS observations, on the contrary, provided evidence that the quiescent solar corona is not isothermal, but it is characterized by a rather narrow temperature structure with a tail that extends to higher temperatures (Warren Brooks 2009, Brooks 2009, Hahn 2011)."
 The aim of the present paper is to assess how reliable such isothermal plasma claims are., The aim of the present paper is to assess how reliable such isothermal plasma claims are.
 These measurements have been carried out either with the Emission Measure (EM) loci technique or with traditional Differential Emission Measured (DEM) techniques (see Feldman Landi 2008 and references therein)., These measurements have been carried out either with the Emission Measure (EM) loci technique or with traditional Differential Emission Measured (DEM) techniques (see Feldman Landi 2008 and references therein).
 The former allows to determine the temperature and the total emission measure (EM) of an isothermal plasma. and it has been the technique used in most of the studies that resulted in the isothermal corona results.," The former allows to determine the temperature and the total emission measure (EM) of an isothermal plasma, and it has been the technique used in most of the studies that resulted in the isothermal corona results."
 This technique. while useful in the case of strictly isothermal plasmas. fails at providing an assessment of the uncertainty of its results.," This technique, while useful in the case of strictly isothermal plasmas, fails at providing an assessment of the uncertainty of its results."
 Traditional DEM techniques rely on several different methods to reconstruct the thermal structure of multithermal plasmas determining the DEM curve., Traditional DEM techniques rely on several different methods to reconstruct the thermal structure of multithermal plasmas determining the DEM curve.
 However. most of them rely on assumptions on the smoothness of the distribution. that make it very difficult to study isothermal or near-isothermal plasmas.," However, most of them rely on assumptions on the smoothness of the distribution, that make it very difficult to study isothermal or near-isothermal plasmas."
 Landi Klimehuk (2010) carried out tests aimed at evaluating the ability of the EM loci technique to discriminate between isothermal and multithermal plasmas. when the spectral line intensities used for the reconstruction included uncertainties.," Landi Klimchuk (2010) carried out tests aimed at evaluating the ability of the EM loci technique to discriminate between isothermal and multithermal plasmas, when the spectral line intensities used for the reconstruction included uncertainties."
 They also developed a quantitative method to determine how multithermal could a plasma distribution be in order to be compatible with a set of observed line intensities that was also compatible with an isothermal condition., They also developed a quantitative method to determine how multithermal could a plasma distribution be in order to be compatible with a set of observed line intensities that was also compatible with an isothermal condition.
 In the present work we extend the Landi Klimehuk (2010) study to the application of DEM diagnostic techniques to spectral line intensities., In the present work we extend the Landi Klimchuk (2010) study to the application of DEM diagnostic techniques to spectral line intensities.
" In particular, we aim at: 1) determining the ability of such techniques to retrieve isothermal plasma from a set of spectral line intensities."," In particular, we aim at: 1) determining the ability of such techniques to retrieve isothermal plasma from a set of spectral line intensities,"
"Bardeen potential ot: here $1, represents a Gaussian random field.",Bardeen potential $\Phi$: here $\Phi_{\rm L}$ represents a Gaussian random field.
" Hereafter we will adopt the so-called large-scale structure (LSS) convention, where ® is linearly extrapolated to the presentepoch|."," Hereafter we will adopt the so-called large-scale structure (LSS) convention, where $\Phi$ is linearly extrapolated to the present."
" Moreover we will consider only the so-called local shape for non-Gaussianity, in which the bispectrum signal is larger on squeezed triangle configurations."," Moreover we will consider only the so-called local shape for non-Gaussianity, in which the bispectrum signal is larger on squeezed triangle configurations."
to For more (2004);details Verdeabout& Matarrese|other possible(2009); shapesBartolo we refer [Bartoloet and al]references therein.," For more details about other possible shapes we refer to \cite{Bartolo2004,Verde2009,Bartolo2009.1} and references therein."
 At present the most stringent constraints on fwr come from the cosmic microwave background (CMB) data., At present the most stringent constraints on $f_{\rm NL}$ come from the cosmic microwave background (CMB) data.
 Their discriminating power derives from the fact that its temperature fluctuations trace the density perturbations before the gravitational non-linearities modify their original distribution., Their discriminating power derives from the fact that its temperature fluctuations trace the density perturbations before the gravitational non-linearities modify their original distribution.
" Whatever is the specific test adopted, all CMB analysis consistently allow only very small deviations from Gaussianity: for instance, analyzing the recent WMAP data, found that fyi varies between -13 and 97, {Smithetal](2009) found fuu<104 (at 95 per cent whileconfidence level); see also ear"," Whatever is the specific test adopted, all CMB analysis consistently allow only very small deviations from Gaussianity: for instance, analyzing the recent WMAP data, \cite{Komatsu2010} found that $f_{\rm NL}$ varies between -13 and 97, while \cite{Smith2009} found $-5<f_{\rm NL}<104$ (at 95 per cent confidence level); see also \cite{Komatsu2009}."
"]AlternativeQ009).. and complementary constraints on fwr can in principle be derived analyzing the LSS (forrecentreviews,see,e.g. 2010).."," Alternative and complementary constraints on $f_{\rm NL}$ can in principle be derived analyzing the LSS \citep[for recent reviews, see, e.g.][]{Verde2010,Desjacques2010}."
" As already evident from the first generation of non-Gaussian N-body simulations [I993), the presence of a positive (negative) skewness in the PDF of the primordial density field tends to favor (disfavor) the formation of cosmic structures, inducing a different timing in the whole process of gravitational instability."," As already evident from the first generation of non-Gaussian N-body simulations \citep{Messina1990,Moscardini1991,Weinberg1992}, the presence of a positive (negative) skewness in the PDF of the primordial density field tends to favor (disfavor) the formation of cosmic structures, inducing a different timing in the whole process of gravitational instability."
" However, to be fully exploited, this approach needs to be complemented with reliable methods to disentangle from the primordial signal the non-Gaussian features introduced by the late non-linear evolution and by the possible presence of a non-linear bias factor."," However, to be fully exploited, this approach needs to be complemented with reliable methods to disentangle from the primordial signal the non-Gaussian features introduced by the late non-linear evolution and by the possible presence of a non-linear bias factor."
" While lensing statistics entirely avoid the latter problem, to attack the former one it is necessary to make use of both analytic techniques (likehigh-orderperturbation and full N-body etsimulations (see,e.g.Kangeta[]007,(Grossietal]D007,[Dalal317008;[et to properly calibrate the theoretical predictions in the non-linear aL.|[2009)regime."," While lensing statistics entirely avoid the latter problem, to attack the former one it is necessary to make use of both analytic techniques \citep[like high-order perturbation theory and the Time-Renormalization Group approach; see, e.g.,][]{Taruya2008,Bartolo2009.2} and full N-body simulations \citep[see, e.g.,][]{Kang2007,Grossi2007,Dalal2008,Viel2009,Desjacques2009,Pillepich2009,Grossi2009.1} to properly calibrate the theoretical predictions in the non-linear regime."
" The large amount of theoretical work recently done in this direction allowed to better understand what is the size of the effects on the abundance of non-linearetaL structures on the halo biasing. (Dalaletal2008;McDonald2008B005),. ,, on the galaxy bispectrum (Scfusati&Komatsul2007,Jeong on the mass density distribution Komatsi2009 the topology (Grossietal12008: on the J2008),,onintegrated Sachs-Wolfe effect 2003}(Afshordi[Hikage&Toley|2008etal]/2008),,eta on the Ly-o fux from low-density intergalactic medium Carbone]2008),et(VielaI]2009)., on the 21 cm fluctuations (Cooray|{2006[PillepichetaI]2007) and on the reionization process (Crocianietal.[2009).."," The large amount of theoretical work recently done in this direction allowed to better understand what is the size of the effects on the abundance of non-linear structures \citep{Matarrese2000,Verde2000,Mathis2004,Kang2007,Grossi2007,Grossi2009.1,Maggiore2009,Roncarelli2009}, on the halo biasing \citep{Dalal2008,McDonald2008,Fedeli2009.1,Carbone2010}, on the galaxy bispectrum \citep{Sefusatti2007.1,Jeong2009,Nishimichi2009.2}, on the mass density distribution \citep{Grossi2008,Lam2009.1,Lam2009.2}, on the topology \citep{Matsubara2003.1,Hikage2008}, on the integrated Sachs-Wolfe effect \citep{Afshordi2008,Carbone2008}, on the $\alpha$ flux from low-density intergalactic medium \citep{Viel2009}, on the 21 cm fluctuations \citep{Cooray2006,Pillepich2007} and on the reionization process \citep{Crociani2009}."
" The first attempts of an application to real observational data gave very encouraging results: [STosaretal.](2008),, combining the bias measurements for two samples of luminous red(2008). galaxies and quasars, found fxn=48153: studying the integrated Sachs- effect (SW) in the NVSS survey, derived fxr,=354+165; all error bars are at 2-0levef?]."," The first attempts of an application to real observational data gave very encouraging results: \cite{Slosar2008}, combining the bias measurements for two samples of luminous red galaxies and quasars, found $f_{\rm NL}=48^{+55}_{-74}$ ; \cite{Afshordi2008}, studying the integrated Sachs-Wolfe effect (ISW) in the NVSS survey, derived $f_{\rm NL}=354\pm 165$; all error bars are at $\sigma$."
 Very recently found 32«fwr152 from the analysis of the auto-correlation of the brightest NVSS sources on angular scales of several degrees., Very recently \cite{Xia2010} found $32 < f_{\rm NL}< 152$ from the analysis of the auto-correlation of the brightest NVSS sources on angular scales of several degrees.
 In this paper we will focus on estimating the weak lensing signals in scenarios with primordial non-Gaussianity., In this paper we will focus on estimating the weak lensing signals in scenarios with primordial non-Gaussianity.
" Being based on the measurement of the shear effect produced by the intervening large-scale structure of the universe on the images of background galaxies, gravitational lensing is a direct probe of the total matter distribution."," Being based on the measurement of the shear effect produced by the intervening large-scale structure of the universe on the images of background galaxies, gravitational lensing is a direct probe of the total matter distribution."
 For this reason it is considered one of the most powerful tools to constrain the main cosmological parameters and many dedicated projects are in progress or under study., For this reason it is considered one of the most powerful tools to constrain the main cosmological parameters and many dedicated projects are in progress or under study.
" A very exciting perspective is certainly opened by the ESA Cosmic Vision project EUCLID currently under study: the goal of its wide survey(Laureijs|2009),, is to obtain the shear measurements for about 40 galaxies per arcmin? on the entire extragalactic sky with galactic latitude b>30 (approximately 20000 deg?)."," A very exciting perspective is certainly opened by the ESA Cosmic Vision project EUCLID \citep{Laureijs2009}, currently under study: the goal of its wide survey is to obtain the shear measurements for about 40 galaxies per $^2$ on the entire extragalactic sky with galactic latitude $b>30$ (approximately 20000 $^2$ )."
 The possibility of using the weak lensing signals to constrain also the amount of primordial non-Gaussianity has been already explored by different authors., The possibility of using the weak lensing signals to constrain also the amount of primordial non-Gaussianity has been already explored by different authors.
"[Amara used a generalized halo model to study the impact&Refregier|(2004) on the estimates of the power spectrum normalization og of primordial non-Gaussianity, modeled assuming various lognormal distributions for the density field."," \cite{Amara2004} used a generalized halo model to study the impact on the estimates of the power spectrum normalization $\sigma_8$ of primordial non-Gaussianity, modeled assuming various lognormal distributions for the density field."
" More recently, computed the power spectrum of the weak cosmic shear for non-Gaussian models with different values of ᾖ[ντ,."," More recently, \cite{Fedeli2010}
 computed the power spectrum of the weak cosmic shear for non-Gaussian models with different values of $f_{\rm NL}$."
" In particular, they improved the halo model including more accurate prescriptions for its ingredients (mass functions, bias and halo profile), calibrated on the last generation of non-Gaussian N-body simulations."," In particular, they improved the halo model including more accurate prescriptions for its ingredients (mass functions, bias and halo profile), calibrated on the last generation of non-Gaussian N-body simulations."
" The application of this model to a survey having the expected characteristics of the EUCLID project showed the possibility of a significant detection of non-Gaussianity at the level of |ντι]= few tens, once the remaining parameters are held fixed."," The application of this model to a survey having the expected characteristics of the EUCLID project showed the possibility of a significant detection of non-Gaussianity at the level of $|f_{\rm NL}|\approx$ few tens, once the remaining parameters are held fixed."
 In this paper we investigate weak lensing statistics in non-Gaussian scenarios using numerical rather than analytical tools., In this paper we investigate weak lensing statistics in non-Gaussian scenarios using numerical rather than analytical tools.
" Specifically, we will create weak lensing maps performing ray-tracing simulations through very deep light-cones extracted from high-resolution N-body simulations."," Specifically, we will create weak lensing maps performing ray-tracing simulations through very deep light-cones extracted from high-resolution N-body simulations."
 The advantage of this approach is twofold., The advantage of this approach is twofold.
" First of all, N-body simulations permit to fully account for the non-linear evolution which is usually modeled less accurately by analytical means."," First of all, N-body simulations permit to fully account for the non-linear evolution which is usually modeled less accurately by analytical means."
" Second, numerical experiments allow us to extract a large set of realistic weak lensing maps that can be used for better evaluating the statistical robustness of the results."," Second, numerical experiments allow us to extract a large set of realistic weak lensing maps that can be used for better evaluating the statistical robustness of the results."
" The main goal of our numerical work is to figure out what are the observational evidences of the presence of some level of primordial non-Gaussianity, as quantified by the fur parameter."," The main goal of our numerical work is to figure out what are the observational evidences of the presence of some level of primordial non-Gaussianity, as quantified by the $f_{\rm NL}$ parameter."
" In particular we will compute a large set of weak lensing statistics in models with different ᾖ[ντ, and we will quantify the deviations from the corresponding results in the Gaussian scenario.", In particular we will compute a large set of weak lensing statistics in models with different $f_{\rm NL}$ and we will quantify the deviations from the corresponding results in the Gaussian scenario.
" This is important not only to address the possibility of a positive detection with future data, but also to establish at which level an"," This is important not only to address the possibility of a positive detection with future data, but also to establish at which level an"
the apsidal plane.,the apsidal plane.
 Similarly. the negative προΠαρ. is understood as the beat between the orbital period and the nodal precession of the disc (Patterson 1999).," Similarly, the negative superhump is understood as the beat between the orbital period and the nodal precession of the disc (Patterson 1999)."
 According to theory. superhumps are developed: when the aceretion cise extends bevond the 3:1. resonance radius ancl becomes elliptical (Whitehurst Wing 1991).," According to theory, superhumps are developed when the accretion disc extends beyond the 3:1 resonance radius and becomes elliptical (Whitehurst King 1991)."
 Osaki (1996) suggested that permanent superhump systems differ from other subclasses of non-magnetic CVs in their relatively short orbital periods and high mass-transfer rates. resulting in accretion disces that are thermally stable but tidally unstable.," Osaki (1996) suggested that permanent superhump systems differ from other subclasses of non-magnetic CVs in their relatively short orbital periods and high mass-transfer rates, resulting in accretion discs that are thermally stable but tidally unstable."
 Retter Navlor (2000) eave observational support to this idea., Retter Naylor (2000) gave observational support to this idea.
 Simulations suggest that superhumps can only occur in binary svstems with small mass ratios (q—Massesο&0.33 AVhitchurst 1988: Whitehurst Wine 1991: Murray 2000)., Simulations suggest that superhumps can only occur in binary systems with small mass ratios – $q=M_{donor}/M_{compact}\la0.33$ (Whitehurst 1988; Whitehurst King 1991; Murray 2000).
 Although this condition is easily met in short orbital period LAINBs. whose primaries are compact massive objects. superhumps have been only seen in a few LMXIUDs in outburst (e.g. ODonoghue Charles 1996).," Although this condition is easily met in short orbital period LMXRBs, whose primaries are compact massive objects, superhumps have been only seen in a few LMXRBs in outburst (e.g. O'Donoghue Charles 1996)."
 Llere we report the detection of permanent superhumps in a persistent LAINRB., Here we report the detection of permanent superhumps in a persistent LMXRB.
 V1405 Aql was first discovered as an X-ray. source showing tvpe-L bursts (Decker. et. al., V1405 Aql was first discovered as an X-ray source showing type-I bursts (Becker et al.
 1977: Doxsey ct al., 1977; Doxsey et al.
 1977) suggesting that its primary star is à neutron star., 1977) suggesting that its primary star is a neutron star.
 White Swank (1982) and Walter ct al. (, White Swank (1982) and Walter et al. (
1982). independentIv found. ~8000-s periodic dips in its X-ray. light. curve.,1982) independently found $\sim$ 3000-s periodic dips in its X-ray light curve.
 The dips are believed. to occur due to obseuration of the disc edge (probably where the accretion stream impacts the disc), The dips are believed to occur due to obscuration of the disc edge (probably where the accretion stream impacts the disc).
 Such a short orbital period. implies a low mass hvdrogen-deficient. secondary. star (Nelson. Rappaport Joss 1986).," Such a short orbital period implies a low mass hydrogen-deficient secondary star (Nelson, Rappaport Joss 1986)."
 The 21-mag optical counterpart of V1405 Ael was identified bv Sehmidtke (1988) and. Grindlay et al. (, The 21-mag optical counterpart of V1405 Aql was identified by Schmidtke (1988) and Grindlay et al. (
LOSS).,1988).
 They also reported a detection of an optical periodicity. about one percent longer than the X-ray period.," They also reported a detection of an optical periodicity, about one percent longer than the X-ray period."
 The cdillerence between the X-ray. ancl optical periods. was confirmed. by further extensive observations of V1405 Aql (Smale οἱ al., The difference between the X-ray and optical periods was confirmed by further extensive observations of V1405 Aql (Smale et al.
 1988: 1989: 1992: Callanan 1993: C'allanan. Crindlay Cool 1995: Yoshida et al.," 1988; 1989; 1992; Callanan 1993; Callanan, Grindlay Cool 1995; Yoshida et al."
 1995: Church ct al., 1995; Church et al.
 1997: 1998: Wo e al., 1997; 1998; Ko et al.
 1999: Morley. et al., 1999; Morley et al.
 1999: Chou. Grindlay Bloser 200 ποιος et al.," 1999; Chou, Grindlay Bloser 2001; Homer et al."
 2001)., 2001).
 The various models olfered so far for the periodicities ound in V1405 Aq) can be grouped into two basic models a triple svstem (Cuinedlay 1986: Caindlav et al., The various models offered so far for the periodicities found in V1405 Aql can be grouped into two basic models – a triple system (Grindlay 1986; Grindlay et al.
 1988) and superhumps (Schmictke 1988: White 1989)., 1988) and superhumps (Schmidtke 1988; White 1989).
 According o the first. model. which was motivated by the detection of a possible 199-d. periodicity in the X-ray. light curve of Vi1405 λα (Prieclhorsky Terrell 1984). the longer 3028-5 »eriod is the binary inner orbital period. while the shorter 3001-8 period is the beat between the binary. period ane rw ded orbital period of a third companion.," According to the first model, which was motivated by the detection of a possible 199-d periodicity in the X-ray light curve of V1405 Aql (Priedhorsky Terrell 1984), the longer 3028-s period is the binary inner orbital period, while the shorter 3001-s period is the beat between the binary period and the $\sim$ 4-d orbital period of a third companion."
 The. 199-4 »eriod. is explained by. the cecentricity of the inner binary orbit (Chou ct al., The 199-d period is explained by the eccentricity of the inner binary orbit (Chou et al.
 2001)., 2001).
 The second. model suggests tha 10 3001-s period is the binary period and that the 3028-5 »eriod is à positive superhump., The second model suggests that the 3001-s period is the binary period and that the 3028-s period is a positive superhump.
 We note that both mocdoels owediet the presence of a beat periodicity of 7 d. Lackece evidence for a period of ~8.9 d has been found in the X-ray light curve of V1405 Aql (Chou et al., We note that both models predict the presence of a beat periodicity of $\sim$ 4 d. Indeed evidence for a period of $\sim$ 3.9 d has been found in the X-ray light curve of V1405 Aql (Chou et al.
 2001: see also Homer et al., 2001; see also Homer et al.
 2001)., 2001).
 The debate on the nature on V1405 Aql has continued Chou et al., The debate on the nature on V1405 Aql has continued (Chou et al.
 2001: Homer et al., 2001; Homer et al.
 2001: Laswell ct al., 2001; Haswell et al.
 2001). although some preference. was given to. the. superhump scenario.," 2001), although some preference was given to the superhump scenario."
 One of the arguments against the triple system model is that the 199-d. period was not confirmed by Further observations. although a second: possible evidence for this periodicity was found bx Smale Lochner (1992).," One of the arguments against the triple system model is that the 199-d period was not confirmed by further observations, although a second possible evidence for this periodicity was found by Smale Lochner (1992)."
 On the other hand. an argument against the superhump. moclel is that the 3028-5. period. is quite stable while superhump periods are unstable (Callanan et al.," On the other hand, an argument against the superhump model is that the 3028-s period is quite stable while superhump periods are unstable (Callanan et al."
 1993: 1995: Chou ct al., 1993; 1995; Chou et al.
 WL: Homer ct al., 2001; Homer et al.
 2001)., 2001).
 There is a strong observational link between positive and negative superhumps in CVs as light curves of several systems show both (Patterson 1999: Arenas et al., There is a strong observational link between positive and negative superhumps in CVs as light curves of several systems show both (Patterson 1999; Arenas et al.
 2000: letter et al., 2000; Retter et al.
 2002b)., 2002b).
 Ln addition. Patterson (1999) found that period deficits in negative superhumps are about half period excesses in positive superhumps: ezθ0ο. where €—(Roane—DosinalfPosnas," In addition, Patterson (1999) found that period deficits in negative superhumps are about half period excesses in positive superhumps: $\epsilon_{-}\approx-0.5\epsilon_{+}$, where $\epsilon=(P_{superhump}-P_{orbital})/P_{orbital}$."
 Thus. we decided to look in available photometric data on V1405 Aq! for negative superhunmps. which would be predicted to have à period near 2086 s. Indeed. we have found a new periodicity in the X- light curve of V1405 αι.," Thus, we decided to look in available photometric data on V1405 Aql for negative superhumps, which would be predicted to have a period near 2986 s. Indeed we have found a new periodicity in the X-ray light curve of V1405 Aql."
 These findings were reported by Retter. Chou Bedding (2002a). and are. presented below.," These findings were reported by Retter, Chou Bedding (2002a), and are presented below."
 We also present a new relation for svstems that have both tvpes of superhunmps. and discuss the mechanisms for superhumps in the X-ray and optical regimes.," We also present a new relation for systems that have both types of superhumps, and discuss the mechanisms for superhumps in the X-ray and optical regimes."
 We have re-analysecl existing N-rav. photometry that was presented by Chou et al. (, We have re-analysed existing X-ray photometry that was presented by Chou et al. (
2001).,2001).
 The clatasct consists of about 190 ksee of observations by the RAPE satellite obtained during 17 occasions between 1900 February to October., The dataset consists of about 150 ksec of observations by the RXTE satellite obtained during 17 occasions between 1996 February to October.
 In Figs., In Figs.
 la & 2a the power spectrum (Scargle 1982) of 10 successive runs in 1996 May is presented., 1a $\&$ 2a the power spectrum (Scargle 1982) of 10 successive runs in 1996 May is presented.
 No de-trencding method. was used., No de-trending method was used.
 Bursts were not observed. during these observations., Bursts were not observed during these observations.
 In addition to the two known periods (3001 & 3028 s. marked as fj & £5) and their + aliases. there is a rd peak (labelled £5) together with its 1-d+ alias pattern.," In addition to the two known periods (3001 $\&$ 3028 s, marked as $_{1}$ $\&$ $_{2}$ ) and their $^{-1}$ aliases, there is a third peak (labelled $_{3}$ ) together with its $^{-1}$ alias pattern."
 After fitting and subtracting the first term of the two known requencies (£j & fs). the third. which corresponds to the »eriodicitv O.O34483+0.000013 α (2979.3-E1.1 s) becomes he strongest peak in the residual power spectrum (Lies.," After fitting and subtracting the first term of the two known frequencies $_{1}$ $\&$ $_{2}$ ), the third, which corresponds to the periodicity $\pm$ 0.000013 d $\pm$ 1.1 s), becomes the strongest peak in the residual power spectrum (Figs."
 Lb & 2b)., 1b $\&$ 2b).
 The error was calculated by 1000 simulations (see e.g. Xettor οἱ al., The error was calculated by 1000 simulations (see e.g. Retter et al.
 2002b)., 2002b).
 The subtraction of nine more terms of he periods (to account for a non-sinusoidal shape) and the 3.9-cl period (whose peak in the raw power spectrum Is not significant) vields similar results., The subtraction of nine more terms of the periods (to account for a non-sinusoidal shape) and the 3.9-d period (whose peak in the raw power spectrum is not significant) yields similar results.
 This point is valid for all ests Listed below., This point is valid for all tests listed below.
 The power spectrum of the whole 1996 dataset resembles the result of the 10 runs in Mas: however. it is noisier as the gaps between the observations are long.," The power spectrum of the whole 1996 dataset resembles the result of the 10 runs in May, however, it is noisier as the gaps between the observations are long."
 In Fig., In Fig.
 3 we show the mean shape of the 2979-5 period, 3 we show the mean shape of the 2979-s period
radio sources.,radio sources.
 ? also found. a significant (~0.3 Gvr) delay between the starburst event and Radio-AGN activity lor the radio galaxy B2 06484-27., \cite{Emonts06} also found a significant $\sim$ 0.3 Gyr) delay between the starburst event and Radio-AGN activity for the radio galaxy B2 0648+27.
 These results are consistent with a scenario in which the radio activity is triggered relatively late in (he merger sequence. following (he major merger-induced starburst.," These results are consistent with a scenario in which the radio activity is triggered relatively late in the merger sequence, following the major merger-induced starburst."
 ILowever. ?.— also discussed the possibility that some radio sources have undergone multiple epochs of activity.," However, \cite{Tadhunter05} also discussed the possibility that some radio sources have undergone multiple epochs of activity."
 Indeed. ? detected low surface brightness extended radio emission surrounding several objects in (their sample of otherwise compact (D.<1 kpe) radio sources. including PIXS13454-12.," Indeed, \cite{Stanghellini05} detected low surface brightness extended radio emission surrounding several objects in their sample of otherwise compact $D < 1$ kpc) radio sources, including PKS1345+12."
. Moreover. this ealaxv shows sone evidence lor continuity between (he compact and extended: emission. which. if confirmed. would suggest that PINS13454-12 has had AGN-Raclio activity over a longer timescale than the estimated age of (he GPS source (0.1 Myr).," Moreover, this galaxy shows some evidence for continuity between the compact and extended emission, which, if confirmed, would suggest that PKS1345+12 has had AGN-Radio activity over a longer timescale than the estimated age of the GPS source $<$ 0.1Myr)."
 Given the evidence for both multiple epochs of star formation. including some YSP with ages less (han 6 Myr. and multiple epochs of radio jet activitv. we conclude that there is no clear evidence Lor a delay between the merger-induced starburst and the start of the AGN phase in PINS13454-12.," Given the evidence for both multiple epochs of star formation, including some YSP with ages less than 6 Myr, and multiple epochs of radio jet activity, we conclude that there is no clear evidence for a delay between the merger-induced starburst and the start of the AGN phase in PKS1345+12."
 In Section 3.3 we presented evidence that (he SSC and their associated IL] regions are moving at high racial velocities relative to (he diffuse ISM detected (ii projection) at similar locations in the halo of the host galaxy. despite the morphological evidence for associations with diffuse ISA] components.," In Section 3.3 we presented evidence that the SSC and their associated HII regions are moving at high radial velocities relative to the diffuse ISM detected (in projection) at similar locations in the halo of the host galaxy, despite the morphological evidence for associations with diffuse ISM components."
 It is interesting to consider how such rapidly moving star clusters might form., It is interesting to consider how such rapidly moving star clusters might form.
 The first possibility recognises (he [act that major galaxy mergers are violent. events. and considerable dissipation must take place before the σας finally settles into a stable configuration in the merged svstem.," The first possibility recognises the fact that major galaxy mergers are violent events, and considerable dissipation must take place before the gas finally settles into a stable configuration in the merged system."
 At all stages in a merger. but particularly the final stages belore the nuclei of the (wo merging galaxies coalesce. it is likely (hat rapidly moving eas sireams/tidal tails will rain down on the central regions of the host galaxy.," At all stages in a merger, but particularly the final stages before the nuclei of the two merging galaxies coalesce, it is likely that rapidly moving gas streams/tidal tails will rain down on the central regions of the host galaxy."
 As individual eas streams collide with the ambient gas or other gas streams. they will generale shocks that may (igger star formation (?)..," As individual gas streams collide with the ambient gas or other gas streams, they will generate shocks that may trigger star formation \citep{Barnes04}."
 The ionizing photons generated in the shocks as the tidal streams interact with the ambient gas may also lead to a local enhancement in the enission line luminosity of (he ambient eas. perhaps explaining the effect seen in Figure 11.," The ionizing photons generated in the shocks as the tidal streams interact with the ambient gas may also lead to a local enhancement in the emission line luminosity of the ambient gas, perhaps explaining the effect seen in Figure 11."
" ""Therefore the formation of the SSC detected in PINS1345+12 may be associated with the overall settling/cdissipation process that must accompany the evolution of the ISAM followiadye a major galaxy merger: il 1s possible (hal we are wilnessing the events (hat will eventually", Therefore the formation of the SSC detected in PKS1345+12 may be associated with the overall settling/dissipation process that must accompany the evolution of the ISM following a major galaxy merger; it is possible that we are witnessing the events that will eventually
from both clireet ancl indirect. collisions for. νο ολ... Mus m100M. and à=1.2.,"from both direct and indirect collisions for $M_{min}$ $_{\odot}$ , $M_{max}$ $_{\odot}$ and $\alpha=1.2$."
 We find that at small raclii the mass loss rate is dominated by direct collisions. and at large racii it is clominatec by indirect collisions.," We find that at small radii the mass loss rate is dominated by direct collisions, and at large radii it is dominated by indirect collisions."
 Mass loss cue to indirect collisions is suppressed in the Galactic center. due to the very fast relative stellar velocities.," Mass loss due to indirect collisions is suppressed in the Galactic center, due to the very fast relative stellar velocities."
 Even though the high velocities (and high densities) in the Galactic center make collisions more frequent. under the impulse approximation. when velocities are very. fast. mass loss is minimized.," Even though the high velocities (and high densities) in the Galactic center make collisions more frequent, under the impulse approximation, when velocities are very fast, mass loss is minimized."
 To illustrate which mass stars contribute the most to the total mass loss rate. we plot dAldinAl. as à function of M in Fig.," To illustrate which mass stars contribute the most to the total mass loss rate, we plot $d\dot M/dlnM_{pd}$ as a function of $M_{pd}$ in Fig."
 11. for several different PDMESs., \ref{fig:dmdot_dlnm_v_m} for several different PDMFs.
 The range of integration we choose for rgar is from 0 to O.06pe (see 85))., The range of integration we choose for $r_{gal}$ is from 0 to 0.06pc (see \ref{sec:constrain}) ).
" We choose Aly, to be 0.05. 0.5 andSAL. (left to right in the figure). and we use a constant M, of 125M.."," We choose $M_{min}$ to be 0.05, 0.5 and$_{\odot}$ (left to right in the figure), and we use a constant $M_{max}$ of $_{\odot}$."
 In each panel. we vary à from 1.5 to 2.5 in equal increments.," In each panel, we vary $\alpha$ from 1.5 to 2.5 in equal increments."
 The figure shows that for ων=0.05 M. changing à has little effect on what mass stars contribute the most to the mass loss rate (although. the total mass loss rate is decreased with increasing a).," The figure shows that for $M_{min}=0.05$ $_{\odot}$, changing $\alpha$ has little effect on what mass stars contribute the most to the mass loss rate (although, the total mass loss rate is decreased with increasing $\alpha$ )."
 For the Ανν= 0.5M. and SAL. cases. increasing à results in lower mass stars contributing more to the mass loss rate.," For the $M_{min}=0.5$ $_{\odot}$ and $5$ $_{\odot}$ cases, increasing $\alpha$ results in lower mass stars contributing more to the mass loss rate."
 This trend makes sense. since PDME profiles with higher values of à have fractionally more lower mass stews.," This trend makes sense, since PDMF profiles with higher values of $\alpha$ have fractionally more lower mass stars."
 To test how our interpolation between the m=1.5 and 3 polvtropic indeces allects the main results of this paper. we consider two extreme cases.," To test how our interpolation between the $n=1.5$ and 3 polytropic indeces affects the main results of this paper, we consider two extreme cases."
" The first case we consider has n—1.5 lor AZ,<LU. and Al,75M.. and n=3 for 1M.xM,SAL..."," The first case we consider has $n=1.5$ for $M_{\star}<1M_{\odot}$ and $M_{\star}>5M_{\odot}$, and $n=3$ for $1M_{\odot}\le M_{\star}\le 5M_{\odot}$."
 This approach has n=1.5 for much of the mass spectrum. and should result in the highest. mass loss rates since (as is evident [rom Fig. 1))," This approach has $n=1.5$ for much of the mass spectrum, and should result in the highest mass loss rates since (as is evident from Fig. \ref{fig:mass_loss}) )"
 collisions with the perturbed star having n=1.5 result in the most mass lost., collisions with the perturbed star having $n=1.5$ result in the most mass lost.
 This is due to the fact that for p=1.5 stars. the mass is less centrally concentrated. aad more mass can therefore escape at large racii which receive a stronger velocity kick.," This is due to the fact that for $n=1.5$ stars, the mass is less centrally concentrated, and more mass can therefore escape at large radii which receive a stronger velocity kick."
" The second case we consider has n=1.5 for M,«0.3M. and Ad,10A/.. and η=3 for OBAL.xM,LOAL.."," The second case we consider has $n=1.5$ for $M_{\star}<0.3M_{\odot}$ and $M_{\star}>10M_{\odot}$, and $n=3$ for $0.3M_{\odot}\le M_{\star}\le 10M_{\odot}$."
 This case should result in the smallest mass loss rates. since it has»=1.5 for a smaller fraction of the mass spectrum.," This case should result in the smallest mass loss rates, since it has $n=1.5$ for a smaller fraction of the mass spectrum."
 Since different mass functions have dillerent fractions of the total mass in the neighborhood of LAL. (where we expect the least mass loss per collision since n= 3). we test the two cases for several dillerent mass functions.," Since different mass functions have different fractions of the total mass in the neighborhood of $1M_{\odot}$ (where we expect the least mass loss per collision since $n=3$ ), we test the two cases for several different mass functions."
 We find that clillerences in dMπμνα) For both cases are relatively minor. dillering at most by ~10% depending on the mass function that we use.," We find that differences in $d\dot{M}/dlnr_{gal} (r_{gal})$ for both cases are relatively minor, differing at most by $\sim10\%$ depending on the mass function that we use."
]t is known through x-ray observations [romChandra. tha the central supermassive black hole in the Galactic center is surrounded by eas donated from stellar winds (e.g. Baganoll 2003).,"It is known through x-ray observations from, that the central supermassive black hole in the Galactic center is surrounded by gas donated from stellar winds (e.g. Baganoff 2003)."
 The x-ray luminosity is due to Dremsstrahlung emission from unbouncl material supplied at a rate of ~10ΝΕ.vet (Najarroetal.1997)., The x-ray luminosity is due to Bremsstrahlung emission from unbound material supplied at a rate of $\sim 10^{-3} \mathrm{M_{\odot}yr^{-1}}$ \citep{najarro:1997}.
. This unbound materia has been studied theoretically by Quataert(2004)... who solved. the equations of hvedrodynamies (under. spherica symmetry) to follow how the gas is accreted onto Ser A," This unbound material has been studied theoretically by \citet{quataert:2004}, who solved the equations of hydrodynamics (under spherical symmetry) to follow how the gas is accreted onto Sgr A*."
" Quataert(2004) finds that his mocdel agrees with the level of cilfuse x-ray emission. measured byCheandre.. ancl predicts . ∕∕ ⋜⋯↓⊔∐∪∖∖⋎∪⇂⊔↓⋜↧⊳∖⊳∖⋜∐∣⋮∠∕⊓∣∿↓⋜∐⋜↧↓⋅⋜⋯⋅∪⇂∿↓∪ ""Above."," \citet{quataert:2004} finds that his model agrees with the level of diffuse x-ray emission measured by, and predicts an inflow of mass at $r_{gal} \sim 1''$at a rate of $\sim 10^{-5} \mathrm{M_{\odot} yr^{-1}}$ ."
 ]t is interesting to note that the mass loss rate profiles due to stellar collisions are consistent with the observed. x- surface brightness profile., It is interesting to note that the mass loss rate profiles due to stellar collisions are consistent with the observed x-ray surface brightness profile.
 From his derived temperature and density profiles. Quatacrt(2004). solves for the x-ray surface brightness as a function of ry... Quatacrt(2004) ," From his derived temperature and density profiles, \citet{quataert:2004} solves for the x-ray surface brightness as a function of $r_{gal}$ \citet{quataert:2004} "
The ATLAS (H-ATLAS) (Eales et al. 2010)),The ATLAS (H-ATLAS) (Eales et al. \cite{eales}) )
" will survey over 550 deg? in 5 wavebands at 100, 160, 250, 350 and 500 um. One field covering ~4°x degrees was observed during the science demonstration phase of the mission, and has produced a catalogue of ~6600 sources with significance >5c above the combined instrumental and confusion noise."," will survey over 550 $^2$ in 5 wavebands at 100, 160, 250, 350 and $500
\ \mu$ m. One field covering $\sim 4^\circ\times4^\circ $ degrees was observed during the science demonstration phase of the mission, and has produced a catalogue of $\sim 6600$ sources with significance $>5
\sigma$ above the combined instrumental and confusion noise."
 This represents roughly 1/30 of the final H-ATLAS data-set., This represents roughly 1/30 of the final H-ATLAS data-set.
 In this paper we present measurements of the angular correlation function of five flux and colour-selected samples of the H-ATLAS sources., In this paper we present measurements of the angular correlation function of five flux and colour-selected samples of the H-ATLAS sources.
" H-ATLAS uses parallel scan mode observations performed with the ESA Space Observatory (Pilbratt et al. 2010)),"," H-ATLAS uses parallel scan mode observations performed with the ESA Space Observatory (Pilbratt et al. \cite{herschel}) ),"
 providing data simultaneously from both the PACS (Poglitsch et al. 2010)), providing data simultaneously from both the PACS (Poglitsch et al. \cite{pacs}) )
 and SPIRE (Griffin et al. 2010)), and SPIRE (Griffin et al. \cite{spire}) )
 instruments., instruments.
 The time-line data are reduced usingHIPE., The time-line data are reduced using.
 Maps are produced from the SPIRE data using a naive mapping technique after removing instrumental temperature variations from the time-line data (Pascale et al., Maps are produced from the SPIRE data using a naive mapping technique after removing instrumental temperature variations from the time-line data (Pascale et al.
 in prep)., in prep).
" Noise maps are generated by using the two cross-scan measurements to estimate the noise per detector pass, and then for each pixel the noise is scaled by Νικος where Nj; is the number of detector passes."," Noise maps are generated by using the two cross-scan measurements to estimate the noise per detector pass, and then for each pixel the noise is scaled by $N_{passes}^{1/2}$, where $N_{passes}$ is the number of detector passes."
" A false-colour image combining the SPIRE 250, 350 and um maps is shown in Fig. 1.."," A false-colour image combining the SPIRE 250, 350 and $\ \mu$ m maps is shown in Fig. \ref{fig:1}."
" The PACS H-ATLAS maps currently yield only a few hundred sources, which is insufficient for us to attempt a clustering measurement."," The PACS H-ATLAS maps currently yield only a few hundred sources, which is insufficient for us to attempt a clustering measurement."
 The clustering in the PACS bands will be investigated in a future paper using more data., The clustering in the PACS bands will be investigated in a future paper using more data.
" Sources were identified in the SPIRE maps using a Multi-band Algorithm for source eXtraction (MADX, Maddox et al."," Sources were identified in the SPIRE maps using a Multi-band Algorithm for source eXtraction (MADX, Maddox et al."
 in prep)., in prep).
 First a local background is estimated from the peak of the histogram of pixel values in 30x blocks of pixels., First a local background is estimated from the peak of the histogram of pixel values in $30\times 30$ blocks of pixels.
" This corresponds to 2.5' for the um map, and 5' for the 350 and um maps."," This corresponds to $2.5'$ for the $\ \mu$ m map, and $5'$ for the 350 and $\ \mu$ m maps."
" The background at each pixel is then estimated using a bi-cubic interpolation between the coarse grid of backgrounds, and this value subtracted from the pixel."," The background at each pixel is then estimated using a bi-cubic interpolation between the coarse grid of backgrounds, and this value subtracted from the pixel."
 The filter scale was chosen to be as large as possible while still following the variations in the cirrus background., The filter scale was chosen to be as large as possible while still following the variations in the cirrus background.
" Since the local background is estimated from the peak of the flux histogram, it is insensitive to the presence of resolved sources within the background block, so long as they do not cover a significant number of pixels in the block."," Since the local background is estimated from the peak of the flux histogram, it is insensitive to the presence of resolved sources within the background block, so long as they do not cover a significant number of pixels in the block."
" This approach should remove the local background without removing flux from the resolved sources, and so should be less susceptible to removing real structure in the source distribution compared to standard Fourier filtering approaches."," This approach should remove the local background without removing flux from the resolved sources, and so should be less susceptible to removing real structure in the source distribution compared to standard Fourier filtering approaches."
" The background subtracted maps are then filtered by the estimated PSF, including a local inverse variance weighting."," The background subtracted maps are then filtered by the estimated PSF, including a local inverse variance weighting."
" The maps from all three bands are then combined with weights set by the local inverse variance, and also the prior expectation of the SED of the galaxies."," The maps from all three bands are then combined with weights set by the local inverse variance, and also the prior expectation of the SED of the galaxies."
" We tried a flat-spectrum prior, where equal weight is given to each band and also um weighting, where only the um band was included."," We tried a flat-spectrum prior, where equal weight is given to each band and also $\ \mu$ m weighting, where only the $\ \mu$ m band was included."
" At the depth of the filtered maps source confusion becomes an issue in the longer wavelength bands, and the higher resolution of the um maps outweighs the signal-to-noise gain from adding in the other bands."," At the depth of the filtered maps source confusion becomes an issue in the longer wavelength bands, and the higher resolution of the $\ \mu$ m maps outweighs the signal-to-noise gain from adding in the other bands."
 The current catalogues use the um-only prior and we will revisit this issue in future data releases., The current catalogues use the $\ \mu$ m-only prior and we will revisit this issue in future data releases.
" All local peaks are identified in the combined PSF filtered map as potential sources, and a Gaussian is fitted to each peak to give estimates of the position at the sub-pixel level"," All local peaks are identified in the combined PSF filtered map as potential sources, and a Gaussian is fitted to each peak to give estimates of the position at the sub-pixel level"
 (seePi-forarecentreview). (Rhoads (seec.g.Deuenuauu1999).. (e.g.Frailetal.2001:CdirlaudaFriedinan&Bloom2005).. (c.g.Sarietal.1998:Rhoads1999," \citep[see][for a recent review]{Piran2005}, \citep{Rhoads1999}. \citep[see e.g.][]{Beuermann1999}, \citep[e.g.][]{Frail2001,Ghirlanda2004, Friedman2005}. \citep[e.g.][]{SarPirNar1998,Rhoads1999},"
).. (Johannessonet2005) Johannessonctal., \citep{Johannesson2006} \citet{Johannesson2006}.
(2005).. Rhoads(1999) afterglow Πο curves where the model parameter values were selected at random frou a normal distribution over a narrow parameter range eiven in table L.., \citet{Rhoads1999} afterglow light curves where the model parameter values were selected at random from a normal distribution over a narrow parameter range given in table \ref{tab:parameterrange}.
 We assum a constant density mterstellur medium for all events., We assume a constant density interstellar medium for all events.
 The standard deviation of the distribution for cach parameter was fixed at L/1 of its range and the distribution was clipped to fit within cach parameter rauge., The standard deviation of the distribution for each parameter was fixed at 1/4 of its range and the distribution was clipped to fit within each parameter range.
 The rauge of paraincters was chosen after fitting several afterelows with our model over a wide range of frequencies (πουex-etal. 2005).," The range of parameters was chosen after fitting several afterglows with our model over a wide range of frequencies \citep[see examples
in][]{Johannesson2006, Postigo2005}."
". Although the fraction of cucrev contained in electrons. ej. and maguetic field. ep. do not directly euter our formalisin. they are included as they can affect the results via the characteristic svuchrotron frequencics. Dy, aud vy. (Sarictal.1998)."," Although the fraction of energy contained in electrons, $\epsilon_e$, and magnetic field, $\epsilon_B$, do not directly enter our formalism, they are included as they can affect the results via the characteristic synchrotron frequencies, $\nu_m$ and $\nu_c$ \citep{SarPirNar1998}."
. The viewing angele of the observer. 6.. can smooth the jet break (Granotctal.2001) and was also included as a parameter.," The viewing angle of the observer, $\theta_v$, can smooth the jet break \citep{Granot2001} and was also included as a parameter."
 The liebt curves were evenly sampled in the logarithm of observer time aud consisted of about 35 points cach (see fie., The light curves were evenly sampled in the logarithm of observer time and consisted of about 35 points each (see fig.
 1 for a typical example)., \ref{fig:exlc} for a typical example).
 To make the svuthetic data more realistic. we added normally distributed random fluctuations with a standard deviation of3%.," To make the synthetic data more realistic, we added normally distributed random fluctuations with a standard deviation of."
. We also assumed a fixed error of for cach data point., We also assumed a fixed error of for each data point.
 We find that our results are not sensitive to the values of these error parameters ancl lowerlng the error estimate only reduces the uuuber of successful fits without reducing the scatter in our results., We find that our results are not sensitive to the values of these error parameters and lowering the error estimate only reduces the number of successful fits without reducing the scatter in our results.
 The results were also tested for the effects of the EATS by turning it off in a test sample., The results were also tested for the effects of the EATS by turning it off in a test sample.
 We found that it had no senificaut effect., We found that it had no significant effect.
 According to standard fireball theory. the sinplest afterelow light curves can be described with the widely adopted analytical approximation (Rhoads1999).. for a given observing frequency. v. m the rauge BOO ," According to standard fireball theory, the simplest afterglow light curves can be described with the widely adopted analytical approximation \citep{Rhoads1999}, for a given observing frequency, $\nu$ , in the range $\nu_m < \nu <
\nu_c$ ."
dero tis the time from the onset of the burst iu the observer's frame. fj is the jet break tine iu the observers frame aud p isthe electron energydistribution index.," Here $t$ is the time from the onset of the burst in the observer's frame, $t_j$ is the jet break time in the observer's frame and $p$ isthe electron energydistribution index."
 Since the jet break is iu geucral uot sharp. eq. (1))," Since the jet break is in general not sharp, eq. \ref{eq:brapp}) )"
 is often replaced with a simoothly joiued. broken power, is often replaced with a smoothly joined broken power
"the lower boundary to 04,=O.Ole,y and thus allow the coronal magnetic field to evolve quasi-statically until it erupts dvnanmically.",the lower boundary to $v_{\rm rise} = 0.01 v_{a0}$ and thus allow the coronal magnetic field to evolve quasi-statically until it erupts dynamically.
" Figures 4 and 5 show snapshots of the 3D coronal magnetic field evolution (as viewed from 2 different perspectives) after the initial stage of relatively [ast emergence has ended al /=L2. and the speed for driving the flux emergence al the lower boundary. has been reduced (to ο=001c,4."," Figures \ref{fig4} and \ref{fig5} show snapshots of the 3D coronal magnetic field evolution (as viewed from 2 different perspectives) after the initial stage of relatively fast emergence has ended at $t=1.2$, and the speed for driving the flux emergence at the lower boundary has been reduced to $v_{\rm rise} = 0.01 v_{a0}$."
 The view shown in Figure 4. corresponds (o the observation perspective at the Gime of the flare. lor which the center of the emerging region (also the center of the simulation lower boundary) is located at 7.1°S and 24°W [rom the solar disk center (or the line-ol-sight).," The view shown in Figure \ref{fig4} corresponds to the observation perspective at the time of the flare, for which the center of the emerging region (also the center of the simulation lower boundary) is located at $7.1^{\circ}$ S and $24^{\circ}$ W from the solar disk center (or the line-of-sight)."
 GIF movies for the evolution shown in Figure + and Figure are available in (he electronic version of (he paper., GIF movies for the evolution shown in Figure \ref{fig4} and Figure \ref{fig5} are available in the electronic version of the paper.
 We see that the emerged coronal [lux rope settles into a cquasi-static rise phase and (hen uidergoes a dynamic eruption., We see that the emerged coronal flux rope settles into a quasi-static rise phase and then undergoes a dynamic eruption.
" Figure shows the evolution of the rise velocity. 0, measured al the apex of the tracked axis of the emerged fIux rope (triangle points). aud also measured at the leading edge of the flux rope (crosses)."," Figure \ref{fig6} shows the evolution of the rise velocity $v_r$ measured at the apex of the tracked axis of the emerged flux rope (triangle points), and also measured at the leading edge of the flux rope (crosses)."
 After the emergence is slowed down al /=1.2. the rise velocity at the apex of the flux rope axis slows down. ancl undergoes some small oscillations as Che [Inx rope settles into a quasi-static rise.," After the emergence is slowed down at $t=1.2$, the rise velocity at the apex of the flux rope axis slows down, and undergoes some small oscillations as the flux rope settles into a quasi-static rise."
 The quasi-static rise phase extends from about /=1.2 until about |—2.5. over a lime period of 1.3. long compared to the dvnamic time scale of ~0.1 for the estimated Allvénn crossing time of the flix rope.," The quasi-static rise phase extends from about $t=1.2$ until about $t=2.5$, over a time period of $1.3$, long compared to the dynamic time scale of $\sim 0.1$ for the estimated Alfvénn crossing time of the flux rope."
 At about /= 2.5. the flux rope axis starts to accelerate significantly and a dynamic eruption ensues.," At about $t=2.5$ , the flux rope axis starts to accelerate significantly and a dynamic eruption ensues."
" The flux emergence is stopped αἱ |—2,8. alter which the flux rope continues to accelerate outward."," The flux emergence is stopped at $t=2.8$, after which the flux rope continues to accelerate outward."
" We are able to follow the acceleration of the axial field line up to c,=0.541050 kin/s at /=3.. when the axial Ποια line undergoes a reconnection and we are subsequently unable to track it."," We are able to follow the acceleration of the axial field line up to $v_r = 0.54 = 1050 $ km/s at $t=3.$, when the axial field line undergoes a reconnection and we are subsequently unable to track it."
" Figure 6 also shows ¢, measured at the leading edge of the low density cavity (as shown in Figure 1)). corresponding to the expanding flux rope."," Figure \ref{fig6} also shows $v_r$ measured at the leading edge of the low density cavity (as shown in Figure \ref{fig7}) ), corresponding to the expanding flux rope."
" Wefind that by the time of about /=3.2. a shock front followed by a condensed sheath has formed ahead of the flix rope cavity. (see Figure 7 ad /= 3.25). and the ve, measured at the front edge of the cavity (or the inner edge of the sheath) reaches a steady speed of about 0.425 or 530 kin/s (see crosses in Figure 6))."," Wefind that by the time of about $t=3.2$, a shock front followed by a condensed sheath has formed ahead of the flux rope cavity (see Figure \ref{fig7} at $t=3.25$ ), and the $v_r$ measured at the front edge of the cavity (or the inner edge of the sheath) reaches a steady speed of about 0.425 or 830 km/s (see crosses in Figure \ref{fig6}) )."
 When the flux rope begins significant. acceleration (at /2 2.5). the decay index which describes the rate of decline of the corresponding potential field Bp with height / is found to be ηzz1.2 at the apex of the flux rope axis. and η21.4 at the apex of the flix rope cavity.," When the flux rope begins significant acceleration (at $t \approx 2.5$ ), the decay index $n \equiv d \ln | {\bf B}_p | / d \ln h$ which describes the rate of decline of the corresponding potential field ${\bf B}_p$ with height $h$ is found to be $n \approx 1.2$ at the apex of the flux rope axis, and $n \approx 1.4$ at the apex of the flux rope cavity."
 These values are smaller (han (hecritical value of nag=1.5 for the onset of the torus instability for a circular current ring (Bateman19758:IXliemand 2010).. although there is a range of variability for the," These values are smaller than thecritical value of $n_{\rm crit} = 1.5$ for the onset of the torus instability for a circular current ring \citep{bateman1978, kliem_toeroek2006,
demoulin_aulanier2010}, , although there is a range of variability for the"
to inerease the > ray flux by c.g. increasing (2) would ouly result in au increase of the reprocessed radiation.,to increase the $\gamma-$ ray flux by e.g. increasing $\lginj$ would only result in an increase of the reprocessed radiation.
 Fig., Fig.
 11. depicts a series of similax ruus with R=Lote cu. à=10? and B=16 G (solid line). B=1.6 G (dashed lino) and B=16 € (dotted lino).," \ref{mw_b} depicts a series of similar runs with $R=10^{16}$ cm, $\delta=10^{1.2}$ and $B=.16$ G (solid line), $B=1.6$ G (dashed line) and $B=16$ G (dotted line)."
 While for low D values there is no quenching. this beegius to appear at higher values of the maguetic field.," While for low $B$ values there is no quenching, this begins to appear at higher values of the magnetic field."
 Note. however. hat he value of B=1.6 € produces au acceptable fit. )ecause the soft reprocessed radiation. docs not violate auv observable nuit (auost notably the N-ravs) aud it can xoduce a imareinally acceptable fit o TeV 3 rave.," Note, however, that the value of $B=1.6$ G produces an acceptable fit, because the soft reprocessed radiation does not violate any observable limit (most notably the X-rays) and it can produce a marginally acceptable fit to TeV $\gamma-$ rays."
 Based ou this. one can define a liniting value of the magnetic Beld Boasx(R.6) with the property that if Bo>Dus he quenching does not allow the internally produced 5 rav huumosity to escape.," Based on this, one can define a limiting value of the magnetic field $\bquemx(R,\delta)$ with the property that if $B>\bquemx$, the quenching does not allow the internally produced $\gamma-$ ray luminosity to escape."
" Therefore Byais is the highest allowed value that the B field cau take for the observed flux of TeV > rays,", Therefore $\bquemx$ is the highest allowed value that the $B-$ field can take for the observed flux of TeV $\gamma-$ rays.
 Figure 12 epicts the locus Baas as a function of the size of the source A for various Doppler factors à., Figure \ref{Blocus} depicts the locus $\bquemx$ as a function of the size of the source $R$ for various Doppler factors $\delta$.
 Each liue separates the parameter space into two regimes: Because of the quenching.i a choice of Π and B above the locus xoduces a fax too low 2. rav flux that camo fit the TeV observations.," Each line separates the parameter space into two regimes: Because of the quenching, a choice of $R$ and $B$ above the locus produces a far too low $\gamma-$ ray flux that cannot fit the TeV observations."
 For example. if one asstues tha à—10 and R=2.5«10/5 an so that it correspouds to a variability iuescale of 1 dav. the inagnetic ficld cannot exceed 0.3 C. However. owing to the steep depeudence of the flux on à (c.£.," For example, if one assumes that $\delta=10$ and $R=2.5\times 10^{16}$ cm – so that it corresponds to a variability timescale of 1 day, the magnetic field cannot exceed 0.3 G. However, owing to the steep dependence of the flux on $\delta$ (c.f."
 Equ 10). the above coucition is ercatly relaxed for ieher values of à.," Eqn 40), the above condition is greatly relaxed for higher values of $\delta$."
 We should also meution that Barus uot ouly αραιά» on ó and HR. but also ou the upper cutoff e.a of he injected power-law iu 5 ravs.," We should also mention that $\bquemx$ not only depends on $\delta$ and $R$, but also on the upper cutoff $\egmx$ of the injected power-law in $\gamma-$ rays."
 For the caleulatious orforiued here. we chose ἕως such as to match the üehbest. observed frequency. in. the data set. i.. NM1079 Tz.," For the calculations performed here, we chose $\egmx$ such as to match the highest observed frequency in the data set, i.e., $10^{26}$ Hz."
 Tf one asunes that the internally produced + rays cau exteud to even higher euergies then the value of Baas is affected as well.," If one assumes that the internally produced $\gamma-$ rays can extend to even higher energies, then the value of $\bquemx$ is affected as well."
" Accordingly we found that if we were to repeat the same calculations with e.4,44: ten times higher than the value we usec. then quenching would affect the flux more and the lines of Fig.12 would have to shift upwards by a factor of 2."," Accordingly we found that if we were to repeat the same calculations with $\egmx$ ten times higher than the value we used, then quenching would affect the flux more and the lines of Fig.12 would have to shift upwards by a factor of 2."
 Therefore. Fig.12 can be taken as a conservative lower limit on D.," Therefore, Fig.12 can be taken as a conservative lower limit on $\bquemx$."
 We have examined the effects of photon quenching on conipact > ray sources., We have examined the effects of photon quenching on compact $\gamma-$ ray sources.
 Automatic photon queuchine. first proposed by Sk. is a non-linear network of two processes that are very conunon in bhiel-enuerev astroplivsies. namely photon-photon amnililation and lepton svuchrotron radiation.," Automatic photon quenching, first proposed by SK, is a non-linear network of two processes that are very common in high-energy astrophysics, namely photon-photon annihilation and lepton synchrotron radiation."
 We verified and expaucded both analytically and ummerically the results of SIN aud we showed that this network. wheu applied to compact ~ rav sources. can give limits ou the Doppler factor. which depend on the source radius and on the magnetic field strength oulv.," We verified and expanded both analytically and numerically the results of SK and we showed that this network, when applied to compact $\gamma-$ ray sources, can give limits on the Doppler factor, which depend on the source radius and on the magnetic field strength only."
 This is a theoretically derived limit aud as such does not depeud ou the soft photon populations that might be present inside the source., This is a theoretically derived limit and as such does not depend on the soft photon populations that might be present inside the source.
 It is iuportaut to note that photon quenching can be investigated ouly when the radiation trausfer problems. is solved selt-cousisteutlv because it is a non-linear process that involves redistribution of energy from the 5 ravs to the ower energv parts of the photon spectrum., It is important to note that photon quenching can be investigated only when the radiation transfer problem is solved self-consistently because it is a non-linear process that involves redistribution of energy from the $\gamma-$ rays to the lower energy parts of the photon spectrum.
 For this we used a svsteni of kinetic equations that treats in the siuuplest case > rav annililation and electrou svuchrotron radiation., For this we used a system of kinetic equations that treats in the simplest case $\gamma-$ ray annihilation and electron synchrotron radiation.
 Using suitable 4 fiction approximatious for he photou-photon annihilation cross section aud for the photon svuchrotrou cluissivity. we were able to verify the results of SIX in the Πατ where svuchrotrom losses can © considered as catastrophic (see Fig. 1)).," Using suitable $\delta-$ function approximations for the photon-photon annihilation cross section and for the photon synchrotron emissivity, we were able to verify the results of SK in the limit where synchrotron losses can be considered as catastrophic (see Fig. \ref{lcr1}) )."
 Ou the other iud. using a more realistic step fiction for the cross section. we were able to find a very good agreement with munerical results. which utilized the full photou-photou aunihilation cross section.," On the other hand, using a more realistic step function for the cross section, we were able to find a very good agreement with numerical results, which utilized the full photon-photon annihilation cross section."
 Furthermore. we extended our iuuerncal study in the case where 5. rays are injected iu he form of a power-law.," Furthermore, we extended our numerical study in the case where $\gamma-$ rays are injected in the form of a power-law."
 We found that quenching can be responsible for wide spectral breaks. which become more xononunced as the injected + rays become stecper (see Fig. 9)).," We found that quenching can be responsible for wide spectral breaks, which become more pronounced as the injected $\gamma-$ rays become steeper (see Fig. \ref{ainj}) )."
 Quenching has some interesting consequences when applied to existing > rav sources., Quenching has some interesting consequences when applied to existing $\gamma-$ ray sources.
" For example. for the 2006 MIACIC observations of the quasar 3C'279, modelliug with combinations of low values of the Doppler factor 4 and high values of the maeuetic field streneth D lead to strong absorption of the injected + ravs."," For example, for the 2006 MAGIC observations of the quasar 3C279, modelling with combinations of low values of the Doppler factor $\delta$ and high values of the magnetic field strength $B$ lead to strong absorption of the injected $\gamma-$ rays."
 The oulv wav to avoid quenching is either to assume typical values of, The only way to avoid quenching is either to assume typical values of
The spherical-collapse model (e.g.2222?) is a fundamental ingredient in the theory of cosmic structure formation.,"The spherical-collapse model \citep[e.g.][]{Wang1998,Engineer2000,Mota2004,Bartelmann2006,Schaefer2008} is a fundamental ingredient in the theory of cosmic structure formation."
" Following the collapse of a slightly overdense, homogeneous sphere, it allows the derivation of two essential parameters; the overdensity A, of a virialised halo compared to the mean or the critical cosmic density, and the critical linear density contrast ó,."," Following the collapse of a slightly overdense, homogeneous sphere, it allows the derivation of two essential parameters; the overdensity $\Delta_\vv$ of a virialised halo compared to the mean or the critical cosmic density, and the critical linear density contrast $\delta_\cc$."
" The former is important because it allows relating sizes to masses of virialised structures, and the latter because it establishes a link between linear structure formation and the population statistics of collapsed haloes."," The former is important because it allows relating sizes to masses of virialised structures, and the latter because it establishes a link between linear structure formation and the population statistics of collapsed haloes."
" Despite fundamental doubts as to the validity of such a simplified model for accurate cosmological predictions, the parameters derived from the spherical-collapse model or variants thereof allow surprisingly far-reaching predictions such as the halo mass function or the correlation properties of haloes, which are confirmed by numerical simulations."," Despite fundamental doubts as to the validity of such a simplified model for accurate cosmological predictions, the parameters derived from the spherical-collapse model or variants thereof allow surprisingly far-reaching predictions such as the halo mass function or the correlation properties of haloes, which are confirmed by numerical simulations."
 The statistics of a Gaussian random field implies that spherical collapse should not occur., The statistics of a Gaussian random field implies that spherical collapse should not occur.
" In fact, the probability distribution for the eigenvalues 4; of the Zel'dovich deformation tensor (?) shows that spherical collapse has a vanishing probability, because it requires 4;=4»A."," In fact, the probability distribution for the eigenvalues $\lambda_i$ of the Zel'dovich deformation tensor \citep{Doroshkevich1970} shows that spherical collapse has a vanishing probability, because it requires $\lambda_1=\lambda_2=\lambda_3$."
" In this equation, 6=Aytap+ Az, 1=AdoΛο:+Λιλ"". and o? denotes the variance of the matter power spectrum smoothed on a scale that corresponds to a halo mass M."," In this equation, $\delta\equiv\lambda_1+\lambda_2+\lambda_3$ , $I\equiv\lambda_1\lambda_2+\lambda_2\lambda_3+\lambda_1\lambda_3$, and $\sigma^2$ denotes the variance of the matter power spectrum smoothed on a scale that corresponds to a halo mass $M$."
 Modifications of the original Press-Schechter mass function (??) motivated by ellipsoidal collapse (?) substantially improve the agreement of analytic predictions on the halo population with numerical simulations (e.g.theMillenniumsimulation3).," Modifications of the original Press-Schechter mass function \citep{Press1974,Bond1991} motivated by ellipsoidal collapse \citep{Sheth2002} substantially improve the agreement of analytic predictions on the halo population with numerical simulations \citep[e.g. the \emph{Millennium simulation} ."
" Ellipsoidal collapse was analysed many times before (see222,for examples).."," Ellipsoidal collapse was analysed many times before \citep[see][for examples]{Bartelmann1993, Eisenstein1995, Bond1996}."
" Several authors have worked with the model by ?, generalising it for different cosmologies and a"," Several authors have worked with the model by \citet{Bond1996}, generalising it for different cosmologies and introducing the scale factor $a$ as a time variable \citep{Monaco1995, Monaco1997, Monaco1998, Sheth2001, Sheth2002, Ohta2004}."
re reconsidering it here for two reasons., We are reconsidering it here for two reasons.
" First, we want to analyse how different assumptions on the treatment of the environment of a halo impact on the parameters A, and ὃς."," First, we want to analyse how different assumptions on the treatment of the environment of a halo impact on the parameters $\Delta_\vv$ and $\delta_\cc$."
" Previous assumptions have been that the principal-axis system of the homogeneous ellipsoid is either identical with that of the external gravitational shear field, or that the two eigensystems do not coincide, introducing rotation and a deviation from the homogeneous mass profile."," Previous assumptions have been that the principal-axis system of the homogeneous ellipsoid is either identical with that of the external gravitational shear field, or that the two eigensystems do not coincide, introducing rotation and a deviation from the homogeneous mass profile."
" We introduce another assumption here, letting the eigensystems of the collapsing ellipsoid and its surrounding shear field follow each other until turn-around of the major principal axis and then decoupling both."," We introduce another assumption here, letting the eigensystems of the collapsing ellipsoid and its surrounding shear field follow each other until turn-around of the major principal axis and then decoupling both."
" Second, we want to stop the collapse along any of the principal axes according to a physically motivated virialisation condition."," Second, we want to stop the collapse along any of the principal axes according to a physically motivated virialisation condition."
" Virialisation must be invoked to prevent the axes from collapsing to zero, and thus to be able to follow the entire collapse of an ellipsoid, i.e. the collapse of its three principal axes."," Virialisation must be invoked to prevent the axes from collapsing to zero, and thus to be able to follow the entire collapse of an ellipsoid, i.e. the collapse of its three principal axes."
" Conventionally, the collapse of each axis is stopped when a;(a)=0.177a, where a; is the scale factor of the i-th axis and a is the background scale factor (??).."," Conventionally, the collapse of each axis is stopped when $a_i(a)=0.177a$, where $a_i$ is the scale factor of the $i$ -th axis and $a$ is the background scale factor \citep{Bond1996, Sheth2001}."
" In this way, A=a?/(a,a2a3)178 for spherical collapse in an Einstein-de Sitter universe at the time when the third axis virialises."," In this way, $\Delta=a^3/(a_1a_2a_3)=178$ for spherical collapse in an Einstein-de Sitter universe at the time when the third axis virialises."
" However, the value 0.177 has no fundamental physical motivation, and there is no guarantee for it not to be different for ellipsoidal rather than spherical collapse, or when cosmologies other than EdS are to be considered."," However, the value $0.177$ has no fundamental physical motivation, and there is no guarantee for it not to be different for ellipsoidal rather than spherical collapse, or when cosmologies other than EdS are to be considered."
 We introduce a general virialisation condition based on the that avoids introducting such an uncalibrated factor., We introduce a general virialisation condition based on the that avoids introducting such an uncalibrated factor.
" We find substantial changes on both A, and 6, from both modifications and point out several discrepancies of our results with earlier studies.", We find substantial changes on both $\Delta_\vv$ and $\delta_\cc$ from both modifications and point out several discrepancies of our results with earlier studies.
" We introduce the ellipsoidal-collapse model including our modifications in Sect. D],"," We introduce the ellipsoidal-collapse model including our modifications in Sect. \ref{sec:model},"
 present our results in Sect., present our results in Sect.
 B] and conclude in Sect. 4]., \ref{sec:results} and conclude in Sect. \ref{sec:conclusions}.
" In this Section, we shall briefly review the ellipsoidal-collapse model of ? for cosmologies with a cosmological constant introducing the scale factor a as time variable."," In this Section, we shall briefly review the ellipsoidal-collapse model of \citet{Bond1996} for cosmologies with a cosmological constant introducing the scale factor $a$ as time variable."
" Furthermore, we shall present a physically motivated virialisation condition to stop the collapse of each axis, and show how to find the proper initial ellipticity and prolaticity as a function of mass and virialisation redshift."," Furthermore, we shall present a physically motivated virialisation condition to stop the collapse of each axis, and show how to find the proper initial ellipticity and prolaticity as a function of mass and virialisation redshift."
galaxies are nearly exponential and have no well-defined breaks as the center isapproached”.,galaxies are nearly exponential and have no well-defined breaks as the center is.
.. Their overall forms are completely dillerent. from the luminous core galaxies and thus they are trivially distinguished [rou them. lxormendsy(1985).. Ixormeudy(LOST)," Their overall forms are completely different from the luminous core galaxies and thus they are trivially distinguished from them. \citet{k85}, , \citet{k87}, ,"
... aud Ixorimeudy showed that ellipticals aud spleroicdals form well defined. cdisjoiut sequeuces iii parameter space.," and \citet{kb94} showed that ellipticals and spheroidals form well defined, disjoint sequences in parameter space."
 A recent analysis of the proliles of low-Iuminosity. ellipticals aud a comparison to a large sample ol Sph galaxies strongly coulirms the E/Sph dichotomy (Ixormendyetal.2007)., A recent analysis of the profiles of low-luminosity ellipticals and a comparison to a large sample of Sph galaxies strongly confirms the E/Sph dichotomy \citep{k07}.
. The resolution of this question. however. is not relevant to the strong physical dichotomy between power-law aud core galaxies advauced tu Faberοἱal.(1997).," The resolution of this question, however, is not relevant to the strong physical dichotomy between power-law and core galaxies advanced in \citet{f97}."
. Cores are a phenomenon of the most luminous galaxies., Cores are a phenomenon of the most luminous galaxies.
 Their likely iuunediate progenitors are the power-law galaxies of slightly. lower luminosity., Their likely immediate progenitors are the power-law galaxies of slightly lower luminosity.
 The distribution of s for Ady;«—19 is the relevant diagnostic [or how this transition occurred., The distribution of $\gamma'$ for $M_V<-19$ is the relevant diagnostic for how this transition occurred.
 This research was supported in part by several grauts provided through STScl associated witli GO programs 5512. 6099. GOST. 7388. 8591. 9106. aud 9107.," This research was supported in part by several grants provided through STScI associated with GO programs 5512, 6099, 6587, 7388, 8591, 9106, and 9107."
 Our team meetines were generously hosted by the National Optical Astronomy Observatory. the Observatories ofthe Carnegie Institution of Washineton. the Aspen Center for Physics. the Leiden. Observatory. aud the University. of California. Santa Cruz Center lor Adaptive Optics.," Our team meetings were generously hosted by the National Optical Astronomy Observatory, the Observatories of the Carnegie Institution of Washington, the Aspen Center for Physics, the Leiden Observatory, and the University of California, Santa Cruz Center for Adaptive Optics."
 We thank Steve Zepf for the WMA software., We thank Steve Zepf for the KMM software.
 We thank the releree for a careful review of the manuscript., We thank the referee for a careful review of the manuscript.
 We thank Dr. Laura Ferrarese for Ler comments on an earliercraft of thispaper., We thank Dr. Laura Ferrarese for her comments on an earlierdraft of thispaper.
eniploved. by Scehródder. Pols. Eeveletou (1997). and Pols et al. (,"employed by Schrödder, Pols, Eggleton (1997) and Pols et al. ("
1997) to test critically their evolutionary calculations (Pols et al.,1997) to test critically their evolutionary calculations (Pols et al.
 1995) aud tune parameters iu their scheme that take account of convection., 1995) and tune parameters in their scheme that take account of convection.
 Stellar radi are alternatively aud independently measured by interferometric techniques (e.g. Bichichi et al., Stellar radii are alternatively and independently measured by interferometric techniques (e.g. Richichi et al.
 1998) and also bv the so-called InfraRed Flux. Method. (Blacksvell LvuasCrav 199, 1998) and also by the so-called InfraRed Flux Method (Blackwell Lynas-Gray 1994).
" The later provides probably the most direct and modelindepeudent estimate of the stellar effective temperature,", The later provides probably the most direct and model-independent estimate of the stellar effective temperature.
 Other parameters involved are the cheiiical couposition aud the age., Other parameters involved are the chemical composition and the age.
" Although it is possible to check the metallicity estimates from the isochrones with the results from spectroscopic micasurements, and the y.cellar ages cau be compared with other methods. such as activity indicators (sec. e.g. Rocha-Pinto Maciel 1998). or the abuudances of radioactive nuclei (see. ο,ον, Coricly Clerbaux 1999). we shall not conceutrate on them here."," Although it is possible to check the metallicity estimates from the isochrones with the results from spectroscopic measurements, and the stellar ages can be compared with other methods, such as activity indicators (see, e.g., Rocha-Pinto Maciel 1998), or the abundances of radioactive nuclei (see, e.g., Goriely Clerbaux 1999), we shall not concentrate on them here."
 Tn this paper. we have uade use of the stellar evolutionary calculations of Bertelli ct al. (," In this paper, we have made use of the stellar evolutionary calculations of Bertelli et al. ("
1991) to study the possibility of retrieving the fundamental stellar paraiucters: radius. ass. and effective teniperature. from the ceuparison of the stars’ position iu the colow-magnitude diagram with couations of stellar evolution for isolated. stars.,"1994) to study the possibility of retrieving the fundamental stellar parameters: radius, mass, and effective temperature, from the comparison of the stars' position in the colour-magnitude diagram with computations of stellar evolution for isolated stars."
 We quantity the degencracy between lnass aud age for a given place in the WR diagram usine error bars in the estimates of these parameters., We quantify the degeneracy between mass and age for a given place in the HR diagram using error bars in the estimates of these parameters.
 Stars in eclipsing binary svstems aud the InfraRed Flux Method are used to test the procedure., Stars in eclipsing binary systems and the InfraRed Flux Method are used to test the procedure.
 And finally. the technique is applied to 17.219 stars identified by the astrometric niüssion within au sphere with radius of 100 pc ceutred in the Sun.," And finally, the technique is applied to 17,219 stars identified by the astrometric mission within an sphere with radius of 100 pc centred in the Sun."
 Calculations of stellar evolution interpolated to fiud the yosition of coeval stars of a given chemical composition but different masses are conuuonlv referred to as isochrones., Calculations of stellar evolution interpolated to find the position of coeval stars of a given chemical composition but different masses are commonly referred to as isochrones.
 Among the published isochrones. those presented by Bertelli et al. (," Among the published isochrones, those presented by Bertelli et al. ("
1991) span all observed stellar masses. luctallicitics from |M/II] 2. 1.7 to 0.LL and stellar ages frou ἐν10° to 2«1029 vrs.,"1994) span all observed stellar masses, metallicities from [M/H] = $-1.7$ to 0.4, and stellar ages from $ 4 \times 10^6$ to $2 \times 10^{10}$ yrs."
 The computations mostly use the OPAL radiative opacities bv Rogers Iglesias (1992) and Iglesias et al. (, The computations mostly use the OPAL radiative opacities by Rogers Iglesias (1992) and Iglesias et al. (
1992).,1992).
 The reader is referred to Bertelli et al. (, The reader is referred to Bertelli et al. (
1991) and references therein for details.,1994) and references therein for details.
 We assume that both the V colour. and the absolute magnitude (aud therefore the istauce to the star) are known with enough precision to avoid siguificaut observational errors affecting derived quantities.," We assume that both the $-$ V colour, and the absolute magnitude (and therefore the distance to the star) are known with enough precision to avoid significant observational errors affecting derived quantities."
 We quantifv this statement below., We quantify this statement below.
 To retrieve he appropriate set of stellar parameters that reproduces the position of a star in the colouranunagnitude diagram. we proceed as follows., To retrieve the appropriate set of stellar parameters that reproduces the position of a star in the colour-magnitude diagram we proceed as follows.
 First we search the ecutive set of isochrones to fud what are. if auv. the V colours hat reproduce the observed absolute V. magnitude (1) within the observational errors.," First we search the entire set of isochrones to find what are, if any, the $-$ V colours that reproduce the observed absolute V magnitude $M_v$ ) within the observational errors."
 Then. the different possible solutions. correspouding to different ages. metallicities. masses. etc.," Then, the different possible solutions, corresponding to different ages, metallicities, masses, etc."
 are averaged to obtain a mean value for every stellar parameter. and an estimate of the uncertaiutv frou the standard deviation.," are averaged to obtain a mean value for every stellar parameter, and an estimate of the uncertainty from the standard deviation."
 The final uncertainty includes three components., The final uncertainty includes three components.
 First of all. the iutvinsic spread. as different evolutionary paths cross the same area of the colom-imaeuitude diagram.," First of all, the intrinsic spread, as different evolutionary paths cross the same area of the colour-magnitude diagram."
 Second. the noise introduced in the translation frou the theoretical loggJig plane to the observational colouranagnuitude diagram.," Second, the noise introduced in the translation from the theoretical $\log g - T_{\rm eff}$ plane to the observational colour-magnitude diagram."
 Also. observational errors in the two considered parameters (D.V aud AL.) coutribute.," Also, observational errors in the two considered parameters (B–V and $M_v$ ) contribute."
 Applvine he procedure along a grid of colours aud absolute magnitudes. we coustruct maps of the combined uucertaimties expected for the different stellar parameters.," Applying the procedure along a grid of colours and absolute magnitudes, we construct maps of the combined uncertainties expected for the different stellar parameters."
 We are imainly iuterested dn radi iuassess and ngu," We are mainly interested in radii, masses, and $T_{\rm eff}$ s:"
and quiescent inflows of hot gas from the galaxy halo (2).,and quiescent inflows of hot gas from the galaxy halo .
 Hydrodynamie simulations also propose that the secular evolution of galaxies provides a large enough reservoir of gas through stellar winds. which could give episodic AGN activity and substantial black hole growth ??).," Hydrodynamic simulations also propose that the secular evolution of galaxies provides a large enough reservoir of gas through stellar winds, which could give episodic AGN activity and substantial black hole growth ."
. To our knowledge this process has hot as yet been implemented in SAMs., To our knowledge this process has not as yet been implemented in SAMs.
 The different SMBH ueling modes above operate by design under different conditions and therefore make specitic predictions for the properties of he galaxies that host AGN. including their star-formation rate. morphology and environment.," The different SMBH fueling modes above operate by design under different conditions and therefore make specific predictions for the properties of the galaxies that host AGN, including their star-formation rate, morphology and environment."
 This offers an observational tool to est different SMBH growth models., This offers an observational tool to test different SMBH growth models.
 The evidence above has motivated efforts to study AGN josts using data from large multiwavelength programs. such as he AEGIS 1). the COSMOS 2). the GOODS (Great Observatories Origins Survey) and the XBOOTES (2).," The evidence above has motivated efforts to study AGN hosts using data from large multiwavelength programs, such as the AEGIS , the COSMOS , the GOODS (Great Observatories Origins Survey) and the XBOOTES ."
 These programs use primarily X-rays to locate AGN and multiwavelength observations to study their hosts., These programs use primarily X-rays to locate AGN and multiwavelength observations to study their hosts.
 As a result observational constraints to different AGN fueling models have become available in the last few vears. although the results are often contradictory and the interpretation may vary among groups.," As a result observational constraints to different AGN fueling models have become available in the last few years, although the results are often contradictory and the interpretation may vary among groups."
 Morphological studies for example. have identified a large fraction of disks (z30 per cent) among X-ray selected AGN at 2zm1. suggesting that major mergers cannot be the dominant mode of SMBH growth at that epoch(2222).," Morphological studies for example, have identified a large fraction of disks $\approx 30$ per cent) among X-ray selected AGN at $z\approx1$, suggesting that major mergers cannot be the dominant mode of SMBH growth at that epoch."
. X-ray AGN at 21 are typically found in small and moderate size groups with dark matter halo masses zzή107AL.(22)...," X-ray AGN at $z\approx1$ are typically found in small and moderate size groups with dark matter halo masses $\rm \approx10^{12}-10^{13}\,M_{\odot}$."
 This environment is proposed to be the most conducive to merger events, This environment is proposed to be the most conducive to merger events.
 ?mainBodyCitationEnd469]Hopkins2008 .C'alacginlteractions. pauédiadaedshlioscusb (07a Nar deer Lype(red)spstem.arcindeedobservedlobemorecommoningroupen em θα (2).," Galaxy interactions, particularly those which include at least one early-type (red) system, are indeed observed to be more common in group environments at $z\approx1$ ."
 The star-formation rate of X-ray AGN. although still not well constrained. is suggested to depend on their accretion luminosity22).," The star-formation rate of X-ray AGN, although still not well constrained, is suggested to depend on their accretion luminosity."
. Below the break of the X-ray luminosity function it is claimed that the star-formation is decoupled to AGN activity. consistent with the stochastic accretion mode scenario.," Below the break of the X-ray luminosity function it is claimed that the star-formation is decoupled to AGN activity, consistent with the stochastic accretion mode scenario."
 Above the knee of the X-ray luminosity function it is argued that the accretion luminosity scales with the star-formation rate. in agreement with the major-merger formation picture.," Above the knee of the X-ray luminosity function it is argued that the accretion luminosity scales with the star-formation rate, in agreement with the major-merger formation picture."
 The observational studies above essentially probe active SMBHs close to the peak of the accretion power of the Universe. 27]. and lack the volume to provide meaningful constraints on the AGN population at low redshift z70.5.," The observational studies above essentially probe active SMBHs close to the peak of the accretion power of the Universe, $z\approx 1$, and lack the volume to provide meaningful constraints on the AGN population at low redshift $z\la 0.5$."
 This is an important gap in AGN evolution studies. as there are suggestions that the dominant mode of SMBH accretion changes with redshift. thereby leading to the observed rapid decline of the accretion power at 2c1(222).," This is an important gap in AGN evolution studies, as there are suggestions that the dominant mode of SMBH accretion changes with redshift, thereby leading to the observed rapid decline of the accretion power at $z<1$."
 for example. propose two main channels of SMBH growth. disk instabilities. which dominate at high redshift and bright AGN luminosities. and hot gas accretion. which becomes important at low redshift and faint luminosities.," for example, propose two main channels of SMBH growth, disk instabilities, which dominate at high redshift and bright AGN luminosities, and hot gas accretion, which becomes important at low redshift and faint luminosities."
 In the models above shifts between different fueling modes with cosmic time should imprint detectable changes in the properties of AGN hosts from 2= Oto >=| and beyond., In the models above shifts between different fueling modes with cosmic time should imprint detectable changes in the properties of AGN hosts from $z=0$ to $z=1$ and beyond.
 The Sloan Digital Sky Survey has identitied the largest sample of low redshift (2= 0.1) AGN todate. using diagnostic emission line ratios(2).," The Sloan Digital Sky Survey has identified the largest sample of low redshift $z\approx0.1$ ) AGN todate, using diagnostic emission line ratios."
. The selection function of that sample however. is very different from that of X-ray AGN in deep surveys. rendering the comparison difficult.," The selection function of that sample however, is very different from that of X-ray AGN in deep surveys, rendering the comparison difficult."
 Recent work has demonstrated that there is only partial overlap between AGN samples selected using diagnostic emission lines and X-rays?2)., Recent work has demonstrated that there is only partial overlap between AGN samples selected using diagnostic emission lines and X-rays.
. This might be due to obscuration effects. which are likely to be more severe at X-rays. dilution of the AGN optical emission lines by stellar continuum or nebular emission lines from HII regions. intrinsic scatter in the X-ray to optical luminosity ratio of AGN or the controversial nature of sources such as LINERs.," This might be due to obscuration effects, which are likely to be more severe at X-rays, dilution of the AGN optical emission lines by stellar continuum or nebular emission lines from HII regions, intrinsic scatter in the X-ray to optical luminosity ratio of AGN or the controversial nature of sources such as LINERs."
 This class of objects are often included in AGN samples selected by diagnostic emission line ratios. although their emission lines may not be dominated by accretion onto a SMBH?).," This class of objects are often included in AGN samples selected by diagnostic emission line ratios, although their emission lines may not be dominated by accretion onto a SMBH."
. The study of AGN hosts as a function of time would benetit from a homogeneous selection of active SMBHs at all redshifts., The study of AGN hosts as a function of time would benefit from a homogeneous selection of active SMBHs at all redshifts.
 A step in this direction is the extended Chandra Tultiwavelength Project?).. which detects serendipitous X-ray sources on archival Chandra images.," A step in this direction is the extended Chandra Multiwavelength Project, which detects serendipitous X-ray sources on archival Chandra images."
 That survey covers an area of 32deg? and therefore probes a large enough volume o detect in a systematic way X-ray AGN at low redshift(2).," That survey covers an area of $\rm
32\,deg^2$ and therefore probes a large enough volume to detect in a systematic way X-ray AGN at low redshift."
. lore recently combined archival XMM observations with the SDSS (XMM/SDSS survey} to compile a sample of X-ray selected AGN at low redshift. +0.1. over an area of 120deg?. much arger than the extended ChaMP.," More recently combined archival XMM observations with the SDSS (XMM/SDSS survey) to compile a sample of X-ray selected AGN at low redshift, $z\approx0.1$, over an area of $\rm 120\,deg^2$, much larger than the extended ChaMP."
 The advantage of serendipitous X-ray surveys like the XMM/SDSS and the ChaMP is that the AGN selection function at low redshift is almost identical to that of X-ray AGN at 2c| and beyond in deep surveys. thereby facilitating the comparison by minimising differential selection effects.," The advantage of serendipitous X-ray surveys like the XMM/SDSS and the ChaMP is that the AGN selection function at low redshift is almost identical to that of X-ray AGN at $z\approx1$ and beyond in deep surveys, thereby facilitating the comparison by minimising differential selection effects."
 demonstrated the power ofthis approach., demonstrated the power of this approach.
 By comparing X-ray AGN in the XMM/SDSS at +z0.1 with those in the AEGIS at 5zc1. they found little evolution of the fraction of AGN in red/blue hosts., By comparing X-ray AGN in the XMM/SDSS at $z\approx0.1$ with those in the AEGIS at $z\approx1$ they found little evolution of the fraction of AGN in red/blue hosts.
 They argued that this is evidence against a change of the dominant SMBH fueling mechanism with redshift., They argued that this is evidence against a change of the dominant SMBH fueling mechanism with redshift.
 In this paper we expand on the results and provide new estimates of the volume density evolution of X-ray AGN in red and blue hosts by combining three different X-ray surveys with QB δα d," In this paper we expand on the results and provide new estimates of the volume density evolution of X-ray AGN in red and blue hosts by combining three different X-ray surveys with median redshifts $z=0.1$, 0.3 and 0.8."
eyermine how the ολ 3xease AGN in the overall galaxy population changes with redshift., We also determine how the fraction of X-ray AGN in the overall galaxy population changes with redshift.
 These results are discussed in the contextof semi-analytic models for the growth of SMBHs and the evolution of the AGN population., These results are discussed in the contextof semi-analytic models for the growth of SMBHs and the evolution of the AGN population.
" Throughout this paper we adopt Ly=]üOkmsMpe +. Ox,=O.3 and Q,=(0.7."," Throughout this paper we adopt $\rm H_{0} = 100 \, km \,
s^{-1} \, Mpc^{-1}$ , $\rm \Omega_{M} = 0.3$ and $\rm \Omega_{\Lambda}
= 0.7$."
" Rest frame quantities (e.g. absolute magnitudes. luminosities. stellar masses) are parametrised by /—//,4/100."," Rest frame quantities (e.g. absolute magnitudes, luminosities, stellar masses) are parametrised by $h=H_{0} / 100$."
 Three different datasets are used to explore the evolution from x=(Ll to z&| of the space density of AGN in blue and red cloud hosts., Three different datasets are used to explore the evolution from $z\approx0.1$ to $z\approx1$ of the space density of AGN in blue and red cloud hosts.
 The XMM/SDSS serendipitous X-ray survey is used to select AGN at z50.1., The XMM/SDSS serendipitous X-ray survey is used to select AGN at $z\approx0.1$.
 The Needles in the Haystack Survey (NHS: Georgantopoulos 2005) is employed to study X-ray AGN at a median redshift of +=0.3., The Needles in the Haystack Survey (NHS; Georgantopoulos 2005) is employed to study X-ray AGN at a median redshift of $z\approx0.3$.
 At 2=0.5. the AEGIS(?).. the GOODS-North and the GOODS-South tields are combined to explore the properties of the X-ray AGN population.," At $z\approx0.8$, the AEGIS, the GOODS-North and the GOODS-South fields are combined to explore the properties of the X-ray AGN population."
 The number of X-ray AGN with Ly(2./—10keV)>ottCLES+ in each of the above three samples is presented in Table I..," The number of X-ray AGN with $L_X (\rm 2 - 10 \, keV)
>10^{41}\, erg \, s^{-1}$ in each of the above three samples is presented in Table \ref{tab_sample}."
 The XMM/SDSS is a serendipitous XMM survey of the SDSS area and is designed to study the properties of X-ray AGN at 2cO.1., The XMM/SDSS is a serendipitous XMM survey of the SDSS area and is designed to study the properties of X-ray AGN at $z\approx0.1$.
 Details on the reduction of the XMM observations. the detection of sources. the estimation of fluxes and hardness ratios. the determination of the survey sensitivity and the identification of X-ray sources with optical counterparts are deseribed in ?..," Details on the reduction of the XMM observations, the detection of sources, the estimation of fluxes and hardness ratios, the determination of the survey sensitivity and the identification of X-ray sources with optical counterparts are described in ."
 The XMM/SDSS survey includes pointings which have clusters of galaxies as their prime targets., The XMM/SDSS survey includes pointings which have clusters of galaxies as their prime targets.
 The overdensity of, The overdensity of
A particular example is the ALENILAT wavelet define by This wavelet. (introduced. by Marr (1980)). is proportiona to the 2D Laplacian of the Gaussian function.,A particular example is the MEXHAT wavelet defined by This wavelet (introduced by Marr (1980)) is proportional to the 2D Laplacian of the Gaussian function.
 Ht has been extensively used in the literature to detect structure on a 2D image (c.g. in astrophysics to detect. point sources in à noisy background), It has been extensively used in the literature to detect structure on a 2D image (e.g. in astrophysics to detect point sources in a noisy background).
 For CAMB analyses we are interested in the extension. of these isotropic wavelets to the sphere., For CMB analyses we are interested in the extension of these isotropic wavelets to the sphere.
 Recenth. Antoine Vanderghevnst (1998). have followed. a group. theory approach το deal with this problem.," Recently, Antoine Vandergheynst (1998) have followed a group theory approach to deal with this problem."
 This extension incorporates four basic properties: a) the basic function is a compensated. filter. b) translations. ο) dilations and. d) Euchdean limit for small angles.," This extension incorporates four basic properties: a) the basic function is a compensated filter, b) translations, c) dilations and d) Euclidean limit for small angles."
 Γον conclude that. the stereographic projection on the sphere is the appropriate one to translate the mentioned properties from the plane to the sphere., They conclude that the stereographic projection on the sphere is the appropriate one to translate the mentioned properties from the plane to the sphere.
" Such a projection is defined by (CF)12(8.6) where (6.6) represent polar coordinates on. S7 and (gy(,=21an5.©) are polar coordinates in the tangent plane to the North pole (sce Figure 1)."," Such a projection is defined by $(\vec{x})\mapsto (\theta ,\phi )$ where $(\theta ,\phi )$ represent polar coordinates on $S^2$ and $(y\equiv 2\tan \frac{\theta}{2}, \phi)$ are polar coordinates in the tangent plane to the North pole (see Figure 1)."
 Therefore. the isotropic wavelet Wor:22) transforms to ]t can be proven that the new wavelet on s? incorporates the basic properties. i. e. à) it is a compensated filter (f(dedosinAW4(6:2)= 0). b) translations are defined by translations along the sphere. i. e. rotations about the center of the sphere. c) the dilations are defined. by. the stereographie projection of dilations on the plane and d) for small angles one recovers the Lucliclean limit.," Therefore, the isotropic wavelet $\Psi (x; R)$ transforms to It can be proven that the new wavelet on $S^2$ incorporates the basic properties, i. e. a) it is a compensated filter $\int d\theta d\phi \sin \theta \Psi_{S} (\theta; R) = 0$ ), b) translations are defined by translations along the sphere, i. e. rotations about the center of the sphere, c) the dilations are defined by the stereographic projection of dilations on the plane and d) for small angles one recovers the Euclidean limit."
 0.3cm. Let us consider a function on the sphere (0.ο].," 0.3cm Let us consider a function on the sphere $f(\theta, \phi)$."
 The continuous wavelet. transform. with respect to Vs(0:2) is defined as the linear operation (0.o:I0)—worRI) ave the wavelet coelfieients dependent on 3 parameters.," The continuous wavelet transform with respect to $\Psi_S (\theta ; R)$ is defined as the linear operation $w(\theta ,\phi ; R)\equiv \tilde{w}(\vec{x}, R)$ are the wavelet coefficients dependent on $3$ parameters."
 O.3em AX straightforward. caleulation based. on the equation (5) leads. after stereographic projection. to the continuous reconstruction formula: where CF.2)=w(o.o: RH).," 0.3cm A straightforward calculation based on the equation (5) leads, after stereographic projection, to the continuous reconstruction formula: where $\tilde{w}(\vec{x}, R)\equiv w(\theta ,\phi ; R)$ ."
 0.3cm. A particular example is the MIZNILNT wavelet. defined by (see Figure 2)) We remark that the normalization constant has been chosen such that {αράώκαθ”(0:I)= 1., 0.3cm A particular example is the MEXHAT wavelet defined by (see Figure \ref{MexHatSphere}) ) We remark that the normalization constant has been chosen such that $\int d\theta d\phi \sin \theta \Psi^2 (\theta; R) = 1$ .
 This is the wavelet we are going to use in this paper to analize Gaussianity associated to dillerent. models., This is the wavelet we are going to use in this paper to analize non-Gaussianity associated to different models.
 We comment that the stereographic projection of the MENXILNT. wavelet has been recently.usedto analize maps, We comment that the stereographic projection of the MEXHAT wavelet has been recentlyusedto analize maps
hat racial velocity measurements have a limit of 0.5 m/s. which is about the highest accuracy that has been achieved fron he ground. and may be at the limit imposed. by stellar variability (2)...,"that radial velocity measurements have a limit of 0.5 m/s, which is about the highest accuracy that has been achieved from the ground, and may be at the limit imposed by stellar variability \citep{but04}."
 We assume that astrometric measurements have an accuracy of 1 praresecond which is the accuracy. which is projected to be achieved by the (?7)..," We assume that astrometric measurements have an accuracy of 1 $\mu$ arcsecond which is the accuracy which is projected to be achieved by the \citep{for03,soz03}."
 Finally. we assume that the transit timing can be measured to an accuracy of LO seconds. which is the highest accuracy of ransit timing measurements of LID 200458 (?)..," Finally, we assume that the transit timing can be measured to an accuracy of 10 seconds, which is the highest accuracy of transit timing measurements of HD 209458 \citep{bro01}."
 We concentrate on LID 209458 since it is the best studied transiting planet., We concentrate on HD 209458 since it is the best studied transiting planet.
 Fhis svstem is at a clistance of 46 pe and has a »eriod. of 3.5 days., This system is at a distance of 46 pc and has a period of 3.5 days.
 Figure 7 shows a comparison of the 10-7 sensitivity of these three techniques., Figure \ref{fig7} shows a comparison of the $\sigma$ sensitivity of these three techniques.
 Phe solid curve is computed ⋅OL Πανκ=6.7.⋅−LO1ΤΑ and mp=LOSAL. and both planets on circular orbits., The solid curve is computed for $m_{trans}=6.7\times 10^{-4} {\rm M_\odot}$ and $m_{pert}=10^{-7} {\rm M_\odot}$ and both planets on circular orbits.
 When not trapped in resonance. the amplitude of the timing veulations scales as mad/mo. so we scale the results to the mass of the perturber to compute where he timine variations are 00 seconds — this determines the sensitivity.," When not trapped in resonance, the amplitude of the timing variations scales as $m_{pert}/m_0$, so we scale the results to the mass of the perturber to compute where the timing variations are 100 seconds – this determines the sensitivity."
 For exact resonance we plot equation (33))., For exact resonance we plot equation \ref{reemeqn}) ).
 “Phe T'TV echnique is more sensitive than the astrometric technique at semi-major axis ratios smaller than about 2., The TTV technique is more sensitive than the astrometric technique at semi-major axis ratios smaller than about 2.
 OL resonance. racial velocity. nieasurenm«nts are the technique of choice for this system. while on resonance the EV is sensitive to much smaller planet masses.," Off resonance, radial velocity measurements are the technique of choice for this system, while on resonance the TTV is sensitive to much smaller planet masses."
 No ethat in Figure 7 the T'EV and astrometric techniques have the same slope at small P544Dos., Note that in Figure \ref{fig7} the TTV and astrometric techniques have the same slope at small $P_{pert}/P_{trans}$.
 ‘This is because the transi timing technique is measuring the reflex motion of the host star due to the inner planet. which is also being measured by astrometry.," This is because the transit timing technique is measuring the reflex motion of the host star due to the inner planet, which is also being measured by astrometry."
 The solid curve is an to the minimum mass detectable in ED. 209458 since a non-zero eccentricity will ead to larger timing variations (Figure 5)) and thus a smaller detectable mass., The solid curve is an to the minimum mass detectable in HD 209458 since a non-zero eccentricity will lead to larger timing variations (Figure \ref{fig5}) ) and thus a smaller detectable mass.
 1n the case that. two planets are discovered to transit their host star. measurement of the transit timing variations can break the degeneracy betw‘con mass and radius needed to derive the physical parameters of the planetary system.," In the case that two planets are discovered to transit their host star, measurement of the transit timing variations can break the degeneracy between mass and radius needed to derive the physical parameters of the planetary system."
 This has been, This has been
10 VLBA stations participated in the experiment. but due to technical problems. data from litt Peak and Los Alamos were not useful,"10 VLBA stations participated in the experiment, but due to technical problems, data from Kitt Peak and Los Alamos were not useful."
 Pie Town had a clock error which was corrected halfway through the experiment., Pie Town had a clock error which was corrected halfway through the experiment.
 The compact source BOT124-103. was observed for calibration. purposes., The compact source B0742+103 was observed for calibration purposes.
 The data were reduced using standard routines in the AIPS data reduction package., The data were reduced using standard routines in the AIPS data reduction package.
 The VLBA system lias a poor respouse at 1163 MHz. so the observations of the z=0.2212 system previded little useful information.," The VLBA system has a poor response at 1163 MHz, so the observations of the $z=0.2212$ system provided little useful information."
 At 1302 MHz the QSO center is lightly resolved into a core anc an extended: component. as shown in Figure 3. and all of the “core” flux from unresolved. observatious is recoverecl.," At 1302 MHz the QSO center is lightly resolved into a core and an extended component, as shown in Figure 3, and all of the “core” flux from unresolved observations is recovered."
 The symunetric. weak. extended lobes. which lie at roughly 30 to the north and south of the cqUasar (Murphy. Brown. Perley 1993). have been resolved away.," The symmetric, weak, extended lobes, which lie at roughly $''$ to the north and south of the quasar (Murphy, Brown, Perley 1993), have been resolved away."
 Spectra from three locatious across he QSO have been extracted., Spectra from three locations across the QSO have been extracted.
 The upper spectrum is offset by oue syuthesized beam (FWHP) to he north: the lower offset is one beam to the south. aud two beans to the east. correspeicliug Oa position along the extension of the QSO.," The upper spectrum is offset by one synthesized beam (FWHP) to the north; the lower offset is one beam to the south, and two beams to the east, corresponding to a position along the extension of the QSO."
 Although nearly of the continuum level in the ipper spectrum is due to contaminating signal from the center sightliue. the lower positiou is uearly completely independent and is uulikely to be contaminated at a detectable level.," Although nearly of the continuum level in the upper spectrum is due to contaminating signal from the center sightline, the lower position is nearly completely independent and is unlikely to be contaminated at a detectable level."
 The total sluft iu j»ositiou on the sky between the center aud lowermost siglitliue is 13.2 mas. which corresponds to 20 hos ο at the redshift of the absorber.," The total shift in position on the sky between the center and lowermost sightline is 13.2 mas, which corresponds to 20 $h^{-1}_{75}$ pc at the redshift of the absorber."
 The Arecibo Radio Telescope was used to take three one-hour on source integrations of the 1665 aud 1667 MHz OH lines iu the two absorption systems ou the BOT38+313 sightline., The Arecibo Radio Telescope was used to make three one-hour on source integrations of the 1665 and 1667 MHz OH lines in the two absorption systems on the B0738+313 sightline.
 Civen tlie uarrowness of the 2icin absorption features at both redshifts (Laue et al., Given the narrowness of the 21cm absorption features at both redshifts (Lane et al.
 1998: Chenealur Ixauekar 1909). we knew that the absorbing clouds must be very cold aud therefore uuelit contain detectable molecular gas.," 1998; Chengalur Kanekar 1999), we knew that the absorbing clouds must be very cold and therefore might contain detectable molecular gas."
 The observatious were carried out by Ix. O'Neill on 1998 November 30 aud 1908 December 1&33., The observations were carried out by K. O'Neill on 1998 November 30 and 1998 December 3.
 Using the post-upgrade system. the L-Baud wide receiver. aid the lew autocorrelation spectrometer. it was possible to observe both OH lines at the redsifts of thi DLA systems simultaneously by placing | bandpass subcorrelators at appropriate frequencies.," Using the post-upgrade system, the L-Band wide receiver, and the new autocorrelation spectrometer, it was possible to observe both OH lines at the redshifts of both DLA systems simultaneously by placing 4 bandpass subcorrelators at appropriate frequencies."
 A bandwidth of 3.125 MHz aud 1021 channels gave a velocity resolution of 0.7 km s! lu two j»olarizations., A bandwidth of 3.125 MHz and 1024 channels gave a velocity resolution of 0.7 km $^{-1}$ in two polarizations.
 The observations were made iu an ou/olf “total power” mode., The observations were made in an on/off “total power” mode.
 Total power spectra or the calibration source 3C236 were processed to form gain templates for the spectral passπε and these templates were applied to the BO7384+313 dillereuce spectra using Analyz softwae.," Total power spectra for the calibration source 3C236 were processed to form gain templates for the spectral passbands, and these templates were applied to the B0738+313 difference spectra using Analyz software."
" The ancl placed arouud the OH 1665 MHz liue at 2,=0.0912 was corrupted by interference. au could 101 be used."," The band placed around the OH 1665 MHz line at $z_1=0.0912$ was corrupted by interference, and could not be used."
 No OH was detected in the other three bauds (OH 1667 MHz at 24=0.0912 aud both OH 1667 aud 1665 MHz at to= 0.2212). which were Hanning smoothed to a velocity resoion of $e—Ld dns ! and had a 3o detection limit of 7=0.0012.," No OH was detected in the other three bands (OH 1667 MHz at $z_1=0.0912$ and both OH 1667 and 1665 MHz at $z_2=0.2212$ ), which were Hanning smoothed to a velocity resolution of $\delta v = 1.4$ km $^{-1}$ and had a $\sigma$ detection limit of $\tau = 0.0012$."
 On 1999 April 21-25. we observed the 2lem liue in the zy=0.0912 system for 15 mi[ustes on source using the ou/olI observation technique. in two polarizations.," On 1999 April 24-25, we observed the 21cm line in the $z_1=0.0912$ system for 45 minutes on source using the on/off observation technique, in two polarizations."
 Placing 2018 channels across, Placing 2048 channels across
The posterior distribution contains all the information available for inference. and different best estimates for the power-law index max be extracted from this distribution.,"The posterior distribution contains all the information available for inference, and different best estimates for the power-law index may be extracted from this distribution."
 The most probable value of> is (he maximum of P(5D). which is called the modal estimate: This corresponds to (he maximun likelihood estimate of > given by Crawlord. Jauncey. Murdoch (1970). ancl Bai (1993).," The most probable value of$\gamma$ is the maximum of $P(\gamma|D)$, which is called the modal estimate: This corresponds to the maximum likelihood estimate of $\gamma$ given by Crawford, Jauncey, Murdoch (1970), and Bai (1993)."
 A simple corresponding estimate of the uncertainty in the most likely value of 5 (for a uniform prior) was given bv Wheatland (2004) using the assumption of Gaussian behavior of P(5|D) in the vicinity of the peak SSivia Skilling 2006)., A simple corresponding estimate of the uncertainty in the most likely value of $\gamma$ (for a uniform prior) was given by Wheatland (2004) using the assumption of Gaussian behavior of $P(\gamma|D )$ in the vicinity of the peak Sivia Skilling 2006).
 In that case the width of the distribution is roughly eus=[L'(5444)]F7. where £(5)=—1n|P(35|D)] and where the prime denotes a derivative.," In that case the width of the distribution is roughly $\sigma_{\rm mod}=[L^{\prime\prime}(\gamma_{\rm mod})]^{-1/2}$, where $L(\gamma)=-\ln \left[ P(\gamma|D )\right]$ and where the prime denotes a derivative."
 This leads (0 Ina., This leads to $\sigma_{\rm mod}=M^{1/2}/\ln\pi$ .
 Using equation (A10)) then gives An alternative (and more general) approach to estimation is provided by moments of the posterior distribution., Using equation \ref{eq:gam_ML}) ) then gives An alternative (and more general) approach to estimation is provided by moments of the posterior distribution.
" Specilicallv. if the A""! moment is defined by then a best estimate and a corresponding uncertainty are provided by and The model probability distribution is with 5> 54. where σ is the upper rollover. ancl where (he normalisation constant is"," Specifically, if the $k^{\rm th}$ moment is defined by then a best estimate and a corresponding uncertainty are provided by and The model probability distribution is with $S\geq S_1$ , where $\sigma $ is the upper rollover, and where the normalisation constant is"
Our paradigm. of. galaxy formation. is. built. upon the foundation that all galaxies form within halos of dark matter (White Rees 1978: Blumenthal et al.,Our paradigm of galaxy formation is built upon the foundation that all galaxies form within halos of dark matter (White Rees 1978; Blumenthal et al.
 1984)., 1984).
 The evolution and growth of galaxies is governed by the merging of their dark matter halos and the conversion. of⋅ gas into. stars., The evolution and growth of galaxies is governed by the merging of their dark matter halos and the conversion of gas into stars.
 sThe massass of a halohi is: therefore ia fundamentalial parameterpare thati ⋠⋯∐⋯⊔⊓∖≀∙⋝≀↓≀∖∖⋰∖∐↓∙⋯⋅∙⋅⋡⋜⊓, The mass of a halo is therefore a fundamental parameter that influences a galaxy's form.
"⋜⋜⇀∙⊳⋟⊳ While] the, physicsVSICS of© haloalo Browerowth in a coldCO dardark matter atter(CDM) universe ≼can be explored ""with 'computer: (e.g..v. Springel al.", While the physics of halo growth in a cold dark matter (CDM) universe can be explored with computer simulations (e.g. Springel al.
. 2005). .actually measuring thetechniques : darkclark mass contentealaxies of galaxies is problematic.," 2005), actually measuring the total dark mass content of galaxies is problematic."
atic. RecentlyV. only . : . ⋏∙≟↓⋅⋖⋅⋜↧↿⊳∖↿↓⋅⊔⇂⋖⋅⊳∖⋯∖⇁⋖⋅∣⋡⋖⊾⋖⋅⊔⊔↓⋯⇂⋖⋅↓⊔⇂↓⊔⊳∖⋜⋯⊾⋜↧∣⋡∙∖⇁⇂∐↓↓↓⊳∖↓⊔⋏∙≟↿↓↥⋖⊾↔⊀ echnique. of⋅ weak gravitationalMN lensing. to infer∢⋅ halo masses ab very [large radi (CEvson et al.," Recently, great strides have been made in this area by utilising the technique of weak gravitational lensing to infer halo masses at very large radii (Tyson et al."
 1984: Llockstra ct al., 1984; Hoekstra et al.
 2005: Alancelbaum et al., 2005; Mandelbaum et al.
 2006: Parker et al., 2006; Parker et al.
 2007)., 2007).
 This echnique statistically combines the lensing signal from a arge number of⋅ halos and. can vicld. general. relationships. jvetween galaxy stellar and total halo masses (MHoekstra et al., This technique statistically combines the lensing signal from a large number of halos and can yield general relationships between galaxy stellar and total halo masses (Hoekstra et al.
 2005: Mancdelbauni et al., 2005; Mandelbaum et al.
 2006)., 2006).
 Such a relation also follows rom analysis of the cconditional luminosity function. which is an observationally-basec model. connecting the stellar mass of galaxies to their hosting halos (Yang. Mo van den Bosch 2003. 2008: van den Bosch et al.," Such a relation also follows from analysis of the “conditional luminosity function”, which is an observationally-based model connecting the stellar mass of galaxies to their hosting halos (Yang, Mo van den Bosch 2003, 2008; van den Bosch et al."
 2007)., 2007).
" Incdiviclual galaxies are expected. to scatter. about. these ""ecneral relations.", Individual galaxies are expected to scatter about these general relations.
 A number of other techniques. exist. for measuring the halo mass of@ndividual galaxies: rotation curves in disk DDe&alaxies (SofueD DENRubin 2001).u-— X-rav NORgas in MNgiant elliptical galaxies COSullivan Ponman 2004). kinematics ⋅ ∪⇂⋏∙≟⇂∩∣⋡⊔↓⋜⊔⋅≼⇍⇂⊔⊳∖∩⋅↓⋅⊳∖⋜⋃∐⇂↓≻↓⋜⋯∢⋅↿⋜⊔⋅∙∖⇁⊔∢⊾∣⋡⊔↓⋯⊾↿∖∐∪⊔↓⋜⋯∪∖∖⋱∖↳∖⇁⋖⋅↿ al.," A number of other techniques exist for measuring the halo mass of galaxies: rotation curves in disk galaxies (Sofue Rubin 2001), X-ray gas in giant elliptical galaxies (O'Sullivan Ponman 2004), kinematics of globular clusters and planetary nebulae (Romanowsky et al."
 ∙⋅2003: Romanowsky et al., 2003; Romanowsky et al.
 ⋅2008) and strong gravitationalDN lensing. (Ferreras.. Prasenjite Williamsna 2005).," 2008) and strong gravitational lensing (Ferreras, Prasenjit Williams 2005)."
 However. such all⋠⋠⋠⋠ have their limitations. ∙e.g. they are often simulations. applicable to. certain galaxy⊾⊲⋅ types. require equilibriumep. alotal ≼∙∪↓⊔⇂⋠⊔⋠↓∪⊔≱∖⋡≼∙⋜⋯∣⋡∢⋅∪∣⋡≱∖∢⋅↓⋅∖⇁⋜↧↿⊲↓∪⊔⋜↧↓∙∖⇁∖⇁⋖⊾↓⋅∙∖⇁∢⊾⇀∖↓≻⋖⋅⊔⊳∖⊲↓∖⇁∢⊾⊳↓≻↓⋅∪∣⋊⋅ D. . ∪⊔↓∙∖⇁⋜↧↓⊔∐↿⋯⇂↓⋅⋯⊔⋜↧⇂∢⋅⇀∖∩⋅⊔∣⋜⋯∠⇂⊳↓⊔↿↓∐⋅≼⇍⋜↧⊳∖∢⊾∪⇂⊳∖⇂↓⋅∪⊔⋏∙≟↓⋖⋅⊔⊳∖↓⊔⋏∙≟⊳ . . is only available for galaxies along random lines of sight.," However, such techniques all have their limitations, e.g. they are often only applicable to certain galaxy types, require equilibrium conditions, can be observationally very expensive, probe only a limited radial extent and, in the case of strong lensing, is only available for galaxies along random lines of sight."
 As a result. only a very small number of individual galaxies have direct measures of their total halo mass available.," As a result, only a very small number of individual galaxies have direct measures of their total halo mass available."
" In this1 avewe present""QR aκ newTU method. forDr estimatingestimatine the total halo mass surrounding individual galaxies based on the total mass contained in their elobular cluster (6€) svstcHis.", In this we present a new method for estimating the total halo mass surrounding individual galaxies based on the total mass contained in their globular cluster (GC) systems.
 All large galaxies. including the Milky Way. are known to host a system of GCs.," All large galaxies, including the Milky Way, are known to host a system of GCs."
 “Phey are among the oldest (>11 Ον) stellar populations known (Stracer et al., They are among the oldest $>11$ Gyrs) stellar populations known (Strader et al.
 2005: Puzia et al., 2005; Puzia et al.
 2005: Proctor et al., 2005; Proctor et al.
 2007) and therefore trace the first stages of ealaxy formation (West et al., 2007) and therefore trace the first stages of galaxy formation (West et al.
 2004: Broclic, 2004; Brodie
while only 254. was assembled (ο halos in this mass range.,while only $2\%$ was assembled to halos in this mass range.
 Even αἱ z~10. more than 1%. of (he cosmic mass was already in sheets more massive than LOMAL... whereas the mass fraction in halos in the same mass range was ten times smaller.," Even at $z\sim 10$, more than $1\%$ of the cosmic mass was already in sheets more massive than $10^{10}M_\odot$, whereas the mass fraction in halos in the same mass range was ten times smaller."
 This suggests that the structure at highoO = is dominated by sheets., This suggests that the structure at high $z$ is dominated by sheets.
" Al any given redshift. we may define a characteristic mass. M.(z). so thato(M,) is equal to the critical overdensitv 9,2)."," At any given redshift, we may define a characteristic mass, $M_*(z)$, so that$\sigma(M_*)$ is equal to the critical overdensity $\delta_{sc}(z)$ ."
" In our ACDM model. the value of AL, drops rapidly with redshift. from just over 10734, at 2=0 (o ~LOMOAZ, at 2=2."," In our $\Lambda$ CDM model, the value of $M_*$ drops rapidly with redshift, from just over $10^{13}M_{\odot}$ at $z=0$ to $\sim 10^{10.5}M_{\odot}$ at $z=2$."
" Objects at 2 whose mass satisfies a(.M)=Ας)NoCM,). where N=1.2.:--. can be thought of as being increasingly unusual compared to objects of the characteristic mass AL,(2)."," Objects at $z$ whose mass satisfies $\sigma(M)=N\delta_{sc}(z)=N\sigma(M_*)$, where $N=1,2,\cdot\cdot\cdot$, can be thought of as being increasingly unusual compared to objects of the characteristic mass $M_*(z)$."
" The top two solid curves in (he first three panels of Figure 3. show the abundance-redshift locii of sheets. filaments and halos whose masses satisfy 0(.M)=o(M,) and o(VM)=26(M,)."," The top two solid curves in the first three panels of Figure \ref{numfracall} show the abundance-redshift locii of sheets, filaments and halos whose masses satisfy $\sigma(M)=\sigma(M_*)$ and $\sigma(M)=2\sigma(M_*)$."
 The panels show that the mass [raction in lo sheets. 0.59. is larger than that in lo filaments (0.32) and than that in lo halos (0.17).," The panels show that the mass fraction in $1\sigma$ sheets, $0.59$, is larger than that in $1\sigma$ filaments (0.32) and than that in $1\sigma$ halos (0.17)."
 A similar scaling applies to the more massive 26 objects. implving that. at anv redshift. the (vpical structure is dominated by sheets.," A similar scaling applies to the more massive $2\sigma$ objects, implying that, at any redshift, the typical structure is dominated by sheets."
 The bottom right panel in Figure 3. compares the umber densiGes of sheets (dotted). filaments (dashed) aud halos (solid) at different redshilts in the three other panels directly (i.e.. Chese were the dashed curves in the other three panels): top-niost curves show our moclel [or masses of LOMAL... bottom-most for 10M... and the curves in between show results in which the mass changes by 1 dex.," The bottom right panel in Figure \ref{numfracall} compares the number densities of sheets (dotted), filaments (dashed) and halos (solid) at different redshifts in the three other panels directly (i.e., these were the dashed curves in the other three panels): top-most curves show our model for masses of $10^{10} M_\odot$, bottom-most for $10^{15} M_\odot$, and the curves in between show results in which the mass changes by 1 dex."
 These curves show that. ab masses larger than LOMAL.. {here are more sheets than filaments. and more filaments (han halos. for almost all redshifts.," These curves show that, at masses larger than $10^{12} M_{\odot}$, there are more sheets than filaments, and more filaments than halos, for almost all redshifts."
 llowever. at small masses and late (times this trencl begins to reverse: al 2=0 there are fewer sheets than filaments with masses smaller than LOMAL.. and there are more halos than sheets of mass LOMA...," However, at small masses and late times this trend begins to reverse: at $z=0$ there are fewer sheets than filaments with masses smaller than $10^{11} M_{\odot}$, and there are more halos than sheets of mass $10^{10} M_{\odot}$."
 Based on these mass functions. we can estimate the average mass ol objects more massive (han some given minimum mass.," Based on these mass functions, we can estimate the average mass of objects more massive than some given minimum mass."
 For instance. when the minimum mass is LOMA. then. at the present time. this average mass is 3.9x10!5M. [or sheets. 21x10!M. for filaments. and 1.6xLOMAL. for halos.," For instance, when the minimum mass is $10^{10} M_{\odot}$ then, at the present time, this average mass is $3.9 \times 10^{11} M_{\odot}$ for sheets, $2.1 \times 10^{11} M_{\odot}$ for filaments, and $1.6 \times 10^{11} M_{\odot}$ for halos."
 The average mass of sheets is about 2.5 limes that of halos. again suggesting that large scale structure is dominated by sheets rather than by virialized halos.," The average mass of sheets is about 2.5 times that of halos, again suggesting that large scale structure is dominated by sheets rather than by virialized halos."
 Talo abundances are expected (to correlate. with the overclensity of their surroundings halos populate dense regions (Mo&White 1996).., Halo abundances are expected to correlate with the overdensity of their surroundings---massive halos populate dense regions \citep{MW96}. .
 However. ourmodel predicts that halo abundances will also correlate. with the morphology of their surroundings.," However, ourmodel predicts that halo abundances will also correlate with the morphology of their surroundings."
 To illustrate. the solid curve in Figure 4. shows the mass fraction of LOMAL. sheets that is in," To illustrate, the solid curve in Figure \ref{cosmicweb} shows the mass fraction of $10^{13} M_{\odot}$ sheets that is in"
"is ~9 4400, 83300, 68800, 37700 years for models A-D, respectively, and the characteristic time scale for the shocked wind, Rso/Vw, which is approximately 2200 years.","is $\sim$ 400, 300, 800, 700 years for models A-D, respectively, and the characteristic time scale for the shocked wind, $R_{\rm SO} / v_{\rm w}$, which is approximately 200 years."
 In Fig., In Fig.
" 10 we compare the radiative cooling and characteristic dynamical time scales of the gas: =,. and = Κον», respectively."," \ref{fig: tau} we compare the radiative cooling and characteristic dynamical time scales of the gas: =, and = / v, respectively."
" For regions with rA<<Tdyn, e.g. the bow shock arc where the gas is strongly decelerated, the gas cools strongly and behaves almost isothermally."," For regions with $\tau_\Lambda << \tau_{\rm dyn}$, e.g. the bow shock arc where the gas is strongly decelerated, the gas cools strongly and behaves almost isothermally."
" The gas in the tail has a much longer cooling time scale, 4300000 years (much greater than the cooling time scale for the shocked ISM) resulting in a largely adiabatic flow there."," The gas in the tail has a much longer cooling time scale, $\sim$ 000 years (much greater than the cooling time scale for the shocked ISM) resulting in a largely adiabatic flow there."
" The shape of the bow shock changes with time: it is initially circular and becomes increasingly *parabolic' until it reaches a steady state, at which point the global morphology is described by Eq. 3,"," The shape of the bow shock changes with time: it is initially circular and becomes increasingly `parabolic' until it reaches a steady state, at which point the global morphology is described by Eq. \ref{eq: shape},"
 the analytic ? solution., the analytic \cite{Wil96} solution.
" The ratio of the shock position at angles 6=0? and 0=90°, r,(0°)/r;(90°), as a function of time, is shown in Fig."," The ratio of the shock position at angles $\theta=0^\circ$ and $\theta=90^\circ$, $r_{\rm s}(0^\circ)/r_{\rm s}(90^\circ)$, as a function of time, is shown in Fig."
 11 [top]., \ref{fig: shape} [top].
 It takes several dynamical time scales for the bow shock to achieve the steady state value R(0°)/R(90°)=1/V3., It takes several dynamical time scales for the bow shock to achieve the steady state value $R(0^\circ)/R(90^\circ)$ $1/\sqrt{3}$.
" The wind takes at least {=R(0)/w, to reach the thin-shell position, ignoring the drag due to the moving environment, thus it takes 160000 years for 0=0° (R(0*)-Rso), while for 6=90° (R(90°) = V3Rso) it takes ~28 0000 years."," The wind takes at least $t = R(\theta)/v_{\rm w}$ to reach the thin-shell position, ignoring the drag due to the moving environment, thus it takes $\sim$ 000 years for $\theta=0^\circ$ $R(0^\circ)$ $R_{\rm SO}$ ), while for $\theta=90^\circ$ $R(90^\circ)$ = $\sqrt{3} R_{\rm SO}$ ) it takes $\sim$ 000 years."
" In reality, the above time scales are longer."," In reality, the above time scales are longer."
" The drag is maximal for 6=0°, so it takes about 300000 years to reach Rso (e.g. Fig. 9)),"," The drag is maximal for $\theta=0^\circ$, so it takes about 000 years to reach $R_{\rm SO}$ (e.g. Fig. \ref{fig: shockpos}) ),"
 and coincidentally it takes about the same time to reach the solution position for, and coincidentally it takes about the same time to reach the solution position for
"For MDI data these correlations are 0.36 and 0.40 respectively,",For MDI data these correlations are 0.36 and 0.40 respectively.
 The correlation with other parameters of the tachocline are found to be very. small., The correlation with other parameters of the tachocline are found to be very small.
 GONG and MDI give somewhat disparate results about the differences in the tachocline between the minimum before evele 23 and that before evele 24., GONG and MDI give somewhat disparate results about the differences in the tachocline between the minimum before cycle 23 and that before cycle 24.
 This could be a result of the fact that MIDI data do not completely span the cycle 23 minimum., This could be a result of the fact that MDI data do not completely span the cycle 23 minimum.
 The only significant change is in the parameter 0€2;., The only significant change is in the parameter $\delta\Omega_5$.
 Even (hat is of opposite sign between GONG ancl MIDI., Even that is of opposite sign between GONG and MDI.
 Thus we cannot sav lor certain whether the tachocline parameters were different between the two minima., Thus we cannot say for certain whether the tachocline parameters were different between the two minima.
 We take a closer look at the changes in tachocline parameters by examining a lew latitudes in more detail., We take a closer look at the changes in tachocline parameters by examining a few latitudes in more detail.
 For each latitude we have determined the correlation coefficient between r;. 0 and w and the 10.7 em radio flux.," For each latitude we have determined the correlation coefficient between $r_t$ , $\delta\Omega$ and $w$ and the 10.7 cm radio flux."
 However. since the 10.7 em flux is an indicator of global activity. while the changes in the tachocline may be correlated with other local changes. which depend on Iatitude. we also fit an oscillatory Form will a period of solar cvcle.," However, since the 10.7 cm flux is an indicator of global activity, while the changes in the tachocline may be correlated with other local changes, which depend on latitude, we also fit an oscillatory form with a period of solar cycle."
 Thus for 0€ we fit where wy is the frequency corresponding to 11.7 vears. the dynamical length of the solar evele as determined by Antia Basu (2010). and the angular brackets denote average over time.," Thus for $\delta\Omega$ we fit where $\omega_0$ is the frequency corresponding to 11.7 years, the dynamical length of the solar cycle as determined by Antia Basu (2010), and the angular brackets denote average over time."
 This form was motivated by the fact that the change of zonal flow velocities al a given ]atitude ancl radius can be fitted wilh thisexpression., This form was motivated by the fact that the change of zonal flow velocities at a given latitude and radius can be fitted with thisexpression.
 We fit similar expressions to r; ancl d., We fit similar expressions to $r_t$ and $w$.
 In Figure | we show (he position of the tachocline. r; at four latitudes plotted as a [unetion of time.," In Figure \ref{fig:rd} we show the position of the tachocline, $r_t$ at four latitudes plotted as a function of time."
 As can be seen. there is a Large spread in the results and there is no clear temporal variation.," As can be seen, there is a large spread in the results and there is no clear temporal variation."
 In the same figure we also show a running mean ol (he results to reduce the scatter., In the same figure we also show a running mean of the results to reduce the scatter.
 The running mean is taken over a period of about 1 vear. i.e.. over 9 sets [or GONG and 5 sets for ALDI.," The running mean is taken over a period of about 1 year, i.e., over 9 sets for GONG and 5 sets for MDI."
 Also shown is the fit to the form given by Eq. (8))., Also shown is the fit to the form given by Eq. \ref{eq:osc}) ).
 There is a good agreement between the results obtained using GONG and MDI data at low latitudes. but there are significant differences at high latitudes as also seen in the time-averaged data.," There is a good agreement between the results obtained using GONG and MDI data at low latitudes, but there are significant differences at high latitudes as also seen in the time-averaged data."
 The r; values are not correlated with the 10.7 «m flux. the correlation coefficients lie between —0.14 and 0.02 for the different latitudes ancl sets of data.," The $r_t$ values are not correlated with the 10.7 cm flux, the correlation coefficients lie between $-0.14$ and $0.02$ for the different latitudes and sets of data."
 There appears to be an oscillatory time variation al the lower latitudes. aud the resulis are consistent for both data sets. implving the results may have some significance.," There appears to be an oscillatory time variation at the lower latitudes, and the results are consistent for both data sets, implying the results may have some significance."
 Figure 2. shows similar results for the width., Figure \ref{fig:wid} shows similar results for the width.
 At low latitudes the width is comparable to errors in individual data points and it is difficult to sav much about the temporal variations., At low latitudes the width is comparable to errors in individual data points and it is difficult to say much about the temporal variations.
 At high latitudes also the temporal variations are unclear {here is no agreement between results obtained with the GONG and MDI data., At high latitudes also the temporal variations are unclear — there is no agreement between results obtained with the GONG and MDI data.
" Aswith r,. the correlation with the 10.7 em flux is small."," Aswith $r_t$ , the correlation with the 10.7 cm flux is small."
and well-known scaling relation. and most of the disks do follow such a relation.,"and well-known scaling relation, and most of the disks do follow such a relation."
" However. some disks in the ""SEP|SDSSDEDLSDSSBING” photometric dataset lic onu au “upper brauch? of large disks that appear to be 2-3 ines larecr at a given disk Iuninosity."," However, some disks in the “SEP+SDSSDEBL+SDSSBKG"" photometric dataset lie on an “upper branch"" of large disks that appear to be 2-3 times larger at a given disk luminosity."
 Plotting the same relation for the «-band showed au increasing bimodality in the sizes of disks., Plotting the same relation for the $u$ -band showed an increasing bimodality in the sizes of disks.
 Visual iuspectiou of these apparently aree disks showed that they were in fact caused by constant «kv backeround residuals left after subtracting he SDSS skv values., Visual inspection of these apparently large disks showed that they were in fact caused by constant sky background residuals left after subtracting the SDSS sky values.
" As explained in Section ??.. our initial procedure asstuned that the SDSS skv values were ""perfect? and that the GIAI2D fits could be performed w fixing the sky value to zero once the SDSS sky values wad been subtracted frou, the input science inages."," As explained in Section \ref{skybkg}, our initial procedure assumed that the SDSS sky values were “perfect"" and that the GIM2D fits could be performed by fixing the sky value to zero once the SDSS sky values had been subtracted from the input science images."
 These constant positive sky residuals were being fit by he algorithm with a very large (aud thus very flat) disk component., These constant positive sky residuals were being fit by the algorithm with a very large (and thus very flat) disk component.
 The false disks lic ou a linear relation jecause of the strong covariance between Iunuinositv aud size errors through the relation PF-—2aNyry for the otal flux F of au exponential disk with ceutral surface xiehtuess Xj aud scale leneth ry., The false disks lie on a linear relation because of the strong covariance between luminosity and size errors through the relation $F = 2\pi\Sigma_0 r_d$ for the total flux $F$ of an exponential disk with central surface brightness $\Sigma_0$ and scale length $r_d$.
 Recomputing local sky backerouud levels following the procedure described in Section ?? corrected. this problem as shown by the “SEP|SDSSDEBLGM2DDRG. photometric dataset in the bottom paucl of Figure 7..," Recomputing local sky background levels following the procedure described in Section \ref{skybkg} corrected this problem as shown by the “SEP+SDSSDEBL+GM2DBKG"" photometric dataset in the bottom panel of Figure \ref{disk-lumsize}."
 The «hand relation was also corrected by the same procedure., The $u$ -band relation was also corrected by the same procedure.
 The disk size-uunmositv relation is therefore a scusitive test of our «kv vackeround level deteriinations., The disk size-luminosity relation is therefore a sensitive test of our sky background level determinations.
 It is worth noting here that this sky background woblem is different from the well-known problem| with he photometry of bright galaxies in SDSS (Abazajiauctal.2009.audreferences therein).., It is worth noting here that this sky background problem is different from the well-known problem with the photometry of bright galaxies in SDSS \citep[][and references therein]{abazajian09}.
 The SDSS photometric xpeliue produces systematic errors m the estimation of he skv near bright (r« 16) galaxies that cause their fluxes ancl scale sizes to be underestimated., The SDSS photometric pipeline produces systematic errors in the estimation of the sky near bright $r < 16$ ) galaxies that cause their fluxes and scale sizes to be underestimated.
 The sky vackeround problem eucouutered here was related to yositive sky residuals., The sky background problem encountered here was related to positive sky residuals.
 These effects of these residuals )eanie especially significant here because sky level errors dominate the systematic errors in bulge|disk decompositions (Simardetal.2002)., These effects of these residuals became especially significant here because sky level errors dominate the systematic errors in bulge+disk decompositions \citep{simard02}.
. It will also be more sjeuificautf i amore crowded environmneuts (ucludiug close pairs of galaxies)., It will also be more significant in more crowded environments (including close pairs of galaxies).
 Figure 8 compares our GIM2D-based total magnitudes with the SDSS Petyosian and model magnitudes as a function of bulee fraction., Figure \ref{dr7_mags_cmp} compares our GIM2D-based total magnitudes with the SDSS Petrosian and model magnitudes as a function of bulge fraction.
 There are two different types of SDSS model magnitudes., There are two different types of SDSS model magnitudes.
" The first type (called ""mnodelMag comes from fitting the object with a pure deVaucouleurs) and a pure exponcutial profile. determining which profile fits the object better and computing the magnitude of this profile."," The first type (called “modelMag"") comes from fitting the object with a pure deVaucouleurs and a pure exponential profile, determining which profile fits the object better and computing the magnitude of this profile."
" The ""modelMag"" magnitudes are widely used throughout the literature because they are the ones explicitly even in the SDSS database."," The “modelMag"" magnitudes are widely used throughout the literature because they are the ones explicitly given in the SDSS database."
" The second type of model magnitudes (called “cumodel”) is computed from the composite fiux eiven by the linear combination of the two profiles. Los ροκο=(FraeDVFay10FracDeVM, (the value offracDer is stored iu the SDSS database)."," The second type of model magnitudes (called “cmodel"") is computed from the composite flux given by the linear combination of the two profiles, i.e., $F_{composite} = (fracDeV)F_{deV}+ (1-fracDeV)F_{exp}$ (the value of is stored in the SDSS database)."
 Although it would be more fair to compare the GIM2D maeguitudes to the vcmodel maguitudes. we decided to compare our photometry with the most widely used set of SDSS model magnitudes to cuphasize that the use of SDSS inodel magnitudes requires sole caution or galaxies of intermediate morphological types.," Although it would be more fair to compare the GIM2D magnitudes to the “cmodel"" magnitudes, we decided to compare our photometry with the most widely used set of SDSS model magnitudes to emphasize that the use of SDSS model magnitudes requires some caution for galaxies of intermediate morphological types."
 This is obvious from looking at the panels in the first colin of Figure 8.., This is obvious from looking at the panels in the first column of Figure \ref{dr7_mags_cmp}.
 These pancls do not depend ou any GIM2D photometry., These panels do not depend on any GIM2D photometry.
" Thev show a “bifurcation” at intermediate πι]οο fraction values because the ""modelMag"" model uagnitudes can only come from either a pure disk or a pure bulee model. aud some objects with a iuixture of »ilee aud disk components become ambiguous enoueh to such a binary classification that they cud up in differeut LS."," They show a “bifurcation"" at intermediate bulge fraction values because the “modelMag"" model magnitudes can only come from either a pure disk or a pure bulge model, and some objects with a mixture of bulge and disk components become ambiguous enough to such a binary classification that they end up in different bins."
 The second aud third cohununus of Figure δ show hat the offset between GIN2D and SDSS maguitudes is inclependenut of bulee fraction., The second and third columns of Figure \ref{dr7_mags_cmp} show that the offset between GIM2D and SDSS magnitudes is independent of bulge fraction.
 This lack of depeudeuce ou the type of galaxy light profiles makes it easier to compare the different sets of magnitudes., This lack of dependence on the type of galaxy light profiles makes it easier to compare the different sets of magnitudes.
 Blantonctal.(2003) reported a systematic trend between their Sérrsic- maguitude aud SDSS Petrosian magnitudes as a function of the Sérrsic index value with Petrosian magnitudes underestimating the g- aud ας faxes by 20% in their hieh (0~ 1) Séórrsic index galaxies., \citet{blanton03} reported a systematic trend between their S\'errsic-based magnitude and SDSS Petrosian magnitudes as a function of the Sérrsic index value with Petrosian magnitudes underestimating the $g$ - and $r$ -band fluxes by $\%$ in their high $n \simeq$ 4) Sérrsic index galaxies.
 The bifurcation at iutermediate bulge fractious is also visible in the comparison between GIM2D and SDSS model, The bifurcation at intermediate bulge fractions is also visible in the comparison between GIM2D and SDSS model
broad.,broad.
 The range of values of r aud j for the ERIIb radio galaxies studied is remarkably small. while the range of values of r ancl j for outflows associated with CDGs is quite broad.," The range of values of $r$ and $j$ for the FRIIb radio galaxies studied is remarkably small, while the range of values of $r$ and $j$ for outflows associated with CDGs is quite broad."
 There are (wo Κον distinctions between these samples that may partially explain these differences., There are two key distinctions between these samples that may partially explain these differences.
 Firstly. the outflow energv measured for à ΕΠΟ source is an estimate of the total outflow enerev (hat will be produced over the entire lifetime of the outflow. whereas the outflow enerev measured [ον the CDGs is an estimate of the outflow enerev produced by the source up to the time of the observation.," Firstly, the outflow energy measured for a FRIIb source is an estimate of the total outflow energy that will be produced over the entire lifetime of the outflow, whereas the outflow energy measured for the CDGs is an estimate of the outflow energy produced by the source up to the time of the observation."
 Catching each CDG outflow at a different phase in the course of its lifetime will cause the distribution of outflow energies for CDGs to be broad: (his is not a factor for the ΕΠΗ sources., Catching each CDG outflow at a different phase in the course of its lifetime will cause the distribution of outflow energies for CDGs to be broad; this is not a factor for the FRIIb sources.
 In. addition. a prior outflow event [rom a CDG could cause the outflow energy of a given event to be overestimated.," In addition, a prior outflow event from a CDG could cause the outflow energy of a given event to be overestimated."
 Secondly. the EFRIIb radio galaxies are selected from the flux limited 8CRR catalogue (Laing. Rilev. Longair 1983). so only the most powerlul sources are detected at each redshilt.," Secondly, the FRIIb radio galaxies are selected from the flux limited 3CRR catalogue (Laing, Riley, Longair 1983), so only the most powerful sources are detected at each redshift."
" Interestingly. the total outflow energv ZZ, measured for powerful 3CRB radio galaxies ΜΕ”... (O'Dea οἱ al."," Interestingly, the total outflow energy $E_*$ measured for powerful 3CRR radio galaxies increases as $(1+z)^{1.8 \pm 0.2}$ (O'Dea et al."
 2008). while the black hole mass rises as (1+2?7*5 (MeLure et al.," 2008), while the black hole mass rises as $(1+z)^{2.07 \pm 0.76}$ (McLure et al."
 2006). so the ratio i: is independent of redshift.," 2006), so the ratio $r$ is independent of redshift."
 The fact that lower mass aud lower energy black holes are not observed at high redshift is due to the fIux limit of the survey., The fact that lower mass and lower energy black holes are not observed at high redshift is due to the flux limit of the survey.
 The fact that the high mass and high energv sources FRIIb sources are not seen at low redshift may suggest that the higher mass sources are more active al high redshift. or that changes in the environment cause (he radio power of the sources to fall below the flux limit of the survey. that is. the sources may be lower power FRII sources or FRI sources. though they could still have large outflow energies (e.g. ILucdee οἱ al.," The fact that the high mass and high energy sources FRIIb sources are not seen at low redshift may suggest that the higher mass sources are more active at high redshift, or that changes in the environment cause the radio power of the sources to fall below the flux limit of the survey, that is, the sources may be lower power FRII sources or FRI sources, though they could still have large outflow energies (e.g. Hardee et al."
 1992: ODonoghue. Eilek. Owen 1993: Bicknell 1995).," 1992; O'Donoghue, Eilek, Owen 1993; Bicknell 1995)."
" A selection bias also seems to be present in the CDG sample: the 6 sources with outflow energies less than 10/62, are all at verv low redshift. ancl the outflow energy. black hole mass. and black hole spin all tend (o increase with redshift."," A selection bias also seems to be present in the CDG sample; the 6 sources with outflow energies less than $10^4 c^2 M_{\odot}$ are all at very low redshift, and the outflow energy, black hole mass, and black hole spin all tend to increase with redshift."
 Very powerful 3CRIB radio galaxies are likely to evolve into the CDGs (Lilly Longair 1984: Best et al., Very powerful 3CRR radio galaxies are likely to evolve into the CDGs (Lilly Longair 1984; Best et al.
 L998: MeLure οἱ al., 1998; McLure et al.
 2004)., 2004).
 Thus. the two samples studied are representative of the early ancl late phases of evolution of these svstems.," Thus, the two samples studied are representative of the early and late phases of evolution of these systems."
 The svstems have black holes with similar mass., The systems have black holes with similar mass.
 Clearly. (he ΕΠΗΠΙΟ sources have a very small range of outflow energy per unil black hole mass. r.," Clearly, the FRIIb sources have a very small range of outflow energy per unit black hole mass, $r$."
 It would appear that there is à larger range of values for this parameter for CDGs. but. given the [actors affecting (he measurements of outflow energy [rom black holes of CDGs. it is possible thatthese sources have a similarly small intrinsic range of + and j.," It would appear that there is a larger range of values for this parameter for CDGs, but, given the factors affecting the measurements of outflow energy from black holes of CDGs, it is possible thatthese sources have a similarly small intrinsic range of $r$ and $j$ ."
develop a description of the equation of state that has the correct critical exponents which are not integers or simple Iractions: in other words. it is not a mean field theory.,"develop a description of the equation of state that has the correct critical exponents which are not integers or simple fractions; in other words, it is not a mean field theory."
 It will also include the critical amplitudes. which are universal.," It will also include the critical amplitudes, which are universal."
 The chiral phase transition is in the same universality class as (he 3D Ising moclel and liquid gas phase Gransition., The chiral phase transition is in the same universality class as the 3D Ising model and liquid gas phase transition.
 Finally. (he parameterization will incorporate knowledge about the zero barvon density equation of state. computed by lattice QCD. as well as knowledge about the hieh densitv behavior of nuclear matter al zero temperature.," Finally, the parameterization will incorporate knowledge about the zero baryon density equation of state, computed by lattice QCD, as well as knowledge about the high density behavior of nuclear matter at zero temperature."
 Perhaps the closest work (hat addressed these issues was ref., Perhaps the closest work that addressed these issues was ref.
[20]. which blended a parameterization of the 3D Ising model equation of state into quark and hadron equations of state.,\cite{attract} which blended a parameterization of the 3D Ising model equation of state into quark and hadron equations of state.
 In that work there was an ambiguity as (o how to relate the magnetization to the barvon chemical potential., In that work there was an ambiguity as to how to relate the magnetization to the baryon chemical potential.
 In (hiis paper there is no such issue., In this paper there is no such issue.
 These {wo parameterizalions can perhaps be viewed as alternatives which provide some idea as to the range of uncertainty in how to describe matter near the chiral critical point., These two parameterizations can perhaps be viewed as alternatives which provide some idea as to the range of uncertainty in how to describe matter near the chiral critical point.
 The outline of the paper is as follows., The outline of the paper is as follows.
 In Sec., In Sec.
 2 we summarize some thermodynamic relations and definitions of the relevant critical exponents., 2 we summarize some thermodynamic relations and definitions of the relevant critical exponents.
 In Sec., In Sec.
 3 we develop a relatively general parameterization of the equation of state in the vicinity of the critical point whose motivation comes [rom mean field theory., 3 we develop a relatively general parameterization of the equation of state in the vicinity of the critical point whose motivation comes from mean field theory.
 In Sec., In Sec.
" 4 we fix the parameters to best match onto what is known about the equation of state at finite temperature bul zero barvon density, and al finite barvon densitv but zero temperature."," 4 we fix the parameters to best match onto what is known about the equation of state at finite temperature but zero baryon density, and at finite baryon density but zero temperature."
 In. Sec., In Sec.
 5 we show the numerical results obtained with the developed. parameterization and parameters., 5 we show the numerical results obtained with the developed parameterization and parameters.
 In Sec., In Sec.
 6 we apply the Landau theory of [actuations away from the stable thermodynamic phases to estimate the magnitude of density fluctuations. which turn out to be surprisingly large.," 6 we apply the Landau theory of fluctuations away from the stable thermodynamic phases to estimate the magnitude of density fluctuations, which turn out to be surprisingly large."
 Concluding comments are made in Sec., Concluding comments are made in Sec.
 7., 7.
 Those wishing to know only the results should read Sec., Those wishing to know only the results should read Sec.
 5 and the Appendix which summarizes (he parameterized equation of state., 5 and the Appendix which summarizes the parameterized equation of state.
 Consider the equation of state of matter in (he vicinity of a critical point which is in (he same universality class as (he liquid-gas phase transition and the 3D Ising model., Consider the equation of state of matter in the vicinity of a critical point which is in the same universality class as the liquid-gas phase transition and the 3D Ising model.
 It. is advantageous to work with the [Helmholtz free enerev clensity / which is a function of the temperature PT and barvon density »., It is advantageous to work with the Helmholtz free energy density $f$ which is a function of the temperature $T$ and baryon density $n$.
 Some useful thermodynamic relations involving the pressure J?.," Some useful thermodynamic relations involving the pressure $P$ ,"
map the extinction distribution. within 1 towards these windows.,map the extinction distribution within $^{\circ}$ towards these windows.
 The reddening distribution in the Bulge area is also fundamental to understand the spatial distribution of elobular clusters and their relation to the Bulge field stellar »opulation itself., The reddening distribution in the Bulge area is also fundamental to understand the spatial distribution of globular clusters and their relation to the Bulge field stellar population itself.
 Barbuy et al. (, Barbuy et al. (
1998) discussed the elobular clusters projected within 5° of the nucleus and which appear o be related to the ίσο. while Barbuy et al. (,"1998) discussed the globular clusters projected within $^{\circ}$ of the nucleus and which appear to be related to the Bulge, while Barbuy et al. ("
1999) discussed the properties of those found in the area covered w the present study.,1999) discussed the properties of those found in the area covered by the present study.
 In this work we use the 2ALASS (1.25j0m) and (2.175). photometric data to build ο) CMDs of Bulge fields in order to map out the interstellar extinction in the central 107. o£ the Galaxy., In this work we use the 2MASS $\mu$ m) and $\mu$ m) photometric data to build ) CMDs of Bulge fields in order to map out the interstellar extinction in the central $^{\circ}$ of the Galaxy.
 Phe extinction. was derived by means of the upper giant branch fitting method. similar to that described in Paper LI. In Sect.," The extinction was derived by means of the upper giant branch fitting method, similar to that described in Paper I. In Sect."
 2 we revise this method and. describe in detail the way we build. the extinction map., 2 we revise this method and describe in detail the way we build the extinction map.
 In Sect., In Sect.
 3 we compare the present results with those from DENIS., 3 we compare the present results with those from DENIS.
 In Sect., In Sect.
 4 we analyze the 2ALASS and DIRBE/IRAS extinction maps. making use of simple models for the optical depth in cillerent directions in the Galaxy. anc discuss the asvmametries between observations in the northern and southern Galactic hemispheres.," 4 we analyze the 2MASS and DIRBE/IRAS extinction maps, making use of simple models for the optical depth in different directions in the Galaxy, and discuss the asymmetries between observations in the northern and southern Galactic hemispheres."
 Finally. the concluding remarks are given in Sect.," Finally, the concluding remarks are given in Sect."
 5., 5.
 In Paper Lowe have built an extinction map in regions of low-extinetion by means of upper giant branch fitting., In Paper I we have built an extinction map in regions of low-extinction by means of upper giant branch fitting.
 We adopted an upper giant branch template from a composite CALD using 7 Bulge fields from E'EIN99., We adopted an upper giant branch template from a composite CMD using 7 Bulge fields from FTK99.
 Using this template we managed to reproduce the mean extinetion in Baade's and Ser E windows., Using this template we managed to reproduce the mean extinction in Baade's and Sgr I windows.
 Aletallicity effects were not. considered in Paper L since Ramirez ct al. (," Metallicity effects were not considered in Paper I, since Ramirez et al. ("
2000). [rom a spectroscopic study of central M giant stars. pointed out that there is no evidence for metallicity gracient. within the inner Bulee(A < 560 pc).,"2000), from a spectroscopic study of central M giant stars, pointed out that there is no evidence for metallicity gradient within the inner Bulge $<$ 560 pc)."
 Besides. E'PIN99. concluded that the amplitude of metallicity variations in the inner Bulge implies very small giant branch slope changes.," Besides, FTK99 concluded that the amplitude of metallicity variations in the inner Bulge implies very small giant branch slope changes."
" Therefore. we assume that the metallicity variations in the inner Bulge do not alfect significantly the extinction estimates and use as reference the same upper giant branch adopted in Paper L: The equation above describes appropriately the upper elant branch locus for stars with 8 —Ay, <12.5."," Therefore, we assume that the metallicity variations in the inner Bulge do not affect significantly the extinction estimates and use as reference the same upper giant branch adopted in Paper I: The equation above describes appropriately the upper giant branch locus for stars with 8 $\le {\it K_0} \le $ 12.5."
" We carried out photometric extractions of stars in the 2MLASS Point Source Catalog available in the Web InterfaceAirsa.ipac.callech.cdu/applicalions/Galor/ lor 61 [lields with ractius —1"" each.", We carried out photometric extractions of stars in the 2MASS Point Source Catalog available in the Web Interface for 61 fields with radius $^{\circ}$ each.
 For most fields. we extracted stars in the range 8S <A.x 11.5. there being typically. 60.000 such stars in cach.," For most fields, we extracted stars in the range 8 $\le {\it K_s} \le $ 11.5, there being typically 60,000 such stars in each."
 For the fields with | «155 we expected: very high extinction (Schultheis et al.," For the fields with $|{\it b}| <$ $^{\circ}$, we expected very high extinction (Schultheis et al."
 1999): thus we extracted fainter stars. down to A.—13.," 1999); thus we extracted fainter stars, down to =13."
 This extra lI mag allowed us to determine the best magnitude range for upper giant branch fitting. taking into account factors such as the increasing photometric errors and contamination by disk stars with fainter magnitude limits. and the decreasing observed. extent of the upper giant branch in. heavily redcened inner Bulge fields.," This extra 1.5 mag allowed us to determine the best magnitude range for upper giant branch fitting, taking into account factors such as the increasing photometric errors and contamination by disk stars with fainter magnitude limits, and the decreasing observed extent of the upper giant branch in heavily reddened inner Bulge fields."
" ""This issue ds cliscussecl in Sect.", This issue is discussed in Sect.
 2.1., 2.1.
 In these low-latituce regions the deeper 2ALASS extractions [led to a considerably larger. number of. stars per feld. = 200.000.," In these low-latitude regions the deeper 2MASS extractions led to a considerably larger number of stars per field, $\approx$ 200,000."
 Considering all fields in the area. we extracted.J and magnitudes for approximately 10? stars.," Considering all fields in the area, we extracted and magnitudes for approximately $\times 10^6$ stars."
 In order to map out the extinction in the 10 Galactic central region. we defined 961 cells with 4.4 in size in cach one of 61 extracted fields.," In order to map out the extinction in the $^{\circ}$ Galactic central region, we defined 961 cells with $^{\prime} \times 4^{\prime}$ in size in each one of 61 extracted fields."
 The extinction was determined, The extinction was determined
As discussed. by several authors. a formal estimate of errors ou the resulting Nuker parameters cau not be performed (e.g.Byunctal.1996).,"As discussed by several authors, a formal estimate of errors on the resulting Nuker parameters can not be performed \citep[e.g.][]{byun96}."
. Since we are mostly interested iu separating galaxies with shallow cusps from those with steep cusps. we explored how the variations of ~ are reflected in the ft aud in the residuals.," Since we are mostly interested in separating galaxies with shallow cusps from those with steep cusps, we explored how the variations of $\gamma$ are reflected in the fit and in the residuals."
 We focus on UGC 7575. one of the two galaxies for which we found an intermediate value of 5=0.33.," We focus on UGC 7575, one of the two galaxies for which we found an intermediate value of $\gamma = 0.33$."
 Its best fit (sce Fie.6)) shows residuals always simaller than 1.5%. quite typical for all fitted galaxies.," Its best fit (see \ref{profiles2}) ) shows residuals always smaller than $\pm
1.5\%$, quite typical for all fitted galaxies."
 In Fig., In Fig.
 8 we present the Nuker fit obtained forcing > to 0.13. a value 0.1 larger with respect to the best fit. leaving all other parameters free to vary.," \ref{ugc7575} we present the Nuker fit obtained forcing $\gamma$ to 0.43, a value 0.1 larger with respect to the best fit, leaving all other parameters free to vary."
 Large and systematic residuals appear toward the ceuter. reaching at (711.," Large and systematic residuals appear toward the center, reaching at 1."
 Although this does not provide us witli a precise error estimate. it indicates that the error on 5 are generally significantly smaller than 0.1. in agreeimneut with the analvsis performed by Byunctal... This implies that except for a few border-line objects our classification is robust.," Although this does not provide us with a precise error estimate, it indicates that the error on $\gamma$ are generally significantly smaller than 0.1, in agreement with the analysis performed by \citeauthor{byun96} This implies that except for a few border-line objects our classification is robust."
 We avo now iu the position to explore how the multiwavelength properties of the galaxies of our sample are related to the brightuess profiles and low they colpare with results obtained on purely optically selected saluples., We are now in the position to explore how the multiwavelength properties of the galaxies of our sample are related to the brightness profiles and how they compare with results obtained on purely optically selected samples.
 Iu Fig., In Fig.
 9 we plot the absolute I baud iaeuitude Mj: derived from 2\LASS with, \ref{mkgamm} we plot the absolute K band magnitude $_K$ derived from 2MASS with
and it is further limited given that a center of rotation needs to be identified.,and it is further limited given that a center of rotation needs to be identified.
 Phe brightest. pixel was used as the center of rotation. which is a reasonable approximation for the optical light. distribution. but the light ancl physical properties do not share well correlated: distributions and the high spatial frequency nature of the physical properties means that asymmetry is probably not the best approach for quantifying these.," The brightest pixel was used as the center of rotation, which is a reasonable approximation for the optical light distribution, but the light and physical properties do not share well correlated distributions and the high spatial frequency nature of the physical properties means that asymmetry is probably not the best approach for quantifying these."
 Moreover. such “patchy” disributions are inherently asvimmetric. and many systems would show an asymmetry close to the maximum possible asymmetry value CXsvmnmetrv-2).," Moreover, such “patchy” distributions are inherently asymmetric, and many systems would show an asymmetry close to the maximum possible asymmetry value (Asymmetry=2)."
 For these reasons asymmetry. too. seems not to be a particularly useful parameter in describing 1e distribution of physical properties in galaxies.," For these reasons asymmetry, too, seems not to be a particularly useful parameter in describing the distribution of physical properties in galaxies."
 Finally. there is the ebumpiness parameter. which would Jearly seem to be a more useful statistic in quantifving these “patchy” distributions of physical properties of galaxies than the previous (wo. as HC was developed to measure this type ol distribution.," Finally, there is the clumpiness parameter, which would clearly seem to be a more useful statistic in quantifying these “patchy” distributions of physical properties of galaxies than the previous two, as it was developed to measure this type of distribution."
 The correlation. cocllicients measured clo not indicate a real correlation. but we serve to distinguish the relative correlation strength. of QAIAL with the CCAS measurements compared to r anc v-bands.," The correlation coefficients measured do not indicate a real correlation, but we serve to distinguish the relative correlation strength of QMM with the CAS measurements compared to $r$ and $u$ -bands."
 The range of values of clumpiness is more likely to be encoding much more physical information about galaxies. rather han the less useful and. potentially biased estimates or concentration and asvmuimetry.," The range of values of clumpiness is more likely to be encoding much more physical information about galaxies, rather than the less useful and potentially biased estimates for concentration and asymmetry."
 “Phe OMM. morphology should aso be sensitive to the “patchy” distributions. of jxel. colours that underlies the clumpy nature of. the »hvsical xoperties., The QMM morphology should also be sensitive to the “patchy” distributions of pixel colours that underlies the clumpy nature of the physical properties.
 ιν could. well be the reason for the vigher correlation values with QMM seen. for. clumpiness compared to the concentration and. asvmmetry as well as he clumpiness of the r-band light distribution., This could well be the reason for the higher correlation values with QMM seen for clumpiness compared to the concentration and asymmetry as well as the clumpiness of the $r$ -band light distribution.
 We also confirmed that this clumpiness is not due to any noise in the v-bane images by showing that the correlation between the clumpiness parameters and QMM output is higher than that between the ρα and the QMSM output., We also confirmed that this clumpiness is not due to any noise in the $u$ -band images by showing that the correlation between the clumpiness parameters and QMM output is higher than that between the $u$ -band and the QMM output.
 This was true for all redshift bins as well., This was true for all redshift bins as well.
 Using other approaches towards quantifving ¢lumpy structure may be helpful in confirming this., Using other approaches towards quantifying clumpy structure may be helpful in confirming this.
 We have compared two separate techniques for analysing galaxies. one using a purely physical approach ZuCCARS) and the other using a purely morphological approach (OMM).," We have compared two separate techniques for analysing galaxies, one using a purely physical approach CAS) and the other using a purely morphological approach (QMM)."
" “Phe purely. physical approach uses the ""uxel-z method to infer the distribution of four. physical oropoertLes (age. star formation rate. colour excess ancl metallicitv). within galaxies."," The purely physical approach uses the Pixel-z method to infer the distribution of four physical properties (age, star formation rate, colour excess and metallicity), within galaxies."
 Vhese distributions were then quantified using the CAS method. where each physical sroperty was assigned a value. for its concentration. asvonmetry ancl elumpiness.," These distributions were then quantified using the CAS method where each physical property was assigned a value for its concentration, asymmetry and clumpiness."
 This is the first. application of CAS to distributed physical properties within galaxies., This is the first application of CAS to distributed physical properties within galaxies.
 The CCAS procedure was also applied to the r-xuxd ancl ρα light. distribution., The CAS procedure was also applied to the $r$ -band and $u$ -band light distribution.
 We did not discover a rivial relationship between the optical light distribution of galaxies and the distribution of the physical properties in CAS space., We did not discover a trivial relationship between the optical light distribution of galaxies and the distribution of the physical properties in CAS space.
 This paces further significance on our objective of using QMM to connect the muliwavelength photometric morphology with the underlving physical properties., This places further significance on our objective of using QMM to connect the multiwavelength photometric morphology with the underlying physical properties.
 This is not. unexpected as CAS is able to be applied only to single-ilter imaging. rather than a full multiwavelength morphology.," This is not unexpected as CAS is able to be applied only to single-filter imaging, rather than a full multiwavelength morphology."
 We also analyzed the morphology of the galaxies using OMM where the images of the galaxies in five filters were cecomposecL using shapelets. followed by a Principal Component Analysis.," We also analyzed the morphology of the galaxies using QMM where the images of the galaxies in five filters were decomposed using shapelets, followed by a Principal Component Analysis."
 To measure the possible correlations between ZuCCAS and QAMAL we carried. out a regression analysis for the CAS parameters of cach physical property against the S principal components and compared these to a similar analysis using the CAS parameters for the ¢-bancl anc," To measure the possible correlations between CAS and QMM, we carried out a regression analysis for the CAS parameters of each physical property against the 8 principal components and compared these to a similar analysis using the CAS parameters for the $r$ -band and"
 To measure the possible correlations between ZuCCAS and QAMAL we carried. out a regression analysis for the CAS parameters of cach physical property against the S principal components and compared these to a similar analysis using the CAS parameters for the ¢-bancl ancl," To measure the possible correlations between CAS and QMM, we carried out a regression analysis for the CAS parameters of each physical property against the 8 principal components and compared these to a similar analysis using the CAS parameters for the $r$ -band and"
"also be a sub-stellar companion. so that we would then have x3 sub-stellar companious around 11 πλατό, ic. στ[ ","also be a sub-stellar companion, so that we would then have $\le 3$ sub-stellar companions around 11 primaries, i.e. $27 \pm 16~\%$."
This should be compared to the frequency and orbit characteristics of other high-mass-ratio binaries with large separations. Le. brown dwarfs iu orbit around a star.," This should be compared to the frequency and orbit characteristics of other high-mass-ratio binaries with large separations, i.e. brown dwarfs in orbit around a star."
 Receutly Cagis et al. (," Recently, Gizis et al. ("
2001) estimated he frequency of wide. visual. ok L- and T-type brown dwarf compauious to normal stars o be ISx114.,"2001) estimated the frequency of wide, visual, old L- and T-type brown dwarf companions to normal stars to be $18 \pm 14~\%$."
 This ΠΡ is consistcat with our estimate eiven above. so that we find no evidence for the ejection scenario.," This number is consistent with our estimate given above, so that we find no evidence for the ejection scenario."
 However. because of the large error bars due to siiuidl-nmuber-statisties. we do not have suffüiient evidence to make a strong statement in this regard.," However, because of the large error bars due to small-number-statistics, we do not have sufficient evidence to make a strong statement in this regard."
 In addition. the surveys discussed by Cagis et al. (," In addition, the surveys discussed by Gizis et al. ("
2001) and our survey have different cvnamical ranges aud detection Iits.,2001) and our survey have different dynamical ranges and detection limits.
 We also ive to refrain from comparing our results with previous survevs for conrpanious around voung low-lnass stars in Chaimacleoun (Reipurth Ziunecker 1993. Brauduer et al.," We also have to refrain from comparing our results with previous surveys for companions around young low-mass stars in Chamaeleon (Reipurth Zinnecker 1993, Brandner et al."
 1996. Chez et al.," 1996, Ghez et al."
 1997. Ioóhhler 2002). because our suuple of prinaries (M6- to M&8-tvpe bona-fide aud. candidate brown dwarfs) is cliffereut from the previously surveved samples (G- to carly M-tvpe T Tauri stars) and because the previous surveys were restricted to stnaller dynamical ranges. up to a magnitude difference of 5 mag between primary aud compauion candidate. while all but one (Cha Πα 5) or two (Cha Ue 2) of oux conipauion candidates have arecr magnitude differences.," 1997, Köhhler 2002), because our sample of primaries (M6- to M8-type bona-fide and candidate brown dwarfs) is different from the previously surveyed samples (G- to early M-type T Tauri stars) and because the previous surveys were restricted to smaller dynamical ranges, up to a magnitude difference of 5 mag between primary and companion candidate, while all but one (Cha $\alpha$ 5) or two (Cha $\alpha$ 2) of our companion candidates have larger magnitude differences."
 Ini our observations. we found one promising colmpanion candidate (Cha Πα 5/cec 1). which could e a brown dwarf (or even elant planet] companion.," In our observations, we found one promising companion candidate (Cha $\alpha$ 5/cc 1), which could be a brown dwarf (or even giant planet) companion."
 The other candidates are all fainter and often detected onlv in L Π they were true companions. they would all be eiut planets. given their magnitudes. asstuuine the same age and distance as towards the primarics.," The other candidates are all fainter and often detected only in I. If they were true companions, they would all be giant planets, given their magnitudes, assuming the same age and distance as towards the primaries."
 However. thev are all at physical separatious co200 AU. which is not expected for planets. but should not be excluded.," However, they are all at physical separations $\ge 200$ AU, which is not expected for planets, but should not be excluded."
 Follow-up 2ud epoch imagine and/or spectroscopy will show. whetler aud which of our companion candidates are trucly cool aud bonuud.," Follow-up 2nd epoch imaging and/or spectroscopy will show, whether and which of our companion candidates are truely cool and bound."
by Fermi (Abdoetal.2010c).,by $\emph{Fermi}$ \citep{abdo10c}.
. The innermost region of Cen A has been resolved with VLBI and shown to have a size of ~3x1016 cm (Kellermannetal.1997;Horiuchi2006).," The innermost region of Cen A has been resolved with VLBI and shown to have a size of $\sim 3 \times 10^{16}$ cm \citep{kellermann97,horiuchi06}."
. Observations at shorter wavelengths also reveal a small core., Observations at shorter wavelengths also reveal a small core.
" VLT infrared interferometry resolves the core size to ~6x10!"" cm (Meisenheimeretal.2007).", VLT infrared interferometry resolves the core size to $\sim 6 \times 10^{17}$ cm \citep{meisen07}.
. The angle of the jet of Cen A to our line of sight is ~50°—80? (Tingayetal.1998) with a preferred value of ~70? (Steinle 2009).., The angle of the sub-parsec jet of Cen A to our line of sight is $\thicksim 50^0-80^0$ \citep{Tingay98} with a preferred value of $\sim 70^o$ \citep{steinle09}. .
" The K-band nuclear flux with starlight subtracted is F(K)~38 mJy, corresponding to vL,~7x10*° erg s! for a distance of 3.5 Mpc (Marconietal.2000)."," The K-band nuclear flux with starlight subtracted is $F(K) \thicksim 38$ mJy, corresponding to $\nu L_{\nu} \thicksim 7 \times 10^{40}$ erg $s^{-1}$ for a distance of $3.5$ Mpc \citep{marconi00}."
". The mid-IR flux of ~1.6 Jy at 11.7 um corresponds to vL,~6x103 erg s~!."," The mid-IR flux of $\thicksim 1.6$ Jy at $11.7 \, \mu$ m corresponds to $\nu L_\nu \thicksim 6 \times 10^{41}$ erg $s^{-1}$ ."
 The broadband SED from radio to y-rays has been fitted with a synchrotron self-Compton model by Abdoetal. (2010a)., The broadband SED from radio to $\gamma$ -rays has been fitted with a synchrotron self-Compton model by \cite{abdo10a}.
". In their model (seeTable2ofAbdoetal.2010a),, a maximum electron Lorentz factor of Ύπιας=1x105 was required in order to produce the observed y-ray emission."," In their model \citep[see Table 2 of ][]{abdo10a}, a maximum electron Lorentz factor of $\gamma_{max}= 1\times10^8$ was required in order to produce the observed $\gamma$ -ray emission."
" However, given the assumed magnetic field of B=6.2 G, this does not seem possible since for y=10? electrons, the synchrotronloss time scale is shorter than their gyro-timescale, which sets the shortest plausible acceleration time scale."," However, given the assumed magnetic field of $B=6.2$ G, this does not seem possible since for $\gamma = 10^8$ electrons, the synchrotronloss time scale is shorter than their gyro-timescale, which sets the shortest plausible acceleration time scale."
" Here, wetherefore present an alternative"," Here, wetherefore present an alternative"
In August 1934 we initiated a program to svstematically monitor the sources in the flix and position-limited Pearson-Beadhead (PearsonandBeadhead.1931.1988). VLBI survey al cm wavelengths using the University of Michigan 26-meter paraboloid operating al 14.5. 8.0. and 4.8 GlIIz. with the primary goal of delineating the rauge of behavior in the total flux and near polarization in a complete ΠακΠίος sample of extragalactic objects.,"In August 1984 we initiated a program to systematically monitor the sources in the flux and position-limited Pearson-Readhead \citep{per81,per88}
 VLBI survey at cm wavelengths using the University of Michigan 26-meter paraboloid operating at 14.5, 8.0, and 4.8 GHz, with the primary goal of delineating the range of behavior in the total flux and linear polarization in a complete flux-limited sample of extragalactic objects."
 The source eroup imcludes representatives of the optical classes QSO. galaxy. and BL Lac object. as well as of several VLBI morphological classes. making it ideally suited for exploring variability properties as a function of source tvpe.," The source group includes representatives of the optical classes QSO, galaxy, and BL Lac object, as well as of several VLBI morphological classes, making it ideally suited for exploring variability properties as a function of source type."
 The limit on source flix density provided by the, The limit on source flux density provided by the
Carretta et al. (,Carretta et al. (
2010b). whose spectroscopic metallicity distribution is similar to tat described above (see their Figure 1).,"2010b), whose spectroscopic metallicity distribution is similar to that described above (see their Figure 1)."
 However. (μον describe the M54. MPP population as having a slightly higher metallicity of [Fe/IJ=—1.55 with a significant (0.19 dex) dispersion.," However, they describe the M54 MPP population as having a slightly higher metallicity of $= -1.55$ with a significant (0.19 dex) dispersion."
 Frinchabov et al. (, Frinchaboy et al. (
in prep.),in prep.)
 explore metallicity distribution along 9grs major and minor axes ancl shows clear evidence of a spatial trend in abundance (about 0.2-0.3 dex out to the radii of our bulge cluster fields)., explore metallicity distribution along Sgr's major and minor axes and shows clear evidence of a spatial trend in abundance (about 0.2-0.3 dex out to the radii of our bulge cluster fields).
 This trend could be eritical to analvsis of the Ser background features., This trend could be critical to analysis of the Sgr background features.
 The derived absolute distance of the Ser CMD features. as measured by comparison to the Paper IV isochrones. is sensitive to which of these seven populations are included in (he analvsis.," The derived absolute distance of the Sgr CMD features, as measured by comparison to the Paper IV isochrones, is sensitive to which of these seven populations are included in the analysis."
 If we include all seven populations. the inferred distances would be many kpe farther than the distances measured [for the Ser core clusters.," If we include all seven populations, the inferred distances would be many kpc farther than the distances measured for the Sgr core clusters."
 However. including all seven populations is likelv inappropriate for the background features. which all lie outside ol the core of Ser (defined as 3.7-3.9 degrees in MO3) and are well wav [rom the central M54 pointing.," However, including all seven populations is likely inappropriate for the background features, which all lie outside of the core of Sgr (defined as 3.7-3.9 degrees in M03) and are well way from the central M54 pointing."
 No study. from either the ground or LST. has found the voungest populations seen in (he Paper IV. analvsis bevond (hese innermost regions of Ser: and none of the five backeround features has the bright blue 5Yng population seen in the M54 field (the brightest turnolL in Figure 1).," No study, from either the ground or HST, has found the youngest populations seen in the Paper IV analysis beyond these innermost regions of Sgr; and none of the five background features has the bright blue SYng population seen in the M54 field (the brightest turnoff in Figure 1)."
 The aforementioned spectroscopic studies of Sgrs inner regions (Clarretta οἱ al., The aforementioned spectroscopic studies of Sgr's inner regions (Carretta et al.
 2010b: Frinchabov οἱ al.," 2010b; Frinchaboy et al.,"
 in prep) would favor using (hie more metal-poor of Sgrs seven populations., in prep) would favor using the more metal-poor of Sgr's seven populations.
 The Frinchabow et al., The Frinchaboy et al.
 study. in particular. (races out the major axis of Ser into regions overlapping our survey.," study, in particular, traces out the major axis of Sgr into regions overlapping our survey."
 The metallicity distribution in (heir outer regions would be best malched by a combination of the Ser MPP population with the more metal poor Slat populations (those with [Fe/II] of 20.5 and —0.7)., The metallicity distribution in their outer regions would be best matched by a combination of the Sgr MPP population with the more metal poor SInt populations (those with [Fe/H] of $-0.5$ and $-0.7$ ).
 Such a choice of isochrones would also be consistent with the approximately 0.2 dex drop in metallicity between the core aud the tidal arms indicated by studies of the tidal arms (LAILOa: Aliud οἱ al., Such a choice of isochrones would also be consistent with the approximately 0.2 dex drop in metallicity between the core and the tidal arms indicated by studies of the tidal arms (LM10a; Alard et al.
 2001: Chou et al., 2001; Chou et al.
 2007. 2010: Giulhida et al.," 2007, 2010; Giuffrida et al."
 2010)., 2010).
 We therefore have chosen these (μου populations Ser AIPP aud the two poorest οἱ (o represent the background Ser debris., We therefore have chosen these three populations – Sgr MPP and the two poorest SInt's – to represent the background Sgr debris.
 The uncertainty. over which populations to use (and the relative strengths of each) introduces some uncertainty into our distauce measures., The uncertainty over which populations to use (and the relative strengths of each) introduces some uncertainty into our distance measures.
 This uncertainty is estimated to be around 0.05 magnitudes based on a comparison of the distance moduli derived for various combinations of stellar populations., This uncertainty is estimated to be around 0.05 magnitudes based on a comparison of the distance moduli derived for various combinations of stellar populations.
 The isochrones were fist placed at M54's distance and NGC 6624s reddening. with the Apa and. psg extinction coefficients set to Chose of Sirianni et al. (," The isochrones were first placed at M54's distance and NGC 6624's reddening, with the $A_{F606W}$ and $A_{F814W}$ extinction coefficients set to those of Sirianni et al. ("
2005).,2005).
 We then varied (he distance ancl reddening until the Ser isochrones overlapped the composite backeround population created from the shifted aud co-added color-magnitude diagrams of all five clusters., We then varied the distance and reddening until the Sgr isochrones overlapped the composite background population created from the shifted and co-added color-magnitude diagrams of all five clusters.
 To check internal consistency. we then fitted the isochrones to the background population of each cluster CMD on an individual basis.," To check internal consistency, we then fitted the isochrones to the background population of each cluster CMD on an individual basis."
 We recover the relative shifts between, We recover the relative shifts between
As expected. (he LPAI [requeney. decreases with the density of the electrons at the star's surface. and [or z>500 Im the effect of the multiple scattering can be neelected.,"As expected, the LPM frequency decreases with the density of the electrons at the star's surface, and for $z\geq 500$ fm the effect of the multiple scattering can be neglected."
 The variation of the LPAI frequency as a funcüon of z in the presence and in the absence of an electric field is represented. for (wo different values of the temperature. in Fie.," The variation of the LPM frequency as a function of $z$ in the presence and in the absence of an electric field is represented, for two different values of the temperature, in Fig."
 3., 3.
 The mean value of the LPAI frequency. (eppay). can be easily obtained in the low temperature limit. 7— 0.," The mean value of the LPM frequency, $\left\langle \omega
_{LPM}\right\rangle $, can be easily obtained in the low temperature limit, $%
T\rightarrow 0."
" In this case. bv neglecting the effect of the external electric field and by taking into account that the effect of multiple scattering has a significant effect on the electron radiation ouly [or a distance of the order τη. we obtain. The value of τη can be obtained as a solution of the algebraic equation 2.37(2,4.1)= L7)."," In this case, by neglecting the effect of the external electric field and by taking into account that the effect of multiple scattering has a significant effect on the electron radiation only for a distance of the order $z_{crit}$, we obtain, The value of $z_{crit}$ can be obtained as a solution of the algebraic equation $z_{crit}n\left( z_{crit},T\right) =1/16\pi
r_{e}^{2}Z^{2}\ln \left( 184Z^{-1/3}\right)$ ."
" The mean frequency of the emitted radiation essentially depends on the value of the electric potential at the quark star surface V, and of the value of the electric field.", The mean frequency of the emitted radiation essentially depends on the value of the electric potential at the quark star surface $V_{q}$ and of the value of the electric field.
 The mean value of the electron energy in the dense laver of the electrosphere. (e)=--hcV(z)dz. can be represented as For the mean value of the electric field in the region 0€z<z4; we obtain The mean value of the parameter /m?.describingeEe/m?=y/2o/3xeV the effect of the electric field on the radiation processes in (the electrosphere of the quark stars is given by," The mean value of the electron energy in the dense layer of the electrosphere, $\left\langle \epsilon
\right\rangle =\left( 1/z_{crit}\right)
\int_{0}^{z_{crit}}V(z)dz$, can be represented as For the mean value of the electric field in the region $0\leq
z\leq z_{crit}$ we obtain The mean value of the parameter }$, describing the effect of the electric field on the radiation processes in the electrosphere of the quark stars is given by"
Gaudi&Petteis(2001) emphasized to have worked out in (2-d) real variables.,\citet{vertical} emphasized to have worked out in (2-d) real variables.
 They made an erroneous conclusion and kindly warned the readers accordinglv. This claàm may imply a correlation between the error and the choice of variables. but that is not the case.," They made an erroneous conclusion and kindly warned the readers accordingly, This claim may imply a correlation between the error and the choice of variables, but that is not the case."
 The missing term in the second equation of SEF.6.17 could have been recovered if one had stared at the two equations of SEF.6.17 and scratched the head about the meanings of the linear terms in comparison to the square root terms., The missing term in the second equation of SEF.6.17 could have been recovered if one had stared at the two equations of SEF.6.17 and scratched the head about the meanings of the linear terms in comparison to the square root terms.
 In our humble belief. (astvo)physically meaninglul quantities have intrinsic values (many (mes as (he form of geometric quantities) ancl offer intuitive understanding.," In our humble belief, (astro)physically meaningful quantities have intrinsic values (many times as the form of geometric quantities) and offer intuitive understanding."
 One advantage of the complex variable is (hat. quasi-canalvlic lens equations are. well. quasi-analvtic.," One advantage of the complex variable is that quasi-analytic lens equations are, well, quasi-analytic."
 The derivatives are completely determined by one analytic function., The derivatives are completely determined by one analytic function.
 There is nothing much great about it except that it is simpler., There is nothing much great about it except that it is simpler.
 The svstem is constrained and the constraint is completely contained within the nature of the variable. and so il is easier {ο eel to know the svstem.," The system is constrained and the constraint is completely contained within the nature of the variable, and so it is easier to get to know the system."
 Especially. for one variable analytic fanction. we even know how to," Especially, for one variable analytic function, we even know how to"
"the JPL (Pickettetal.1998) and CDMS (Mülleretal.2001,2005) databases and cross-checked those with the online Splatalogue compilation.","the JPL \citep{pickett98} and CDMS \citep{mueller01,mueller05} databases and cross-checked those with the online Splatalogue compilation."
" Since all of the lines are narrow little line-blending occurs, in contrast with high-mass star-forming regions."," Since all of the lines are narrow little line-blending occurs, in contrast with high-mass star-forming regions."
" Also, all the assigned lines are centered within 0.1-0.2 oof the systemic velocity of 7.0 oof IRAS4B. Most prominent in the spectrum are lines of dimethyl ether, CH30CHs3, with 5 identified transitions."," Also, all the assigned lines are centered within 0.1–0.2 of the systemic velocity of 7.0 of IRAS4B. Most prominent in the spectrum are lines of dimethyl ether, $_3$ $_3$, with 5 identified transitions."
" In addition, lines of sulfur dioxide, SO2, and water, H80, are clear detections with a fainter line of ethyl cyanide, C2H5CN, also present."," In addition, lines of sulfur dioxide, $_2$, and water, $_2^{18}$ O, are clear detections with a fainter line of ethyl cyanide, $_2$ $_5$ CN, also present."
 Concerning the confidence of the assignment of the H;80 line: according to the Splatalogue compilation no other transitions fall within +1 oof the location of the H;80 line., Concerning the confidence of the assignment of the $_2^{18}$ O line: according to the Splatalogue compilation no other transitions fall within $\pm$ 1 of the location of the $_2^{18}$ O line.
" Offset by 1-2 aare transitions of CH;CHCN, ?CH4CH5;CN and (CH3);CO: the two former can be ruled out because of the lack of additional components which should have been observable at larger velocity offsets, whereas the latter has a very low intrinsic line strength."," Offset by 1–2 are transitions of $^{13}$ $_2$ CHCN, $^{13}$ $_3$ $_2$ CN and $_3$ $_2$ CO: the two former can be ruled out because of the lack of additional components which should have been observable at larger velocity offsets, whereas the latter has a very low intrinsic line strength."
" For each of the detected lines we fit a circular Gaussian in the (u,v) plane to the line emission integrated over the widths of the lines to zero intensity (about +1 ))."," For each of the detected lines we fit a circular Gaussian in the $(u,v)$ plane to the line emission integrated over the widths of the lines to zero intensity (about $\pm 1$ )."
" The derived fluxes, peak positions and sizes are given in Table 3 as well."," The derived fluxes, peak positions and sizes are given in Table \ref{line_id} as well."
" AII the detected lines are marginally resolved with deconvolved sizes of 0.2""— 0.6"", corresponding to radii of only 25 to 75 AU."," All the detected lines are marginally resolved with deconvolved sizes of $''$ $''$, corresponding to radii of only 25 to 75 AU."
" Interestingly, most of the detected lines are offset by about 0.10—-0.15""from the continuum position in right ascension — with the exception of CoH5CN."," Interestingly, most of the detected lines are offset by about from the continuum position in right ascension – with the exception of $_2$ $_5$ CN."
" The peaks of the molecular emission still fall within the deconvolved extent of the continuum emission, though."," The peaks of the molecular emission still fall within the deconvolved extent of the continuum emission, though."
 The images of the complex organic molecules and SO» will be discussed elsewhere., The images of the complex organic molecules and $_2$ will be discussed elsewhere.
 The high spectral and spatial resolution offered by the IRAM data reveals a tentative velocity gradient in the H580 emission., The high spectral and spatial resolution offered by the IRAM data reveals a tentative velocity gradient in the $_2^{18}$ O emission.
 Fig., Fig.
 3 shows a moment-1 (velocity) map of this emission indicating a change in velocity of a few x 0.1 oover the extent of the source emission., \ref{vel_gradient} shows a moment-1 (velocity) map of this emission indicating a change in velocity of a few $\times$ 0.1 over the extent of the source emission.
" To estimate the magnitude of the velocity gradient the centroid of the emission was determined channel-by-channel in the (u, v)-plane."," To estimate the magnitude of the velocity gradient the centroid of the emission was determined channel-by-channel in the $(u,v)$ -plane."
 The derived offsets along and perpendicular to the largest velocity gradient are then plotted as function of velocity and a linear fit performed (Fig. 3))., The derived offsets along and perpendicular to the largest velocity gradient are then plotted as function of velocity and a linear fit performed (Fig. \ref{vel_gradient}) ).
 The velocity gradient found in this way is 9.4 arcsec! or 7.7x10? pc! at a 3.5-4σ confidence level., The velocity gradient found in this way is 9.4 $^{-1}$ or $\times 10^3$ $^{-1}$ at a $\sigma$ confidence level.
 The velocity gradient is about 3 orders of magnitude larger than the typical gradients observed on arcminute scales in pre-stellar cores (Goodmanetal.1993) — supporting a scenario in which the H;80 emission has its origin in a more rapidly rotating structure such as a central disk — although it is not possible to address whether the velocity for example is Keplerian in nature., The velocity gradient is about 3 orders of magnitude larger than the typical gradients observed on arcminute scales in pre-stellar cores \citep{goodman93} – supporting a scenario in which the $_2^{18}$ O emission has its origin in a more rapidly rotating structure such as a central disk – although it is not possible to address whether the velocity for example is Keplerian in nature.
" The position angle of the largest gradient is ~75° measured East of North (with an accuracy of +5—10° from the gradient measured in the (u,v) plane), which, interestingly, is nearly perpendicular to the axis of the protostellar outflow (upper panel of Fig. 3))"," The position angle of the largest gradient is $\approx$ $^\circ$ measured East of North (with an accuracy of $\pm$ $^\circ$ from the gradient measured in the $(u,v)$ plane), which, interestingly, is nearly perpendicular to the axis of the protostellar outflow (upper panel of Fig. \ref{vel_gradient}) )"
 traced by water masers in the North-South direction with a position angle of —29° and thermal line emission on larger scales at a position angle of 0? ," traced by water masers in the North-South direction with a position angle of $-$ $^\circ$ \citep[][]{marvel08,desmurs09} and thermal line emission on larger scales at a position angle of $^\circ$ "
equation (26)) around pr=0.,equation \ref{EQ2}) ) around $x=0$ .
 For |.r|«1. we have As2(diAsiu(ora).," For $\left| x \right|
\ll 1$ , we have ${\cal A}_2 \simeq (d{\cal A}_2/dr)_{r=r_M}(r-r_M)$."
" Here we assume that h—ef(082)~Hron the parameters 6). bo. οι. d, and do are at most of order unity/dr) and 3~1."," Here we assume that $h \equiv c/ \left( r \Omega \right) \sim H/r \ll
1$ , the parameters $b_1$, $b_2$, $c_1$, $d_1$ and $d_2$ are at most of order unity and $\beta \sim 1$ ."
 This amounts to assuming that only e varies in the expression (27)) of de., This amounts to assuming that only $\sigma$ varies in the expression \ref{coefA2}) ) of ${\cal A}_2$.
 For |r“lowe then have: where the subscript AZ indicates that the corresponding quantities are taken at r=ry. and εξ| or |1 depending on whether ray=rug o£ rog," For $\left| x \right| \ll 1$ we then have: where the subscript $M$ indicates that the corresponding quantities are taken at $r=r_M$ , and $\epsilon = -1$ or $+1$ depending on whether $r_M=r_{IMR}$ or $r_{OMR}$."
 We now approximate i. 45 and So by their values at r=riy.," We now approximate ${\cal A}_1$, ${\cal A}_0$ and ${\cal S}_0$ by their values at $r=r_M$ ."
 This requires [e|«fay., This requires $\left| x \right| \ll h_M$.
" HE this condition is not satisfied. QQO, cannot be replaced by its value at r=ru."," If this condition is not satisfied, $\Omega-\Omega_p$ cannot be replaced by its value at $r=r_M$ ."
 To the first non zero order in fay and allowing for mhyp. we canthen rewrite equation (26)) under the form: where we have defined the dimensionless quantity νε—f./rap and where νὰ and iS are given by: introduce the new variable v= 2yVakr. so that equation (42)) becomes: gin 1 Note that is real outside the magnetic resonances. where Ar>0. whereas it is imaginary inside the resonances. where alr<0.," To the first non zero order in $h_M$ and allowing for $m \sim h_M^{-1}$, we canthen rewrite equation \ref{EQ2}) ) under the form: where we have defined the dimensionless quantity $\tilde{\xi}_{mr} \equiv \xi_{mr}/r_M$ and where ${\cal A}$ and ${\cal
S}$ are given by: We next introduce the new variable $u=2 \sqrt{{\cal A}x}$ , so that equation \ref{EQ2b}) ) becomes: Note that $u$ is real outside the magnetic resonances, where ${\cal
A}x>0$, whereas it is imaginary inside the resonances, where ${\cal
A}x<0$."
" The solutions of this equation can be written as: where Jy and Yo are the Bessel functions of first and second kind. respectively. and €, and C7 are (complex) constants wich depend on y. fay andm."," The solutions of this equation can be written as: where $J_0$ and $Y_0$ are the Bessel functions of first and second kind, respectively, and $C_\epsilon$ and $C'_\epsilon$ are (complex) constants which depend on $\beta_M$, $h_M$ and$m$."
 The factor 7/2 has been introduced to simplify the asymptotic expression of £ below., The factor $\pi/2$ has been introduced to simplify the asymptotic expression of $\tilde{\xi}$ below.
 When w+0 (ie. Le]< hi). Jot)~ land Yo(u)~οπ)απ (Abramowitz Stegun 1972).," When $u \rightarrow 0$ (i.e. $\left| x \right| \ll h_M$ ), $J_0(u) \sim
1$ and $Y_0(u) \sim (2 / \pi) \ln u$ (Abramowitz Stegun 1972)."
" To deal with the singularity abn=0. we use the Landau prescription. te. we clisplace the pole slightly from the real axis by. replacing QO, with QODO,IE. where E is a small positive real constant."," To deal with the singularity at $u=0$, we use the Landau prescription, i.e. we displace the pole slightly from the real axis by replacing $\Omega - \Omega_p$ with $\Omega - \Omega_p - i \Gamma$, where $\Gamma$ is a small positive real constant."
 This amounts to multiplying all the perturbed quantities by exp(L7). and therefore to considering the response of the disc to a slowly increasing perturbation.," This amounts to multiplying all the perturbed quantities by ${\rm{exp}}\left(\Gamma t
\right)$, and therefore to considering the response of the disc to a slowly increasing perturbation."
 This procedure can also be thought of as introducing a viscous force. or as shifting the resonances progressively with time.," This procedure can also be thought of as introducing a viscous force, or as shifting the resonances progressively with time."
 Note that such a shift occurs when the planet drifts through the disc. since the material in the resonances then changes with time.," Note that such a shift occurs when the planet drifts through the disc, since the material in the resonances then changes with time."
" Shifting QO, below the real axis is equivalent to shifting .7 above this axis in equation (42)). since only QO, was allowed to vary in the expression of Ab."," Shifting $\Omega - \Omega_p$ below the real axis is equivalent to shifting $x$ above this axis in equation \ref{EQ2b}) ), since only $\Omega-\Omega_p$ was allowed to vary in the expression of ${\cal A}_2$."
" ""Therefore. we replace  with |is. where + is a small positive real constant."," Therefore, we replace $x$ with $x+ i \gamma$, where $\gamma$ is a small positive real constant."
" For |r|Char. the solutions of equation (45)) can then be written as: so that. near c=0: Note that thereis à phase shift of z on passing through .r—0.Since"" S5/.À has opposite. sign. at r—riyg and kr—roug (ON,∙∕∙0.changessign. on passing. throughcorotation. whereas Wi, does not) and is real.the real part of the solutions of equation (45)) should haveopposite signal r=rap and r=roug whereas their imaginaryparts should have the same sign."," For $\left| x \right| \ll h_M$, the solutions of equation \ref{EQ2c}) ) can then be written as: so that, near $x=0$: Note that thereis a phase shift of $\pi$ on passing through $x=0$.Since ${\cal S}/ {\cal A}$ has opposite sign at $r=r_{IMR}$ and $r=r_{OMR}$ $\partial \Psi'_m / \partial x$changessign on passing throughcorotation whereas $ \Psi'_m $ does not) and is real,the real part of the solutions of equation\ref{EQ2c}) ) should haveopposite signat $r=r_{IMR}$ and $r=r_{OMR}$ whereas their imaginaryparts should have the same sign."
" Therefore. from equation (48)) we ect €,=ο 4. where the asterisk denotes the complex conjugate."," Therefore, from equation \ref{sol00}) ) we get $C_{+1} = -C_{-1}^{\ast}$ , where the asterisk denotes the complex conjugate."
" Equations (32)) ancl (34)) can now be used to evaluate 05,"" and W7, for fe]« ha: where Ou, is Ὁ taken atr= ry.", Equations \ref{Wm2}) ) and \ref{vphim2}) ) can now be used to evaluate $v'_{m \varphi}$ and $W'_m$ for $\left| x \right| \ll h_M$ : where $\Omega_M$ is $\Omega$ taken at$r=r_M$ .
" Finally. (5,=Amet; vields. for |r]"," Finally, $v'_{mr} = i m \sigma \xi_{mr}$ yields, for $\left| x \right| \ll
h_M$ :"
relation predicted by self-enrichment models is a sensitive function of the assumed protoglobular cloud mass-radius relation.,relation predicted by self-enrichment models is a sensitive function of the assumed protoglobular cloud mass-radius relation.
 We have investigated how the cluster-forming core mass-radius relation influences the resilience of clusters to an external tidal tield during their violent relaxation. namely. while they dynamically respond to the expulsion of their residual star-forming gas.," We have investigated how the cluster-forming core mass-radius relation influences the resilience of clusters to an external tidal field during their violent relaxation, namely, while they dynamically respond to the expulsion of their residual star-forming gas."
" Owing to their high volume densities (152vIdOen 4). observed Galactic eluster-forming cores are virtually ""immune? o the external tidal field arising from an isothermal galactic xo with a circular velocity V.220&nrs|."," Owing to their high volume densities $n_{H2} \simeq 10^4-10^5\,cm^{-3}$ ), observed Galactic cluster-forming cores are virtually `immune' to the external tidal field arising from an isothermal galactic halo with a circular velocity $V_c \simeq 220\,km.s^{-1}$."
 Following gas expulsion. however. clusters experience a severe density decrease hrough spatial expansion and star-loss (cluster infant weight-loss).," Following gas expulsion, however, clusters experience a severe density decrease through spatial expansion and star-loss (cluster infant weight-loss)."
 This renders expanded gas-free clusters more vulnerable to an external tidal field. possibly enhancing their infant weight-loss and ikelihood of dissolution compared to clusters evolving in a tield-free environment.," This renders expanded gas-free clusters more vulnerable to an external tidal field, possibly enhancing their infant weight-loss and likelihood of dissolution compared to clusters evolving in a tidal-field-free environment."
" To assess the impact of the core mass-radius relation on this issue. we have built on the N-body model grid of(20073... which provides the bound fraction F5, of stars in clusters at the end of violent relaxation as a function of core SFE. gas expulsion time-scale CGgsp/Goss (n. units of a core crossing-time Trove) and tidal field impact ryfr. where ry and r, are the half-mass radius and tidal radius of the embedded cluster. respectively."," To assess the impact of the core mass-radius relation on this issue, we have built on the $N$ -body model grid of, which provides the bound fraction $F_{bound}$ of stars in clusters at the end of violent relaxation as a function of core SFE, gas expulsion time-scale $\tau_{GExp}/\tau_{cross}$ (in units of a core crossing-time $\tau_{cross}$ ) and tidal field impact $r_h/r_t$, where $r_h$ and $r_t$ are the half-mass radius and tidal radius of the embedded cluster, respectively."
 Assuming a given external tidal field. SFE and σας expulsion time-scale. we have estimated the bound fraction Fp as a function of the mass of cluster progenitor cores for various core mass-radius relations rio;=Ya$e," Assuming a given external tidal field, SFE and gas expulsion time-scale, we have estimated the bound fraction $F_{bound}$ as a function of the mass of cluster progenitor cores for various core mass-radius relations $r_{core} = \chi m_{core}^{\delta}$."
 We have considered the cases of constant core surface density (6= 1/2). constant core volume density (6= 1/3) and constant core radius (9= 0). which we refer to the Leare. Peore and río; models.," We have considered the cases of constant core surface density $\delta = 1/2$ ), constant core volume density $\delta = 1/3$ ) and constant core radius $\delta = 0$ ), which we refer to the $\Sigma_{core}$, $\rho_{core}$ and $r_{core}$ models."
" For each slope ὃ. we have tested two different normalizations 7%. which we refer to the ""compact"" and ‘loose’ models (see Table |)."," For each slope $\delta$, we have tested two different normalizations $\chi$, which we refer to the `compact' and `loose' models (see Table \ref{tbl:mrr}) )."
 Loose models describe molecular cores mapped in CO emission line., Loose models describe molecular cores mapped in $C^{18}O$ emission line.
 These cores are sometimes void of YSOs., These cores are sometimes void of YSOs.
 In contrast. compact models describe dense molecular cores selected for their star-formation activity.," In contrast, compact models describe dense molecular cores selected for their star-formation activity."
 The compact core mass-radius relations therefore constitute a better proxy to cluster-formation conditions than their loose counterpars., The compact core mass-radius relations therefore constitute a better proxy to cluster-formation conditions than their loose counterparts.
 We have shown tha constant surface density cores and constant radius cores resut in the preferential removal of high- and low-mass clusters. resyectively.," We have shown that constant surface density cores and constant radius cores result in the preferential removal of high- and low-mass clusters, respectively."
 In contrast. constant volume density cores robustly lead to mass-independent cluster infant weight-loss (Fig. 33).," In contrast, constant volume density cores robustly lead to mass-independent cluster infant weight-loss (Fig. \ref{fig:fb}) )."
 This is because constant surface density cores are characterised by a decreasing volume density — hence growing vulnerability to an external tidal field — with increasing core mass. while the oppposite is true for constant radius cores (Fig. +)).," This is because constant surface density cores are characterised by a decreasing volume density – hence growing vulnerability to an external tidal field – with increasing core mass, while the oppposite is true for constant radius cores (Fig. \ref{fig:crr}) )."
" Building on the compact mass-radius relations (X,=¢.cml ngo—6.IO!em* and reore—0.3pe). these results are further quantified."," Building on the compact mass-radius relations $\Sigma_{core}=0.5\,g.cm^{-2}$ , $n_{H2}=6.10^4\,cm^{-3}$ and $r_{core}=0.3\,pc$ ), these results are further quantified."
" Ata galactocentric distance D,tkkpe in an isothermal galactic halo with a circular velocity V.—220&n.s.|. constant surface density cores preclude the formation of bound clusters more massive than a few 10AZ... as a result of the greater vulnerability of more massive cores to the external tidal field."," At a galactocentric distance $D_{gal} \simeq 4$ kpc in an isothermal galactic halo with a circular velocity $V_c=220\,km.s^{-1}$, constant surface density cores preclude the formation of bound clusters more massive than a few $10^4\,M_{\sun}$ as a result of the greater vulnerability of more massive cores to the external tidal field."
 In contrast. the tidal field impact for constant volume density cores is independent of their mass and weak. that is. these cores respond to gas expulsion as if evolving in a tidal-field-free environment (see Figs 2.. 3. and 4».," In contrast, the tidal field impact for constant volume density cores is independent of their mass and weak, that is, these cores respond to gas expulsion as if evolving in a tidal-field-free environment (see Figs \ref{fig:rhrt}, \ref{fig:fb} and \ref{fig:crr}) )."
 The presence of massive star clusters pu106 AZ...) in the old haloes of elliptical and spiral galaxies. as well as instarbursts and galaxy mergers. may therefore constitute a hint that cluster-progenitor cores are not characterised by a constant surface density.," The presence of massive star clusters $\gtrsim 10^6\,M_{\sun}$ ) in the old haloes of elliptical and spiral galaxies, as well as instarbursts and galaxy mergers, may therefore constitute a hint that cluster-progenitor cores are not characterised by a constant surface density."
 Besides. almost all of the cluster-forming cores in the middle and bottom panels of Fig.," Besides, almost all of the cluster-forming cores in the middle and bottom panels of Fig."
 |. have ΊοIQem.7.," \ref{fig:obs} have $n_{H2} > 10^4\,cm^{-3}$."
 Very massive molecular cores whose mean surface density is X4;=0.5¢.cm* are not that dense.," Very massive molecular cores whose mean surface density is $\Sigma_{core} = 0.5\,g.cm^{-2}$ are not that dense."
" For instance. the mean number density of a core of mass 10)ΑΗ. and surface density X4,=0.5¢.cm* is nyo~750c0m3 only."," For instance, the mean number density of a core of mass $10^7\,M_{\sun}$ and surface density $\Sigma_{core} = 0.5\,g.cm^{-2}$ is $n_{H2} \simeq 750\,cm^{-3}$ only."
 Therefore. such cores may no result in the formation of massive clusters because of inefticien star formation in the first place.," Therefore, such cores may not result in the formation of massive clusters because of inefficient star formation in the first place."
 The shape of the mass function of young clusters is reported to be invariant through violent relaxation (albeit with a declining amplitude). which points to mass-independent cluster. infan weight-loss.," The shape of the mass function of young clusters is reported to be invariant through violent relaxation (albeit with a declining amplitude), which points to mass-independent cluster infant weight-loss."
 In that respect. our analysis supports the hypothesis that cluster-forming cores have a constant volume density.," In that respect, our analysis supports the hypothesis that cluster-forming cores have a constant volume density."
" When the SFE and gas expulsion time-scale are constant. cores of constant surface density produce mass-independent infant weight-loss for specitie conditions only: the core mass range upper limi must be such that ry/r,2;0.05."," When the SFE and gas expulsion time-scale are constant, cores of constant surface density produce mass-independent infant weight-loss for specific conditions only: the core mass range upper limit must be such that $r_h/r_t \lesssim 0.05$."
 This can be obtained either through a weak tidal field (i.e. large 7) and/or a high surface density (to decrease rj. and/or a small upper limit on the core mass range as for observed molecular cores in our Galaxy (Fig. H9).," This can be obtained either through a weak tidal field (i.e. large $r_t$ ) and/or a high surface density (to decrease $r_h$ ), and/or a small upper limit on the core mass range as for observed molecular cores in our Galaxy (Fig. \ref{fig:obs}) )."
 Figure + constitutes an efficient and straightforward tool to assess whether molecular cores produce clusters affected by an external tidal field or not., Figure \ref{fig:crr} constitutes an efficient and straightforward tool to assess whether molecular cores produce clusters affected by an external tidal field or not.
 Using Eq. 10.," Using Eq. \ref{eq:rhrt-n},"
" which provides the tidalfield impact ry,/r, as a function of the core number density 7jj» and of the strength of an external tidal field (e. a galactic halo circular velocity V. and a galactocentric distance D,,/). one can compare a given core mass-radius relation to lines of constant η."," which provides the tidalfield impact $r_h/r_t$ as a function of the core number density $n_{H2}$ and of the strength of an external tidal field (i.e. a galactic halo circular velocity $V_c$ and a galactocentric distance $D_{gal}$ ), one can compare a given core mass-radius relation to lines of constant $r_h/r_t$ ."
 As long as the core mass-radius relation stands below ry/r= 0.05. the tidal field impact is weak to non-existent.," As long as the core mass-radius relation stands below $r_h/r_t=0.05$ , the tidal field impact is weak to non-existent."
" AS rj/r, increases. cluster infant weight-loss becomes larger (see also Figs 2. and 3)."," As $r_h/r_t$ increases, cluster infant weight-loss becomes larger (see also Figs \ref{fig:rhrt} and \ref{fig:fb}) )."
2009).,.
". To elucidate the pathway to the first galaxies, an essential ingredient is to characterize the conditions in the gas clouds out of which the second generation of stars form (e.g.,Brommetal.2009)."," To elucidate the pathway to the first galaxies, an essential ingredient is to characterize the conditions in the gas clouds out of which the second generation of stars form \citep[e.g.,][]{bromm09}."
. We in particular need to quantify the metal content in such clouds., We in particular need to quantify the metal content in such clouds.
" This applies to the first galaxy itself, but also to neighboring minihalos which collapse during the assembly process."," This applies to the first galaxy itself, but also to neighboring minihalos which collapse during the assembly process."
" It is generally believed that minihalos can only host Pop III stars, as self-enrichment is unlikely, and transport of metals out of the IGM into pre-condensed halos might be severely limited (Cen&Riquelme2008)."," It is generally believed that minihalos can only host Pop III stars, as self-enrichment is unlikely, and transport of metals out of the IGM into pre-condensed halos might be severely limited \citep{cr08}."
". Here, we can address these questions within their proper cosmological setting."," Here, we can address these questions within their proper cosmological setting."
" We note that we use the terminology to classify early stellar populations, in particular the distinction between Pop IIL1 and IIL2, recently introduced by McKee&Tan(2008).."," We note that we use the terminology to classify early stellar populations, in particular the distinction between Pop III.1 and III.2, recently introduced by \citet{mt08}."
 The former refers to metal-free stars whose formation is entirely unaffected by any stellar, The former refers to metal-free stars whose formation is entirely unaffected by any stellar
Orientation-owed: unified: schemes for. racio-louck quasars and. powerftd radio ealaxies play ai Κον role. in our uncerstanding of these objects.,Orientation-based unified schemes for radio-loud quasars and powerful radio galaxies play a key role in our understanding of these objects.
 Ever since the conception of the idea tiat powerful radio galaxies are simply quasars with heir j«4 axes oriented away from our line of sight. such that tl1 nuclear continuum ancl broac-lince regions are 0ISCULCE by a clusty torus or warped dise. (Scheuer LOST: Barthe 1950). evidence has been sought to prove that these unificaion schemes are correct.," Ever since the conception of the idea that powerful radio galaxies are simply quasars with their jet axes oriented away from our line of sight, such that the nuclear continuum and broad-line regions are obscured by a dusty torus or warped disc (Scheuer 1987; Barthel 1989), evidence has been sought to prove that these unification schemes are correct."
 Reviews of unification schemes [or active galactic nuclei have been presented. by Antonucci (993) and Urrv Padovani (1995)., Reviews of unification schemes for active galactic nuclei have been presented by Antonucci (1993) and Urry Padovani (1995).
 Although the majority ol observations are consistent with them. some observations have been used to cast doubt on their viability over all ranges of radio luminosities ancl redshifts.," Although the majority of observations are consistent with them, some observations have been used to cast doubt on their viability over all ranges of radio luminosities and redshifts."
 The narrow line emission in radio sources is observed to be emitted largely from bevond the obscuring material (c.g. MeCarthy. Spinrad van Breugel 1995: Les. Barthel Fosbury. 1996) and therefore is independent of the jet axis orientation.," The narrow line emission in radio sources is observed to be emitted largely from beyond the obscuring material (e.g. McCarthy, Spinrad van Breugel 1995; Hes, Barthel Fosbury 1996) and therefore is independent of the jet axis orientation."
 Hence in this model one would expect radio galaxies and. quasars to have similar narrow line luminosities., Hence in this model one would expect radio galaxies and quasars to have similar narrow line luminosities.
 The strong positive correlation between the extended: radio luminosities and narrow line luminosities of radio sources (Baum Lleckman 1989: Rawlings et al., The strong positive correlation between the extended radio luminosities and narrow line luminosities of radio sources (Baum Heckman 1989; Rawlings et al.
 1989: Willott et al., 1989; Willott et al.
 1999) means that the samples of radio galaxies ancl quasars to be compared must be matched. in extended: radio luminosity., 1999) means that the samples of radio galaxies and quasars to be compared must be matched in extended radio luminosity.
 At low redshift (z< 0.5} there have been claims that cuaseu’s are observed to have OLLI] line luminosities a factor of," At low redshift $z<0.8$ ), there have been claims that quasars are observed to have [OIII] line luminosities a factor of"
at radio wavelengths (Bhattacharya1990:; G99).,at radio wavelengths \cite{bha90}; G99).
 This argument is supported by recent observations of eemission in the region (Dubneretal.2002))., This argument is supported by recent observations of emission in the region \cite{dgg+02}) ).
 X-ray emission from this system. also displayed in Figure .. has revealed a different but similarly complicated picture.," X-ray emission from this system, also displayed in Figure \ref{fig_g320_most_rosat}, has revealed a different but similarly complicated picture."
 Apart from a pulsed point-source corresponding to the pulsar itself (Seward&HarndenJr. 1982)). X-ray emission from the system is dominated by an elongated non-thermal source centered on the pulsar. and an extended thermal source coincident with the northern radio component of the SNR and with the optical nebula (Sewardetal.1983:; Trussontetal. 1996)).," Apart from a pulsed point-source corresponding to the pulsar itself \cite{sh82}) ), X-ray emission from the system is dominated by an elongated non-thermal source centered on the pulsar, and an extended thermal source coincident with the northern radio component of the SNR and with the optical nebula \cite{shmc83}; \cite{tmc+96}) )."
 Faint diffuse X-rays extend over the entire extent of the radio SNR. and potentially correspond to thermal emission from the SNR blast wave (Trussonietal. 1996)).," Faint diffuse X-rays extend over the entire extent of the radio SNR, and potentially correspond to thermal emission from the SNR blast wave \cite{tmc+96}) )."
 The central non-thermal source has been interpreted as a PWN powered by the pulsar (Sewardetal.1984:: Trussonietal. 1996))., The central non-thermal source has been interpreted as a PWN powered by the pulsar \cite{shss84}; \cite{tmc+96}) ).
 As mentioned above. no radio emission from this PWN has been seen. presumably because of the low density into which the pulsar wind expands (Bhattacharya1990:; G99).," As mentioned above, no radio emission from this PWN has been seen, presumably because of the low density into which the pulsar wind expands \cite{bha90}; G99)."
 High-resolution observations with sshow the X-ray PWN to be significantly collimated to the SE of the pulsar (Fig. 1: Greiveldingeretal. 1995))., High-resolution observations with show the X-ray PWN to be significantly collimated to the SE of the pulsar (Fig. \ref{fig_g320_most_rosat}; \cite{gcm+95}) ).
 Several authors have interpreted this region as corresponding to a jet produced by the pulsar (Tamuraetal.1996:: Brazier&Becker 1997))., Several authors have interpreted this region as corresponding to a jet produced by the pulsar \cite{tkyb96}; \cite{bb97}) ).
" Using PPSPC data. Greiveldinger ((1995)) have analyzed the morphology of X-ray emission close to the pulsar. and concluded that the PWN has two components: a large region extending over many areminutes. and a compact disc of emission of radius 30""—45"" centered on the pulsar."," Using PSPC data, Greiveldinger \nocite{gcm+95}) ) have analyzed the morphology of X-ray emission close to the pulsar, and concluded that the PWN has two components: a large region extending over many arcminutes, and a compact disc of emission of radius $30''-45''$ centered on the pulsar."
 Brazier Becker (1997)) considered data taken with the HHRI., Brazier Becker \nocite{bb97}) ) considered data taken with the HRI.
 They found no evidence for the component PWN proposed by Greiveldinger ((1995)). but pointed out that the X-ray emission around. the pulsar had à cross-shaped morphology.," They found no evidence for the two-component PWN proposed by Greiveldinger \nocite{gcm+95}) ), but pointed out that the X-ray emission around the pulsar had a cross-shaped morphology."
 They interpreted this as indicating that the pulsar powers an equatorial torus. as 1s also seen in X-rays for the Crab pulsar (Aschenbach&Brinkmann 1975:: Hesteretal. 1995)).," They interpreted this as indicating that the pulsar powers an equatorial torus, as is also seen in X-rays for the Crab pulsar \cite{ab75}; \cite{hss+95}) )."
 This morphology would imply that the jet (and presumably the pulsar spin axis) lies largely in the plane of the sky. inclined at ~70° from the line-of-sight.," This morphology would imply that the jet (and presumably the pulsar spin axis) lies largely in the plane of the sky, inclined at $\sim70^\circ$ from the line-of-sight."
 An interpretation for the thermal X-ray source ~6/ to the north of the pulsar is unclear., An interpretation for the thermal X-ray source $\sim6'$ to the north of the pulsar is unclear.
 It hàs been repeatedly proposed that this region represents an interaction between a collimated outflow from the pulsar and the surrounding SNR (Seward 1983:: Manchester&Durdin 1983:: Manchester 1987:; &Becker 1997))., It has been repeatedly proposed that this region represents an interaction between a collimated outflow from the pulsar and the surrounding SNR \cite{shmc83}; \cite{md83}; \cite{man87}; \cite{bb97}) ).
 Tamura ((1996)) showed evidence for a bridge of non-thermal emission connecting the pulsar to the thermal X-ray emission to the north. but it remained uncertain as to whether this was the direct counterpart to the SE jet-like feature.," Tamura \nocite{tkyb96}) ) showed evidence for a bridge of non-thermal emission connecting the pulsar to the thermal X-ray emission to the north, but it remained uncertain as to whether this was the direct counterpart to the SE jet-like feature."
 At high spatial resolution. a collection of X-ray clumps is seen at the center of the northern X-ray region. the morphologies and positions of which closely match those of clumps seen at radio wavelengths (Brazier&Becker 1997:; G99).," At high spatial resolution, a collection of X-ray clumps is seen at the center of the northern X-ray region, the morphologies and positions of which closely match those of clumps seen at radio wavelengths \cite{bb97}; G99)."
 It has not been established whether the proposed pulsar outflow powers the entire rregion (Tamuraetal. 1996)) or just the central collection of clumps (G99)., It has not been established whether the proposed pulsar outflow powers the entire region \cite{tkyb96}) ) or just the central collection of clumps (G99).
 But in either case. the evidence for interaction provides a convincing argument that the pulsar and SNR are physically associated.," But in either case, the evidence for interaction provides a convincing argument that the pulsar and SNR are physically associated."
 Clearly PSR παπά SNR pprovide an important example of the interaction between a pulsar and its environment., Clearly PSR and SNR provide an important example of the interaction between a pulsar and its environment.
 Accordingly. we have obtained sensitive high-resolution observations of this system. with theObservatory.," Accordingly, we have obtained sensitive high-resolution observations of this system with the."
 These data provide the opportunity to clarify the nature of the emission close to the pulsar. to determine the overall structure of the PWN. and to further investigate the apparent interaction between the pulsar and its surrounding SNR.," These data provide the opportunity to clarify the nature of the emission close to the pulsar, to determine the overall structure of the PWN, and to further investigate the apparent interaction between the pulsar and its surrounding SNR."
" In ine pectra, andine refsecgiscusswediscussandinterprettheseresults,"," In \\ref{sec_obs} we describe our observations, in \\ref{sec_results} we describe the resulting images and spectra, and in \\ref{sec_discuss}
 we discuss and interpret these results."
 wwas observed on 2000 Aug 14 (ObsID 754) with the Advanced CCD Imaging Spectrometer (ACIS) aboard the ((Weisskopfetal.2000b))., was observed on 2000 Aug 14 (ObsID 754) with the Advanced CCD Imaging Spectrometer (ACIS) aboard the \cite{wtvo00}) ).
 ACIS is an array of ten « 1024-pixel CCDs fabricated by MIT Lincoln Laboratory., ACIS is an array of ten $\times$ 1024-pixel CCDs fabricated by MIT Lincoln Laboratory.
 These X-ray sensitive devices have large detection efficiency (10-90% )) and moderate energy resolution (10-50) over a 0.2-10.0 keV passband., These X-ray sensitive devices have large detection efficiency $-$ ) and moderate energy resolution $-$ 50) over a 0.2–10.0 keV passband.
 Coupled with the 10-m optics ofChandra.. the CCDs have a scale of 07492 per pixel. well-matched to the on-axis point-spread function (FWHM 1).," Coupled with the 10-m optics of, the CCDs have a scale of $0\farcs492$ per pixel, well-matched to the on-axis point-spread function (FWHM $\la1''$ )."
 ACIS and its calibration are discussed in detail by Burke ((1997)) and Bautz ((1998))., ACIS and its calibration are discussed in detail by Burke \nocite{bgb+97}) ) and Bautz \nocite{bpb+98}) ).
" A single exposure of length. 20 ks was made in the standard “Timed Exposure"" mode.", A single exposure of length 20 ks was made in the standard “Timed Exposure” mode.
 The data discussed here were recorded from the four CCDs of the ACIS-I. array., The data discussed here were recorded from the four CCDs of the ACIS-I array.
 After standard processing had been carried out at the Chandra Science Center (ASCDS version number RACUSUPDI4.1). we analyzed the resultant events list using CIAO v2.1.3.," After standard processing had been carried out at the Chandra Science Center (ASCDS version number R4CU5UPD14.1), we analyzed the resultant events list using CIAO v2.1.3."
" We first applied the observation-specific list of bad pixels to the data. and then corrected for the ~1"" offset present in ACIS-Ivations!!.. After removing small intervals in which data were not recorded. the final exposure time for this observation was 19039 s. The effective area of the detector is a function of both the energy of incident photons and their position on the sky."," We first applied the observation-specific list of bad pixels to the data, and then corrected for the $\sim1''$ offset present in ACIS-I. After removing small intervals in which data were not recorded, the final exposure time for this observation was 19039 s. The effective area of the detector is a function of both the energy of incident photons and their position on the sky."
 To generate an exposure-corrected image. we computed exposure maps at five energies: 0.5. 1.5. 3.0. 5.0 and 7.0 keV. We also computed images of the counts per pixel in the five separate energy bands 0.3-0.8. 0.8—2.0. 2.0-4.0. 4.0-6.0 and 6.0-8.0 keV. For each energy band. we divided the map of raw counts by the exposure map corresponding to that band. and then normalized by the exposure time to produce a final image with units of photons em s.," To generate an exposure-corrected image, we computed exposure maps at five energies: 0.5, 1.5, 3.0, 5.0 and 7.0 keV. We also computed images of the counts per pixel in the five separate energy bands 0.3–0.8, 0.8–2.0, 2.0–4.0, 4.0–6.0 and 6.0–8.0 keV. For each energy band, we divided the map of raw counts by the exposure map corresponding to that band, and then normalized by the exposure time to produce a final image with units of photons $^{-2}$ $^{-1}$."
 For compact sources of emission (<5” in extent). spectra were extracted from the data using the sseript within CIAO. using a small extraction region. immediately surrounding the source of interest.," For compact sources of emission $\le5''$ in extent), spectra were extracted from the data using the script within CIAO, using a small extraction region immediately surrounding the source of interest."
 For extended sources we made use of v1.08 of the and provided by Alexey Vikhlinin. which generate the response matrix and effective area for a region of the CCD by computing an appropriately weighted mean over small sub-regions.," For extended sources we made use of v1.08 of the and provided by Alexey Vikhlinin, which generate the response matrix and effective area for a region of the CCD by computing an appropriately weighted mean over small sub-regions."
 In all cases. the resulting spectra were rebinned to ensure a minimum of 15 counts per spectral channel.," In all cases, the resulting spectra were rebinned to ensure a minimum of 15 counts per spectral channel."
 The background spectrum for each region was measured by considering an annular region, The background spectrum for each region was measured by considering an annular region
hereafter RLWO2) we identified the most probable optical IDs for 50 out of 51 fF(s.5Cllz)=7 900 µν sources. detected in three deep VLA surveys.,"hereafter RLK02) we identified the most probable optical IDs for 50 out of 51 $F(8.5~\rm GHz)=7$ –900 $\rm \mu Jy$ sources, detected in three deep VLA surveys."
 These. galaxies were imaged in BRL using the Ixeck Low Resolution Imaging Spectrograph (LIIS). and in fy with the Near Infra-red Camera (NIC).," These galaxies were imaged in $BRI$ using the Keck Low Resolution Imaging Spectrograph (LRIS), and in $K$ with the Near Infra-red Camera (NIRC)."
 Spectra and redshifts were obtained. for lY. using LRIS in spectroscopic mode.," Spectra and redshifts were obtained for 17, using LRIS in spectroscopic mode."
 Of these. Lt were founcl to be emission-line galaxies at redshifts up to zI. and the other 3 were QSOs at zz1.," Of these, 14 were found to be emission-line galaxies at redshifts up to $z\sim 1$, and the other 3 were QSOs at $z>1$."
 A further 9 of the radio LD galaxies had known redshifts from. previous observations. giving à total of 26 with known recdshifts.," A further 9 of the radio ID galaxies had known redshifts from previous observations, giving a total of 26 with known redshifts."
 1111909 concluded that more than half of these 26 were disk ealaxies with enhanced radio luminosities. resulting [roni major starbursts (SE~1OOAL.vr +).," RLK02 concluded that more than half of these 26 were disk galaxies with enhanced radio luminosities, resulting from major starbursts $\rm SFR\sim 100 M_{\odot}yr^{-1}$ )."
 The remainder were apparently normal. non-starburst spirals and. ellipticals at lower redshifts (2< 0.4) (~20 per cent). QSOs (~15 per cent). or giant raciioluminous cllipticals suspected to contain obsceured AGN (ο8 per cent).," The remainder were apparently normal, non-starburst spirals and ellipticals at lower redshifts $z<0.4$ ) $\sim 20$ per cent), QSOs $\sim 15$ per cent), or giant radioluminous ellipticals suspected to contain obscured AGN $\sim 8$ per cent)."
 The LRIS spectra of the 14. non-QSO ealaxics showed the expected. emission lines. ΟΠΗΑΤΟΤΑ. 13 anc OLLS007A. but the line luminosities typically correspon to SERs an order of magnitude lower than the racio Iuminosities imply.," The LRIS spectra of the 14 non-QSO galaxies showed the expected emission lines, $\rm [OII]3727\AA$, $H\beta$ and $\rm [OIII]5007\AA$, but the line luminosities typically correspond to SFRs an order of magnitude lower than the radio luminosities imply."
 Other surveys of raclio-sclectecl galaxies have found a similar discrepancy between emission-line anc radio [luxes (e.g. Smith et al., Other surveys of radio-selected galaxies have found a similar discrepancy between emission-line and radio fluxes (e.g. Smith et al.
 1996: Beck. Furner and Ixovo 2000: Serjeant et al.," 1996; Beck, Turner and Kovo 2000; Serjeant et al."
 2000)., 2000).
 The simplest explanation is à high dust extinction. (elyo2 3 mag) of the starburs regions., The simplest explanation is a high dust extinction $A_V\simeq 2$ –3 mag) of the starburst regions.
 We also found that 11 of the 26 galaxies were in close pairs and several others looked. cisturbed., We also found that 11 of the 26 galaxies were in close pairs and several others looked disturbed.
 However. it was clear that to adequately investigate the interaction status of these galaxies. much higher resolution imaging. e.& WEPC2. would be required.," However, it was clear that to adequately investigate the interaction status of these galaxies, much higher resolution imaging, e.g WFPC2, would be required."
 Serjeant et al. (, Serjeant et al. (
2000) performed an WEDPC2 imaging survey of raclio-selected star-forming galaxies. ancl describe the first four.,"2000) performed an WFPC2 imaging survey of radio-selected star-forming galaxies, and describe the first four."
 These galaxies are at lower recdshifts. 2o0.2. than the RLWO2 sample but. have. comparable radio luminosites.," These galaxies are at lower redshifts, $z\sim 0.2$, than the RLK02 sample but have comparable radio luminosites."
" They found at least one galaxy to be interacting ancl all 4 to be disturbed to some degree. ane quantified morphological disturbance in terms ofa rotationa asvmumetrv parameter zl,s,, (Consclice. Bershacky ane Jangren 2000)."," They found at least one galaxy to be interacting and all 4 to be disturbed to some degree, and quantified morphological disturbance in terms of a rotational asymmetry parameter $A_{asym}$ (Conselice, Bershady and Jangren 2000)."
 They concluded that racio-selected galaxies are generally more asymmetric than. opticallv-selectec ealaxies at similar magnitudes., They concluded that radio-selected galaxies are generally more asymmetric than optically-selected galaxies at similar magnitudes.
 Similarly. and at higher redshifts. Fomatont et al. (," Similarly, and at higher redshifts, Fomalont et al. ("
2002). identified 37/63 fain ](8.561111)c1.55Jv radio sources with /x23.3 galaxies. and using WEPC2 imaging found a high proportion. 46 per cent. to be multiple or interacting.,"2002) identified 37/63 faint $\rm F(8.5~GHz)\geq7.5\mu \rm Jy$ radio sources with $I\leq 23.3$ galaxies, and using WFPC2 imaging found a high proportion, 46 per cent, to be multiple or interacting."
 In this paper we perform a similar WEPC? study of a subsample of the IRLINO2 radio IDs., In this paper we perform a similar WFPC2 study of a subsample of the RLK02 radio IDs.
 Our LST program was allocated a total of 24 orbits. which was devoted to ΑΕΡΟΣ I-band imaging of four fields. containing a total of 9 of the 11903 galaxies.," Our HST program was allocated a total of 24 orbits, which was devoted to WFPC2 $I$ -band imaging of four fields, containing a total of 9 of the RLK02 galaxies."
 For the first two fields we imaged in parallel with STIS. and obtained images in the broad STIS V-band of two of the galaxies.," For the first two fields we imaged in parallel with STIS, and obtained images in the broad STIS $V$ -band of two of the galaxies."
 Section 2 of this paper describes the observations ancl data reduction., Section 2 of this paper describes the observations and data reduction.
 Section 3. catalogs the observed galaxies and presents LIST images., Section 3 catalogs the observed galaxies and presents HST images.
 In Section 4. radial intensity profiles are presented. model profiles fitted. ancl surface brightness compared with local galaxies.," In Section 4, radial intensity profiles are presented, model profiles fitted, and surface brightness compared with local galaxies."
 In Section 5 the evidence for interactions is investigated., In Section 5 the evidence for interactions is investigated.
" Asvnimetry parameters are evaluated ancl model. profiles subtracted from the galaxies to highlight ""residual. features.", Asymmetry parameters are evaluated and model profiles subtracted from the galaxies to highlight `residual' features.
 In Section 6. galaxy colours are interpreted.," In Section 6, galaxy colours are interpreted."
 Section 7 is a discussion of the nature of these galaxies and the source of their racio luminosity., Section 7 is a discussion of the nature of these galaxies and the source of their radio luminosity.
 The galaxies studied here are an unbiased: subsaniple of the galaxies identificd by RLINO2 on Ixeck LIIS images as probable optical counterparts of the radio. detections on two VLA surveys., The galaxies studied here are an unbiased subsample of the galaxies identified by RLK02 on Keck LRIS images as probable optical counterparts of the radio detections on two VLA surveys.
" The first survey is of the SAGS Ποια. centred at IA. 23""59'157. Dee 14:55:00 (equinox 2000.0 used throughout). in which observations at 5.0 GllIz detected sources to integrated [Lux limits of 60/5]v. or 40. 455v at 8.44 Gllz (Weistrop et al. ("," The first survey is of the SA68 field, centred at R.A. $23^h 59^m
15^s$ Dec 14:55:00 (equinox 2000.0 used throughout), in which observations at 5.0 GHz detected sources to integrated flux limits of $60\rm \mu Jy$, or $\sim 40$ $45\rm \mu Jy$ at 8.44 GHz (Weistrop et al. ("
1987).,1987).
 The WEPC2 field of view covers only one of these sources at a timo., The WFPC2 field of view covers only one of these sources at a time.
" The second survey. smaller in area but much deeper. is ofthe ‘Lynx?’ or LGW field IN. OS’45""04"" Doe 44:34:05). observed at S44 CGllz for 68 hours in December 1989 ane January 1990."," The second survey, smaller in area but much deeper, is of the `Lynx2' or `16V' field ( R.A. $08^h 45^m 04^s$ Dec 44:34:05), observed at 8.44 GHz for 63 hours in December 1989 and January 1990."
 This gave a noise level of e=3.21yJv/beam and a point-source completeness limit at (8.44Cilz)=14.55. — Windhorst et al. (, This gave a noise level of $\sigma=3.21\rm \mu Jy/beam $ and a point-source completeness limit at $F(8.44~{\rm GHz})\simeq 14.5\rm \mu Jy$ – Windhorst et al. (
1993) catalog the detections.,1993) catalog the detections.
 A WEDPC? Ποιά may include several of these fainter sources., A WFPC2 field may include several of these fainter sources.
 The UST observations. in. 14 time slots from September 2000 to September 2001. consisted of two pointings within the SAGS VLA survey. (Weistrop et al.," The HST observations, in 14 time slots from September 2000 to September 2001, consisted of two pointings within the SA68 VLA survey (Weistrop et al."
 1987) ancl two in the Lvnx2/16V. (Windhorst et., 1987) and two in the Lynx2/16V (Windhorst et.
 al 1993) field., al 1993) field.
 Our program originally called for. observations with both WEPC2 and STIS. but the temporary failure of STIS partway through our observations meant that we used only WEDPC?2 for the targets in the Lynx? field.," Our program originally called for observations with both WFPC2 and STIS, but the temporary failure of STIS partway through our observations meant that we used only WFPC2 for the targets in the Lynx2 field."
 Imaging was performed in four telescope pointings as. described below. (, Imaging was performed in four telescope pointings as 	described below. (
"i) The optical ID. for the radio source SAGS:10 (hereafter. ""SIO) was centered in the WEPC2 and imaged in the /-band |CESIAW) for 5 orbits. giving LO exposures of 1200s or 13008. for a total of 12.4 Kksec.","i) The optical ID for the radio source SA68:10 (hereafter 	`S10') was centered in the WFPC2 and imaged in the $I$ -band 	(F814W) for 5 orbits, giving 10 exposures of 1200s or 1300s 	for a total of 12.4 ksec."
 Each orbit. was split into two exposures for cosmic-ray rejection. and the telescope was. dithered a few arcseconds between successive orbits to improve [at-fielding and bad pixel rejection.," Each orbit was split into two 	exposures for cosmic-ray rejection, and the telescope was 	dithered a few arcseconds between successive orbits to 	improve flat-fielding and bad pixel rejection."
" STIS. operating in parallel. simultaneously imaged another ealaxy from the same survey. SA68:4 (hereafter ""S4). for 5 exposures οἱ 400 sec each totalling 2000 sec in a wide V -band filter. Cmirvis)."," STIS, 	operating in parallel, simultaneously imaged another galaxy 	from the same survey, SA68:4 (hereafter `S4'), for 5 exposures 	of 400 sec each totalling 2000 sec in a wide $V$ -band filter 	('mirvis')."
 For the remainder of the available time at that. pointing. STIS was used in spectroscopic mode. but the resulting spectra were of poor signal-to-noise and will not be cliseussecl here further. (," For the remainder of the available time at that 	pointing, STIS was used in spectroscopic mode, but the 	resulting spectra were of poor signal-to-noise and will not be discussed 	here further. ("
ii) Phe telescope was rotated 150 degrees at the same as in (1) to switch instruments. so that the WETC2 was centered on S4 and STIS on SLO.,"ii) The telescope was rotated 180 degrees at the same 	pointing as in (i) to switch instruments, so that the WFPC2 	was centered on S4 and STIS on S10."
 Again we obtained LO, Again we obtained 10
Array (SIXKA‘)) with collecting area LO?m= and a FoV much greater than Ideg?. with these specifications the SILA surveys niav detect 10 galaxies.,"Array ) with collecting area $\sim 10^{6}\,{\rm m}^2$ and a FoV much greater than $1\,{\rm deg}^2$, with these specifications the SKA surveys may detect $\sim 10^{9}$ galaxies."
 Before the full SIX comes into operation (2020) another large radio telescope will have been constructed which also has great potential for HE surveys., Before the full SKA comes into operation $\sim 2020$ ) another large radio telescope will have been constructed which also has great potential for HI surveys.
 The Five-huncdred-ametre Aperture Spherical Telescope (PAST) will be an AXrecibo-like telescope with a significantly larger aperture ( 500m) and FoV (Nan2006)..," The Five-hundred-metre Aperture Spherical Telescope (FAST) will be an Arecibo-like telescope with a significantly larger aperture $\sim 500\,{\rm m}$ ) and FoV \citep{KARST}."
 Phe aim of this paper is to discuss the scientifie potential of a low-recshift LUE survey of galaxies asa PAST Κον science programme., The aim of this paper is to discuss the scientific potential of a low-redshift HI survey of galaxies as a FAST key science programme.
 We will show that. if a focal plane array with 100 instantaneous beams can be developed. ~10' galaxies can be detected. within a realistic survey period. (2 vears) and that the. [ow-redshift galaxy power spectrum can be measured much more accurately than is currently possible with optical surveys.," We will show that, if a focal plane array with 100 instantaneous beams can be developed, $\sim 10^{7}$ galaxies can be detected within a realistic survey period $\sim 2$ years) and that the low-redshift galaxy power spectrum can be measured much more accurately than is currently possible with optical surveys."
 Since hvdrogen is the fundamental barvonic constituent of the universe such a survey would provide an important. view of the large-scale structure of the universe with dilferent biases and be an important stepping stone to the larger LIL galaxy surveys which will become possible with the SIA., Since hydrogen is the fundamental baryonic constituent of the universe such a survey would provide an important view of the large-scale structure of the universe with different biases and be an important stepping stone to the larger HI galaxy surveys which will become possible with the SKA.
 In the next section we will brielly discuss some relevant observational parameters of FAST and then in subsequent sections. we will estimate the number counts. make forecasts for errors on the power spectrum ancl constraints on cosmological parameters that could. be possible with such @ SULVON.," In the next section we will briefly discuss some relevant observational parameters of FAST and then in subsequent sections, we will estimate the number counts, make forecasts for errors on the power spectrum and constraints on cosmological parameters that could be possible with such a survey."
 The LEAST telescope is. funded. and when completed in around 2012-2013 it will be the largest |single clish telescope in the world., The FAST telescope is funded and when completed in around 2012-2013 it will be the largest single dish telescope in the world.
 The expected. specifications of the telescope (Nan2006) relevant to our discussion here are: The expected. thermal noise for adual polarization single beam can be computed using for an observing time of / and a frequenes bancwiclt of Av. where &=1380JvmIKLis the Boltzmann constant.," The expected specifications of the telescope \citep{KARST} relevant to our discussion here are: The expected thermal noise for adual polarization single beam can be computed using for an observing time of $t$ and a frequency bandwidth of $\bw$, where $k=1380\,{\rm Jy}\,{\rm m}^2\,{\rm K}^{-1}$ is the Boltzmann constant."
 Lf we assume a bandwidth of Av=1MllIz which corresponds to à velocity linewidth of zz2O0kms at 2lem then the instantaneous sensitivity of each beam of the FAST system will be ejua&Lindyst?," If we assume a bandwidth of $\Delta\nu=1\,{\rm MHz}$ which corresponds to a velocity linewidth of $\approx 200\,{\rm km}\,s^{-1}$ at 21cm then the instantaneous sensitivity of each beam of the FAST system will be $\sigma_{\rm inst}\approx 1\,{\rm mJy}\,s^{1/2}$."
 The focal ratio of the telescope (ff=0.47) restricts np to x19 for a conventional close-packed horn-based array. similar to that used for LUPASS and operated at the prime [ocus.," The focal ratio of the telescope $f/D=0.47$ ) restricts $n_{\rm B}$ to $\le 19$ for a conventional close-packed horn-based array, similar to that used for HIPASS and operated at the prime focus."
 Larger arravs require some of the horns to be far enough oll-axis that aberration losses become unacceptable., Larger arrays require some of the horns to be far enough off-axis that aberration losses become unacceptable.
" A l19-beam array gives an instantaneous FOV of 3,42200aremin?."," A 19-beam array gives an instantaneous FoV of $\Omega_{\rm inst}\approx 200\,{\rm arcmin}^2$ ."
 We will use this I9-beam system as the fiducial benchmark for our basic caleulations., We will use this 19-beam system as the fiducial benchmark for our basic calculations.
 A receiver array with np100 may well be possible., A receiver array with $n_{\rm B}\sim 100$ may well be possible.
 This requirement could only be met with close-packed phased arrays. of small antenna elements: these would not take up much more area than the 19-beam horn system., This requirement could only be met with close-packed phased arrays of small antenna elements; these would not take up much more area than the 19-beam horn system.
 The required phased array technology is actively being developed by several groups working within the international Square Kilometre Array R&DD and hence we can be confident that by the time EAST comes into operation in 2012-3 the I00-beam receiver capability will be available., The required phased array technology is actively being developed by several groups working within the international Square Kilometre Array D and hence we can be confident that by the time FAST comes into operation in 2012-3 the 100-beam receiver capability will be available.
" ""his would allow for Oa&(ng/19)200aremin."," This would allow for $\Omega_{\rm inst}\approx (n_{\rm B}/19)200\,{\rm arcmin}^2$."
 Such a system will be at least twice as sensitive per beam as that on the recibo Telescope and have a survey. speed (Quia(ζωεν 1} whieh is at least 6605/19) times faster than the 7-beam system presently available there., Such a system will be at least twice as sensitive per beam as that on the Arecibo Telescope and have a survey speed $\Omega_{\rm inst}(A_{\rm eff}/T_{\rm sys})^2$ ) which is at least $6(n_{\rm B}/19)$ times faster than the 7-beam system presently available there.
 EAST will also have more frequency flexibility due to the excellent REL enviroment at the PAST site anc will be able to observe twice the area of the sky than does the recibo Telescope., FAST will also have more frequency flexibility due to the excellent RFI enviroment at the FAST site and will be able to observe twice the area of the sky than does the Arecibo Telescope.
 The HE mass. Mg. of a galaxy at redshift z is given in terms of the observed Dux. S. and line width. AV. bv (Roberts where di(z) is the luminosity distance to the galaxy.," The HI mass, $M_{\rm HI}$, of a galaxy at redshift $z$ is given in terms of the observed flux, $S$ , and line width, $\Delta V_{\rm o}$, by \citep{Roberts}
 where $d_{\rm L}(z)$ is the luminosity distance to the galaxy."
 We note that this formula is usually quoted in terms of the [Euclidean distance and without the (1|2)ii this is the correct formula in an EIU universe., We note that this formula is usually quoted in terms of the Euclidean distance and without the $(1+z)^{-1}$; this is the correct formula in an FRW universe.
 One can estimate the number of galaxies above some limiting mass Mi(2) bv computing where the sky area covered is AQ and the size of the redshift bin is Ac., One can estimate the number of galaxies above some limiting mass $M_{\rm lim}(z)$ by computing where the sky area covered is $\Delta\Omega$ and the size of the redshift bin is $\Delta z$ .
 dVατα is the comoving volume clement for the ERN universe and δαναλ is the comoving number density. ofgalaxies per unit mass which we shall assume can be approximated by a Schechter function, $dV/dzd\Omega$ is the comoving volume element for the FRW universe and $dN/dVdM$ is the comoving number density ofgalaxies per unit mass which we shall assume can be approximated by a Schechter function
this Poy regime.,this $_{\rm eff}$ regime.
 We note that a dip in the LE observec in the sequences of both open clusters and. the field a Ale&7 (the Wielen gap: c.g Wielen. Jahreiss Kruger 1983: Delikov at al.," We note that a dip in the LF observed in the sequences of both open clusters and the field at $_{V}\approx7$ (the Wielen gap; e.g Wielen, Jahreiss Kruger 1983; Belikov at al."
 1998). has been attributed to a loca steepening of the Iuminosity-mass relation due to strong opacity in the atmospheres of objects in this T regime (Ixroupa. Tout. Gilmore 1990).," 1998), has been attributed to a local steepening of the luminosity-mass relation due to strong $^{-}$ opacity in the atmospheres of objects in this $_{\rm eff}$ regime (Kroupa, Tout Gilmore 1990)."
" Theoretical coo model atmospheres predict. that dust. grains of AlsO,. be and MeSsiO, begin to condense in the outermost layers of a brown cdwarl (or very low mass star) at T150019. corresponding to a ‘Pa consistent with the AIT-ALS LE dip (e.g Allard et al."," Theoretical cool model atmospheres predict that dust grains of $_{2}$ $_{3}$, Fe and $_{4}$ begin to condense in the outermost layers of a brown dwarf (or very low mass star) at $\sim1800$ K, corresponding to a $_{\rm
eff}$ consistent with the M7-M8 LF dip (e.g Allard et al."
 2001. Tsuji 2001).," 2001, Tsuji 2001)."
 Indeed. spectroscopic observations of late-M. field chwarls lend. support to these predictions (Jones Vsuji 1997).," Indeed, spectroscopic observations of late-M field dwarfs lend support to these predictions (Jones Tsuji 1997)."
 However. the DUSTY evolutionary models of Chabrier et al. (," However, the DUSTY evolutionary models of Chabrier et al. ("
2000) which include a treatment for the formation of dust. do not. predict any discontinuitv in the [uminositv-niass relation here.,"2000) which include a treatment for the formation of dust, do not predict any discontinuity in the luminosity-mass relation here."
 In the synthetic atmospheres of Allard. et al.," In the synthetic atmospheres of Allard et al.,"
 which provide the outer boundary conditions for the DUSTY calculations. cust is treated as erains with sizes in the range 0.00625— 0.2454m. In this case. dust does not contribute significantly to the opacities in the near-infrared where the water hands are dominant. (Allarc et al.," which provide the outer boundary conditions for the DUSTY calculations, dust is treated as grains with sizes in the range $0.00625-0.24\mu$ m. In this case, dust does not contribute significantly to the opacities in the near-infrared where the water bands are dominant (Allard et al."
 2001)., 2001).
 However. as illustrated by Figure 4 of Allard ct al.," However, as illustrated by Figure 4 of Allard et al."
 if the &rain size is increased. into the regime described by Mie scattering theory. the opacity provided: by dust. is dramatically enhanced. (we note that the opacity of dust. increases with decreasing wavelength).," if the grain size is increased into the regime described by Mie scattering theory, the opacity provided by dust is dramatically enhanced (we note that the opacity of dust increases with decreasing wavelength)."
 While current observations argue against &rains as large as l0074n. the most detailed: current cloud. formation models predict sizes in the range 5.— 205m (Cooper et al.," While current observations argue against grains as large as $100\mu$ m, the most detailed current cloud formation models predict sizes in the range $5-20\mu$ m (Cooper et al."
 2002)., 2002).
 Lt might well be possible to reproduce the observed dips in the LEs with a moderate increase in the model grain size., It might well be possible to reproduce the observed dips in the LFs with a moderate increase in the model grain size.
 With the recent. advent of wide-field near-infrared imagers (ee ΕΕΑΛΛΗΝΟ) and the ever. increasing availability. of wice-ficlel optical imagers (c.g WEIL) the results of deeper. larger area surveys for low mass nicmbers of rich. voung clusters such as NGC2547. NGC2516 and 1€2391. should. soon become available.," With the recent advent of wide-field near-infrared imagers (e.g FLAMINGOS) and the ever increasing availability of wide-field optical imagers (e.g WFI) the results of deeper, larger area surveys for low mass members of rich young clusters such as NGC2547, NGC2516 and IC2391 should soon become available."
 Phe improved statistics of these surveys will confirm or otherwise the reported dip in the LE between ALT-AMIS ancl allow a more thorough examination of its implications for mass determinations at Tr£ 2700E. 1., The improved statistics of these surveys will confirm or otherwise the reported dip in the LF between M7-M8 and allow a more thorough examination of its implications for mass determinations at $_{\rm eff}\simless2700$ K. 1.
 We have presented evidence for a deficit of very-Low. mass stars/brown dwarfs between spectral types M7-MS and have argued that this arises from a sharp local change in the shape of the luminosity-mass relation., We have presented evidence for a deficit of very-low mass stars/brown dwarfs between spectral types M7-M8 and have argued that this arises from a sharp local change in the shape of the luminosity-mass relation.
 2., 2.
 Structure near the bottom end of the Trapezium LAL recently proposed by Muench et al. (, Structure near the bottom end of the Trapezium IMF recently proposed by Muench et al. (
2002) is more likely a manifestation of the feature in the luminositv-mass relation reported here.,2002) is more likely a manifestation of the feature in the luminosity-mass relation reported here.
 3., 3.
 We find that brown cdwarfs with spectral types later than ADT-MS may have masses greater than predicted by current theoretical models., We find that brown dwarfs with spectral types later than M7-M8 may have masses greater than predicted by current theoretical models.
 4., 4.
 We speculate that the onset of dust formation in the upper atmosphere of objects in this Tir regime is responsible for the change in the shape of the Iuminositv-mass relation., We speculate that the onset of dust formation in the upper atmosphere of objects in this $_{\rm eff}$ regime is responsible for the change in the shape of the luminosity-mass relation.
 This might be mocdelled by assuming larger grain sizes than are currently used in the DUSTY calculations., This might be modelled by assuming larger grain sizes than are currently used in the DUSTY calculations.
 PDD. DIP and STL acknowledge the financial support of PPARC.," PDD, DJP and STH acknowledge the financial support of PPARC."
 We thank Gilles Chabrier ancl Isabelle: Daralfe for useful discussions on their evolutionary models. David Doarrado ν Navascudieis for making his latest results available to us and the referee. Nigel Hambly. for a swift ancl useful response.," We thank Gilles Chabrier and Isabelle Baraffe for useful discussions on their evolutionary models, David Barrado y Navascués for making his latest results available to us and the referee, Nigel Hambly, for a swift and useful response."
"refsec:discuss,, that the CO gas is located in a ring around the star which has constant temperature and column density.",", that the CO gas is located in a ring around the star which has constant temperature and column density."
" These assumptions are reasonable since the hottest CO component, which arises close to the star where usually the disk density is highest, dominates the total CO spectrum completely."," These assumptions are reasonable since the hottest CO component, which arises close to the star where usually the disk density is highest, dominates the total CO spectrum completely."
 For details on the CO band modeling we refer to Krausetal.(2000) and Kraus(2009).," For details on the CO band modeling we refer to \citet{Kraus00}
 and \citet{Kraus09}."
". To constrain the range of the !?C/!?C abundance ratio in the disk, we took expected !?C/'?C surface abundance ratios from the stellar evolution models of Meynet&Maeder(2005) for massive stars with LMC metallicity and initial rotation velocities of ss!"," To constrain the range of the $^{12}$ $^{13}$ C abundance ratio in the disk, we took expected $^{12}$ $^{13}$ C surface abundance ratios from the stellar evolution models of \citet{MeynetMaeder05}
 for massive stars with LMC metallicity and initial rotation velocities of $^{-1}$."
 reffig:ratio shows. how this ratio decreases with time for stars of different initial mass., \\ref{fig:ratio} shows how this ratio decreases with time for stars of different initial mass.
 The stars 112 and 773 have evolved from progenitors with Mis;=25...30 (see reftab:parameters)) whose surface abundanceMo should have dropped to a value !?C/P?C <20 at the turn-off from the main sequence and decreases down to a value of ~4 at late evolutionary phases.," The stars 12 and 73 have evolved from progenitors with $M_{\rm ini} = 25\ldots 30\,M_{\sun}$ (see \\ref{tab:parameters}) ) whose surface abundance should have dropped to a value $^{12}$ $^{13}$ C $\la 20$ at the turn-off from the main sequence and decreases down to a value of $\sim 4$ at late evolutionary phases."
" We computed synthetic CO band spectra for temperatures in the range 2000 — KK, !?C/P?C ratios varying from 4 to 20, and CO column densities higher than 10?! ccm? to account for the optical depth effects."," We computed synthetic CO band spectra for temperatures in the range 2000 – K, $^{12}$ $^{13}$ C ratios varying from 4 to 20, and CO column densities higher than $10^{21}$ $^{-2}$ to account for the optical depth effects."
" We did not find any indication for additional line broadening from either Keplerian rotation or outflow signatures, meaning that the CO gas has no significant line-of-sight velocity component, in agreement with an (almost) pole-on orientation of the disks."," We did not find any indication for additional line broadening from either Keplerian rotation or outflow signatures, meaning that the CO gas has no significant line-of-sight velocity component, in agreement with an (almost) pole-on orientation of the disks."
" Therefore, we used the thermal velocity as the intrinsic gas velocity and convolved the resulting spectrum with the resolution of SINFONI."," Therefore, we used the thermal velocity as the intrinsic gas velocity and convolved the resulting spectrum with the resolution of SINFONI."
" Pure CO models that match the observations best are shown in red in the top panels of reffig:fits,, with the parameters listed in reftab:bestfit.."," Pure CO models that match the observations best are shown in red in the top panels of \\ref{fig:fits}, with the parameters listed in \\ref{tab:bestfit}."
" The presence of some additional emission features, especially at the red end of the spectrum, can be explained by the (if weak) recombination lines of the hydrogen Pfund series."," The presence of some additional emission features, especially at the red end of the spectrum, can be explained by the (if weak) recombination lines of the hydrogen Pfund series."
" Following the description of Krausetal.(2000), we modeled the Pfund series assuming a constant electron density and an electron temperature of T,= 10*KK. The lack of Pfund lines higher than 334 773) and 331 112) means that due to pressure ionization in a high-density gas, recombination into such high-energy states is suppressed."," Following the description of \citet{Kraus00}, we modeled the Pfund series assuming a constant electron density and an electron temperature of $T_{e} = 10^{4}$ K. The lack of Pfund lines higher than 34 73) and 31 12) means that due to pressure ionization in a high-density gas, recombination into such high-energy states is suppressed."
" This is in good agreement with the electron densities of neZ10? ccm? needed for the lines to be formed in LTE conditions, as is obvious from their intensity ratios."," This is in good agreement with the electron densities of $n_{e} \ga 
10^{9}$ $^{-3}$ needed for the lines to be formed in LTE conditions, as is obvious from their intensity ratios."
" As in the case of the CO bands, no additional line broadening is required to match the width of the Pfund lines."," As in the case of the CO bands, no additional line broadening is required to match the width of the Pfund lines."
" This could indicate that the Pfund emission is generated either in the disk but much closer to the star than the CO bands, or in the polar wind regions so close to the star that the wind velocity is still very small compared to the velocity resolution of SINFONI."," This could indicate that the Pfund emission is generated either in the disk but much closer to the star than the CO bands, or in the polar wind regions so close to the star that the wind velocity is still very small compared to the velocity resolution of SINFONI."
" The resulting synthetic Pfund series are included as the blue lines in the top panels of reffig:fits,, and the combined model spectra of CO bands and Pfund"," The resulting synthetic Pfund series are included as the blue lines in the top panels of \\ref{fig:fits}, , and the combined model spectra of CO bands and Pfund"
discussed. above is too uch to explain the observed + process abuudances in metal-poor stars.,discussed above is too much to explain the observed $r$ -process abundances in metal-poor stars.
 This can be secu by considering mixing of the ejecta from cach eveut with the ISM., This can be seen by considering mixing of the ejecta from each event with the ISM.
 Rosswoe et al (, Rosswog et al. (
1999) showed that the total energv of the NSAL ejecta is at most comparable to the SN explosion cucrey (1075 ere).,1999) showed that the total energy of the NSM ejecta is at most comparable to the SN explosion energy $\sim 10^{51}$ erg).
 Therefore. when its energv/Amonmoeutuni is dispersed in the ISM. this ejecta can wun with at most the same amount of material as swept up by a SN vemmant. Maisστον101AL. (ee. Thornton et al.," Therefore, when its energy/momentum is dispersed in the ISM, this ejecta can mix with at most the same amount of material as swept up by a SN remnant, $M_{\rm mix}\approx 3\times 10^4\,M_\odot$ (e.g., Thornton et al."
 1998)., 1998).
 Consequently. if all of this ejecta were r-process material as required to account for the present Galactic + process inventory. an ISAL enriched by a single NSM eveut would have a total i-process mass fraction The solar r-process mass fractious of clemenuts with 100cL4z130 aud A130 are ystud.fyVitexLs10 (Küpppeler ct al.," Consequently, if all of this ejecta were $r$ -process material as required to account for the present Galactic $r$ -process inventory, an ISM enriched by a single NSM event would have a total $r$ -process mass fraction The solar $r$ -process mass fractions of elements with $100\lesssim A\leq 130$ and $A>130$ are $X_{\odot,r}^{100\lesssim A\leq 130}\approx X_{\odot,r}^{A>130}\approx
4\times 10^{-8}$ (Käpppeler et al."
 1989)., 1989).
" According to equation (1)). NES, would be approximately equal NT m.NOY|uu.A130Vieteven for theto lowest AIS interest."," According to equation \ref{xnsm}) ), $X_{r,{\rm NSM}}^{\rm tot}$ would be approximately equal to $X_{\odot,r}^{\rm tot}\approx X_{\odot,r}^{100\lesssim A\leq 130} +
X_{\odot,r}^{A>130}$ even for the lowest $M_{\rm ej}^{\rm NSM}$ of interest."
" Therefore. whether the NSM ejecta were composed. of r-process clemeuts with 100 xA<130or Ao 130.or both eroups iu solar proportion. equation (19) predicts that abundance ratios of e$. Ag CY ~107) and/or Eu CY —153) with respect to livdrogen in an ISM enriched by a single event would be at least comparable to the correspondiug solar r-process values (Ae/ID).., aud (ιν."," Therefore, whether the NSM ejecta were composed of $r$ -process elements with $100\lesssim A\leq 130$ or $A>130$, or both groups in solar proportion, equation \ref{xnsm}) ) predicts that abundance ratios of e.g., Ag $A\sim 107$ ) and/or Eu $A\sim 153$ ) with respect to hydrogen in an ISM enriched by a single event would be at least comparable to the corresponding solar $r$ -process values $_{\odot,r}$ and $_{\odot,r}$."
 This predicted level of r-process euricluneut is in disagrecmeut with recent observations of r-process abundances in metal-poor stars. as Ag/IL in stars with ο ~1.3to. 2.2 are ~ 1010? times lower than (ΑΠ (Crawford et al.," This predicted level of $r$ -process enrichment is in disagreement with recent observations of $r$ -process abundances in metal-poor stars, as Ag/H in stars with ${\rm [Fe/H]}\sim -1.3$ to $-2.2$ are $\sim 10$ $10^2$ times lower than $_{\odot,r}$ (Crawford et al."
" 1998) while Eu/II in stars with |Fo/II| ~3are ~ 30101 lower than (Eu/II)..,. (AIcWilliam et al."," 1998) while Eu/H in stars with ${\rm [Fe/H]}\sim -3$ are $\sim 30$ $10^3$ lower than $_{\odot,r}$ (McWilliam et al."
 1995: Suedeu et al., 1995; Sneden et al.
 1996. 1998).," 1996, 1998)."
 The following interpretation of equation (1)) gives some insights iuto how NSMSs would explain the preseut Galactic r-process duventorv and may help appreciate why this explanation is disfavored by observations., The following interpretation of equation \ref{xnsm}) ) gives some insights into how NSMs would explain the present Galactic $r$ -process inventory and may help appreciate why this explanation is disfavored by observations.
 The Calaxy cau he divided iuto ~3«109 wnits cach laving a mass Mj73 1047...," The Galaxy can be divided into $\sim 3\times 10^6$ units each having a mass $M_{\rm mix}\approx 3\times 10^4\,M_\odot$ ."
 This division has a plivsieal meaning as Miis is the iiaxinmun mass within which the ejecta from au individual NSAL event can be distributed., This division has a physical meaning as $M_{\rm mix}$ is the maximum mass within which the ejecta from an individual NSM event can be distributed.
 Due to the rather low rate of NSAIs in the Galaxy. on average at most one NSM event occurred in a uuit over Calactic history.," Due to the rather low rate of NSMs in the Galaxy, on average at most one NSM event occurred in a unit over Galactic history."
 Therefore. in order to account for the present Galactic r-process mveutorv. cach unit would have to be enriched with an approximately solar r-process mass fraction bv a single NSM. event.," Therefore, in order to account for the present Galactic $r$ -process inventory, each unit would have to be enriched with an approximately solar $r$ -process mass fraction by a single NSM event."
 As shown above. this picture of +-process chrichment is inconsistent with the observed ;-process abundances iu uetal-poor stars.," As shown above, this picture of $r$ -process enrichment is inconsistent with the observed $r$ -process abundances in metal-poor stars."
 Furthermore. due to the high rate of SNe iu the Galaxy. niu SNe occured in a unit where a single NSM event also occurred at sometime in Galactic listory (sec 811).," Furthermore, due to the high rate of SNe in the Galaxy, many SNe occurred in a unit where a single NSM event also occurred at sometime in Galactic history (see 4)."
 As Fe euriclineut of this unit was provided bv these SNe. stars formed at different times in this unit would have varying |[Fe/T| but either zero or approximately solar r-process mass fraction if NSMs were the major rprocess sources.," As Fe enrichment of this unit was provided by these SNe, stars formed at different times in this unit would have varying [Fe/H] but either zero or approximately solar $r$ -process mass fraction if NSMs were the major $r$ -process sources."
 This is in disagreement with the observed correlation between abundances of r-process elements and Fe at [Fe/T]|FS2.5 (Caatton οποίοι 1991: Crawford et al., This is in disagreement with the observed correlation between abundances of $r$ -process elements and Fe at ${\rm [Fe/H]}\gtrsim -2.5$ (Gratton Sneden 1994; Crawford et al.
 1998: sec also McWilliam et al., 1998; see also McWilliam et al.
 1995)., 1995).
 Of course. the above discussion of mixing of tle NSM ejecta and the ISAL is oversimplified.," Of course, the above discussion of mixing of the NSM ejecta and the ISM is oversimplified."
 For example. oue could imagine that a smaller than average fraction of + Xocess ejecta was mixed into the ISAL near the edge of a NSM remmnant.," For example, one could imagine that a smaller than average fraction of $r$ -process ejecta was mixed into the ISM near the edge of a NSM remnant."
 Iu this case stars formed near the edge of he remnaut would have r-process lass fractions sinaller hau that eiven by equation (1))., In this case stars formed near the edge of the remnant would have $r$ -process mass fractions smaller than that given by equation \ref{xnsm}) ).
 However. one would also expect that less than average eurichnieut was not uuduly pervasive and a siguificaut fraction of the metal»oor stars would have the +-process enrichment indicated w equation (1)).," However, one would also expect that less than average enrichment was not unduly pervasive and a significant fraction of the metal-poor stars would have the $r$ -process enrichment indicated by equation \ref{xnsm}) )."
" The fact that no such stars lave con, observed sugeests a difficulty of the NSAL irprocess uodel in explaining the observations at low ictallicitics.", The fact that no such stars have been observed suggests a difficulty of the NSM $r$ -process model in explaining the observations at low metallicities.
 Furthermore. even if the observed r-process abundances n uctal-poor stars could be attributed to pervasive less tla average enricluueut by NSALs. one would still face the other difficulty in explaining the observed correlation between abundances of r-process eleiieuts and Feat |Fe/II|=2.5.," Furthermore, even if the observed $r$ -process abundances in metal-poor stars could be attributed to pervasive less than average enrichment by NSMs, one would still face the other difficulty in explaining the observed correlation between abundances of $r$ -process elements and Fe at ${\rm [Fe/H]}\gtrsim -2.5$."
 As Fe enrichment was controlled by SNe occurring at a much ligher rate. widely-varving deerees of mixing of the r-process ejecta in an already existing NSM reimmnaut with Fe produced by fresh SNe would result iu large scatter iun process abuudauces over a broad rauge of {Fe/TI].," As Fe enrichment was controlled by SNe occurring at a much higher rate, widely-varying degrees of mixing of the $r$ -process ejecta in an already existing NSM remnant with Fe produced by fresh SNe would result in large scatter in $r$ -process abundances over a broad range of [Fe/H]."
 This is in disagreement with the rather carly onset of the correlation between abundances of r-process elements ane Fe at |Fo/HI|z2.5., This is in disagreement with the rather early onset of the correlation between abundances of $r$ -process elements and Fe at ${\rm [Fe/H]}\approx -2.5$.
 Iu sunny. observations of iietal-poor stars would— be difficult to explain if NSMs were the Major r-process sources.," In summary, observations of metal-poor stars would be difficult to explain if NSMs were the major $r$ -process sources."
 Iu contrast to the case of NSMs. observations of poor stars as well as iieteoritie data ou ΕΤΠΕ aud. !IP?T iu the early solar svsteni support a self-consisteut picture of rprocess enriclineut by SNe (Qian Wasserburg 2000. hereafter QW: Wasserbure Qiau 2000: hereafter WQJ.," In contrast to the case of NSMs, observations of metal-poor stars as well as meteoritic data on $^{182}$ Hf and $^{129}$ I in the early solar system support a self-consistent picture of $r$ -process enrichment by SNe (Qian Wasserburg 2000, hereafter QW; Wasserburg Qian 2000; hereafter WQ)."
 Iu this picture the ejecta from cach SN event is mixed with an average mass Mags©ον105A of ISAL swept up by the SN remuaut (c5. Thornton et al.," In this picture the ejecta from each SN event is mixed with an average mass $M_{\rm mix}\approx 3\times 10^4\,M_\odot$ of ISM swept up by the SN remnant (e.g., Thornton et al."
 1998)., 1998).
 It is assumed that the SN rate per unit mass of gas RSMia Is approximately constaut over Galactic historv (this seenis reasonable as the star formation rate is proportional to the eas nis)., It is assumed that the SN rate per unit mass of gas $f_{\rm G}^{\rm SN}/M_{\rm gas}$ is approximately constant over Galactic history (this seems reasonable as the star formation rate is proportional to the gas mass).
" Consequently. au average ISAL in the Calaxyv is enriched by newh-svuthesized material frou SNe at a frequency where ReΑς, Is estimated using quantities for the present Galaxy."," Consequently, an average ISM in the Galaxy is enriched by newly-synthesized material from SNe at a frequency where $f_{\rm G}^{\rm SN}/M_{\rm gas}$ is estimated using quantities for the present Galaxy."
 Meteoritic data on ΟΠΕ (with a lifetime r4=1.3« vr) and DT (roy=2.3«10* vr) in the early solar svsteii shed iuportaut lieht on the r-process aud its association with SNe.," Meteoritic data on $^{182}$ Hf (with a lifetime $\bar\tau_{182}
=1.3\times 10^7$ yr) and $^{129}$ I $\bar\tau_{129}=2.3\times 10^7$ yr) in the early solar system shed important light on the $r$ -process and its association with SNe."
" Wasserbure. Dusso. Calline (1996: sce also. QW) showed that the abuudauce ratio (πεΟΠ=2.010. (Παρα Jacobsen 1996: Lee Halliday 1995, 1998) is consistent with common SNe injecting ΟΠΕ iuto the ISM at a hieh frequency"," Wasserburg, Busso, Gallino (1996; see also QW) showed that the abundance ratio $({^{182}{\rm Hf}}/{^{180}{\rm Hf}})_{\rm SSF}
=2.4\times 10^{-4}$ (Harper Jacobsen 1996; Lee Halliday 1995, 1998) is consistent with common SNe injecting $^{182}$ Hf into the ISM at a high frequency"
we [find a value of9.3%.. which is close to the Cummingetal.(2008) value of.,"we find a value of, which is close to the \citet{cumming} value of."
. The slight dilfereuce is probably not significant. but might be viewed as upward bias iu the value due to the inclusion of the hot Jupiters.," The slight difference is probably not significant, but might be viewed as upward bias in the \citet{cumming} value due to the inclusion of the hot Jupiters."
 In any. case we would uot have obtained very different constraints if we µας used the Cumunineetal.(2008) normalization in our Monte Carlo simulations., In any case we would not have obtained very different constraints if we had used the \citet{cumming} normalization in our Monte Carlo simulations.
 For comparison. amoug the other papers reporting Monte Carlo sunulations similar to ours. ]xasperetal.(2007) used a normalization of for planets with semimajor axes of 1-3 AU and masses greater than 1Mjap.," For comparison, among the other papers reporting Monte Carlo simulations similar to ours, \citet{kasper} used a normalization of for planets with semimajor axes of 1-3 AU and masses greater than 1."
. This is close to our value of for a similar range., This is close to our value of for a similar range.
 and Nielsenetal.(2008). fixed α aud 9 in their simulations. aud let the normalization be a [ree parameter.," \citet{GDPS} and \citet{nielsen} fixed $\alpha$ and $\beta$ in their simulations, and let the normalization be a free parameter."
 Chauvinetal.(2010) obtaiued their uormalizatiou from Cummingetal.(2008).. aud Nielsen&Close(2009) obtained theirs from Fischer&Valenti(," \citet{chauvin}
 obtained their normalization from \citet{cumming}, and \citet{nielsenclose} obtained theirs from \citet{carnp}."
2005).. Juric&Tremaine(2008) provide a helpful mathematical description of the ecceutricity distribution of known RV planets: where P(e) is the probability of a given extrasolar planets having orbital eccentricity €. € is the root of the natural logarithin. aud o=0.3.," \citet{juric} provide a helpful mathematical description of the eccentricity distribution of known RV planets: where $P(\epsilon)$ is the probability of a given extrasolar planet's having orbital eccentricity $\epsilon$, $e$ is the root of the natural logarithm, and $\sigma = 0.3$."
 We find that this mathematical form provides au excellent fit to the distribution of real exoplanet eccentricities from. Table 2 of (2005).. so we Lave used it as our probability distribution to generate random eccentricities for the Monte Carlo simulations we describe in Section 3. below. Burrow," We find that this mathematical form provides an excellent fit to the distribution of real exoplanet eccentricities from Table 2 of \citet{carnp}, so we have used it as our probability distribution to generate random eccentricities for the Monte Carlo simulations we describe in Section \ref{sec:tmod} below."
setal.(2003) present Ligh resolution. flux-calibrated theoretical spectra of. giaut planets or brown dwarfs [or ages ranging from 0.1-5.0 Gyr aud masses from 1 to 20 (these are availablefor download from burrows/)).," \citet{bur} present high resolution, flux-calibrated theoretical spectra of giant planets or brown dwarfs for ages ranging from 0.1-5.0 Gyr and masses from 1 to 20 (these are availablefor download from )."
 We have integrated these spectra to give absolute magnitudes[n] in the £/ and M filters used in Clio (see Tables P. aud 2)). and have found that the results can be reasonably interpolated to give the £ or AL band magnitudes for all planets of interest for our survey.," We have integrated these spectra to give absolute magnitudes in the $L'$ and $M$ filters used in Clio (see Tables \ref{tab:burl} and \ref{tab:burm}) ), and have found that the results can be reasonably interpolated to give the $L'$ or $M$ band magnitudes for all planets of interest for our survey."
 Baralleetal.(2003) also. present 1uodels of giant planets aud brown cdwarls. pre-integrated into maguitucles in the popular infrared. bauds.," \citet{bar} also present models of giant planets and brown dwarfs, pre-integrated into magnitudes in the popular infrared bands."
 These models predict slightly better sensitivity to low mass planets in the { band and slightly poorer sensitivity in the AZ banc. relative to the Burrowsetal.(2003) models.," These models predict slightly better sensitivity to low mass planets in the $L'$ band and slightly poorer sensitivity in the $M$ band, relative to the \citet{bur} models."
 We cannot say if the clifference is due to the slightly different filter sets used (AMUSO for Clio vs. Johuson-Cilass aud Joliusou for Baralleetal. (2003))). or ifit is intrinsic to the clilferent moclel spectra used iu Burrowsetal.(2003) aud Baralleetal. (2003).," We cannot say if the difference is due to the slightly different filter sets used (MKO for Clio vs. Johnson-Glass and Johnson for \citet{bar}) ), or if it is intrinsic to the different model spectra used in \citet{bur} and \citet{bar}. ."
. We have chosen to use the Burrowsetal.(2003) moclels exclusively herein. to avoid auy errors due to the slight filter differences.," We have chosen to use the \citet{bur} models exclusively herein, to avoid any errors due to the slight filter differences."
 Since the, Since the
Jiao for pointing out the simulated annealing method applied in this paper.,Jiao for pointing out the simulated annealing method applied in this paper.
 Thanks to Jorge Horvath for careful reading of the manuscript., Thanks to Jorge Horvath for careful reading of the manuscript.
to describe divergent lensing by the Ellis wormhole.,to describe divergent lensing by the Ellis wormhole.
 However. negative mass is not a physical enitity.," However, negative mass is not a physical enitity."
 As the lensing by the Ellis wormhole is convergent al the center. lensing al some other place must be divergent because (he wormhole has zero asvimptotic mass.," As the lensing by the Ellis wormhole is convergent at the center, lensing at some other place must be divergent because the wormhole has zero asymptotic mass."
" For "">LO. the light curve of the wormhole has a basin at {ῃ and no peak."," For $\hat{\beta_o} > 1.0$, the light curve of the wormhole has a basin at $t_0$ and no peak."
 Using these features. discrimination from Schwarzschild lensing can be achieved.," Using these features, discrimination from Schwarzschild lensing can be achieved."
 Equations (7)) and (20)) indicate (hat plivsical parameters (12). e. and ep) are degenerate in / and cannot be derived by fitting the light-curve data.," Equations \ref{eqn:re}) ) and \ref{eqn:te}) ) indicate that physical parameters $D_L$, $a$ , and $v_T$ ) are degenerate in $t_E$ and cannot be derived by fitting the light-curve data."
 This situation is the same as that lor Sehwarzschilcl lensing., This situation is the same as that for Schwarzschild lensing.
 To obtain or constrain these values. observations of the finite-source effect. Wickramasinghe1994) or parallax (Alcocketal.1995) are necessary.," To obtain or constrain these values, observations of the finite-source effect \citep{nem94} or parallax \citep{alc95} are necessary."
 The detectability of the magnilication of the star brightness depends on the timescale., The detectability of the magnification of the star brightness depends on the timescale.
 The Einstein radius crossing time { depends on the transverse velocity vy., The Einstein radius crossing time $t_E$ depends on the transverse velocity $v_T$.
 There is no reliable estimate of ey for wormlbholes., There is no reliable estimate of $v_T$ for wormholes.
 Here we assume that the velocity of the wornhole is approximately equal to the rotation velocity of stars (ep=220/n/s) if il is bound to the Galaxy., Here we assume that the velocity of the wormhole is approximately equal to the rotation velocity of stars $v_T = 220 km/s$ ) if it is bound to the Galaxy.
 If the wormhole is not bound to our Galaxy. the transverse velocity would be much higher.," If the wormhole is not bound to our Galaxy, the transverse velocity would be much higher."
 We assume ve;=5000/0/s (SalonovaTorres&Romero2002) [or the unbound swornhbole., We assume $v_T = 5000 km/s$ \citep{saf01} for the unbound wormhole.
 Table 2 shows the Einstein radius crossing times of the Ellis wormhole lensings for the Galactie bulge ancl LMC in both bound and unbound scenarios., Table 2 shows the Einstein radius crossing times of the Ellis wormhole lensings for the Galactic bulge and LMC in both bound and unbound scenarios.
 As the frequencies of current microlensing observations are limited to once every lew hours. an event lor which the timescale is less than one day is difficult to detect.," As the frequencies of current microlensing observations are limited to once every few hours, an event for which the timescale is less than one day is difficult to detect."
 To find very long timescale events ({> 1000days). long-term monitoring of events is necessary.," To find very long timescale events $t_E \geq 1000 days$ ), long-term monitoring of events is necessary."
 The realistic period of observation is <I0gears., The realistic period of observation is $\leq 10 years$.
 Thus. (he realistic size of the throat that we can search for is limited to 1004<a<10*Am both for the Galactic bulge and LMC if wormholes arebound (o our Galaxy.," Thus, the realistic size of the throat that we can search for is limited to $100 km \le a \le 10^7 km$ both for the Galactic bulge and LMC if wormholes arebound to our Galaxy."
 If wormholes are unbouncl. the detection is limited to 10hmxa10? flam.," If wormholes are unbound, the detection is limited to $ 10^5 km \le a \le 10^9 km$ ."
have used the non-detection of planetary companions to the WDs in our sample together with the estimated. detection sensitivities determined. from our simulations. to place upper limits to the frequency. of substellar ancl planetary companions to white charts.,"have used the non-detection of planetary companions to the WDs in our sample together with the estimated detection sensitivities determined from our simulations, to place upper limits to the frequency of substellar and planetary companions to white dwarfs."
 While no useful limits can be placed on the Likely frequeney of Mereury-sizecdl or smaller companions. and only weak constraints on the frequency of Earth-sized objects in the closest. orbits. slightly stronger constraints can be placed on the frequency. of larger bodies in very short-period orbits.," While no useful limits can be placed on the likely frequency of Mercury-sized or smaller companions, and only weak constraints on the frequency of Earth-sized objects in the closest orbits, slightly stronger constraints can be placed on the frequency of larger bodies in very short-period orbits."
" For example. brown cdwarfs and gas giants radius =lt, with periods <0.0.2 clays. similar to the known | BD binary WDOI8T349 (?).. must. certainly be relatively rare C10%)."," For example, brown dwarfs and gas giants radius $\approx \Rjup$ with periods $<0.1-0.2$ days, similar to the known $+$ BD binary $0137-349$ \citep{Maxted06}, must certainly be relatively rare $\la10\%$ )."
 Of course. this limit needs to be compared with those derived. [rom other sources. ee. infra-red sky surveys.," Of course, this limit needs to be compared with those derived from other sources, e.g. infra-red sky surveys."
 For example. ?. estimated that <0.54 of WDs have Lo cdwarf. companions. while Steele ct al. (," For example, \citet{Farihi05} estimated that $<0.5\%$ of WDs have L dwarf companions, while Steele et al. ("
in prep.),in prep.)
 tentatively suggests the fraction of unresolved: brown dwarf companions (including “Po cwarls) may be slightlv higher. between 12%.," tentatively suggests the fraction of unresolved brown dwarf companions (including T dwarfs) may be slightly higher, between $1-2\%$."
 From Spitzer photometry ? suggests that <4% of white dwarfs have unresolvecl substellar companions 10Mj. although the limits at lower masses (c.g. «6 are considerably weak.," From Spitzer photometry \citet{Farihi08a} suggests that $<4\%$ of white dwarfs have unresolved substellar companions $>10\Mjup$, although the limits at lower masses (e.g. $<6\Mjup$ ) are considerably weak."
 To place more stringent constraints on close substellar and gas giant companions to WDs. and similarly stringent constraints on Earth size bodies in close orbital separations likely requires significantly larger WD samples.," To place more stringent constraints on close substellar and gas giant companions to WDs, and similarly stringent constraints on Earth size bodies in close orbital separations likely requires significantly larger WD samples."
 In addition. our simulations and analvsis of WD light curves in the WASP archive suggests the degree of photometric precision is of somewhat secondary. importance compared to a high cadence and continuous Coverage.," In addition, our simulations and analysis of WD light curves in the WASP archive suggests the degree of photometric precision is of somewhat secondary importance compared to a high cadence and continuous coverage."
" The short curation of eclipses and transits of WDs (ος520 mins for companions radius = lu,:2:1.5 mins for terrestrial bodies) compared to the zS min cadence of WASP observations. appears to be the main factor limiting the transit detection rate in a survey optimised for planetary transits of main sequence stars."," The short duration of eclipses and transits of WDs $\approx5-20$ mins for companions radius $\approx \Rjup$; $\approx1-5$ mins for terrestrial bodies) compared to the $\approx
8$ min cadence of WASP observations, appears to be the main factor limiting the transit detection rate in a survey optimised for planetary transits of main sequence stars."
 Future surveys such as Pan-SVTARRS and LSST will be capable of detecting tens of thousands of WDs., Future surveys such as Pan-STARRS and LSST will be capable of detecting tens of thousands of WDs.
 However. we emphasise that observations of high cadence ancl long xwseline are of greatest benefit when attempting to detect he signature of close. eclipsing and transiting substellar and planetary companions to WDs.," However, we emphasise that observations of high cadence and long baseline are of greatest benefit when attempting to detect the signature of close, eclipsing and transiting substellar and planetary companions to WDs."
 Space missions such asCOROT.Nepler (see ?2)) and. especially.PLATO may herefore be better suited. to a survey of WDs as they deliver uninterrupted coverage at high cadence and exquisite ποιοιτὶς precision (~101o[O0 7%) and could at least in winciple detect the transits of asteroid-sized bodies across a WD.," Space missions such as, (see \citealt{DiStefano}) ) and, especially, may therefore be better suited to a survey of WDs as they deliver uninterrupted coverage at high cadence and exquisite photometric precision $\sim 10^{-4}-10^{-5}$ ) and could at least in principle detect the transits of asteroid-sized bodies across a WD."
 FE acknowledges funding from the European Commission under the. Marie. Curie Llost Fellowship for IZarly Stage Research Training SPARTAN. Contract No ALEST-C'p-2004-007512. University of Leicester. Ul.," FF acknowledges funding from the European Commission under the Marie Curie Host Fellowship for Early Stage Research Training SPARTAN, Contract No MEST-CT-2004-007512, University of Leicester, UK."
 AIRB acknowledges the support of an STEC Advanced Fellowship during part of this research., MRB acknowledges the support of an STFC Advanced Fellowship during part of this research.
 Phe WASP Consortium consists of astronomers primarily from. the Queen's. University Belfast. Ixeele.. Leicester. Phe Open University. ancl St Andrews. the Isaac Newton Croup (La Palma). the Instituto de Astrofisica de Canarias (Tenerife) and the South African Astronomical Observatory.," The WASP Consortium consists of astronomers primarily from the Queen's University Belfast, Keele, Leicester, The Open University, and St Andrews, the Isaac Newton Group (La Palma), the Instituto de Astrofisica de Canarias (Tenerife) and the South African Astronomical Observatory."
 Phe SuperWASP-N and \WASDP-S Cameras were constructed and operated with funds made available from Consortium Universities and the Ulv's Science and Technology Facilities Council, The SuperWASP-N and WASP-S Cameras were constructed and operated with funds made available from Consortium Universities and the UK's Science and Technology Facilities Council.
 WASP-South is hosted by the South African. Astronomical Observatory (SAAQO) and we are grateful for their support and assistance., WASP-South is hosted by the South African Astronomical Observatory (SAAO) and we are grateful for their support and assistance.
 The author also thank Professor Andrew Collier Cameron and the anonymous referee for helpful. comments on this paper., The author also thank Professor Andrew Collier Cameron and the anonymous referee for helpful comments on this paper.
some type of active galactic nucleus phenomena.,some type of active galactic nucleus phenomena.
" As mentioned above, event 23 now appears to be a variable star (Glicenstein2004;Jetzeretal.2004:Tisserand&Milsztajn2005;Bennett,Becker,Tomaney2005)., so 2 of the A00 events can be considered to be rejected."," As mentioned above, event 23 now appears to be a variable star \citep{glicens-hawaii,jetz-mil-tiss,tiss-mil,lmc-conf}, so 2 of the A00 events can be considered to be rejected."
" However, only event 23 was selected by the more restrictive cut A, which put tighter constraints on the quality of the light curve fit to à standard llight curve shape."," However, only event 23 was selected by the more restrictive cut A, which put tighter constraints on the quality of the light curve fit to a standard light curve shape."
 The drawback of criteria A is that it does reject events like the binary caustic crossing event 9 (Bennettetal.1996)., The drawback of criteria A is that it does reject events like the binary caustic crossing event 9 \citep{macho-lmc9}.
". The observed light curve features of this event are unique to binary caustic crossing events, and this means that this event is also almost certainly a true microlensing event."," The observed light curve features of this event are unique to binary caustic crossing events, and this means that this event is also almost certainly a true microlensing event."
" However, these same features also mean that this event has a poor fit to the standard light curve model, and this causes this event to fail the cuts for selection criteria A. [ will return to the effect of this systematic error in Sec. 4.."," However, these same features also mean that this event has a poor fit to the standard light curve model, and this causes this event to fail the cuts for selection criteria A. I will return to the effect of this systematic error in Sec. \ref{sec-conclude}."
 Our focus on the events selected by criteria A also means that we will not consider events 20 and 27 have the lowest S/N of all the A00 microlensing candidates., Our focus on the events selected by criteria A also means that we will not consider events 20 and 27 have the lowest S/N of all the A00 microlensing candidates.
 My classification of the 13 MACHO LMC microlensing candidates selected by criteria A of A00 is listed in Table 5.., My classification of the 13 MACHO LMC microlensing candidates selected by criteria A of A00 is listed in Table \ref{tab-that}.
 There are three dilferent categories of confirmed microlensing candidates listed in this table: follow-up. clump giant. and lens ID.," There are three different categories of confirmed microlensing candidates listed in this table: follow-up, clump giant, and lens ID."
 Event 5 is the only candidate confirmed with a direct lens identification (Alcocketal.2001b).., Event 5 is the only candidate confirmed with a direct lens identification \citep{macho-hstlmc5}.
" The final analysis of this event shows (Drake,Cook,&Keller2004:Gould.Bennett,Alves2004). that the microlensing data predicts a lens mass and distance that are completely consistent with the stellar brightness, color, and parallax."," The final analysis of this event shows \citep{dck-lmc5,lmc5-mass} that the microlensing data predicts a lens mass and distance that are completely consistent with the stellar brightness, color, and parallax."
" So, this event has passed three independent tests, and its microlensing interpretation is considered to be essentially certain."," So, this event has passed three independent tests, and its microlensing interpretation is considered to be essentially certain."
" Events 4, 13, 14. 15 all have high precision follow-up photometry [rom the CTIO 0.9m telescope obtained in seeing that is typically much better than in the MACHO survey images (AlcocketTomaney 2005).."," Events 4, 13, 14, 15 all have high precision follow-up photometry from the CTIO 0.9m telescope obtained in seeing that is typically much better than in the MACHO survey images \citep{macho-96lmc2,lmc-conf}."
" The improved seeing and longer exposure ümes lor the CTIO data yields photometry that has estimated uncertainties 2.5-5 Umes smaller than the original MACHO data, and in each case the microlensing interpretation is confirmed."," The improved seeing and longer exposure times for the CTIO data yields photometry that has estimated uncertainties 2.5-5 times smaller than the original MACHO data, and in each case the microlensing interpretation is confirmed."
" Events 16 and 17, were also discovered by the MACHO alert system (Alcocketal.1996b) as alerts 97-LMC-4 and 98-LMC-1, but these were classified as supernovae by the MACHO 5.7 year analysis."," Events 16 and 17, were also discovered by the MACHO alert system \citep{macho-alert}
 as alerts 97-LMC-4 and 98-LMC-1, but these were classified as supernovae by the MACHO 5.7 year analysis."
" MACHO Alerts 97-LMC-2, 3 and 5 were also in the MACHO 5.7 year data set, but these light curves failed other cuts and made neither the microlensing or supernovae samples."," MACHO Alerts 97-LMC-2, 3 and 5 were also in the MACHO 5.7 year data set, but these light curves failed other cuts and made neither the microlensing or supernovae samples."
 MACHO Alerts 97-LMC-6 and 7 occurred in fields not included in the 5.7 year analysis., MACHO Alerts 97-LMC-6 and 7 occurred in fields not included in the 5.7 year analysis.
" The only microlensing candidates from the 5.7 year analysis that occurred when the MACHO alert system was operaüng are events 18, 20, 22, 23, and 25."," The only microlensing candidates from the 5.7 year analysis that occurred when the MACHO alert system was operating are events 18, 20, 22, 23, and 25."
" However, events 22, 23 and 25 occurred at times when the alert system was not running in their respective fields."," However, events 22, 23 and 25 occurred at times when the alert system was not running in their respective fields."
 Event 18 reached maximum magnification in the middle of a 40-day period when the MACHO alert system was not operating due to a camera upgrade and the subsequent need to modily photometry code input files., Event 18 reached maximum magnification in the middle of a 40-day period when the MACHO alert system was not operating due to a camera upgrade and the subsequent need to modify photometry code input files.
" Event 20 was idenufied by the MACHO alert system, but it did not pass the subjecuive selecüon criteria for an alert."," Event 20 was identified by the MACHO alert system, but it did not pass the subjective selection criteria for an alert."
" However, event 20 did not pass selection criteria A, and so this is not relevant to our analysis."," However, event 20 did not pass selection criteria A, and so this is not relevant to our analysis."
 The only criteria A events that occurred, The only criteria A events that occurred
The range of £ is carefully chosen so that we cover the same values for the mic-laver plasma beta (the ratio of the gas pressure to the magnetic pressure) that were covered in the single-Iaver case.,The range of $F$ is carefully chosen so that we cover the same values for the mid-layer plasma beta (the ratio of the gas pressure to the magnetic pressure) that were covered in the single-layer case.
 Initially. we set f=0.0001. which corresponds to a," Initially, we set $F=0.0001$, which corresponds to a"
hreshold values (3.5. L0. 1.5. 5.0. 5.5 and 6 respectively).,"threshold values (3.5, 4.0, 4.5, 5.0, 5.5 and 6 respectively)."
 The choice of [5 as threshold in Sp is a compromise jetween a low flux limit aud the necessity to keep the iiber of spurious sources as low as possible., The choice of 4.5 as threshold in $_{\rm P}$ is a compromise between a low flux limit and the necessity to keep the number of spurious sources as low as possible.
 Frou: the analvsis of the negative peaks. we are confident that at nost three or four spurious sources are present iu the salple.," From the analysis of the negative peaks, we are confident that at most three or four spurious sources are present in the sample."
 The sources were extracted using the task SAD (Search Aud Destrov) which selects all the sources with peaks xiehter than a eiven level., The sources were extracted using the task SAD (Search And Destroy) which selects all the sources with peaks brighter than a given level.
 For cach selected source the Hux. the position aud the size are estimated using a least square Coussin ft.," For each selected source the flux, the position and the size are estimated using a least square Gaussian fit."
 However the Gaussian fit may be unreliable both for faint aud bright sources., However the Gaussian fit may be unreliable both for faint and bright sources.
 In fact the 101se map could iufiueuce the fit for faint sources (Coudon. 1997). while the differences between the CGaussiau beam and the real svuthesized beam could in principle ifluc1ος he goveness of a Gaussian fit even for sources with a aree signal to noise raio (Wiudhorst et al.," In fact the noise map could influence the fit for faint sources (Condon, 1997), while the differences between the Gaussian beam and the real synthesized beam could in principle influence the goodness of a Gaussian fit even for sources with a large signal to noise ratio (Windhorst et al.,"
 195D)., 1984).
 Thus. we uscd SAD to extrac all the sources whose peak ftX Sp Was orcaer than 3 tiucs the local riis value.," Thus, we used SAD to extract all the sources whose peak flux $_{\rm P}$ was greater than 3 times the local rms value."
 Subsequeutly. we derived the peal ftx for all the sources using a second deeree interpolation (ask MAXFIT).," Subsequently, we derived the peak flux for all the sources using a second degree interpolation (task MAXFIT)."
 Ouly the sources witha MAXFIT peak fux deusitv 2 1.50 were included iu the final sample., Only the sources with a MAXFIT peak flux density $\geq$ 4.5 $\sigma$ were included in the final sample.
 Texifter we will use the AIANFIT peak. flux density as the )eak flux density Sp of the sources., Hereafter we will use the MAXFIT peak flux density as the peak flux density $_{\rm P}$ of the sources.
 For mregular aud extuded sources the total flux density was determined by suunius the values of all the pixels covering the source., For irregular and extended sources the total flux density was determined by summing the values of all the pixels covering the source.
 With this proceure we selected 63 radio sources (2 of which have imultiple components for a total of 68 conmouents) over a otal area of 0087. dee., With this procedure we selected 63 radio sources (2 of which have multiple components for a total of 68 components) over a total area of 0.087 $^2$.
 Tn Fig., In Fig.
" 5 we plot the ratio between the total (ST) and peal flux density (Sp) as a function of the peak flux density for all he radio sources,", \ref{F3} we plot the ratio between the total $_{\rm T}$ ) and peak flux density $_{\rm P}$ ) as a function of the peak flux density for all the radio sources.
" Since the ratio of he total flux to the peak flux is a direct measure of the extension of a radio source. it can be used to discriminate vetwween resolved or extended sources 1ιο, larger than the ρα). and unresolved. sources."," Since the ratio of the total flux to the peak flux is a direct measure of the extension of a radio source, it can be used to discriminate between resolved or extended sources (i.e. larger than the beam) and unresolved sources."
" To select the extended SOULCCS, we have determined the lower cuveope of the Hux ratio distribution of Fig."," To select the extended sources, we have determined the lower envelope of the flux ratio distribution of Fig."
 3 and we have mirrored it above the Sp/Sp-l value (upper envelope of Fig. 3))., \ref{F3} and we have mirrored it above the $_{\rm T}$ $_{\rm P}$ =1 value (upper envelope of Fig. \ref{F3}) ).
 We have considere extended the 12 sources lavi18o above the upper euvelope that can be characterized by log (Sp/Sp) — log (1-15 / SU?), We have considered extended the 12 sources laying above the upper envelope that can be characterized by log $_{\rm T}$ $_{\rm P}$ ) = $-$ log (1-0.15 / $_{\rm P}^{0.3}$ ).
 Towever we shotld note that only 5 of the 12 sources classified as exteuded lio significantly above 1ο upper envelope (see Fie. 3)), However we should note that only 5 of the 12 sources classified as extended lie significantly above the upper envelope (see Fig. \ref{F3}) )
" and can be assumed as ""genuine"" extended sources.", and can be assumed as “genuine” extended sources.
 The other 7T sources just above the upper euvolope might also be unresolved sotrees whose total flux has been overestimated due to tlic! noise effect (see Windhorst et al. (, The other 7 sources just above the upper envelope might also be unresolved sources whose total flux has been overestimated due to the noise effect (see Windhorst et al. (
"198D) for a detailed disc""ussonj.",1984) for a detailed discussion).
 Throughout the paper we ave used the iuteeratec flux for the 12 formally extended sources and the peak flrx for the unresolved sources in all he calculations involving the radio flux., Throughout the paper we have used the integrated flux for the 12 formally extended sources and the peak flux for the unresolved sources in all the calculations involving the radio flux.
 The catalogue with the 63 sources (68 coniponents) is given iu Table 1., The catalogue with the 63 sources (68 components) is given in Table 1.
 For cach source we report the mame. he peak fux density οι» (and relative error 93.) in imd. he total flux. deusitv Sy du imuJw and error. the RA aud DEC (J200) aux corresponcdiue errors.," For each source we report the name, the peak flux density $_{\rm P}$ (and relative error $\sigma_{\rm S_{\rm P}}$ ) in mJy, the total flux density $_{\rm T}$ in mJy and error, the RA and DEC (J2000) and corresponding errors."
" Moreover. or resolved sources we also report fjio full width half uaxiuuni (FWIIM) of the non deconvolved major aud uinor axes (Oa; aud 6, in aresec) aud the position angele PA of the major axis (iu degrees. iicasured east to north)."," Moreover, for resolved sources we also report the full width half maximum (FWHM) of the non deconvolved major and minor axes $\theta_M$ and $\theta_m$ in arcsec) and the position angle PA of the major axis (in degrees, measured east to north)."
 All the Crrors du the source piarineters were determined following t10 Yecipes givon wv Condon (1997)., All the errors in the source parameters were determined following the recipes given by Condon (1997).
 Also in the calculation of the flux error we used the iuteerated flux for extended sources aud the peak flix for, Also in the calculation of the flux error we used the integrated flux for extended sources and the peak flux for
we did for the central position. using the method of line fitting in CLASS.,"we did for the central position, using the method of line fitting in CLASS."
 The results of these fits are listed in Table 5.., The results of these fits are listed in Table \ref{offlratio}.
" We also conducted followup observations in early 2004 (o detect additional lines of formaldehyde. also using the IRAM 30m. We used the A230 and B230 receivers to observe the 3244—245. 3n3—Zoo and 245—2y, transitions of formaldehyde. with rest frequencies of 225697.775 MllIz. respectively."," We also conducted followup observations in early 2004 to detect additional lines of formaldehyde, also using the IRAM $30$ m. We used the A230 and B230 receivers to observe the $3_{13}-2_{12}$, $3_{03}-2_{02}$ and $3_{12}-2_{11}$ transitions of formaldehyde, with rest frequencies of $211211.468$, $218222.192$, and $225697.775\,$ MHz, respectively."
 We also used an experimental setup of the 100 receiver in an unsuccessful attempt to detect the 15;—0o; transition of formaldehyde at a rest. frequency of 72327.948MIIZ.., We also used an experimental setup of the B100 receiver in an unsuccessful attempt to detect the $1_{01}-0_{00}$ transition of formaldehyde at a rest frequency of $72837.948$ MHz.
 We again used (he wobbler switching mode. with the same period and beam throw. and we used the same calibration procedures as for our initial observations.," We again used the wobbler switching mode, with the same period and beam throw, and we used the same calibration procedures as for our initial observations."
 On each receiver. we used one 256 channel 1 MIIz filterbank in parallel with the Vespa correlator svstem. set to a resolution of 320 kHz and a bandwidth ol 160 MIIz (100. GIIz receiver) or 320 MIIz (for the 230 GlIIz receivers).," On each receiver, we used one $256\,$ channel $1\,$ MHz filterbank in parallel with the Vespa correlator system, set to a resolution of $320\,$ kHz and a bandwidth of $160\,$ MHz (100 GHz receiver) or $320\,$ MHz (for the 230 GHz receivers)."
 For all of our observations. lines were placed in the lower sideband (LSB).," For all of our observations, lines were placed in the lower sideband (LSB)."
 Data reduction was (he same as for our initial observations. except as noted in discussions below.," Data reduction was the same as for our initial observations, except as noted in discussions below."
" Our lollowup observation spectra are plotted in Figure 7 with line fits: we discuss each spectrum separately,", Our followup observation spectra are plotted in Figure \ref{folspec} with line fits; we discuss each spectrum separately.
 Our 72.838 GHz observations were conducted using a new receiver setup which involved (tuning the DI00 receiver to frequencies below the nominal limits of the svstem. and using the LO svstem Irom the B230 receiver.," Our $72.838\,$ GHz observations were conducted using a new receiver setup which involved tuning the B100 receiver to frequencies below the nominal limits of the system, and using the LO system from the B230 receiver."
 The spectrum plotted is from the 1 MIs filterbank.," The spectrum plotted is from the $1\,$ MHz filterbank."
 Owing to operational difficulties with the svstem. we obtained data at only one LO setting and we did not obtain a spectrum of the image banc.," Owing to operational difficulties with the system, we obtained data at only one LO setting and we did not obtain a spectrum of the image band."
 The system was still undergoing commissioning tests during our observations and (hax calibrations may not be entirely reliable., The system was still undergoing commissioning tests during our observations and flux calibrations may not be entirely reliable.
 The single observed line in (his spectrum. the WCyN J=8—7 line. has been observed only once before. by Morrisetal.(1975) ancl our observation is consistent with their flix measurements. bul owing to the relatively large flix uncertainties of the Morrisetal.(1975). observation. we are unable to obtain a better calibration. and estimate that our svstematic uncertaintv may be as high as 50%.," The single observed line in this spectrum, the $_3$ N $J=8-7$ line, has been observed only once before, by \citet{mor75} and our observation is consistent with their flux measurements, but owing to the relatively large flux uncertainties of the \citet{mor75} observation, we are unable to obtain a better calibration, and estimate that our systematic uncertainty may be as high as $50\%$."
" The vertical dashed lines in this spectrum indicate the expected position ol the lo,—Oo transition of formaldehyde.", The vertical dashed lines in this spectrum indicate the expected position of the $1_{01}-0_{00}$ transition of formaldehyde.
 The spectrum plotted around 211.211 GIHz represents data taken at (wo LO settings aud is from the 1 MIIz filterbank.," The spectrum plotted around $211.211\,$ GHz represents data taken at two LO settings and is from the $1\,$ MHz filterbank."
 The spectrum shows only (wo features., The spectrum shows only two features.
 We identify Cie line al tose=—26.5kms! as the 344—24» transition of formaldehyde.," We identify the line at $v_{LSR}=-26.5\,{\rm km\,s^{-1}}$ as the $3_{13}-2_{12}$ transition of formaldehyde."
 The remaining feature is due to a blend of the 545—44; 11/2—9/2 6—5 and 5—4 transitions of c-C'll.," The remaining feature is due to a blend of the $5_{15}-4_{14}$ $11/2-9/2$ $6-5$ and $5-4$ transitions of $_3$ H. The $218.222\,$ GHz spectrum was obtained using a wider bandwidth than the rest of our"
"o that ""expected"" from the Universal barvon fraction: fo=ALAL).",to that “expected” from the Universal baryon fraction: $f_* \equiv M_*/(f_b M_{\rm vir})$.
 We make the (rather loose) demaud hat f£. lies within the range of plausible values f.=OL 1f., We make the (rather loose) demand that $f_*$ lies within the range of plausible values $f_* = 0.01 - 1 f_b$ .
 For each chosen value of ej; and acopted disk orlmation scenario (AC or uou-AC') the chosen 155 constraint defines the rotation curve completely. aud thus xovides an inplied shear rate at every racius.," For each chosen value of $c_{\rm vir}$ and adopted disk formation scenario (AC or non-AC) the chosen $V_{2.2}$ constraint defines the rotation curve completely, and thus provides an implied shear rate at every radius."
 The three xuels of Figures 5. and 6 show the results of this exercise for ESO 582-CG12 aud IC 2522 respectively., The three panels of Figures \ref{eso582_m} and \ref{ic2522_m} show the results of this exercise for ESO 582-G12 and IC 2522 respectively.
 In cach panel. open syaubols are for the non-AC model aud he solid sviubols reflect the AC assumption.," In each panel, open symbols are for the non-AC model and the solid symbols reflect the AC assumption."
 Each point represents a distinct iuput combination of h. κε. aud (MfL).," Each point represents a distinct input combination of $h$, $L_{disk}$, and $(M/L)$."
 The measured shear rate is illustrated by a solid vertical line aud the 1006 range in the observe shear rate for cach ealaxy is shown by the two vertical dashed hues in cach panel., The measured shear rate is illustrated by a solid vertical line and the $\pm 1-\sigma$ range in the observe shear rate for each galaxy is shown by the two vertical dashed lines in each panel.
 Consider first the left laud panel of Figure 5.., Consider first the left hand panel of Figure \ref{eso582_m}.
 Here we plot the dark matter halo concentration parameter versus the shear rate measured at 10 kpc., Here we plot the dark matter halo concentration parameter versus the shear rate measured at $10$ kpc.
 More concentrated halos generally produce higher shear rates. as expected.," More concentrated halos generally produce higher shear rates, as expected."
" It cau be seen that for a given NEW concentration. 6,5, several values of shear are possible."," It can be seen that for a given NFW concentration, $c_{vir}$, several values of shear are possible."
 This is due to changes oei the barvon coutrbution (1.60. the disk mass απο cüsk scaleleusth) to the rotation curve., This is due to changes in the baryon contrbution (i.e. the disk mass and disk scalelength) to the rotation curve.
 Whether an merease in the barvou contribution causes the shear to increase or decrease depeuds on the size of the disk (6. the isk scaleleusth)., Whether an increase in the baryon contribution causes the shear to increase or decrease depends on the size of the disk (i.e. the disk scalelength).
 The same is true for Figure 6., The same is true for Figure 6.
 In the AC inodel (solid). eq refers to the halo conceutratiou the halo is adiabatically contracted.," In the AC model (solid), $c_{\rm vir}$ refers to the halo concentration the halo is adiabatically contracted."
 This is why the AC points tend to have higher shear values at fixed ey cColmpared to the non-ÀC case., This is why the AC points tend to have higher shear values at fixed $c_{\rm vir}$ compared to the non-AC case.
 Every point (or h. Lana. (ALPE) input combination) has au associated stellar biarvou fraction f. aud dark halo virial mass My.," Every point (or $h$, $L_{disk}$, $(M/L)$ input combination) has an associated stellar baryon fraction $f_*$ and dark halo virial mass $M_{\rm vir}$ ."
 These values are plotted versus shear rate in the iuiddle and right-hand panels in Figure 5.., These values are plotted versus shear rate in the middle and right-hand panels in Figure \ref{eso582_m}.
 Observe that with oulv the V5.2 constraiut iniposed (all poiuts). then a wide ranee of dark matter halo properties are allowed.," Observe that with only the $V_{2.2}$ constraint imposed (all points), then a wide range of dark matter halo properties are allowed."
 Ouce we constrain the models by forcing the predicted shear to be consistent the observed rauge S=O.LF—0.57. we favor models that include adiabatic contraction.," Once we constrain the models by forcing the predicted shear to be consistent the observed range $S=0.47-0.57$, we favor models that include adiabatic contraction."
 The dark matter of halo of ESO 582-C12 is constrained to be of relatively low mass. Mag25.8&10H ML... and to have a αν ligh NEW concentration. eq2716.," The dark matter of halo of ESO 582-G12 is constrained to be of relatively low mass, $M_{\rm vir} \simeq 5-8 \times 10^{11}$ $_{\odot}$, and to have a fairly high NFW concentration, $c_{\rm vir} > \sim 16$."
" This is consistent with the ligh-end ofthe the «4, distribution predicted for LCDM (Bullocketal.20015).", This is consistent with the high-end of the the $c_{\rm vir}$ distribution predicted for LCDM \citep{bullock01b}.
.. The stellar mass in the ealaxy is between ~15 aud 10:4 of the the barvonie mass associated with its dark matter halo., The stellar mass in the galaxy is between $\sim 15$ and $40$ of the the baryonic mass associated with its dark matter halo.
 Figure 6 shows analogous results for IC 2522., Figure \ref{ic2522_m} shows analogous results for IC 2522.
 Uulike ESO 582 C12. this ealaxy prefers a more massive. concentration halo: Mag21.5.6& ΤΟΝΙ with 2. N. af the low end of the e distribution (Bulloc-etal. 2001b).," Unlike ESO 582 G12, this galaxy prefers a more massive, low-concentration halo: $M_{\rm vir} \simeq 1.5-6 \times 10^{12}$ $_{\odot}$ with $c_{\rm vir}
\simeq 2-8$ , at the low end of the $c_{\rm vir}$ distribution \citep{bullock01b}."
. The implied stellar mass is quite siall: <10% of the universal barvou budget has cuded wp on stars., The implied stellar mass is quite small: $< 10$ of the universal baryon budget has ended up on stars.
 Furthermore. the measured shear rate of thi« ealaxy favors our non-AC model.," Furthermore, the measured shear rate of this galaxy favors our non-AC model."
 This is in qualitative aerectnent with results presented by Ditton oct ((2005)., This is in qualitative agreement with results presented by Dutton et (2005).
 Model results for the central mass concentration iu cach galaxy are shown in Figures 7 and 8., Model results for the central mass concentration in each galaxy are shown in Figures \ref{eso582} and \ref{ic2522}.
" Were e; and Cay, correspond to the fraction of the total and dar- latter mass contained within central 10 kpe of each system.", Here $c_{\rm tot}$ and $c_{\rm dm}$ correspond to the fraction of the total and dark matter mass contained within central $10$ kpc of each system.
" We see that ESO 582-CGH12 contains roughly 20% of ifs nass wilin 10 kpc. while IC 2552 is quite diffuse with euον, "," We see that ESO 582-G12 contains roughly $20$ of its mass within $10$ kpc, while IC 2552 is quite diffuse with $c_{\rm tot} < 5$."
What drives these differences?, What drives these differences?
 Tale IC 2522 for example. which has a very low concentration.," Take IC 2522 for example, which has a very low concentration."
 This arises because of the combination of its (small) disk scalelenetland the fact that the rotation curve is rising at 10 kpc (auplied by a shear of 0.3940.1)., This arises because of the combination of its (small) disk scalelength the fact that the rotation curve is rising at 10 kpc (implied by a shear of $\pm$ 0.1).
 The rotation curve of the stellar disk alone peaks at 2.2 scalelengths (8.5 kpc). aud will be falling at LO kpe Gvhere the shear is measured).," The rotation curve of the stellar disk alone peaks at 2.2 scalelengths $\sim 8.8$ kpc), and will be falling at 10 kpc (where the shear is measured)."
 The measured shear of 0.39-40.01 at 10 kpc iuplies that the total rotation curve is rising at 10 Xpc. which means the dark matter must be very extended (or of low concentration).," The measured shear of $\pm$ 0.01 at 10 kpc implies that the total rotation curve is rising at 10 kpc, which means the dark matter must be very extended (or of low concentration)."
 ESO 582-C12. on the other haud. has a falling rotation curve at LO kpe. but aleryer disk scaleleusth.," ESO 582-G12, on the other hand, has a falling rotation curve at 10 kpc, but a disk scalelength."
 This requires a more concentrated halo to explain the observed shear rate., This requires a more concentrated halo to explain the observed shear rate.
 The two galaxies chosen for this modeling are. in any Ware. quite comparable.," The two galaxies chosen for this modeling are, in many ways, quite comparable."
 Their Hubble types. B/D ratios and disk Dpmuuinosities are very simular.," Their Hubble types, B/D ratios and disk luminosities are very similar."
 The oulv sienificaut difference is m their measured rates of shear., The only significant difference is in their measured rates of shear.
 The differeuces in their favored formation models (AC for ESO 582-CG12 versus uon-AC for IC 2522). halo lüasses and concentrations. conies down to a different measurement for their shear aud the sizes of their disks.," The differences in their favored formation models (AC for ESO 582-G12 versus non-AC for IC 2522), halo masses and concentrations, comes down to a different measurement for their shear and the sizes of their disks."
 It is quite interesting that two ealaxies. which scci so siuilar in size aud IIubble type. ποσα to inhabit very different types of dark matter halos.," It is quite interesting that two galaxies, which seem so similar in size and Hubble type, seem to inhabit very different types of dark matter halos."
 It is tempting to speculate that the shear rate itselfprovides an uuderlviug physical driver to push galaxies in drastically differeut dark matter halos towards similar Iuninuosities., It is tempting to speculate that the shear rate itself provides an underlying physical driver to push galaxies in drastically different dark matter halos towards similar luminosities.
 However. the problem would then be what stops such a dive aud makes the ealaxy huninosities so similar. when the halo masses show large differences.," However, the problem would then be what stops such a drive and makes the galaxy luminosities so similar, when the halo masses show large differences."
 This clearly demonstrates that the shear rate adds an important coustraint on galaxy formation moclels compared to what can be learned. from standard Tully-Fisher constraints alone., This clearly demonstrates that the shear rate adds an important constraint on galaxy formation models compared to what can be learned from standard Tully-Fisher constraints alone.
 Iudeed. the results for IC 2522 and ESO 582-G12 are tantalizing.," Indeed, the results for IC 2522 and ESO 582-G12 are tantalizing."
 The first. with a low shear rate. favors a non-adiabaticallv contracted halo. with a low halo concentration aud a massive. extended dark matter halo.," The first, with a low shear rate, favors a non-adiabatically contracted halo, with a low halo concentration and a massive, extended dark matter halo."
 The second. with a high shear rate. favors an adiabatically contracted halo with a lugh NEW concentration. and relatively low virial mass.," The second, with a high shear rate, favors an adiabatically contracted halo with a high NFW concentration, and relatively low virial mass."
 These two cases motivate the application of the shear-rate constraint to a larger sample of galaxies., These two cases motivate the application of the shear-rate constraint to a larger sample of galaxies.
 Since shear is related to the spiral aria pitch angle aud. spiral armis are clearly detectable in disk galaxies at :~I (e.g. Ehueereen. Ehugreen Ist 2001). it is possible to constrain models of disk formation and dark matter halo structure back to a look-back time of ~7 Cevrs. and evaluate how the mass distribution in disk galaxies changes with time.," Since shear is related to the spiral arm pitch angle and spiral arms are clearly detectable in disk galaxies at $z\sim 1$ (e.g. Elmegreen, Elmgreen Hirst 2004), it is possible to constrain models of disk formation and dark matter halo structure back to a look-back time of $\sim 7$ Gyrs, and evaluate how the mass distribution in disk galaxies changes with time."
 An important test of our modeling is how well the modeled rotation curves fit the data., An important test of our modeling is how well the modeled rotation curves fit the data.
 Figure 9 shows the observed rotation curves from Perse Salucei (1995) overlaid with our modeled rotation curves., Figure 9 shows the observed rotation curves from Persic Salucci (1995) overlaid with our modeled rotation curves.
 Iu both cases the rotation curve is modeled well iu the outer parts. where we are measuring the shear.," In both cases the rotation curve is modeled well in the outer parts, where we are measuring the shear."
 In the inner regions. where the bulge is inportaut. both rotation curve models overestimate the observed rotation curve. althoughthe model for ESO 582-Cl2 presents a mach better case than that for IC 2522.," In the inner regions, where the bulge is important, both rotation curve models overestimate the observed rotation curve, althoughthe model for ESO 582-G12 presents a much better case than that for IC 2522."
 For IC 2522. the solid line represeuts a model rotation curve. with a bulge-to-disk ratio of 0.161 (see Table 1) aud a bulge effective radius of 0.2," For IC 2522, the solid line represents a model rotation curve, with a bulge-to-disk ratio of 0.164 (see Table 4) and a bulge effective radius of 0.2"
polarized than the upper level (seeBelluzzietal..2007).. in this scandium line the upper and lower levels carry the same amount of atomic (see Fig. 4)).,"polarized than the upper level \citep[see][]{Bel07}, in this scandium line the upper and lower levels carry the same amount of atomic (see Fig. \ref{fig:atomic-pol}) )."
 Although the contribution of dichroism ts thus expected to be more important in this line than in the Ba D» line. we start our investigation by neglecting it.," Although the contribution of dichroism is thus expected to be more important in this line than in the Ba $_2$ line, we start our investigation by neglecting it."
 We therefore calculate the polarization of the radiation emerging from the slab by means of the equation Some results obtained taking into account dichroism effects (1.e.. applying Eq. (9)))," We therefore calculate the polarization of the radiation emerging from the slab by means of the equation Some results obtained taking into account dichroism effects (i.e., applying Eq. \ref{eq:dichroism}) ))"
 will be shown in Sect., will be shown in Sect.
 3.4., 3.4.
 We recall that the emission coefficient given in Eq. (5)), We recall that the emission coefficient given in Eq. \ref{eq:epsilon}) )
 includes only line processes., includes only line processes.
 To reproduce. albeit qualitatively. the observed profile. we need to add the contribution of the continuum.," To reproduce, albeit qualitatively, the observed profile, we need to add the contribution of the continuum."
" Assuming that the continuum is constant across the line. we have where the superseripts ο and ""( denote that the corresponding quantities refer to continuum and line processes. respectively."," Assuming that the continuum is constant across the line, we have where the superscripts $c$ ” and $\ell$ ” denote that the corresponding quantities refer to continuum and line processes, respectively."
 The quantities εν and «71 are considered to be free parameters in the problem that are adjusted to reproduce the observed polarization profile.," The quantities $\varepsilon_{X}^{\, c}$ and $\varepsilon_I^{\, c}$ are considered to be free parameters in the problem that are adjusted to reproduce the observed polarization profile."
 We point out that the three-peak structure of the Q/7 profile shown by the Ba D» line could be reproduced by Belluzzietal.(2007) within the constraints of the same modeling assumptions and approximations outlined above., We point out that the three-peak structure of the $Q/I$ profile shown by the Ba $_2$ line could be reproduced by \citet{Bel07} within the constraints of the same modeling assumptions and approximations outlined above.
 However. we note that this ionized scandium line is a rather strong spectral line: as for Ba D». the wings of this line originate in the photosphere. while the line core originates in the high photosphere/low chromosphere.," However, we note that this ionized scandium line is a rather strong spectral line: as for Ba $_2$, the wings of this line originate in the photosphere, while the line core originates in the high photosphere/low chromosphere."
 The optically thin slab model considered in this paper ts therefore just a first order approximation., The optically thin slab model considered in this paper is therefore just a first order approximation.
 Nevertheless. as previously stated. it allows us to take into account in a very rigorous way the relevant atomic aspects of the problem. avoiding the complications produced by radiative transfer effects.," Nevertheless, as previously stated, it allows us to take into account in a very rigorous way the relevant atomic aspects of the problem, avoiding the complications produced by radiative transfer effects."
 As discussed in Belluzzietal.(2007).. the presence of barium isotopes both with and without HFS is understood to be important in explaining the three-peak structure of the Q/7 profile shown in the second solar spectrum by the Ba D» line.," As discussed in \citet{Bel07}, the presence of barium isotopes both with and without HFS is understood to be important in explaining the three-peak structure of the $Q/I$ profile shown in the second solar spectrum by the Ba $_2$ line."
" In. particular. the central peak is due to the isotopes without HFS. while the two lateral peaks seem to be caused by the depolarizing effect of the isotopes with HFS (""Ba and Ba). combined with the effect of the continuum. which depolarizes in the wings of the line."," In particular, the central peak is due to the isotopes without HFS, while the two lateral peaks seem to be caused by the depolarizing effect of the isotopes with HFS $^{135}$ Ba and $^{137}$ Ba), combined with the effect of the continuum, which depolarizes in the wings of the line."
 The fractional polarization profile plotted in the left panel of Fig. 5..," The fractional polarization profile plotted in the left panel of Fig. \ref{fig:frac-pol},"
 clearly shows the depolarizing effect of the isotopes with HFS. as well as the polarization enhancement at line center due to the isotopes without HFS.," clearly shows the depolarizing effect of the isotopes with HFS, as well as the polarization enhancement at line center due to the isotopes without HFS."
 This profile was obtained by applying Eq. (10) (, This profile was obtained by applying Eq. \ref{eq:QsuI}) ) (
1.e.. neglecting the continuum) to the case of the Batt D» line. assuming for #. w. and Ay the values corresponding to a slab of barium ions at 6000 K. located 1000 km above the solar surface.,"i.e., neglecting the continuum) to the case of the Ba $_2$ line, assuming for $\bar{n}$, $w$, and $\Delta \lambda_D$ the values corresponding to a slab of barium ions at 6000 K, located 1000 km above the solar surface."
 Starting from this fractional polarization profile. and adding the contribution of the continuum. Belluzzi were able to obtain a good fit to the observed three-peak Q/7 profile.," Starting from this fractional polarization profile, and adding the contribution of the continuum, \citet{Bel07} were able to obtain a good fit to the observed three-peak $Q/I$ profile."
 The fundamental role of the isotopes without HFS in producing the higher central peak of the observed linear polarization profile can be clearly appreciated from the middle panel of Fig. 5.., The fundamental role of the isotopes without HFS in producing the higher central peak of the observed linear polarization profile can be clearly appreciated from the middle panel of Fig. \ref{fig:frac-pol}.
 Here the same fractional polarization profile as in the left panel of Fig., Here the same fractional polarization profile as in the left panel of Fig.
 5. (without continuum) is plotted assuming that only the isotope Ba (with HFS) is present in abundance), \ref{fig:frac-pol} (without continuum) is plotted assuming that only the isotope $^{137}$ Ba (with HFS) is present in abundance).
" Comparing the two profiles. it is noticeable that although a polarization enhancement can. be observed at line center also when only ""Ba is present. this is much larger when the isotopes without HFS are also taken into account."," Comparing the two profiles, it is noticeable that although a polarization enhancement can be observed at line center also when only $^{137}$ Ba is present, this is much larger when the isotopes without HFS are also taken into account."
 In particular. if these isotopes are considered. the polarization at line center reaches the same value as in the far wings.," In particular, if these isotopes are considered, the polarization at line center reaches the same value as in the far wings."
 This is an important peculiarity. since in this case there is no way. by adding the continuum. to obtain three peaks of the same amplitude (such as those observed in Se line).," This is an important peculiarity, since in this case there is no way, by adding the continuum, to obtain three peaks of the same amplitude (such as those observed in Sc line)."
" On the other hand. this would be possible if. hypothetically. only the isotope ""Ba were present."," On the other hand, this would be possible if, hypothetically, only the isotope $^{137}$ Ba were present."
 As previously pointed out. scandium has a single stable isotope. with HFS. which is consistent with the interpretation of this scandium signal. showing three peaks of the same amplitude. in terms of HFS.," As previously pointed out, scandium has a single stable isotope, with HFS, which is consistent with the interpretation of this scandium signal, showing three peaks of the same amplitude, in terms of HFS."
 Applying Eq. (10)).," Applying Eq. \ref{eq:QsuI}) ),"
 and using the values previously calculated for 5. w. and Ap. we find the fractional polarization profile shown in the right panel of Fig. 5..," and using the values previously calculated for $\bar{n}$, $w$, and $\Delta \lambda_D$, we find the fractional polarization profile shown in the right panel of Fig. \ref{fig:frac-pol}."
" The profile clearly shows the depolarizing effect ofHFS*.. but the polarization enhancement at line-center. which produces the central peak. is by far less evident than in the case of ""Ba."," The profile clearly shows the depolarizing effect of, but the polarization enhancement at line-center, which produces the central peak, is by far less evident than in the case of $^{137}$ Ba."
" On the one hand. this is because the HFS components of ""Ba are gathered into (wo well separated groups (because of the large HFS splitting of the ground level). while the HFS components of scandium are gathered into a single group at line center. spreading out over à narrow wavelength interval of about 13 (see the middle and right panels of Fig. 5.."," On the one hand, this is because the HFS components of $^{137}$ Ba are gathered into two well separated groups (because of the large HFS splitting of the ground level), while the HFS components of scandium are gathered into a single group at line center, spreading out over a narrow wavelength interval of about 15 (see the middle and right panels of Fig. \ref{fig:frac-pol},"
 and Fig. 2)), and Fig. \ref{fig:hfs-comp}) ).
 On the other hand. the smaller Doppler width assumed for barium (much heavier than scandium) contributes to making the Ime-core enhancement of the corresponding fractional polarization profile far more evident.," On the other hand, the smaller Doppler width assumed for barium (much heavier than scandium) contributes to making the line-core enhancement of the corresponding fractional polarization profile far more evident."
 As 1n the case of barium. to reproduce the observed profile. it is necessary to include the effect of the continuum.," As in the case of barium, to reproduce the observed profile, it is necessary to include the effect of the continuum."
 By so doing. the value of the fractional polarization decreases in the wings of the line. and a two-peak profile with a small hump at line center is obtained.," By so doing, the value of the fractional polarization decreases in the wings of the line, and a two-peak profile with a small hump at line center is obtained."
 The left panel of Fig., The left panel of Fig.
 6. shows the, \ref{fig:cont} shows the
"uunber of detectious of extended objects (EO). aud in which the central galaxies lave similar redshifts (five of the sample galaxies have v, between 2266 and 3132 liu L ONGC 3226/7 is at 1151 km Pj.","number of detections of extended objects (EO), and in which the central galaxies have similar redshifts (five of the sample galaxies have $_{\rm r}$ between 2266 and 3132 km $^{-1}$, NGC 3226/7 is at 1151 km $^{-1}$ )."
" Thus. a subsample of fields around six galaxies mamoly NCC 520. NGC 772, NGC 3226/7. NGC 3656. Arp 111 and Arp 299 - is analvzed in this paper."," Thus, a subsample of fields around six galaxies – namely NGC 520, NGC 772, NGC 3226/7, NGC 3656, Arp 141 and Arp 299 - is analyzed in this paper."
 Amone these. Arp 299 was observed in one filter onlv. but the high uuuber of detections of EO iu this field lead to its inclusion in this study.," Among these, Arp 299 was observed in one filter only, but the high number of detections of EO in this field lead to its inclusion in this study."
 Iu the following. we describe the observations and the detection procedure in Sect.," In the following, we describe the observations and the detection procedure in Sect."
 2., 2.
 In Sect., In Sect.
 3. we outline the properties of known TDC within the framework of a dvuaimic evolution scenario and describe the statistical analvsis that was performed iu order to quantity the possible number of TDG.," 3, we outline the properties of known TDG within the framework of a dynamic evolution scenario and describe the statistical analysis that was performed in order to quantify the possible number of TDG."
 Couclusious are given im section Sect., Conclusions are given in section Sect.
 L., 4.
" Thnages of the fields were taken in December 1992 with the 2.510 Isaac Newton Telescope at the Observatorio del Roque de los Muchachos. La Palma. with the EEVS CCD-calucra dn prime focus. giving au ciective field of about «10.5' and 0.55 "" *1 resolution."," Images of the fields were taken in December 1992 with the 2.5m Isaac Newton Telescope at the Observatorio del Roque de los Muchachos, La Palma, with the EEV5 CCD-camera in prime focus, giving an effective field of about $~\arcmin \times 10.5~\arcmin$ and 0.55 $\arcsec$ $^{-1}$ resolution."
 The observations took place in bright-ereyv time which affected the limiting magnitude., The observations took place in bright-grey time which affected the limiting magnitude.
" Atmospheric transparency conditious were very good and seeiug ranged from 0.9 to 1.5 "".", Atmospheric transparency conditions were very good and seeing ranged from 0.9 to 1.5 $\arcsec$.
 Tages were taken with the V aud A Cousins filters. with exposure times as eiven in Table 1.," Images were taken with the $V$ and $R$ Cousins filters, with exposure times as given in Table 1."
 In order to determine the stellar PSF. frames with short 20 8 exposures were also taken. iu which kuown xieht field stars are not saturated.," In order to determine the stellar PSF, frames with short 20 s exposures were also taken, in which known bright field stars are not saturated."
 Laucdolt (1992) standards were observed for flux calibration., Landolt (1992) standards were observed for flux calibration.
 Table 1 lists relevant xwanmeters of the target list and of the observations., Table 1 lists relevant parameters of the target list and of the observations.
 Raw πμασοςOo were zero-subtracted aux flatflelcled iu he usual manner., Raw images were zero-subtracted and flatfielded in the usual manner.
 The bias frame used was taken froui averaging all bias frames in the observatious., The bias frame used was taken from averaging all bias frames in the observations.
 Also. Hatfelds were obtained by averaging the individual sky Hatfields taken at the eutire observing run.," Also, flatfields were obtained by averaging the individual sky flatfields taken at the entire observing run."
 This was yossible. since the stability of the flatflelds between üghts was better than," This was possible, since the stability of the flatfields between nights was better than."
 An absolute flux calibration was performed based on the observations of Laudolt standards at clifferent adrniasses., An absolute flux calibration was performed based on the observations of Landolt standards at different airmasses.
 The standards were ueasured using fixec-aperture photometry., The standards were measured using fixed-aperture photometry.
 For luore details about observatious. calibration oxocedure aud classification nie)0dology. sce Paper I. A procedure was developed to identify objects over the background of the CCD frames. discriminating sinall ealaxy-like objects iid star-like objects (see the flow diagram iu Paper I).," For more details about observations, calibration procedure and classification methodology, see Paper I. A procedure was developed to identify objects over the background of the CCD frames, discriminating small galaxy-like objects and star-like objects (see the flow diagram in Paper I)."
 Ideuti&ed objects were then paruneterize by the fitting of elliptical Gaussians. as is described later in the ext.," Identified objects were then parameterized by the fitting of elliptical Gaussians, as is described later in the text."
 The total error in the photometric calibration is about 0.15 mae., The total error in the photometric calibration is about 0.15 mag.
 Objects with a size close to the PSF were considered stellar (SO) whereas those with an extent of more than 1.3 times the PSF in either cé or y direction were classified as extended (EO)., Objects with a size close to the PSF were considered stellar (SO) whereas those with an extent of more than 1.3 times the PSF in either $x$ or $y$ direction were classified as extended (EO).
 Objects whose EWIIM ranges between land 1.3 times the PSF were classifies as almbieuous., Objects whose FWHM ranges between 1 and 1.3 times the PSF were classified as ambiguous.
 All saturated objects; also those having widths larger than the PSF. can safely be classified as stellar. too.," All saturated objects, also those having widths larger than the PSF, can safely be classified as stellar, too."
 Table 2 gives the umber of EO«s. SOs and ambiguous identifications found. aro the galaxies of the guple.," Table 2 gives the number of EOs, SOs and ambiguous identifications found around the galaxies of the sample."
 The selection method used certainly does not exclude that the lists of ealaxies still includei some seiiesolved double or uultiple stars., The selection method used certainly does not exclude that the lists of galaxies still include some semi-resolved double or multiple stars.
 Also the lists of stars may coutain a substantial muuber of galaxies witha nearstellar PSF., Also the lists of stars may contain a substantial number of galaxies with a near–stellar PSF.
 Iu galaxystar separations done in a wav comparable to ours. Motcealfe et al. (," In galaxy–star separations done in a way comparable to ours, Metcalfe et al. ("
1991) found that later spectroscopy revealed of the starlike objects to be compact galaxies. aud of the galaxy ideutificatious turued out to be stars.,"1991) found that later spectroscopy revealed of the star–like objects to be compact galaxies, and of the galaxy identifications turned out to be stars."
 It might xo argued that some of the detected objects are elobular clusters formed during the interaction., It might be argued that some of the detected objects are globular clusters formed during the interaction.
 In fact. Zepf Aslinan (1993) aud Ashunan Zepf (1992) reported two-aged populations o elobular clusters around merecr remnant candidates.," In fact, Zepf Ashman (1993) and Ashman Zepf (1992) reported two-aged populations of globular clusters around merger remnant candidates."
 The colors of our objects would be compatible with those of typical elobular clusters (Iarris 1996). thus opening the possibility hat they are globular clusters tha had originated during the interaction.," The colors of our objects would be compatible with those of typical globular clusters (Harris 1996), thus opening the possibility that they are globular clusters that had originated during the interaction."
 Wowever. globular clusters should be detected as polut-like sources at the distances of the interacting ealaxies of our sample.," However, globular clusters should be detected as point-like sources at the distances of the interacting galaxies of our sample."
 In addition. ouly the faintest EOs iu our sample would have huninosities comparable to those of globular clusters.," In addition, only the faintest EOs in our sample would have luminosities comparable to those of globular clusters."
 Figure 1l shows calibrated Π baud miages of the sample fields. in which the EO. SO and ambiguous objects that have been detected are marked.," Figure 1 shows calibrated $R$ band images of the sample fields, in which the EO, SO and ambiguous objects that have been detected are marked."
 Below. we eive some coments on the galaxy fields:," Below, we give some comments on the galaxy fields:"
ihe Mach number).,the Mach number).
" Thev also assume that the characteristic GMC mass is determined by (he Jeans mass in the galaxy and it relates (ο X, by: llere. we propose a model for star formation which describes the formation of stars in dense protocluster forming clumps which are themselves embedded in lareer GAICS."," They also assume that the characteristic GMC mass is determined by the Jeans mass in the galaxy and it relates to $\Sigma_{g}$ by: Here, we propose a model for star formation which describes the formation of stars in dense protocluster forming clumps which are themselves embedded in larger GMCs."
 Star formation in (hese clumps is regulated by stellar winds feedback from the newly formed OB stars (of masses Z5 M)., Star formation in these clumps is regulated by stellar winds feedback from the newly formed OB stars (of masses $\gtrsim 5$ $_{\odot}$ ).
 In this model. dense cores form in the clump with a given core formation efficiency per Iree-[all time of the clump (CFE;;).," In this model, dense cores form in the clump with a given core formation efficiency per free-fall time of the clump $CFE_{ff}$ )."
 The level and driving scale of turbulence and the strength. of the magnetic fields influence the value of the CFEEjy;., The level and driving scale of turbulence and the strength of the magnetic fields influence the value of the $CFE_{ff}$.
 Dib et al. (, Dib et al. (
"2010a) showed that the CEEy, in a clump/cloud can vary from ~0.06 to 0.33 when the mass-to-magnetic flix ratio changes from joy=2.2 (o0 8.8 (n units of the critical value for collapse).",2010a) showed that the $CFE_{ff}$ in a clump/cloud can vary from $\sim 0.06$ to $0.33$ when the mass-to-magnetic flux ratio changes from $\mu_{B}=2.2$ to $8.8$ (in units of the critical value for collapse).
 This strong dependence on the strength: of (he magnetic field shows that a regulation of the star. formation efficiency which is based solely on turbulence. as postulated by IKMTOO is not entirely satisfactory., This strong dependence on the strength of the magnetic field shows that a regulation of the star formation efficiency which is based solely on turbulence as postulated by KMT09 is not entirely satisfactory.
 Dib et al. (, Dib et al. (
"2011) varied the value of the CEEj, between 0.1 and 0.3 to account [or variations due to different levels of turbulence and magnetic Ποια strengths in the clump.",2011) varied the value of the $CFE_{ff}$ between $0.1$ and $0.3$ to account for variations due to different levels of turbulence and magnetic field strengths in the clump.
 The cores have lifetimes of a few times their Iree-fall time after which they are turned into stus will a one core-to-one star efficiency conversion of 1/3., The cores have lifetimes of a few times their free-fall time after which they are turned into stars with a one core-to-one star efficiency conversion of 1/3.
 Once the cumulated effective kinetic energy [rom the stellar winds of the newly formed OD stars is larger (han (he gravitational binding energy of the clump. gas is expelled from the chimp aud subsequent core aud star formation are quenched.," Once the cumulated effective kinetic energy from the stellar winds of the newly formed OB stars is larger than the gravitational binding energy of the clump, gas is expelled from the clump and subsequent core and star formation are quenched."
 Dib et al. (, Dib et al. (
"2011) measured the final star formation efficiency. 5£E.,,. when gas is expelled from the protocluster chip and found that the value of 5FE, depends weakly on the value of the CFEy; but shows a slrong dependence on metallicity.","2011) measured the final star formation efficiency, $SFE_{exp}$, when gas is expelled from the protocluster clump and found that the value of $SFE_{exp}$ depends weakly on the value of the $CFE_{ff}$ but shows a strong dependence on metallicity."
 This dependence on metallicity is related to (he strong dependence of the wind power on metallicitv., This dependence on metallicity is related to the strong dependence of the wind power on metallicity.
 Stronger winds (ie. associated with higher metallicities) will result in (he gas being expelled from (he protocluster clump at an earlier epoch and thus reduce the value of SFE...," Stronger winds (i.e., associated with higher metallicities) will result in the gas being expelled from the protocluster clump at an earlier epoch and thus reduce the value of $SFE_{exp}$."
 The relevant mass-scale in our model is the characteristic mass of the protocluster clumps in the GAICs and not the characteristic mass of the GAICS as in IKMTO09., The relevant mass-scale in our model is the characteristic mass of the protocluster clumps in the GMCs and not the characteristic mass of the GMCs as in KMT09.
 The characteristic clump mass is given by: where ολο is given by Eq. 3.., The characteristic clump mass is given by: where $M_{GMC}$ is given by Eq. \ref{eq3}.
 M4; is the minimum mass of a protocluster chunp which we take to be 2.5x10? AL. (this guarantees. for final SFEs in the range of 0.05-0.3 a minimum mass for the stellar cluster of ~50 M...," $M_{cl,min}$ is the minimum mass of a protocluster clump which we take to be $2.5 \times 10^{3}$ $_{\odot}$ (this guarantees, for final SFEs in the range of 0.05-0.3 a minimum mass for the stellar cluster of $\sim 50$ $_{\odot}$ )."
" There might be a dependence of M,,,, on the"," There might be a dependence of $M_{cl,max}$ on the"
Extending these techniques to the full survey area will indicate whether these conclusions apply to the inner galaxy as a whole. and also give a much better determination of the distribution of cool cloud temperatures.,"Extending these techniques to the full survey area will indicate whether these conclusions apply to the inner galaxy as a whole, and also give a much better determination of the distribution of cool cloud temperatures."
 We are grateful to the ATNF stall for their hospitality and assistance while the SGPS data were taken., We are grateful to the ATNF staff for their hospitality and assistance while the SGPS data were taken.
 In particular we thank Robin Wark. John Reynolds. Mark Wieringa. and Dob Sault for their support and help.," In particular we thank Robin Wark, John Reynolds, Mark Wieringa, and Bob Sault for their support and help."
 We are grateful to Butler Burton. Carl IHeiles. Dain ]xavars. Michael Ixolpals. Harvey. Liszt. E. J. Lockman. Peter Martin. Simon Strasser. and {ο ihe anonvmous referee for valuable criticism and suggestions.," We are grateful to Butler Burton, Carl Heiles, Dain Kavars, Michael Kolpak, Harvey Liszt, F. J. Lockman, Peter Martin, Simon Strasser, and to the anonymous referee for valuable criticism and suggestions."
 We are grateful to Carl ILleiles and Tom Troland lor giving us their Arecibo data prior to publication., We are grateful to Carl Heiles and Tom Troland for giving us their Arecibo data prior to publication.
 This research was supported in part by the National Science Foundation (grant AST 97-32695 to the University of Minnesota) and NASA (GSRP erant NGT 5-50250 to N.AI-G.)., This research was supported in part by the National Science Foundation (grant AST 97-32695 to the University of Minnesota) and NASA (GSRP grant NGT 5-50250 to N.M-G.).
 Allen. R.J.. 2001. inEvolution. ASP Conlerence Proceedings. Vol.," Allen, R.J., 2001, in, ASP Conference Proceedings, Vol."
 240. eds.," 240, eds."
 J. E. Hibbard. M. Rupen. and J. IH. van Gorkom [San Francisco: Astronomical Society of the Pacific] p. Baker. P.L. and Burton. W.D. 1975. Ap.," J. E. Hibbard, M. Rupen, and J. H. van Gorkom [San Francisco: Astronomical Society of the Pacific] p. Baker, P.L. and Burton, W.B. 1975, Ap."
 J.. 198. Baker. P.L. and Burton. W.D.. 1979. A. and A. Supp..," J., 198, Baker, P.L. and Burton, W.B., 1979, A. and A. Supp.,"
 35.," 35,"
- Latest published papers or preprints. or conference - First-hand updated data ou professional websites.,"- Latest published papers or preprints, or conference - First-hand updated data on professional websites."
 Thev presently are: Anelo-Australian Planet Search aplavs.unsw.edthauscethtihhs California Planet Survey αἲ Geneva Extrasolar Planet Search Progrinnmes a/obswwwaige.ch/exoplaucts. Isepler candidates atlitial.. SuperWASP atSupcerwasp.org. an University of Texas - Dept.," They presently are: Anglo-Australian Planet Search at$\sim$, California Planet Survey at, Geneva Extrasolar Planet Search Programmes at, Kepler candidates at, SuperWASP at, and University of Texas - Dept."
 of Astronomy a/woww.as.utexas.edu. -, of Astronomy at. -
 Tn each individual pages there is a flag “ref for each data.," In each individual pages there is a flag ""ref"" for each data."
 By clicking on that flag the user is directly connected to the reference paper giving the corresponding value., By clicking on that flag the user is directly connected to the reference paper giving the corresponding value.
 These references are updated as they appear in the literature. -, These references are updated as they appear in the literature. -
 For planets detected by racial velocity and timing. ouly the product Misi. where tis the orbit inclination. is known in general.," For planets detected by radial velocity and timing, only the product $M.sin {\it i}$, where is the orbit inclination, is known in general."
 For transiting plauets.ὁ and therefore AM ds known from the fttiug of the transit liehteurve.," For transiting planets, and therefore $M$ is known from the fitting of the transit lightcurve."
 For planets detected by astrometryὁ is directlv iufered frou the pareut star orbit., For planets detected by astrometry is directly infered from the parent star orbit.
 For planets detected by radial velocity in inultiplauct svstenis. if can sometimes be iufored from the dyaizuuical analvsis of the planet-plauct interaction (deviation from purely kepleriau orbits - e.g.Yo Correia et al. (," For planets detected by radial velocity in multiplanet systems, it can sometimes be infered from the dynamical analysis of the planet-planet interaction (deviation from purely keplerian orbits - e.g. Correia et al. ("
2010) for the CJ στο svstem) and im a few vears it will be αίσχος from direct inagiug of some plauets. -,2010) for the GJ 876 system) and in a few years it will be infered from direct imaging of some planets. -
 When the seiiauajor axis e is not eiveu in a detection paper. it is derived from the published orbital period aud from the mass of the parent star through the Ixepler law P=28a?GAL. -magnitudes: ," When the semi-major axis $a$ is not given in a detection paper, it is derived from the published orbital period and from the mass of the parent star through the Kepler law $P=2\pi \sqrt{a^3/GM_{\star}}$. -"
If not eiven in a detection paper. they are taken from SIMBAD. -," If not given in a detection paper, they are taken from SIMBAD. -"
 The precise definition of A is the angle between the sky projections of the perpendicular to the planet orbit and of the star rotation axis (if the planet orbit is iu the star equatorial plane. A = 0j.," The precise definition of $\lambda$ is: the angle between the sky projections of the perpendicular to the planet orbit and of the star rotation axis (if the planet orbit is in the star equatorial plane, $\lambda$ = 0)."
 It is infered from the measurement of the Rossiter-MeLbaugllin effect., It is infered from the measurement of the Rossiter-McLaughlin effect.
 Some authors eive . instead of A. with §= A.," Some authors give $\beta$ instead of $\lambda$, with $\beta = -\lambda$ ."
 We then systematically convert the published ο iuto A=3j -, We then systematically convert the published $\beta$ into $\lambda = -\beta$ .-
the (2.2) to the (1.1) optical depths usiug: This formulation for rotation temperature requires modchne both the (1.1) aud (2.2) spectra.,"the (2,2) to the (1,1) optical depths using: This formulation for rotation temperature requires modeling both the (1,1) and (2,2) spectra."
 However. the area where we detect (2.2) emission is smaller than that with (1.1) cinission (see ??)).," However, the area where we detect (2,2) emission is smaller than that with (1,1) emission (see \ref{sec:morphology}) )."
" To estimate the colhuun density in regions with onlv (1.1) enussion we use the value of τω, from the regions where (2.2) cussion is at the threshold of detection (0. Z5,78.6 IX)."," To estimate the column density in regions with only (1,1) emission we use the value of $T_{rot}$ from the regions where (2,2) emission is at the threshold of detection (i.e., $T_{rot} \sim 8.6$ K)."
 Cohuun deusitv is obtained by integrating the opacity over velocity space., Column density is obtained by integrating the opacity over velocity space.
 From Manectal. (1992).. the column density of the (1.1) aud (2.2) transitions 1s: where Ae is the line width iu kins1. aud » is the frequency of the hyperfine transition iu GIIz.," From \citet{mangum1992}, the column density of the (1,1) and (2,2) transitions is: where $\Delta v$ is the line width in $\rm{km} \rm{s}^{-1}$, and $\nu$ is the frequency of the hyperfine transition in GHz."
 The total column density is then calculated with where g; (=2 for Az0 or 30) is the nuclear spin degeneracy. gy (=2 for Az (0) is the Is-degeneracy. and £C.IN) is the energy of the inversion state above the eround state.," The total column density is then calculated with where $g_I$ $= 2$ for $K \ne 0$ or $3n$ ) is the nuclear spin degeneracy, $g_K$ $= 2$ for $K \ne 0$ ) is the K-degeneracy, and $E(J,K)$ is the energy of the inversion state above the ground state."
 The rotational partition function is given by (Cordy&Cook 1970): where the rotational coustauts D aud C are 208117 and 186726 MITz. respectively.," The rotational partition function is given by \citep{gordy1970}: where the rotational constants B and C are $298117$ and $186726$ $\rm{MHz}$, respectively."
" The temperature of the envelope is considerably loss than Z,=11.5 Ik. reaching temperatures in excess of 20 EK only iu regions very near the central source."," The temperature of the envelope is considerably less than $T_o = 41.5$ K, reaching temperatures in excess of $20$ K only in regions very near the central source."
 In this low temperature regime we can solve the equation from Swiftetal.(2005). to obtain the kinetic temperature., In this low temperature regime we can solve the equation from \citet{swift2005} to obtain the kinetic temperature.
 An inteerated intensity (0th moment) map of the ceutral (1.1) cunission is shown in Figure 2..," An integrated intensity (0th moment) map of the central (1,1) emission is shown in Figure \ref{fig:intensity_all}."
 The iutensitv contours are shown at various detection strengths above the noise. separated by 26 and are superimposed ou the IT; v=10 5(1) line at 2.122 (narrow-band filter) map of ITirauoetal.(2006).. which shows the IIIT 211 protostellar jet.," The intensity contours are shown at various detection strengths above the noise, separated by $2 \sigma$ and are superimposed on the $\mbox{H}_2$ $\nu=1-0$ $S(1)$ line at 2.122 (narrow-band filter) map of \citet{hirano2006}, which shows the HH 211 protostellar jet."
 The NID; (1.1) cinission traces an elongated structure that is perpendicular to the jet outflow axis and shows a peak in inteerated intensity approximately 2.57 (~5/8 of a svuthesized beam width) from he protostellar source IIII 11αι. consistent with the NIT; (1.1) integrated inteusitv map of the same region shown by Wiseman (2001).," The $_3$ (1,1) emission traces an elongated structure that is perpendicular to the jet outflow axis and shows a peak in integrated intensity approximately 2.5"" $\sim 5/8$ of a synthesized beam width) from the protostellar source HH 211-mm, consistent with the $_3$ (1,1) integrated intensity map of the same region shown by Wiseman (2001)."
 We observe züuumnonia emission tracing au exteusion of gas eastward from the southern cud of the, We observe ammonia emission tracing an extension of gas eastward from the southern end of the
(e=0.75. 2=360 7=287) 2000).. (Schaefer&Schaefer.1988.2000:SchaeferTourtellotte2001).. (Schaeler&Seliaefer.1988.2000:Williams.Jones.Dus.&Larson.1989).. ~0.4(Schaefer&Schaefer2000).. (Duratti.Goguen.&Mosher1997:Brown.Webster.19093).. (Schaefer&Tourtellotte2001).. (Dobrovolskis1995).. (Ixlavetter.1989).,"$e=0.75$ $P=360$ $i=28 \degr$ \citep{scs95,scs00}. \citep{scs88,scs00,sct01}. \citep{scs88,scs00,wjt91,bul89}. $\sim 0.4$\citep{scs00}, \citep{bgm97,brw98}. \citep{sct01}. \citep{dob95}, \citep{kla89}."
. >154 (Schaeler&Tourtellotte|2001).., $>1\%$ \citep{sct01}.
is added.,is added.
 | show in Table | the mass of the systemM. the companion and the compact object (including the mass of the accretion disk) as a function of the value off.," I show in Table 1 the mass of the system, the companion and the compact object (including the mass of the accretion disk) as a function of the value of."
 Table | has been composed assuming that the material in the ring of fire is orbiting the centre of mass of the binary at 250 km s!., Table 1 has been composed assuming that the material in the ring of fire is orbiting the centre of mass of the binary at 250 km $^{-1}$.
 This value is model dependent: to the extent that the model is acctrate the radiating He ts orbiting slower than H. It is most unlikely that bulk material radiating 1n the inner rim of the circunibinary disk is orbiting slower than 200 km ∣8., This value is model dependent; to the extent that the model is accurate the radiating He is orbiting slower than H. It is most unlikely that bulk material radiating in the inner rim of the circumbinary disk is orbiting slower than 200 km $^{-1}$.
 For this speed. the system mass is 40 M. for à value of f of 2.3. as assuned in Blundell. Bowler Schmidtobreick (2008); material orbiting in the ring of fire at 250 km s would lie further inwards. f of 1.5.," For this speed, the system mass is 40 $M_\odot$ for a value of $f$ of 2.3, as assumed in Blundell, Bowler Schmidtobreick (2008); material orbiting in the ring of fire at 250 km $^{-1}$ would lie further inwards, $f$ of 1.5."
 Perhaps the Ha radiation is coming from icreasing amourts of material close to jotung the circumbinary disk rather than in stable orbits., Perhaps the $\alpha$ radiation is coming from increasing amounts of material close to joining the circumbinary disk rather than in stable orbits.
 In any event. the evidence from the circumbinary disk is that the systen Is massive. almost certainly exceeding 40 M. and the compact object is a rather massive stellar black hole.," In any event, the evidence from the circumbinary disk is that the system is massive, almost certainly exceeding 40 $M_{\odot}$, and the compact object is a rather massive stellar black hole."
 The Doppler shifts of a number of lines reported by Cherepashchuk et al. (, The Doppler shifts of a number of lines reported by Cherepashchuk et al. (
"2005) anc attributed to the companon yield an orbital velocity of the companion. about the binary tre of mass of 132 km s! and hence a system mass of 42 ο,",2005) and attributed to the companion yield an orbital velocity of the companion about the binary centre of mass of 132 km $^{-1}$ and hence a system mass of 42 $M_{\odot}$.
 This ts not in disagreement with data on the circumbinary κ. although f£ would have to be a little lessthan 1.7 for diδκ material moving as fast as 250 km s!.," This is not in disagreement with data on the circumbinary disk, although $f$ would have to be a little lessthan 1.7 for disk material moving as fast as 250 km $^{-1}$."
 The observations of Hillwig & Gies (2008). Kubota et al (2010). interpreted as absorption in the atmosphere of the companion. are not colsistent with the circumbinary disk modelled in this paper.," The observations of Hillwig $\&$ Gies (2008), Kubota et al (2010), interpreted as absorption in the atmosphere of the companion, are not consistent with the circumbinary disk modelled in this paper."
 They infer an orbital velocity for the companion of 58 km s~∣ alϱ ∣↴⊜∏∁⊜∐⋋⋝⇁⋋↾⊜⋯⋯⋅≏↧⋋⋋⋂↑↴∐⇂↴⋂⋯⊔↕∖∕↕⋮↜∣, They infer an orbital velocity for the companion of 58 km $^{-1}$ and hence a system mass of about 17 $_{\odot}$ .
∶⋯⋪↳∐⋋↥∖⊳⋯∐↾⊜∣⋪⇪∐∣ rb ≣⊓∏∶↔⊺∐↾⊐≤≺⊔∖⊳∏⊔⊣⋔≣⋋⋝⇁≣⊜∐⋋⋅≏↧∖⇁⋅≏↧∣⇂↧⊖⋂↑↴⊀∄∠⊽⋂↑↴∩↜⊱⊰≤− believably close in because the L2 radius corresponds to 1.26. (, For disk material orbiting at 250 km $^{-1}$ this yields a value of $f$ of 0.85 - unbelievably close in because the L2 radius corresponds to $f=1.26$. (
For material at 220 km s! f would be 1.1 ϱ even for 200 km s! an f of 1.33.),For material at 220 km $^{-1}$ $f$ would be 1.1 and even for 200 km $^{-1}$ an $f$ of 1.33.)
" However. it seems pc""sible that the absorption lines of Hillwig Gies (2008) alq Kubota et al (2010) are in fact produced in cireumbinary material."," However, it seems possible that the absorption lines of Hillwig Gies (2008) and Kubota et al (2010) are in fact produced in circumbinary material."
 The arguments against a cireumbinary origin for such absorption lines set out in Hillwig et al (2004) do not obviously apply to material in an orbiting disk., The arguments against a circumbinary origin for such absorption lines set out in Hillwig et al (2004) do not obviously apply to material in an orbiting disk.
 The observations of a sinusoidal oscillation of amplitude ~ 60 km s7!. sharing the orbital phase of the companion. are explained quantitatively if the origin of these lines is absorption of continuum light from the companion in eircumbinary material orbiting at ~ 240 km 67'," The observations of a sinusoidal oscillation of amplitude $\sim$ 60 km $^{-1}$, sharing the orbital phase of the companion, are explained quantitatively if the origin of these lines is absorption of continuum light from the companion in circumbinary material orbiting at $\sim$ 240 km $^{-1}$."
" The reason is that the companion presents an orbital radius projected on the sky of Resinó. where here ὁ is the orbital phase. and at this elongation is viewed through disk material moving with a radial component of velocity V, given by where the radius at which the eirceumbinary material orbits at speed Vy is Ry."," The reason is that the companion presents an orbital radius projected on the sky of $R_{\rm C}{\sin\phi}$, where here $\phi$ is the orbital phase, and at this elongation is viewed through disk material moving with a radial component of velocity $V_{\rm r}$ given by where the radius at which the circumbinary material orbits at speed $V_D$ is $R_D$."
 The radius of the companion orbit ts A/2 and the disk material orbits at ~2A., The radius of the companion orbit is $\sim A/2$ and the disk material orbits at $\sim 2A$.
 This is most easily illustrated by considering the configuration at an orbital phase of 0.25. when the companion is receding.," This is most easily illustrated by considering the configuration at an orbital phase of 0.25, when the companion is receding."
 At extreme elongatiot any circumbinary material through which it is seen has. in the centre of mass system of the binary. a recessional velocity of ~ 60 km s!.," At extreme elongation any circumbinary material through which it is seen has, in the centre of mass system of the binary, a recessional velocity of $\sim$ 60 km $^{-1}$."
 Perhaps it is just a coincidence. but the numerical agreement is as good as it could be.," Perhaps it is just a coincidence, but the numerical agreement is as good as it could be."
 The absorptior line data of Hillwig Gies (2008) and Kubota et al (2010) night thus be reconciled with the data attributed to the circunbinary disk which. as analysed in this paper. imply an orbital velocity for the companion exceeding 130 km s7!.," The absorption line data of Hillwig Gies (2008) and Kubota et al (2010) might thus be reconciled with the data attributed to the circumbinary disk which, as analysed in this paper, imply an orbital velocity for the companion exceeding 130 km $^{-1}$ ."
 If those absorption Ines do in fact originate in the circumbinary disk the identification of the spectral type of the companion as mid A ispremature., If those absorption lines do in fact originate in the circumbinary disk then identification of the spectral type of the companion as mid A ispremature.
2004).,.
. For instance. CH:OH ice has a strong feature at about 9.7 jim (see also Fig. 8)).," For instance, $_{3}$ OH ice has a strong feature at about 9.7 $\mu$ m (see also Fig. \ref{fig:sst18}) )."
 Since the spectra strong absorption bands at about 6 jm and 6.85 jm. caused by water ice and CH3:OH ice together with NH; ice respectively (see Fig. 2)).," Since the spectra strong absorption bands at about 6 $\mu$ m and 6.85 $\mu$ m, caused by water ice and $_{3}$ OH ice together with $_{4}^{+}$ ice respectively (see Fig. \ref{fig:spectra1}) ),"
 we expect the 9.7 um silicate features to be affected much by ice features., we expect the 9.7 $\mu$ m silicate features to be affected much by ice features.
 For the same reason we expecta strong H»O libration band at about 12 jm (Boogertetal..2008).., For the same reason we expecta strong $_{2}$ O libration band at about 12 $\mu$ m \citep{2008ApJ...678..985B}.
 Moreover. the long wavelength side of the 9.7 jm silicate features (from about 10.5 uim onward) in our molecular sightlines is very similar to that of the diffuse sightlines indicating that thereis indeed no H2O libration band present in any of the spectra.," Moreover, the long wavelength side of the 9.7 $\mu$ m silicate features (from about 10.5 $\mu$ m onward) in our molecular sightlines is very similar to that of the diffuse sightlines indicating that thereis indeed no $_{2}$ O libration band present in any of the spectra."
 In Fig. 12..," In Fig. \ref{fig:mol_galcent},"
 only shows a featurethat peaks at about 9.7 jim. that may be compatible with the CH;OH ice band.," only shows a featurethat peaks at about 9.7 $\mu$ m, that may be compatible with the $_{3}$ OH ice band."
 Figure 15. shows the difference between the 9.7 um feature of and the Galactic feature together with the absorption profile of a mixture of CH3OH and ΝΗ ice.," Figure \ref{fig:SSTc2dJ18..minGal_ice}
 shows the difference between the 9.7 $\mu$ m feature of and the Galactic feature together with the absorption profile of a mixture of $_{3}$ OH and $_{3}$ ice."
 In this case the observed difference could be explained by the presence of ices., In this case the observed difference could be explained by the presence of ices.
 On the other hand the spectrum of does not show an absorption band at 6.85 jm. caused by CH:OH ice together with NH; ice.," On the other hand the spectrum of does not show an absorption band at 6.85 $\mu$ m, caused by $_{3}$ OH ice together with $_{4}^{+}$ ice."
 For3-13. however. the difference between its 9.7 jim feature and the feature peaks at about 9 pm even if we correct for the photospheric gas phase SiO band.," For, however, the difference between its 9.7 $\mu$ m feature and the feature peaks at about 9 $\mu$ m even if we correct for the photospheric gas phase SiO band."
 Moreover. it is a relatively broad feature covering about 2—3 jm in wavelength.," Moreover, it is a relatively broad feature covering about 2–3 $\mu$ m in wavelength."
 There are. to the best of our knowledge. no ice species that produce such an absorption profile.," There are, to the best of our knowledge, no ice species that produce such an absorption profile."
 The shape of the 9.7 um silicate absorption feature depends on the dust properties of the amorphots silicates. such as the size. shape and chemical composition of the dust grains.," The shape of the 9.7 $\mu$ m silicate absorption feature depends on the dust properties of the amorphous silicates, such as the size, shape and chemical composition of the dust grains."
 The observed increase in absorption between 7.5 ànd 10.5 pm (1.8 the short wavelength side of the feattre) could. for instance. be the result of a decrease in theO/Si ratio in the anorphous silicates. gonπα from an olivine to pyroxehe composition 2008).," The observed increase in absorption between 7.5 and 10.5 $\mu$ m (i.e the short wavelength side of the feature) could, for instance, be the result of a decrease in theO/Si ratio in the amorphous silicates, going from an olivine to pyroxene composition ."
in order to take into account (some) possible reduction of the tidal torque.,in order to take into account (some) possible reduction of the tidal torque.
" In many previous studies, the gap-opening criterion is derived under the 2D approximation for discs."," In many previous studies, the gap-opening criterion is derived under the 2D approximation for discs."
" 'TTWO02, however, showed that the tidal torque in 3D is a factor of 2-3 smaller than that in 2D. Since the gap-opening mass involves the tidal torque, the mass for 3D discs can increase by at least a factor of 2 or 3."," TTW02, however, showed that the tidal torque in 3D is a factor of 2-3 smaller than that in 2D. Since the gap-opening mass involves the tidal torque, the mass for 3D discs can increase by at least a factor of 2 or 3."
 We therefore take C—3 here., We therefore take $C=3$ here.
 We leave a more detail study on it to a future publication., We leave a more detail study on it to a future publication.
" For viscous discs, a gap-opening mass is where a is the strength of turbulence."," For viscous discs, a gap-opening mass is where $\alpha$ is the strength of turbulence."
" The presence of viscous torque which plays a counteractive role in opening, that is, it closes a gap."," The presence of viscous torque which plays a counteractive role in gap-opening, that is, it closes a gap."
 Fig., Fig.
 13 (Top) shows the gap-opening mass as a function of the distance from the star., \ref{fig13} (Top) shows the gap-opening mass as a function of the distance from the star.
" In general, the mass derived from viscous discs is larger than that from inviscid discs."," In general, the mass derived from viscous discs is larger than that from inviscid discs."
" Both criteria for the dust settling case are smaller than those of the well mixed case, as expected."," Both criteria for the dust settling case are smaller than those of the well mixed case, as expected."
 This is a consequence of having flatter discs which arise from dust settling., This is a consequence of having flatter discs which arise from dust settling.
" We note that, even for the dust settling case, all planets we have considered (2, 5, 10 Ma) satisfy the criterion of type I migration (i.e. not open-up a gap) for r=1 autor=10 au."," We note that, even for the dust settling case, all planets we have considered (2, 5, 10 $M_{\oplus}$ ) satisfy the criterion of type I migration (i.e. not open-up a gap) for $r=1$ au to $r= 10$ au."
" We found that dust settling reduces the gap-opening mass at 10 au, respectively, for inviscid discs by about a factor of 2.5, for viscous discs by about a factor of 2."," We found that dust settling reduces the gap-opening mass at 10 au, respectively, for inviscid discs by about a factor of 2.5, for viscous discs by about a factor of 2."
" In the above analysis, there is no time-dependency."," In the above analysis, there is no time-dependency."
" simulated the interaction of gas and dust in discs in the presence of massive planets, using a two-fluid approach."," \citet{pm06b} simulated the interaction of gas and dust in discs in the presence of massive planets, using a two-fluid approach."
" Although they cannot include dust settling, since they considered 2D discs, they demonstrated that dust-gaps are more readily formed than gas-gaps."," Although they cannot include dust settling, since they considered 2D discs, they demonstrated that dust-gaps are more readily formed than gas-gaps."
" Furthermore, they found that dust removed from the location of planets is accumulated at mean-motion resonances."," Furthermore, they found that dust removed from the location of planets is accumulated at mean-motion resonances."
" Combined with such simulations, gap opening masses for dust settling are likely to be much lower."," Combined with such simulations, gap opening masses for dust settling are likely to be much lower."
" Once a planet opens up a gap in the disc, it undergoes Type II migration (?).."," Once a planet opens up a gap in the disc, it undergoes Type II migration \citep{ward97}."
 This type of migration is controlled by the viscous evolution of the gas in discs., This type of migration is controlled by the viscous evolution of the gas in discs.
" T'he timescale is generally written as where 2, is the angular frequency at r=ry (?)..", The timescale is generally written as where $\Omega_p$ is the angular frequency at $r=r_p$ \citep{t03b}.
" Thus, type II migration is also affected by dust settling via the reduced disc scaleheight (the bottom panel of Fig. 13))."," Thus, type II migration is also affected by dust settling via the reduced disc scaleheight (the bottom panel of Fig. \ref{fig13}) )."
 We found that dust settling increases the timescale by about a factor of 2 at 10 au., We found that dust settling increases the timescale by about a factor of 2 at 10 au.
" As a result, dust settling somewhat alleviates the problem of planetary migration in two respects: the reduction of the gap-opening mass and the slowing of type II migration."," As a result, dust settling somewhat alleviates the problem of planetary migration in two respects: the reduction of the gap-opening mass and the slowing of type II migration."
" In discs with dust settling, low mass planets undergo faster type I migration."," In discs with dust settling, low mass planets undergo faster type I migration."
" At the same time, such planets more readily open-up a gap, and switch to type II migration which is generally slower than type I. We have carefully examined Lindblad torques in radiatively heated discs."," At the same time, such planets more readily open-up a gap, and switch to type II migration which is generally slower than type I. We have carefully examined Lindblad torques in radiatively heated discs."
" In order to make our disc models realistic, we"," In order to make our disc models realistic, we"
Precise measurements of liming variations of strictly periodic| events have been successfully used. to infer the existence bodies orbiting around distant stars.,Precise measurements of timing variations of strictly periodic events have been successfully used to infer the existence bodies orbiting around distant stars.
 Perhaps the best-known examples are the pulsar planets of Wolszezan&Frail (1992)... à system of three extremely low-mass planets orbiting the milisecond. pulsar 8111257|12.," Perhaps the best-known examples are the pulsar planets of \citet{wf92}, a system of three extremely low-mass planets orbiting the millisecond pulsar PSR1257+12."
 Pulsating dwarf D stars have also been found to host planets. e.g. V391 g (Silvottietal.2007) and LIW Vir (Leeetal.2009).," Pulsating sub-dwarf B stars have also been found to host planets, e.g. V391 Peg \citep{silvotti07} and HW Vir \citep{lee09}."
". This same technique has recently been applied to eclipsing »olars. using t10 egress of the small. bright accretion spot as 8 precise ""clock."""," This same technique has recently been applied to eclipsing polars, using the egress of the small, bright accretion spot as a precise “clock.”"
 Unseen orbiting bodies can cause small shifts in the timing of ecipses (ranging [rom 10 sec o a few minues) due to the light-travel time cilferences imposed by the gravitational inlluence of the orbiting bodies on the svstem baryvcentre., Unseen orbiting bodies can cause small shifts in the timing of eclipses (ranging from $\sim$ 10 sec to a few minutes) due to the light-travel time differences imposed by the gravitational influence of the orbiting bodies on the system barycentre.
 Qianetal.(2010) discovered he first such planet orbiting the eclipsing polar DP Leo » combining their observed. eclipse timings with a data set. from Schwopeetal.(2002): the combined data set revealed a sinusoical variation indicativo of a planet with a period of 23.8 vears., \citet{qian10} discovered the first such planet orbiting the eclipsing polar DP Leo by combining their observed eclipse timings with a data set from \citet{schwope02}; the combined data set revealed a sinusoidal variation indicative of a planet with a period of 23.8 years.
 The timing method. has shown that planets can orbit stars which are wildly cdillerent. from. the main-sequence solar-tv stars most commonly targeted. by Doppler ancl transit. pe.planet search programs., The timing method has shown that planets can orbit stars which are wildly different from the main-sequence solar-type stars most commonly targeted by Doppler and transit planet search programs.
" ln a recent letter, Qianetal.(2011) announced the discovery of two (and. potentially more) giant. planets orbiting the eclipsing polar LIU Aquarii (hereafter LLU Λα)."," In a recent letter, \citet{Qian2011} announced the discovery of two (and potentially more) giant planets orbiting the eclipsing polar HU Aquarii (hereafter HU Aqr)."
 The authors provided. fits to the orbits of those planets. placing them at orbital radii of 3.6 and 5.4 AU. from the system barvcentre. and ascribed minimum masses of 5.9 and 45 Morep(respectively) to the two bodies.," The authors provided fits to the orbits of those planets, placing them at orbital radii of 3.6 and 5.4 AU, from the system barycentre, and ascribed minimum masses of 5.9 and 4.5 (respectively) to the two bodies."
 In this Letter. we perform a detailed: cvnamical analysis of the LIU αν planetary system in order to assess the stability of. the proposed. planet. candidates.," In this Letter, we perform a detailed dynamical analysis of the HU Aqr planetary system in order to assess the stability of the proposed planet candidates."
 In their study of the eclipsing polar svstem HU Aqr. Qianet consider the temporal variation in the observed," In their study of the eclipsing polar system HU Aqr, \citet{Qian2011} consider the temporal variation in the observed"
cluperatures. particularly bluewiurd ofwicron.,"temperatures, particularly blueward of."
. This is simply because less dust condenses at higher temperatures. so it makes less difference whether that stavs in suspension in the atinospliere or is oxecipitated out.," This is simply because less dust condenses at higher temperatures, so it makes less difference whether that stays in suspension in the atmosphere or is precipitated out."
 The differeuces are enbliauced at the lower effective temperature TIN) for the converse reason. nuuore condensation of eas-phase species (aud their oxecipitation iu the coud models).," The differences are enhanced at the lower effective temperature K) for the converse reason, more condensation of gas-phase species (and their precipitation in the cond models)."
 The same bands as described above for IIS cau ο used at 210011. but now the cloud signatures are reduced by a factor of four iu the jl aud j2 bands.," The same bands as described above for K can be used at K, but now the cloud signatures are reduced by a factor of four in the j1 and j2 bands."
 For a cloud coverage. this is probably below the level at which systematic errors can be controlled frou eromud-based observations.," For a cloud coverage, this is probably below the level at which systematic errors can be controlled from ground-based observations."
 However. at IKK. the amplitude of the cool spot signature is still relatively large. aud aüucnable to detection.," However, at K, the amplitude of the cool spot signature is still relatively large, and amenable to detection."
 At IIs. the cloud signaturesoO are C»eveathly eubauced and are mich lavecr than the cool spot signatures. such iat the hl colour has an amplitude of Linmags due o a cloud aud ouly uunags due to a cool spot (both covering of the surface).," At K, the cloud signatures are greatly enhanced and are much larger than the cool spot signatures, such that the $-$ h1 colour has an amplitude of mags due to a cloud and only mags due to a cool spot (both covering of the surface)."
 However. it ust be recalled iat below hs UCDs are ecucrally better fit by coud tinospheres. auc dusty clouds on such atinosphleres slow much weaker signature at ISIS (Fie. 1).," However, it must be recalled that below K UCDs are generally better fit by cond atmospheres, and dusty clouds on such atmospheres show a much weaker signature at K (Fig. \ref{var4}) )."
 Moreover. 16 cloud signature is probably scusitive to the of aust precipitation at these intermediate temperatures aud Us is not accounted for in the preseut models (see sections ] aud 3.1)).," Moreover, the cloud signature is probably sensitive to the of dust precipitation at these intermediate temperatures and this is not accounted for in the present models (see sections \ref{intro} and \ref{models}) )."
 I have focused on signatures in the NIR region as that is the region observed and analysed in the rest of this paper., I have focused on signatures in the NIR region as that is the region observed and analysed in the rest of this paper.
 However. there are strong signatures in the optical spectruii. as can be seeu in rofvirb- d..," However, there are strong signatures in the optical spectrum, as can be seen in \\ref{var5}- \ref{var4}."
 There is a trade-off here. of course. because ultra cool dwarfs are much fainter iu the optical.," There is a trade-off here, of course, because ultra cool dwarfs are much fainter in the optical."
 The region bbehliaves in a similar fashion to the jl baud and the region sshows an even strouger siguature across a range of effective temperatures., The region behaves in a similar fashion to the j1 band and the region shows an even stronger signature across a range of effective temperatures.
 Very narrow bands centered on the nicron)) aud uicron)) atomic absorption Lines are in principle a oeood discriminant vetween clouds and cool spots when comibined with a baud in the rangeuicron.. but the narrowness required (~0.02 wicron)) to avoid the opposite signature in the line wines nay preclude their use frou a signal-to-noise ratio (SNR) »oiut of view.," Very narrow bands centered on the ) and ) atomic absorption lines are in principle a good discriminant between clouds and cool spots when combined with a band in the range, but the narrowness required $\sim 0.02$ ) to avoid the opposite signature in the line wings may preclude their use from a signal-to-noise ratio (SNR) point of view."
 The 1 oeifrared appears to contain no wavelength regions Which are of particular advantage over the near infrared for cistineuishine cool spots from dust-rclated clouds. at least not at around TIS (Fig. 53).," The mid infrared appears to contain no wavelength regions which are of particular advantage over the near infrared for distinguishing cool spots from dust-related clouds, at least not at around K (Fig. \ref{var6}) )."
 However. at lower teimiperatures iiethane forms. aud there remaius the interesting possibility of looking for methane clouds either in the near or nud infrared (Bailer-Joues. iu preparation).," However, at lower temperatures methane forms, and there remains the interesting possibility of looking for methane clouds either in the near or mid infrared (Bailer-Jones, in preparation)."
 The analyses of the previous section were motivated by 16 desire to interpret infrared spectra of the L1.5 dwart QATLLLS., The analyses of the previous section were motivated by the desire to interpret infrared spectra of the L1.5 dwarf 2M1145.
 Spectra were obtained using the Omega Cass infrared dager/spectrograph mounted ou the Casscerain ocus of the 3.5m telescope at Calar Alto. Spain.," Spectra were obtained using the Omega Cass infrared imager/spectrograph mounted on the Cassegrain focus of the 3.5m telescope at Calar Alto, Spain."
 The detector 1:8 a «1021 WeCdTe ILAWAILET focal plane uray sensitive hetween 0.8 aud wicron., The detector is a $\times$ 1024 HgCdTe HAWAII-1 focal plane array sensitive between 0.8 and .
". The low resolution camera was used. eiviug iu oenage scale oC 0,37 Apis."," The low resolution camera was used, giving an image scale of $''$ /pix."
 The resolution 120 (= A/AA) eris was used in combination with one of two filters: With a JI filter it provided the waveleneth range minoverlapped in first order at a dispersion of /pix (plus a second order which was not used): with an II filter the wavelength range wwas obtained in first order at the same dispersion., The resolution 420 $=\lambda/\Delta \lambda$ ) grism was used in combination with one of two filters: With a JH filter it provided the wavelength range unoverlapped in first order at a dispersion of /pix (plus a second order which was not used); with an HK filter the wavelength range was obtained in first order at the same dispersion.
" The slit had a width of 0.9"" and leneth of 111"".", The slit had a width of $''$ and length of $''$.
" The mstrunient owas oriented such that the spectroeraph slit was at an augle 10.57East of North. to include both 2NI1115 and the brieht reference star 79"" to"," The instrument was oriented such that the spectrograph slit was at an angle $10.8^{\circ}$East of North, to include both 2M1145 and the bright reference star $''$ to"
of20%... the halo dominates the spheroid [rom 2—Ry to R=dkpe. which latter is about the limit to which the local spheroid density. profile can be reliably extrapolated.,"of, the halo dominates the spheroid from $R=R_0$ to $R=4\,\kpc$, which latter is about the limit to which the local spheroid density profile can be reliably extrapolated."
 However. this domination is not overwhelming: at =4kpc it is only a [actor of 5 aud even at fà=7kpc (where the halo optical depth has fallen by. a factor 5) the halo only dominates by a factor 10.," However, this domination is not overwhelming: at $R=4\,\kpc$ it is only a factor of 5 and even at $R=7\,\kpc$ (where the halo optical depth has fallen by a factor 5) the halo only dominates by a factor 10."
" This means that 2 or 3 halo lenses would have to be idenüfied to constitute a reliable “MACTIO detection"".", This means that 2 or 3 halo lenses would have to be identified to constitute a reliable “MACHO detection”.
 Otherwise. there would be a significant possibility that spheroid lenses were responsible.," Otherwise, there would be a significant possibility that spheroid lenses were responsible."
 Since onemust restrict attention to D;«4kpce. the total available halo optical depth is reduced by a [actor 0.4 relative to the 1.3x10* caleulated by Griestetal.(1991).," Since onemust restrict attention to $D_l<4\,\kpc$, the total available halo optical depth is reduced by a factor $0.4$ relative to the $1.3\times 10^{-7}$ calculated by \citet{griest91}."
". If we further assume f=20%. the available halo optical depth is further reduced to LO. about of the observed optical depth of 7~2x10"" (Alonsoοἱal.2003b:etal. 2005)."," If we further assume $f=20\%$, the available halo optical depth is further reduced to $10^{-8}$, about of the observed optical depth of $\tau\sim 2\times 10^{-6}$ \citep{afonso03b,popow05,sumi05}."
. Hence. assuming for the moment that the event rates are in proportion to the optical depths. roughly 200 measurements would be required to identify a single halo lens.," Hence, assuming for the moment that the event rates are in proportion to the optical depths, roughly 200 measurements would be required to identify a single halo lens."
 Thus. if theS/A/ mission were extended [rom 5 to 10 vears (as is currently envisioned) ihen one might expect to find about 2 halo lenses.," Thus, if the mission were extended from 5 to 10 years (as is currently envisioned) then one might expect to find about 2 halo lenses."
 As noted above this is just ad the margin ol a viable detection., As noted above this is just at the margin of a viable detection.
 The Einstein timescales of these halo events are given by. where ο is the transverse lens velocity relative to the observer-source line of sight.," The Einstein timescales of these halo events are given by, where $v_\perp$ is the transverse lens velocity relative to the observer-source line of sight."
 Thus. for the mass range advocated by Aleockοἱal.(2000).. the tvpical event timescales will be fairly short.," Thus, for the mass range advocated by \citet{alcock00}, the typical event timescales will be fairly short."
 This is important. because the event must be identified ancl alerted to the satellite well before peak in order (ο measure πι (Gould&Salim1999)., This is important because the event must be identified and alerted to the satellite well before peak in order to measure $\pi_\e$ \citep{gs99}.
. IIence. a fairly aggressive posture is required to keep the halo events in the sample.," Hence, a fairly aggressive posture is required to keep the halo events in the sample."
 llowever. the fact that these halo events are somewhat shorter (han (vpical bulge events means (hat (μον are also more frequent than would be inclicated by their optical depth alone.," However, the fact that these halo events are somewhat shorter than typical bulge events means that they are also more frequent than would be indicated by their optical depth alone."
 That is. the event rate Dx7//i. so the rate is inversely proportional to the timescale.," That is, the event rate $\Gamma\propto \tau/t_\e$, so the rate is inversely proportional to the timescale."
 llence. (he shorter timescales enhances (the viability of a given experiment relative to what was discussed in 3.. provided that not too many halo events are lost because they are too short.," Hence, the shorter timescales enhances the viability of a given experiment relative to what was discussed in \ref{sec:spheroid}, , provided that not too many halo events are lost because they are too short."
routine*.,.
. | consider the cases rp=0 and np=—5/2 as representative of ultra-violet and infra-red fields respectively., I consider the cases $n_B=0$ and $n_B=-5/2$ as representative of ultra-violet and infra-red fields respectively.
 Grid dependences were tested employing a range of numbers of samples in the chains: convergence was found to occur at approximately 2~10° for ij=Oand à~107 for ng=—5/2., Grid dependences were tested employing a range of numbers of samples in the chains; convergence was found to occur at approximately $n\sim 10^6$ for $n_B=0$ and $n\sim 10^7$ for $n_B=-5/2$ .
 Results were found using à=5x10° for ig=O and n=5x10! for ng=—5/2., Results were found using $n=5\times 10^6$ for $n_B=0$ and $n=5\times 10^7$ for $n_B=-5/2$.
 The exception is the (ry) correlation where 7=5xI0? samples were needed for convergence across the entire plane due to the complexity of the integrand., The exception is the $\av{\tau_T^3}$ correlation where $n=5\times 10^8$ samples were needed for convergence across the entire plane due to the complexity of the integrand.
 The grid the bispectra are integrated on also varies with 725 but is relatively sparse in &. particularly for ng=—5/2 where the most important area is k«| which is expected to exhibit a scaling behaviour.," The grid the bispectra are integrated on also varies with $n_B$ but is relatively sparse in $k$, particularly for $n_B=-5/2$ where the most important area is $k\ll 1$ which is expected to exhibit a scaling behaviour."
 This sparse sampling ts introduced to speed the numeries: for low & the integration volume is large and integrations take quite some time. especially at low 75.," This sparse sampling is introduced to speed the numerics; for low $k$ the integration volume is large and integrations take quite some time, especially at low $n_B$."
 The angle ὁ is sampled once every five degrees except near the colinear and squeezed planes where it is sampled once every half a degree. and around the equilateral line for (i) which for 52—5/2 contains a feature of particular interest.," The angle $\phi$ is sampled once every five degrees except near the colinear and squeezed planes where it is sampled once every half a degree, and around the equilateral line for $\av{\tau_S^3}$ which for $n_B=-5/2$ contains a feature of particular interest."
 I consider r€[1.5].," I consider $r\in[1,5]$."
 Previous studies of the magnetic bispectra have typically been of scalar correlations. not least as it is only recently that a clear formalism for generating the CMB angular bispectra from vector or tensor perturbations has emerged ??..," Previous studies of the magnetic bispectra have typically been of scalar correlations, not least as it is only recently that a clear formalism for generating the CMB angular bispectra from vector or tensor perturbations has emerged \cite{Shiraishi:2010sm,Shiraishi:2010kd}."
 They have also been in {k.p.q1 space.," They have also been in $\{k,p,q\}$ space."
 Employing instead (4.r.0] space has some benefits. the chief of which is the extent to which it highlights the scalings in Κ.," Employing instead $\{k,r,\phi\}$ space has some benefits, the chief of which is the extent to which it highlights the scalings in $k$."
 In particular. the bispectra sampled along lines of constant [r.ó] are expected to obey a power-law ink for kxkeon.," In particular, the bispectra sampled along lines of constant $\{r,\phi\}$ are expected to obey a power-law in $k$ for $k\lesssim k_\mathrm{Coh}$."
 The bispectrum can then be modelled in a manner similar to the power spectra in. for example. BIO as a plane at constant & and a given scaling relation.," The bispectrum can then be modelled in a manner similar to the power spectra in, for example, B10 as a plane at constant $k$ and a given scaling relation."
 It also simplifies the construction of the complete bispectrum somewhat. in that we can integrate cleanly along lines of constant {7} and assemble these lines into a 3D bispectrum.," It also simplifies the construction of the complete bispectrum somewhat, in that we can integrate cleanly along lines of constant $\{r,\phi\}$ and assemble these lines into a 3D bispectrum."
 The other main benefit of employing this coordinate system is that it results in a fine sampling around the 72| plane: from the results of SS09 and CFPRO9 and the decaying of the bispectrum às r>cc. the degenerate line should dominate and other configurations in and close to the »=1 plane will be more significant than those further away.," The other main benefit of employing this coordinate system is that it results in a fine sampling around the $r=1$ plane; from the results of SS09 and CFPR09 and the decaying of the bispectrum as $r\rightarrow\infty$, the degenerate line should dominate and other configurations in and close to the $r=1$ plane will be more significant than those further away."
 To get the full bispectrum in {k.p.q] space one can transform the results.," To get the full bispectrum in $\{k,p,q\}$ space one can transform the results."
 Since they obey the triangle relations. bispectra in {k.p.d] space are wedge-shaped.," Since they obey the triangle relations, bispectra in $\{k,p,q\}$ space are wedge-shaped."
 Transforming to this space yields a cone filling one third of the wedge., Transforming to this space yields a cone filling one third of the wedge.
 The other two thirds of the bispectrum in {k.p.q] space can be found from permuting across &. p and q.," The other two thirds of the bispectrum in $\{k,p,q\}$ space can be found from permuting across $k$, $p$ and $q$."
 There are. however. drawbacks to the coordinates [K.7.0]. the chief being a lack of immediate clarity.," There are, however, drawbacks to the coordinates $\{k,r,\phi\}$, the chief being a lack of immediate clarity."
 In particular. the q77— sealing found near to the degenerate line for ng<—1 in CFPRO9 is obscured.," In particular, the $q^{2n_B+3}$ scaling found near to the degenerate line for $n_B<-1$ in CFPR09 is obscured."
 Further. reconstructing the bispectrum in [K.p.q] space will be slightly complicated when considering cross-correlations such as (r(or(pors(qQ)) since the anisotropy along the q vector must also be taken into account whenperforming the permutations.," Further, reconstructing the bispectrum in $\{k,p,q\}$ space will be slightly complicated when considering cross-correlations such as $\av{\tau(\mathbf{k})\tau(\mathbf{p})\tau_S(\mathbf{q})}$ since the anisotropy along the $\mathbf{q}$ vector must also be taken into account whenperforming the permutations."
 However. I feel that the convenience with which the bispectra can be found and. more importantly. the clean scalings in & which should be recovered outweigh these disadvantages.," However, I feel that the convenience with which the bispectra can be found and, more importantly, the clean scalings in $k$ which should be recovered outweigh these disadvantages."
 The bispectra are assembled with the following procedure:, The bispectra are assembled with the following procedure:
the + compoucut of the dipole.,the $z$ component of the dipole.
 This then sugeests to use a 47 method to fit a dipolar distributionalong the + direction to ect an unbiased estimate of o., This then suggests to use a $\chi^2$ method to fit a dipolar distributionalong the $\hat z$ direction to get an unbiased estimate of $a_z$.
 To fit a dipole sigual aloug a generic direction d. oue cau just take nv. bius of cos5 (e.g. n.=10 is a reasonable choice. aud we checked that using a larger value does not improve significantly: the results). aud compute the umber of events in each bin IN; aud the correspouding expected values where AQ; is the solid angle forwhich coss! falls in the j-th bin. with LN/dO eiven by Eq. (3)).," To fit a dipole signal along a generic direction $\hat d'$ , one can just take $n_\gamma$ bins of $\cos\gamma'$ (e.g. $n_\gamma=10$ is a reasonable choice, and we checked that using a larger value does not improve significantly the results), and compute the number of events in each bin $N_j$ and the corresponding expected values where $\Delta\Omega_j$ is the solid angle forwhich $\cos\gamma'$ falls in the $j$ -th bin, with ${{\rm d}N}/{{\rm d}\Omega}$ given by Eq. \ref{dndom}) )."
 We can then write the u function as where the sum is clearly restricted to the with JN;>0., We can then write the $\chi^2$ function as where the sum is clearly restricted to the with ${\bar N}_j>0$.
" One cau then determine the value 6,, for which v) is uinimized. aud obtain Ad such that ΓΕ”..."," One can then determine the value $a_m$ for which $\chi^2(a)$ is minimized, and obtain $\Delta a$ such that $\chi^2(a_m+\Delta a)=\chi^2(a_m)+1$ ."
 We hence apply this method along the : direction to obtain first ανν, We hence apply this method along the $\hat z$ direction to obtain first $a_z$.
 Concerning he dipole component orthogonal to the + axis. the natural strategy is to apply Ravleigl’s method.," Concerning the dipole component orthogonal to the $\hat z$ axis, the natural strategy is to apply Rayleigh's method."
" Tn particular. it has been noticed in 10] that this method eives, at least in the case of partial «kv coverage. a better determination of he dipole’s right asceusiou than the SAP method."," In particular, it has been noticed in \cite{au05} that this method gives, at least in the case of partial sky coverage, a better determination of the dipole's right ascension than the SAP method."
 Oue can use then Ravleigh’s ucthod to get the best estimate of the dipoles right ascension aud first harmonic ;uuplitude., One can use then Rayleigh's method to get the best estimate of the dipole's right ascension and first harmonic amplitude.
 This consists simply of computing the quautities (iere a is the right ascension) and then eet the first harmonic amplitude + and phase V through Iu addition. as shown in ref. |5]..," This consists simply of computing the quantities (here $\alpha$ is the right ascension) and then get the first harmonic amplitude $r$ and phase $\Psi$ through In addition, as shown in ref. \cite{au05}, ,"
" there is asimple relation between r aud the original dipole coumpoucuts e; aud e,=«cos iu the case inwhich the exposure is independent of à. which is"," there is asimple relation between $r$ and the original dipole components $a_z$ and $a_\perp\equiv a\cos\delta$ in the case inwhich the exposure is independent of $\alpha$ , which is"
slit. aud achieved typical resolution of 1-0 FFWHM. from 1615 to 6960À.,"slit, and achieved typical resolution of 4-5 FWHM from 4645 to 6960."
. On earlier observing runs. we took spectra of comparison lamps whenever the telescope was moved to achieve accurate wavelength calibration. but on later runs we used the uieht-sky spectrum to shift the zero point of a master wavelength calibration obtained during twilight.," On earlier observing runs, we took spectra of comparison lamps whenever the telescope was moved to achieve accurate wavelength calibration, but on later runs we used the night-sky spectrum to shift the zero point of a master wavelength calibration obtained during twilight."
 When the sky was clear. we observed bright early-type stars to derive corrections for the telluric absorption features. and flux standards to calibrate the instrument response.," When the sky was clear, we observed bright early-type stars to derive corrections for the telluric absorption features, and flux standards to calibrate the instrument response."
 Our spectrophotometry suffers from. uncalibratable errors due to losses at the spectrograph slit. but experience suggests that these typically amount to ~0.2 mag.," Our spectrophotometry suffers from uncalibratable errors due to losses at the spectrograph slit, but experience suggests that these typically amount to $\sim 0.2$ mag."
 Thorsteusen.(2001.hereafterTETOI). discuss the observational protocols in more detail., \citet[hereafter TFT04]{thfe04} discuss the observational protocols in more detail.
 We measurecl centroids for the emission lines using convolution techuiques explained by aud Shafter(1983).. in which the line profile is couvolved with an autisyinnetric [uuction and the zero of the convolution is taken as the line center.," We measured centroids for the emission lines using convolution techniques explained by \citet{sy80} and \citet{shafter83}, in which the line profile is convolved with an antisymmetric function and the zero of the convolution is taken as the line center."
 For the most part we used a convolution fuuction consisting of a positive Gaussian auc a negative Gatssian separated by au adjustable amount., For the most part we used a convolution function consisting of a positive Gaussian and a negative Gaussian separated by an adjustable amount.
 We usually set the separation of the Caussiaus to maximize sensitivity to the steep sides of the line profile., We usually set the separation of the Gaussians to maximize sensitivity to the steep sides of the line profile.
 The Roche geometry dictates that longer-period CVs will have larger and. generally more uiassive secoudary stars than sliorter-period CVs. so their secondary stars olteu contribute siguilicautly to the total light.," The Roche geometry dictates that longer-period CVs will have larger and generally more massive secondary stars than shorter-period CVs, so their secondary stars often contribute significantly to the total light."
 When possible. we measured absorption liue velocities using the cross-correlatiou radial velocity packageπόδας (Ixurtz&Mink19985).," When possible, we measured absorption line velocities using the cross-correlation radial velocity package \citep{kurtzmink}."
. In most. cases the template spectrum was a composite of 76 observations of late-type LAU velocity standard stars. which had been shifted iudividually to zero velocity before averaging: for later-type secoucaries we used a similar template constructed from M dwarls.," In most cases the template spectrum was a composite of 76 observations of late-type IAU velocity standard stars, which had been shifted individually to zero velocity before averaging; for later-type secondaries we used a similar template constructed from M dwarfs."
 The spectral region used in tle cross-correlation was eitlier 6000 to 6500A.. or a window from 5100 to 6500 A.. excluding the complicated region near Hel A5876 aud Nab. Uncertaiuties. based ou the fi-statistic of Toury&Davis(1979).. were typically ou the order of 20 km |.," The spectral region used in the cross-correlation was either 6000 to 6500, or a window from 5100 to 6500 , excluding the complicated region near HeI $\lambda5876$ and NaD. Uncertainties, based on the $R$ -statistic of \citet{tondav}, were typically on the order of 20 km $^{-1}$."
" We omitted tuplysically large velocities (which result when the ""wrong! cross-correlation peak is selected) aud. velocities with large estimated uncertainties.", We omitted unphysically large velocities (which result when the `wrong' cross-correlation peak is selected) and velocities with large estimated uncertainties.
 We analyzed the velocity time series with a residualgram' period-search algorithun 1996) which fits general least-squares sinusoids of the form. to the time series over a deusely-spaced erid of trial frequencies w., We analyzed the velocity time series with a `residualgram' period-search algorithm \citep{tpst} which fits general least-squares sinusoids of the form to the time series over a densely-spaced grid of trial frequencies $\omega$.
 The results are presented as plots of 1/47. where 47—Ye; e(/)/o;: here ej is a measured velocity. e(/;) is the sinusoikl evaluated at the correspoucing time /;; amd σι is the formal uncertainty of the velocity. obtained by propagating the estimated couuting-statisties uncertainties of the individual wavelength channels.," The results are presented as plots of $\chi^2$, where $\chi^2 = \sum[(v_i - v(t_i))/\sigma_i]^2$ ; here $v_i$ is a measured velocity, $v(t_i)$ is the best-fitting sinusoid evaluated at the corresponding time $t_i$, and $\sigma_i$ is the formal uncertainty of the velocity, obtained by propagating the estimated counting-statistics uncertainties of the individual wavelength channels."
 Ideally. a graph of 1/4? as a function of à showsa strong spike at the orbital [requeney. but theimperfect sampling of the time series inevitably results in aliasing.," Ideally, a graph of $1/\chi^2$ as a function of $\omega$ showsa strong spike at the orbital frequency, but theimperfect sampling of the time series inevitably results in aliasing."
 We used a, We used a
we note that synchrotron self-absorption has been roughly computed supposing an homogeneous emitter as wide as long.,we note that synchrotron self-absorption has been roughly computed supposing an homogeneous emitter as wide as long.
" The magnetic field where most of the emission «1 GHz is produced, at a distance of several AU from the star, is ~1 G, in agreement with the values inferred in Sect. ??,,"," The magnetic field where most of the emission $<
1$ GHz is produced, at a distance of several AU from the star, is $\sim 1$ G, in agreement with the values inferred in Sect. \ref{abs},"
 and also compatible with those in the wind at few AU from a massive star (Benaglia Romero 2003))., and also compatible with those in the wind at few AU from a massive star (Benaglia Romero \cite{benaglia03}) ).
" Concerning the SE and NW components, the larger vaa, would point to a different physical structure than that producing the radio core, perhaps a real jet."," Concerning the SE and NW components, the larger $v_{\rm adv}$ would point to a different physical structure than that producing the radio core, perhaps a real jet."
 Small changes in the orientation of ~12? of this jet-like structure have been found (R08)., Small changes in the orientation of $\approx 12^{\circ}$ of this jet-like structure have been found (R08).
" This might be explained, for instance, by an asymmetry introduced by the secondary pairs transported in the stellar wind (B08), the propagation of which might not be spherical, or by jet-stellar wind interactions (Perucho Bosch-Ramon 2008))."," This might be explained, for instance, by an asymmetry introduced by the secondary pairs transported in the stellar wind (B08), the propagation of which might not be spherical, or by jet-stellar wind interactions (Perucho Bosch-Ramon \cite{perucho08}) )."
" About a of the radiation at 5 GHz comes from a jet- structure of projected size ~10—1000 AU detected with VLBA, MERLIN and EVN (PAO; Paredes et al. 2002;;"," About a of the radiation at 5 GHz comes from a jet-like structure of projected size $\sim 10-1000$ AU detected with VLBA, MERLIN and EVN (PA0; Paredes et al. \cite{paredes02};"
 RO8)., R08).
" These extended structures may be associated with the 234 MHz high-state emission, due to the size constraints for this region inferred above."," These extended structures may be associated with the 234 MHz high-state emission, due to the size constraints for this region inferred above."
" The seen at 234 MHz may be related. to changes variabilityin the outflow properties (e.g. via hydrodynamical instabilities), but not to long-term variations in the primary power supply, since it should affect as well the core radio emission."," The variability seen at 234 MHz may be related to changes in the outflow properties (e.g. via hydrodynamical instabilities), but not to long-term variations in the primary power supply, since it should affect as well the core radio emission."
" In this framework, the 234 MHz high state should have associated flux variations at higher frequencies visible outside the radio core."," In this framework, the 234 MHz high state should have associated flux variations at higher frequencies visible outside the radio core."
" Given that t,gy depends on the size and vagq,, the MERLIN and EVN ~100—1000 AU structures at 5 GHz would show variability at timescales of at least months."," Given that $t_{\rm adv}$ depends on the size and $v_{\rm adv}$, the MERLIN and EVN $\sim
100-1000$ AU structures at 5 GHz would show variability at timescales of at least months."
 This 234 MHz/large scale jet link can be tested with simultaneous radio VLBI observations at low and high frequencies taken in two different epochs., This 234 MHz/large scale jet link can be tested with simultaneous radio VLBI observations at low and high frequencies taken in two different epochs.
" We note that a radio emitter of r~1000 AU varying with ~1 yr timescales would have associated a vag,210? cm s!, compatible with the vaqy derived for the VLBA SE and NW components."," We note that a radio emitter of $r\sim 1000$ AU varying with $\sim 1$ yr timescales would have associated a $v_{\rm adv}\ga
10^9$ cm $^{-1}$, compatible with the $v_{\rm adv}$ derived for the VLBA SE and NW components."
" A plausible scenario for the radio emission in LS 5039 would be radiation from secondary pairs created via gamma-ray absorption (see B08), for the compact emitter or core (R few AU), plus a collimated jet, for the resolved emission."," A plausible scenario for the radio emission in LS 5039 would be radiation from secondary pairs created via gamma-ray absorption (see B08), for the compact emitter or core $R\sim$ few AU), plus a collimated jet, for the resolved emission."
" The jet radio emission from the smaller spatial scales would be probably flat, as typically expected in low-hard state microquasars (e.g. Fender 2001)), but the optically thin radiation from secondary pairs would overcome this component, at least up to few AU."," The jet radio emission from the smaller spatial scales would be probably flat, as typically expected in low-hard state microquasars (e.g. Fender \cite{fender01}) ), but the optically thin radiation from secondary pairs would overcome this component, at least up to few AU."
 This would explain why the radio spectrum of LS 5039 does not look flat (M98)., This would explain why the radio spectrum of LS 5039 does not look flat (M98).
" As observed, this compact radio emission would not present significant changes along the orbit, provided first that the system is not particularly eccentric and, second, that secondary pairs emit radio emission in a region big enough (~ few AU) to smooth out any possible modulation of the primary gamma- injection."," As observed, this compact radio emission would not present significant changes along the orbit, provided first that the system is not particularly eccentric and, second, that secondary pairs emit radio emission in a region big enough $\sim$ few AU) to smooth out any possible modulation of the primary gamma-ray injection."
" On the other hand, the VLBA, MERLIN and EVN extended radio emission would come from collimated"," On the other hand, the VLBA, MERLIN and EVN extended radio emission would come from collimated"
the galactic emission.,the galactic emission.
 In the case of KK 246. its HI spectrum was near the edge of the GMRT observing band. hence the flux could not be reliably estimated.," In the case of KK 246, its HI spectrum was near the edge of the GMRT observing band, hence the flux could not be reliably estimated."
 The GMRT integrated HI emission of the sample galaxies. obtained from the eoarsest resolution data cubes (see Table 23). overlayed on the optical Digitized Sky Survey (DSS) images are shown in Figure 6..," The GMRT integrated HI emission of the sample galaxies, obtained from the coarsest resolution data cubes (see Table \ref{tab:obs}) ), overlayed on the optical Digitized Sky Survey (DSS) images are shown in Figure \ref{fig:mom0}."
 The HI morphological inclinations (igi). for our sample galaxies were estimated from the integrated HI maps by fitting elliptical annuli to the HI images at various resolutions., The HI morphological inclinations ${\rm{_{HI}}}$ ) for our sample galaxies were estimated from the integrated HI maps by fitting elliptical annuli to the HI images at various resolutions.
 For sample galaxies which have HI disks less extended than 2 synthesised beams (across the diameter of the galaxy) at the lowest HI resolution. could in principle be derived from the higher resolution HI maps.," For sample galaxies which have HI disks less extended than 2 synthesised beams (across the diameter of the galaxy) at the lowest HI resolution, could in principle be derived from the higher resolution HI maps."
 However. for most sample galaxies ellipse fitting to the high resolution HI maps is not reliable because of clumpiness in the central high column density regions.," However, for most sample galaxies ellipse fitting to the high resolution HI maps is not reliable because of clumpiness in the central high column density regions."
 The derived inclination (without applying any correction for the intrinsic thickness of the HI disk) is given in ColumntlO) in Table 3.., The derived inclination (without applying any correction for the intrinsic thickness of the HI disk) is given in Column(10) in Table \ref{tab:hiresult}.
 Figure 7. shows a comparison between the morphological inclination derived from the optical and HI isophotes of the galaxy., Figure \ref{fig:incl_figgs} shows a comparison between the morphological inclination derived from the optical and HI isophotes of the galaxy.
 No correction has been applied for the intrinsic thickness of the disk in both cases., No correction has been applied for the intrinsic thickness of the disk in both cases.
 The solid line shows a ease when both inclinations are the same., The solid line shows a case when both inclinations are the same.
 We, We
"removal results in a slelit loss of seusitivity }) im tle central 16"".",removal results in a slight loss of sensitivity ) in the central $16\arcsec$.
 Pixels with ceuters less than 17.25 from the readout trail of the uucleus were flagged and excluded from analysis as were pixels containing less thaw sky exposure near the edge of the CCD., Pixels with centers less than $1\arcsec.25$ from the readout trail of the nucleus were flagged and excluded from analysis as were pixels containing less than sky exposure near the edge of the CCD.
 Sources were located using a nünnmuun variance estimator that compares the data to the telescope Point Spreac Fuuction (PSF)., Sources were located using a minimum variance estimator that compares the data to the telescope Point Spread Function (PSF).
 This method is uost suited. to finding point sources that roughlv match the telescopes PSF., This method is most suited to finding point sources that roughly match the telescope's PSF.
 The PSF was approximated bv a circular gaussian with the same width as the measured PSF at the ainrpoiut aud iucreasing quadratically with distance frou the amipoiut to approximately match the core aud axis broadening of the true PSF., The PSF was approximated by a circular gaussian with the same width as the measured PSF at the aimpoint and increasing quadratically with distance from the aimpoint to approximately match the core and off-axis broadening of the true PSF.
 We define the signal-to-noise ratio (S/N) to be the estimated source counts divided by the la uncertaimity in this nuuber., We define the signal-to-noise ratio (S/N) to be the estimated source counts divided by the $\sigma$ uncertainity in this number.
 To set a minimi threshold. we considered the ummber of sources found per S/N interval down to à minimum S/N of 1.0.," To set a minimum threshold, we considered the number of sources found per S/N interval down to a minimum S/N of 1.0."
 This distribution shows a clear gaussian noise peak at low S/N that indicates a single false detection (ou $3) is to be expected for a minima S/N of 3.0 aud less than 0.1 false detections at a threshold of 3.5., This distribution shows a clear gaussian noise peak at low S/N that indicates a single false detection (on S3) is to be expected for a minimum S/N of 3.0 and less than 0.1 false detections at a threshold of 3.5.
 We include all sources with S/NO3U iu our analysis with the exception of the spectral analysis 1)) aud asnalvsis 5)) for which only 8$/N23.5 sources are considered., We include all sources with $>$ 3.0 in our analysis with the exception of the spectral analysis \ref{s:spectral}) ) and analysis \ref{s:photometry}) ) for which only $>$ 3.5 sources are considered.
 A second pass. designed to mitigate source coufusion. was made through the data.," A second pass, designed to mitigate source confusion, was made through the data."
 For this purpose. all pixels within 20 of cach source position were flageed aud the source-finding process repeated.," For this purpose, all pixels within $\sigma$ of each source position were flagged and the source-finding process repeated."
 Iu this way. the variance in the local backeround due to nearby sources is ercatly reduced aud otherwise confused sources isolated.," In this way, the variance in the local background due to nearby sources is greatly reduced and otherwise confused sources isolated."
 Au additional 6 sources. [with 23.5. were detected.," An additional 6 sources, 4 with $>$ 3.5, were detected."
 The final source list contained 81 sources with S/N3.5 with an additional 16 with 3.0€ «3.5., The final source list contained 81 sources with $\ge$ 3.5 with an additional 16 with $\le$ $<$ 3.5.
 These source positions are identified iu the ACTS data show in figure La and superimposed οi the second generation digital skv SHPVOYA- (DSS) mace in figure Ib., These source positions are identified in the ACIS data shown in figure 1a and superimposed on the second generation digital sky survery (DSS) image in figure 1b.
 Sua] spatial regiois centered on each source were theu selected. for further analvsis., Small spatial regions centered on each source were then selected for further analysis.
 The size of these spatial regions also varies quadratically with off-axis position and euconasses ronehly the eucircled energy radii of the PSF at 1.5 keV. No correction for the low backerounud. 0.04 lods made.," The size of these spatial regions also varies quadratically with off-axis position and encompasses roughly the encircled energy radii of the PSF at 1.5 keV. No correction for the low background, $\sim$ 0.04 $^{-1}$, is made."
 For spectral analysis L)). the total counts within each source region were binned into soft (012 2.0 keV) aud hard (2.0 8.0 keV) bands.," For spectral analysis \ref{s:spectral}) ), the total counts within each source region were binned into soft (0.2 – 2.0 keV) and hard (2.0 – 8.0 keV) bands."
 Ilieher resolution spectra. at 16 PI chaunel binning out to —10 keV. were examined for spurious effects such as anomolously high backgrouud.," Higher resolution spectra, at 16 PI channel binning out to $\sim$ 10 keV, were examined for spurious effects such as anomolously high background."
 No anomolies were eucountered., No anomolies were encountered.
" Finally, we distinguish between the ""bulee aud disk? defined here as those regions internal and external. respectively, to a. L/70«2/35 ellipse ceutered on. the uucleus with major axis position angle 119° (cf."," Finally, we distinguish between the `bulge' and `disk' defined here as those regions internal and external, respectively, to a $4.\arcmin70 \times 2.\arcmin35$ ellipse centered on the nucleus with major axis position angle $^{\circ}$ (cf."
 figure 1)., figure 1).
 This cllipse corresponds to a circle in the plane of the ealaxy of radius 2.35 (~2.5 kpc), This ellipse corresponds to a circle in the plane of the galaxy of radius $2.\arcmin35$ $\sim$ 2.5 kpc).
 There is exteuded. resolved x-ray flux in the bulge that is above backerouud and is not due to the PSF wines of the nucleus nor to other point sources.," There is extended, unresolved x-ray flux in the bulge that is above background and is not due to the PSF wings of the nucleus nor to other point sources."
 This commponcut will be referred to as excess bulee emission., This component will be referred to as `excess bulge emission.'
 The bulge coutaius [1 sources at a surface density of 5.1 ~ aud the clisk contains 56 sources (on $3) at a density. of. —1.0 4 source-arcininP7., The bulge contains 41 sources at a surface density of 5.4 $^{-2}$ and the disk contains 56 sources (on S3) at a density of $\sim$ 1.0 $^{-2}$.
 Among the latter. 2] are located within + LOO oc of spiral arius. (as defined by Matouick Fesen (1997) by tracing the Πα emissiou peak. their figure 31b).," Among the latter, 21 are located within $\pm$ 400 pc of spiral arms (as defined by Matonick Fesen (1997) by tracing the $\alpha$ emission peak, their figure 31b)."
 The surface deusitv of spiral avin sources mi our saniple is slightIv less than. but statistically cousisteut with. the surface density of disk sources not associated with spiral arius.," The surface density of spiral arm sources in our sample is slightly less than, but statistically consistent with, the surface density of disk sources not associated with spiral arms."
" The surface density distribution of bulee sources. inclinatiou-corrected and centered on the uucleus; can be deseribed by au exponeutial radial distribution. o(r)«Xέrita with s,~0.92£0.26 kpc."," The surface density distribution of bulge sources, inclination-corrected and centered on the nucleus, can be described by an exponential radial distribution, $\sigma(r) \propto e^{-r/r_o}$ , with $r_o \sim 0.92\pm0.26$ kpc."
" After subtracting the backerouud component aud the coutzibution from the wines of the PSF of the nucleus. the surface brightuess profile of the excess bulge emission can also be described by an exponential with scale leneth kc,7O.8340.20 kpe for οxrp5/5,"," After subtracting the background component and the contribution from the wings of the PSF of the nucleus, the surface brightness profile of the excess bulge emission can also be described by an exponential with scale length $r_o \sim 0.83\pm0.20$ kpc for $0.\arcmin8\le r\le5.\arcmin5$."
 The excess bulee euission is nach steeper at naller radii though: a residual contribution from the wines of the unclear PSF cannot be definitively excluded., The excess bulge emission is much steeper at smaller radii though a residual contribution from the wings of the nuclear PSF cannot be definitively excluded.
 Bevoud r—1l kpc. these profiles are roughly consisteut with the V. (Ceorgiey Cotov 1991). g (Frei 11996). and B aud 7 (Ehuegreen Ehuceereeu 1981) profiles sugecsting the bulge x-ray cinission follows the distribution of starlight.," Beyond $r \sim 1$ kpc, these profiles are roughly consistent with the $V$ (Georgiev Getov 1991), $g$ (Frei 1996), and $B$ and $I$ (Elmegreen Elmegreen 1984) profiles suggesting the bulge x-ray emission follows the distribution of starlight."
 A search of the literature for known objects spatially coincident with ssources found no matches within the bulee. with the exception of the nucleus.," A search of the literature for known objects spatially coincident with sources found no matches within the bulge, with the exception of the nucleus."
 Among the disk sources. 5 are located within 3” of supernova reimnuauts (SNR: Matonick Feseu 1997) iucliding SN 19933 and the bright ssource X-6 (Fas).," Among the disk sources, 5 are located within $3\arcsec$ of supernova remnants (SNR; Matonick Fesen 1997) including SN 1993J and the bright source X-6 (F88)."
 All 5 SNR candidates lie ou spiral arms as expected for vouns. massive star. Type ID supernova progenitors.," All 5 SNR candidates lie on spiral arms as expected for young, massive star, Type II supernova progenitors."
 Source N-G is a point source located within a OU0-pc-dianeter SNR. (Matouick Feseu 1997)., Source X-6 is a point source located within a 90-pc-diameter SNR (Matonick Fesen 1997).
 The lieh luminosity and hard spectrum 10) of N-6 is uncharacteristic of SNRs., The high luminosity and hard spectrum \ref{s:spectral}) ) of X-6 is uncharacteristic of SNRs.
 Six x-ray sources. all located within spiral arius. are within LO” of known TT reeions. half of which are eiaut radio IIIT regions (IXaufiuau 11987).," Six x-ray sources, all located within spiral arms, are within $10\arcsec$ of known HII regions, half of which are giant radio HII regions (Kaufman 1987)."
 Oue of these x-ray sources. though weak. has the hardest spectimm of the entire source sample £)).," One of these x-ray sources, though weak, has the hardest spectrum of the entire source sample \ref{s:spectral}) )."
 No x-ray sources were found to coincide with kuowu elobular clusters. novae. or stars.," No x-ray sources were found to coincide with known globular clusters, novae, or stars."
 Roughly of globular clusters are expected to contain accreting neutron stars detectable above our threshold aud salaxies like AISI typically contain of order 200 elobular clusters (Pereliuuter Racine 1995)., Roughly of globular clusters are expected to contain accreting neutron stars detectable above our threshold and galaxies like M81 typically contain of order 200 globular clusters (Perelmuter Racine 1995).
 However. there are only 25. confirmed elobular clusters in M81 (Perehuuter. Brodie IIuchlira. 1995) aud only Lowithin the $3 field.," However, there are only 25 confirmed globular clusters in M81 (Perelmuter, Brodie Huchra, 1995) and only 4 within the S3 field."
 The nova vate in MSL is probably of order 20 1 (A. W. Shatter. private communication).," The nova rate in M81 is probably of order 20 $^{-1}$ (A. W. Shafter, private communication)."
 The x-ray signatures and light curves of novae are only sparsely docuneuted (Osecehuan.ee Krautte Deuenuauu 1987) making novae difficult to ideutifv iu our source sample.," The x-ray signatures and light curves of novae are only sparsely documented (Öggelman, Krautte Beuermann 1987) making novae difficult to identify in our source sample."
 There are two sources. ssource N-2 anda ROSAT--detected object. which ive near stars or star clusters;," There are two sources, source X-2 and a -detected object, which are near stars or star clusters."
 N-2 is 1.8 distaut froma star cluster with a spectimm cousisteut with C-type stars (Sholukliova et al., X-2 is $1.\arcsec8$ distant from a star cluster with a spectrum consistent with G-type stars (Sholukhova et al.
 1998)., 1998).
 Zickeraf Thuuphlrevs (1991) list the optical object as a probable foreerotud F-tvpe dwarf., Zickgraf Humphreys (1991) list the optical object as a probable foreground F-type dwarf.
 The ssource is coluckdlent with au uudocuueuted ~Ls- object visible im he Digital Sky Survey image., The source is coincident with an undocumented $\sim$ 18-magnitude object visible in the Digital Sky Survey image.
"mass is finally assembled, AGN feedback prevent gas from overcooling and from accumulating in the core.","mass is finally assembled, AGN feedback prevent gas from overcooling and from accumulating in the core."
" This is well illustrated by the sequence of temperature maps shown in Figure 2,, just after the strong AGN outburst occurring at a~0.62 (see Fig. 1))."," This is well illustrated by the sequence of temperature maps shown in Figure \ref{fig:TempMaps}, just after the strong AGN outburst occurring at $a \simeq 0.62$ (see Fig. \ref{fig:AgnGrowth}) )."
" The first image show the mass-weighted projected temperature within the whole cluster, just at the time of the outburst."," The first image show the mass-weighted projected temperature within the whole cluster, just at the time of the outburst."
" Slightly after (in the second frame), the whole cluster has been significantly heated, with buoyantly driven plumes of hot gas escaping the cluster core (???7).."," Slightly after (in the second frame), the whole cluster has been significantly heated, with buoyantly driven plumes of hot gas escaping the cluster core \citep{Chandran:2007p2323, Cattaneo:2007p420, Rasera:2008p2324, Sharma:2009p2272}."
" In the next frame, a strong shock, visible as a sharp temperature discontinuity, develops close to and beyond the virial radius."," In the next frame, a strong shock, visible as a sharp temperature discontinuity, develops close to and beyond the virial radius."
" The last frame shows the cluster back to hydrostatic equilibrium, waiting for the next AGN outburst."," The last frame shows the cluster back to hydrostatic equilibrium, waiting for the next AGN outburst."
" Note that in the core region, where the density is high enough for X-ray detections, only buoyantly rising bubbles are seen."," Note that in the core region, where the density is high enough for X-ray detections, only buoyantly rising bubbles are seen."
" In our simulation, shock waves form only in the outer part of the cluster, with a Mach number of a few."," In our simulation, shock waves form only in the outer part of the cluster, with a Mach number of a few."
" Because of the rather low gas density, we believe that their detection is quite challenging, explaining why there is so little evidence of their presence (?).."," Because of the rather low gas density, we believe that their detection is quite challenging, explaining why there is so little evidence of their presence \citep{Bourdin:2004p444}."
" Although the feedback recipe doesn't explicitly account for it, the AGN energy deposition typically proceeds in two modes: a strong, energetic mode or ""quasar mode"", which in our case corresponds to the Eddington luminosity and reaches 5x1076 erg s!, and a quiescent mode or ""radio mode"", with a Bondi-Hoyle-limited luminosity of 5x1031 erg s ο."," Although the feedback recipe doesn't explicitly account for it, the AGN energy deposition typically proceeds in two modes: a strong, energetic mode or ""quasar mode"", which in our case corresponds to the Eddington luminosity and reaches $5 \times 10^{46}$ erg $^{-1}$, and a quiescent mode or ""radio mode"", with a Bondi-Hoyle-limited luminosity of $5 \times 10^{41}$ erg $^{-1}$ ."
" If we now take into account the coupling efficiency parametere.=0.15 in our energy estimate, using,"," If we now take into account the coupling efficiency parameter$\epsilon_c=0.15$ in our energy estimate, using,"
implies /Q4i?<0.004(c/df)? confidence). where Ojd is the density relative to the critical densitv in the frequency range (f.f+df).,"implies $f\Omega_f h^2 \la 0.004\, (c/d\,f)^2$ confidence), where $\Omega_f {\rm d} f$ is the density relative to the critical density in the frequency range $(f,f+{\rm d}f)$."
 In the presence of an undiscovered massive companion. the known solar svstem barvcenter would experience an acceleration which could be detected in observations of stable astronomical clocks.," In the presence of an undiscovered massive companion, the known solar system barycenter would experience an acceleration which could be detected in observations of stable astronomical clocks."
 No such acceleration has been detected so far., No such acceleration has been detected so far.
 The best limits on the acceleration are provided by the sources lor which the intrinsic period change rate can be somehow estimated and then compared will the observed period change., The best limits on the acceleration are provided by the sources for which the intrinsic period change rate can be somehow estimated and then compared with the observed period change.
 Examples of such svstems are pulsars in binary svstems in which the orbital period decavs due (to emission of gravitational waves. and white dwarls in which the pulsation periods change as the star cools.," Examples of such systems are pulsars in binary systems in which the orbital period decays due to emission of gravitational waves, and white dwarfs in which the pulsation periods change as the star cools."
 Alternatively. statistical limits on (he acceleration can be provided by a group of objects of (hie sime type distributed across the sky. even if (he intrinsic period change rate of these objects is unknown.," Alternatively, statistical limits on the acceleration can be provided by a group of objects of the same type distributed across the sky, even if the intrinsic period change rate of these objects is unknown."
 An example is milliseconcl pulsars. which provide sensitivitv to the acceleration similar to that. obtained using pulsars in binary svstems and white dwarls.," An example is millisecond pulsars, which provide sensitivity to the acceleration similar to that obtained using pulsars in binary systems and white dwarfs."
 The limits on the acceleration of the solar svstem barycenter that we obtained by different methods are summarized in Table 2., The limits on the acceleration of the solar system barycenter that we obtained by different methods are summarized in Table 2.
" The acceleration exerted bv the companion on the solar svstem is @./¢=CanyfCed""). where d, is the current distance to the companion and my is its mass."," The acceleration exerted by the companion on the solar system is $a_{\odot}/c=Gm_p/(c d_p^2)$, where $d_p$ is the current distance to the companion and $m_p$ is its mass."
 Therefore. the upper limit on the mass of the companion as a Iunction of the upper limit on the acceleration is ≼↕⊲∪∐↓↕↽≻≀↧↴∐↕∪∐⋝∖⊽∖∖⊽↕⊔⊔⊔≀↧⊔∖⇁⊳∖⊽≼↲⊳∖⇁∐↕↖≺↽↔↴∐≼↲↕⋅⊔⋯↴∐⊔∏⊳∖⇁≺∢≀↧↴∐∣↽≻≼↲↕⋅∏↥≼↲≺⇂∪∏↥≀↧↴↥≀↧↴⊋⊔≺∢∪∐∐≼⇂≼↲∐≺∢≼↲∪∖↽≼↲↕⋅≀↧," Therefore, the upper limit on the mass of the companion as a function of the upper limit on the acceleration is Companions with masses higher than this can be ruled out at a $\sigma$ confidence over about of the sky (Table 2)."
↴∣↽≻∪∏↥ ⊤≺≨↖∕⋡∣↙∪↓⋟⊔∐↲⋟∖⊽↳↽∡∖↽≼∫⊺≀↧↴∣↽≻↥≼↲⊋↕⋝⋅↻∏↕⋅∐∐↕∐∪∐⊔∐↲≀↧↴≺∢≺∢≼↲↥≼↲↕⋅≀↧↴∐∪∐↕⋟∖⊽≀↧↴∣↽≻∪⋯≀↧↴∐∪↕⋅≼⇂≼↲↕⋅∪↓⋟∐⋯↴↖≺↽↔↴∐∐⋯⇂↩∣↽≻≼↲⊔≼↲↕⋅ ⊔⋯∐⊔∐↲∪∐≼↲≀↧↴∖⇁≀↧↴∐≀↧↴∣↽≻↥≼↲⊥↱≻⋡∖↽≼↲≀↧↴↕⋅⊳∖⊽≀↧↴≸≟∪⋜∫⊺↕⋅≼↲↕∐≀↧↴∐∐↲⊥≤∍≤⋯⇄⋝∣↽≻≼↲≺∢≀↧↴∏⊳∖⇁≼↲∪↓≯⊔∐↲∐∐↕↽≻↕⋅∪∖↽≼↲≼⇂≀↧↴≺∢≺∢⋯⋅≀↧, Our limit on the acceleration is about an order of magnitude better than the one available 15 years ago \citep{trem90} because of the improved accuracy of the timing observations of PSR B1913+16 and because several other objects became available.
↴≺↕∖↽∪↓≯⊔∐↲ ∐∐↓↕↥≸≟∪∣↽≻⋝∖⊽≼↲↕⋅∖⇁≀↧↴∐∪∐⊳∖⇁∪↓≯↥↴⊳↔⊲∐∐⊥≤∐⊑↽⊰⊹∐≨≀↧↴∐≼⇂∣↽≻≼↲≺∢≀↧↴∏⋟∖⇁≼↲⋟∖⇁≼↲∖⇁≼↲↕⋅≀↧↴↥∪⊔∐↲↕⋅∪∣↽≻∙↿≼↲≺∢↥⊳∖⇁∣↽≻≼↲≺∢≀↧↴∐∐↲≀↧↴∖↽≀↧↴∐≀↧↴∣↽≻↥≼↲⋅ e ↴⊺↕∐∐∐∩∪∫↽⋟∏↕⋟∖⊽≀↧↴↕∎⋟∖⊽∐≀↧↪∖⊽∣↽≻≼↲≼↲∐∖∖⊽↕≼∣≼↲↕∖⊽∏⋟∖⊽≼↲≼∣∩≽↽⋟↕≀↧↴∏↲≺⇃∪∐⋟∖⊽∐∎≀↧↴↕∐∟∖⊽∪∐∐∐↲≼↲∐≼↲↕∎∩∖⊽≼∣≼↲∐⋟∖⊽∐∖⊽∪∫ e low-Irequenev Sogravitational waves., Timing of pulsars has been widely used to place constraints on the energy density of low-frequency gravitational waves.
 In Section 4 we summarized these methods ancl outlined the dilferences between (he effects of a companion aud of the gravitational waves on pulsar (imine., In Section \ref{sec_waves} we summarized these methods and outlined the differences between the effects of a companion and of the gravitational waves on pulsar timing.
 The observational accuracy will improve with time for many objects. orbital period decay will likely be measured [or many pulsars in binary svstems. ancl more millisecond pulsars will be discovered.," The observational accuracy will improve with time for many objects, orbital period decay will likely be measured for many pulsars in binary systems, and more millisecond pulsars will be discovered."
 For some objects used in our analysis further observations are, For some objects used in our analysis further observations are
or negligible (Shieldsetal. evolution in the normalisation of the mpu—ox relation.,"or negligible \citep{shie03,gask09}
 evolution in the normalisation of the $\mbh-\sigmas$ relation."
" The predicted evolution does, however, agree with other simulation studies (Robertsonetal.2006) and semi-analytic models (Malbonetal.2007)."," The predicted evolution does, however, agree with other simulation studies \citep{robe06} and semi-analytic models \citep{malb07}."
". Taken together, the simulation predicts that the mpH—m. and ""pH—Ox relations evolve such that the relation between mpg and stellar binding energy (οςMo? ) is independent of redshift (as=0.05+ 0.06)."," Taken together, the simulation predicts that the $\mbh-\ms$ and $\mbh-\sigmas$ relations evolve such that the relation between $\mbh$ and stellar binding energy $\propto\ms\sigmas^2$ ) is independent of redshift $\alphas=\asus$ )."
 It is tempting to conclude from the finding that the ratio mpu/(ms+o2) does not evolve that the BH mass is determined by the binding energy of the galaxy., It is tempting to conclude from the finding that the ratio $\mbh/(\ms\sigmas^2)$ does not evolve that the BH mass is determined by the binding energy of the galaxy.
" However, we demonstrated explicitly in BS10 that the BH mass is instead controlled by the binding energy of the DM halo."," However, we demonstrated explicitly in BS10 that the BH mass is instead controlled by the binding energy of the DM halo."
" This implies that the binding energy of the galaxy tracks the binding energy of the halo, which we will confirm and explain below."," This implies that the binding energy of the galaxy tracks the binding energy of the halo, which we will confirm and explain below."
 We now turn our attention to the relations between the mass of the BH and the DM halo in which it resides., We now turn our attention to the relations between the mass of the BH and the DM halo in which it resides.
 In BS10 we argued that a BH grows until it has injected an amount of energy into its surroundings that scales with the binding energy of its host DM halo., In BS10 we argued that a BH grows until it has injected an amount of energy into its surroundings that scales with the binding energy of its host DM halo.
" We therefore do not expect the mH—Una, relation to evolve.", We therefore do not expect the $\mbh-U_{\rm halo}$ relation to evolve.
" Indeed, in the simulation the amplitude of this relation is independent of redshift, with as=0.03+0.05, as would be expected for a fundamental link between mpu and the binding energy of the host DM halo."," Indeed, in the simulation the amplitude of this relation is independent of redshift, with $\alphas=\asuh$, as would be expected for a fundamental link between $\mbh$ and the binding energy of the host DM halo."
" We do, however, expect the relation between mpg and Mnalo to evolve."," We do, however, expect the relation between $\mbh$ and $\mhalo$ to evolve."
" At higher redshift, haloes of a given mass are more compact than their redshift zero counterparts and are thus more strongly gravitationally bound."," At higher redshift, haloes of a given mass are more compact than their redshift zero counterparts and are thus more strongly gravitationally bound."
" This means that, at a fixed halo mass, more energy is required to eject gas from haloes at high redshift and in order to self-regulate, BHs must grow to be more massive."," This means that, at a fixed halo mass, more energy is required to eject gas from haloes at high redshift and in order to self-regulate, BHs must grow to be more massive."
 Fig., Fig.
" 2 shows the normalisation of the mag—Mhalo relation as a function of redshift and confirms that the simulation predicts the amplitude of this relation to increase with redshift (red, dashed curve), with as=0.65+0.06."," \ref{fig:mh} shows the normalisation of the $\mbh-\mhalo$ relation as a function of redshift and confirms that the simulation predicts the amplitude of this relation to increase with redshift (red, dashed curve), with $\alphas = \asmh$."
" As we already noted, the idea that the binding energy of the dark halo controls the mass of the BH explains our finding that the ΠΙΒΗ—Una relation is independent of redshift."," As we already noted, the idea that the binding energy of the dark halo controls the mass of the BH explains our finding that the $\mbh-U_{\rm halo}$ relation is independent of redshift."
" We will now show that the analytic model of BS10 also reproduces the evolution of the mpm—mnaio relation and that it can explain the observed evolution in the scaling relations with the stellar properties if the observed galaxies evolve predominantly through dry mergers, as predicted by the simulation."," We will now show that the analytic model of BS10 also reproduces the evolution of the $\mbh-\mhalo$ relation and that it can explain the observed evolution in the scaling relations with the stellar properties if the observed galaxies evolve predominantly through dry mergers, as predicted by the simulation."
" If the energy injected by a BH is proportional to the halo gravitational binding energy, Unaio, then, for a DM halo with an NFW (Navarroetal.1997) density profile (BS10) where maj is the halo mass, c(mhalo,Z) is the halo concentration, x is defined to be x=re/rnai, rej is the physical scale on which BH self-regulation takes place, and f(c,x) is the function Simulations have shown that c is a function of both redshift and halo mass, and scales approximately Combining Eqs. 3"," If the energy injected by a BH is proportional to the halo gravitational binding energy, $U_{\rm halo}$, then, for a DM halo with an NFW \citep{nava97} density profile (BS10) where $\mhalo$ is the halo mass, $c(\mhalo,z)$ is the halo concentration, $x$ is defined to be $x\equiv\reject/\rhalo$, $\reject$ is the physical scale on which BH self-regulation takes place, and $f(c,x)$ is the function Simulations have shown that $c$ is a function of both redshift and halo mass, and scales approximately \citep{duff08}
 Combining Eqs. \ref{eq:be}"
euuission being delayed with respect to the peak in IR emission.,emission being delayed with respect to the peak in IR emission.
 Lu this model. flaring at a giveu," In this model, flaring at a given"
four independent fields.,four independent fields.
 Then we compare the inferred galaxy merger rate with the observed number density evolution of massive (M.>10!!M ..). red galaxies from z—1 to the present day. which we obtain by converting Brownetal.(2007) LFs to stellar-mass functions.," Then we compare the inferred galaxy merger rate with the observed number density evolution of massive $M_*>10^{11}M_{\sun}$ ), red galaxies from $z\sim 1$ to the present day, which we obtain by converting \citet{brown} LFs to stellar–mass functions."
" We assume Hy=70 knys. Q,,=0.3 and Q420.7."," We assume $H_0=70$ km/s, $\Omega_{m}=0.3$ and $\Omega_{\Lambda}=0.7$."
 The bulk of our sample is drawn from the —2ssq. ddegree COSMOS survey (Scovilleetal.2007).., The bulk of our sample is drawn from the $\sim$ degree COSMOS survey \citep{scoville}.
 We use photometric redshifts calculated from 30-band photometry by Ilbertetal. (2009):: comparison with spectroscopic redshifts shows excellent accuracy., We use photometric redshifts calculated from 30-band photometry by \citet{ilbert}; comparison with spectroscopic redshifts shows excellent accuracy.
 We use those redshifts to derive rest-frame quantities and stellar masses by using the observed broad-band photometry in conjunction with a non-evolving template library derived using Péggase stellar population model (seeΡίου&Rocca- and à Chabrier(2003). stellar initial mass function (IMF)., We use those redshifts to derive rest-frame quantities and stellar masses by using the observed broad-band photometry in conjunction with a non-evolving template library derived using Péggase stellar population model \citep[see][]{fioc99} and a \cite{chabrier03} stellar initial mass function (IMF).
 The use of a Kroupaetal.(1993) or a Kroupa(2001)4 IMF would. yieldn similar↼↼⋅⊲ stellar NENmasses to withinHP ~10%.," The use of a \cite{kroupa93}
 or a \cite{kroupa01} IMF would yield similar stellar masses to within $\sim$."
. The reddest templates are produced through single exponentially—formationqns episodes. ⊲↴⊀↼intermediate templates↼⊀ also⋅≏⋯∪↾∣↴⊜∣⋪≝∐∣∐⇀↸∖⇁∣⋋∪∣∶∣≀∣∣↲↱≦⋖∣⋅⋟∣∡∏∕⋅∖∖⇁∣↴⊜∣⋪⊜⊔∣⋋⋔⊖∏∐⋯⇂↴⊜∣⋪declining = : have a low-level constant star formation rate (SFR) and the bluer templates have superimposed a recent burst of star formation.," The reddest templates are produced through single exponentially--declining star formation episodes, intermediate templates also have a low-level constant star formation rate (SFR) and the bluer templates have superimposed a recent burst of star formation."
 A full description of the stellar masses will be provided in a future paper: comparison with Pannellaetal.(2009) masses shows agreement at the ddex level., A full description of the stellar masses will be provided in a future paper; comparison with \citet{pan} masses shows agreement at the dex level.
 To combat the sample variance of a single 2 ddegree field. we augment the COSMOS sample with a sample drawn from three widely-separated 0.25 ddegree fields from COMBO-17.," To combat the sample variance of a single 2 degree field, we augment the COSMOS sample with a sample drawn from three widely-separated 0.25 degree fields from COMBO-17."
 COMBO-17 photo-z's. colors and stellar masses have been extensively described in Wolfetal.(2003).. Wolfetal.(2004). Borchetal.(2006) and Grayetal.(2009).," COMBO–17 photo-z's, colors and stellar masses have been extensively described in \citet{wolf03}, \citet{wolf04}, \citet{borch} and \citet{gray}."
. Given the different depths of the two surveys we only include galaxies from the COMBO-17 catalog at z<0.8. where it is complete for ov mass limit.," Given the different depths of the two surveys we only include galaxies from the COMBO-17 catalog at $z<0.8$, where it is complete for our mass limit."
 The final sample comprises ~18000 galaxies with M.5«10M. over an area of ~2.75 sq.," The final sample comprises $\sim 18000$ galaxies with $M_*\ge 5\times
10^{10}M_\sun$ over an area of $\sim 2.75$ sq."
 deg., deg.
 in the redshift range 0.2«z1.2., in the redshift range $0.2<z<1.2$.
 We use the fraction of galaxies with a companion closer than 30 kpe (close pairs) as a proxy for the merger fraction. as those systems are very likely to merge in a few hundredMyr!.," We use the fraction of galaxies with a companion closer than 30 kpc (close pairs) as a proxy for the merger fraction, as those systems are very likely to merge in a few hundred."
. As redshift errors translate into line-of-sight (Lo.s.), As redshift errors translate into line-of-sight (l.o.s.)
 distance uncertainties of the order of ~50—100 Mpe we use projected 2-point correlation functions (2pcf) to find the number of projected close pairs and then deproject into the 3D space., distance uncertainties of the order of $\sim 50-100$ Mpc we use projected 2-point correlation functions (2pcf) to find the number of projected close pairs and then deproject into the 3D space.
" The projected correlation function w(rp) is the integral along the line of sight of the real-space correlation function: where is the distance between the two galaxies projected on the plane7, of sky and π the line-of-sight separation.", The projected correlation function $w(r_P)$ is the integral along the line of sight of the real-space correlation function: where $r_{p}$ is the distance between the two galaxies projected on the plane of sky and $\pi$ the line-of-sight separation.
 A convenient and simple estimator of the 2pef at small scales Is ο)=A(DD/RR—1) (e.g.Belletal.2006a;Robaina2009).. where A is the path length being integrated over. DD(rp) is the histogram of separations between real galaxies and RRG'p) ts the histogram of separations between galaxies in a distributed catalog.," A convenient and simple estimator of the 2pcf at small scales is $w(r_{p})=\Delta(DD/RR-1)$ \citep[e.g.][]{bell06a,robaina09}, where $\Delta$ is the path length being integrated over, $DD(r_P)$ is the histogram of separations between real galaxies and $RR(r_P)$ is the histogram of separations between galaxies in a randomly-distributed catalog."
 As shown in the literature (e.g.Davis&Peebles1983;Lietal. 2006).. the real-space 2pef can be reasonably well fit by a power-law.," As shown in the literature \citep[e.g.][]{davis, li}, the real-space 2pcf can be reasonably well fit by a power-law."
" Assuming £67)2Gra). then wry,2Crn,. with CJ/z(T[G—D/2]/T(G/2))."," Assuming $\xi(r)=(r/r_0)^{- \gamma}$, then $w(r_p)=Cr_{0}^{\gamma}r_{0}^{1-\gamma}$, with $C\sqrt{\pi}\{\Gamma[(\gamma-1)/2]/\Gamma(\gamma/2)\}$."
 We fit the latter expression to our data to find the parameters 5 and ry and use them in the real-space 2pef to find the fraction of galaxies in close pairs., We fit the latter expression to our data to find the parameters $\gamma$ and $r_0$ and use them in the real-space 2pcf to find the fraction of galaxies in close pairs.
 As we wish to preserve S/N we did not integrate. along the entire Lo.s., As we wish to preserve S/N we did not integrate along the entire l.o.s.
 when caleulating w(rp)., when calculating $w(r_P)$ .
" Instead we allowed galaxies to form a pair only if the redshift difference was smaller than 2V2«o-. with σ. being the redshift error of the primary2,5; galaxy."," Instead we allowed galaxies to form a pair only if the redshift difference was smaller than $\sigma_{pair}=\sqrt{2}\times \sigma_z$, with $\sigma_z$ being the redshift error of the primary galaxy."
 As the photo-z errors follow a Gaussian distribution with width o- (Wolfetal.2003.2004) a fraction of the pairs would be missing by simply imposing the Lo.s.," As the photo-z errors follow a Gaussian distribution with width $\sigma_z$ \citep{wolf03,wolf04} a fraction of the pairs would be missing by simply imposing the l.o.s."
 criteria., criteria.
 Thus. following Belletal.(2006a) the fraction of galaxy pairs is Then. w(rp) is multiplied by 1/f in order to account for missing pairs: in our case a correction factor of 1.19.," Thus, following \citet{bell06a} the fraction of galaxy pairs is Then, $w(r_P)$ is multiplied by $1/f$ in order to account for missing pairs; in our case a correction factor of 1.19."
 GivenI. the 3D4 correlationκι function £(7). the differential.a: probability ∪⊓↴⋯↿⋯∶↔⊺∐∶↔⊺∐∣∐↖⋝⇁⋂∁∁∐∣⊃⋝⇁⋯∶↔⊺⋅≏↧∖⇁∪∣∐⋯⊜↙∏∕⋅≏∐⋅≏↧⋔⋋↾∐∏∁⊖∣⋅⋂↑↴↘↼ . - . ↴↙ star lp - . density of secondary galaxies.," Given the 3D correlation function $\xi(r)$, the differential probability of finding a galaxy occupying a volume $\delta V$ at a distance $r$ of another galaxy is $\delta P=n[1+\xi(r)]\delta V$, where $n$ is the number density of secondary galaxies."
" Then. by a simple integration of this expression. we obtain the probability of à galaxy being within a distance +, of any other galaxy (Pattonetal.2000:Belletal.2006a:Masjedi 2006):: Because £67)2(r/rg9) and £(7)2| at the small scales. we obtain: where is the fraction ofgalaxies in close pairs."," Then, by a simple integration of this expression, we obtain the probability of a galaxy being within a distance $r_f$ of any other galaxy \citep{patton00, bell06a, masjedi06}: Because $\xi(r)=(r/r_0)^{-\gamma}$ and $\xi(r)\gg 1$ at the small scales, we obtain: where $f_{pair}$ is the fraction of in close pairs."
 As galaxy interactionsfj are completely decoupled from the Hubble flow. we calculate probabilities as a function of the proper (physical) separation between the two galaxies.," As galaxy interactions are completely decoupled from the Hubble flow, we calculate probabilities as a function of the proper (physical) separation between the two galaxies."
 Errors in the correlation function are calculated by means of bootstrap resampling., Errors in the correlation function are calculated by means of bootstrap resampling.
 This work aims to an understanding of the impact of galaxy mergers on the evolution of massive galaxies and the evolution of the red sequence at its massive end., This work aims to an understanding of the impact of galaxy mergers on the evolution of massive galaxies and the evolution of the red sequence at its massive end.
 To that end we apply the above method to our sample of galaxies with ΜΗ.>5&10'M ... regardless of their optical colors. to find the fraction of galaxies in close physical pairs. which we will use as à proxy for the fraction of galaxies undergoing a galaxy interaction.," To that end we apply the above method to our sample of galaxies with $M_*>5\times10^{10}M_\sun$ , regardless of their optical colors, to find the fraction of galaxies in close physical pairs, which we will use as a proxy for the fraction of galaxies undergoing a galaxy interaction."
 Hereafter. all reference to close pairs will refer to galaxy pairs with inferred separations smaller than 30 kpe in 3D space.," Hereafter, all reference to close pairs will refer to galaxy pairs with inferred separations smaller than 30 kpc in 3D space."
 Fig., Fig.
 | and Table 1. show the result of our 2pcf analysis., \ref{fig:cors} and Table \ref{tab:cor} show the result of our 2pcf analysis.
 The projected correlation functions in Fig., The projected correlation functions in Fig.
" | are shown over the range 10—r,/kpe<1000. in physical (not comoving) units. as is appropriate for a process such as galaxy merging that has decoupled from the Hubble flow."," \ref{fig:cors} are shown over the range $10<r_p/kpc<1000$, in physical (not comoving) units, as is appropriate for a process such as galaxy merging that has decoupled from the Hubble flow."
 One can see that the fitted power laws are good descriptions of the run of ο) with radius. down to the resolution limit of ~15 kkpe.," One can see that the fitted power laws are good descriptions of the run of $w(r_p)$ with radius, down to the resolution limit of $\sim 15$ kpc."
 Below this radius. we extrapolate the power law fit.," Below this radius, we extrapolate the power law fit."
 We have checked the accuracy of this extrapolation by comparison with extremelyclose pairs resolved by HST in the study of two of the four fields used in this paper by Robainaetal. (2009).., We have checked the accuracy of this extrapolation by comparison with extremelyclose pairs resolved by HST in the study of two of the four fields used in this paper by \citet{robaina09}. .
 An extrapolation, An extrapolation
velocity (RV) measurements were made between 2009 June 19 and 2009 September 22 (Table 1)).,velocity (RV) measurements were made between 2009 June 19 and 2009 September 22 (Table \ref{rv-data}) ).
 The resulting radial velocity curve Is shown in Fig. 2.., The resulting radial velocity curve is shown in Fig. \ref{RV}.
 The absence of any correlation with RV of the line bisector spans (ων) in Fig., The absence of any correlation with RV of the line bisector spans $V_{span}$ ) in Fig.
 3. indicates that the RV variations are not due to an unresolved eclipsing binary (Quelozetal.2001)., \ref{RV-BS} indicates that the RV variations are not due to an unresolved eclipsing binary \citep{2001A&A...379..279Q}.
 WASP-26 was observed photometrically on 2009 November 18 using the 2.0-m Faulkes Telescope South (FTS. Siding Spring. Australia) and on 2009 December2 using the2.0-m Faulkes Telescope North (FTN. Maut. Hawar1). both telescopes being operated by LCOGT.," WASP-26 was observed photometrically on 2009 November 18 using the 2.0-m Faulkes Telescope South (FTS, Siding Spring, Australia) and on 2009 December 2 using the 2.0-m Faulkes Telescope North (FTN, Maui, Hawai'i), both telescopes being operated by LCOGT."
 In both cases. a new Spectral CCD was used along with a Pan-STARRS-: filter.," In both cases, a new Spectral CCD was used along with a $z$ filter."
 The Spectral instrument contains a Fairchild. CCD486. back-illuminated 4096x pixel CCD which was binned 2x giving pixels and a field of view of 10’x10’., The Spectral instrument contains a Fairchild CCD486 back-illuminated $4096 \times 4096$ pixel CCD which was binned $2 \times 2$ giving pixels and a field of view of $10\arcmin \times 10\arcmin$.
" The telescope was defocussed a small amount during the observations to prevent saturation in the core of the PSF but not by enough to cause blending problems with the close (.15"")) fainter companion.", The telescope was defocussed a small amount during the observations to prevent saturation in the core of the PSF but not by enough to cause blending problems with the close $\sim 15$ ) fainter companion.
 The frames were pre-processed through the WASP Pipeline (Pollaecoetal.2006) to perform overscan correction. bias subtraction and flat-fielding.," The frames were pre-processed through the WASP Pipeline \citep{2006PASP..118.1407P} to perform overscan correction, bias subtraction and flat-fielding."
 The DAOPHOT photometry package within was used to perform object detection and aperture photometry using aperture radi of 9 and 10 pixels for the FTS and FTN data. respectively.," The DAOPHOT photometry package within was used to perform object detection and aperture photometry using aperture radii of 9 and 10 pixels for the FTS and FTN data, respectively."
 Differential photometry was performed relative to several comparison stars within the field of view., Differential photometry was performed relative to several comparison stars within the field of view.
 Figure + shows the transit photometry.," Figure \ref{FT-phot}
 shows the transit photometry."
 The large scatter in the FTN observations suggests that they were somewhat affected by cloud., The large scatter in the FTN observations suggests that they were somewhat affected by cloud.
 The individual CORALIE spectra of WASP-26 were co-added to produce a single spectrum with an average S/N of around 70:1., The individual CORALIE spectra of WASP-26 were co-added to produce a single spectrum with an average S/N of around 70:1.
 The standard pipeline reduction products were used in the analysis., The standard pipeline reduction products were used in the analysis.
 The analysis was performed using the methods given in Gillonetal.(2009). and Smalley(2005).., The analysis was performed using the methods given in \citet{2009A&A...496..259G} and \citet{2005MSAIS...8..130S}.
" The H, line was used to determinethe effective temperature (Tay). while the D and b lines were used as surface gravity (loge) diagnostics."," The $_\alpha$ line was used to determinethe effective temperature $T_{\rm eff}$ ), while the D and b lines were used as surface gravity $\log g$ ) diagnostics."
 The parameters obtained from the analysis are listed in Table2.., The parameters obtained from the analysis are listed in Table \ref{WASP26-params}.
 The elemental abundances were determined from equivalent width measurements of several clean and unblended lines., The elemental abundances were determined from equivalent width measurements of several clean and unblended lines.
 The quoted error estimates include that given by the uncertainties in Tay.logg and &. as well as the scatter due to measurement and atomic data uncertainties.," The quoted error estimates include that given by the uncertainties in $T_{\rm eff}$ ,$\log
g$ and $\xi_{\rm t}$ , as well as the scatter due to measurement and atomic data uncertainties."
"AL,10M... intermediate to that of Pop III and Pop ILI stars. ancl by extremely low metallicities. below the criGeal value that would normally lead to the formation of VMSs.","$M_{\ast}\sim 10\,\msun$, intermediate to that of Pop III and Pop II stars, and by extremely low metallicities, below the critical value that would normally lead to the formation of VMSs."
" This possibility relies on the gas in the vicinity of the SN being compressed in a radiative shock. leading to a significant increase in densitv. and thus possibly (o a reduction of the fragment mass. MpxpL7. to values much lower than the one realized in primordial eas,"," This possibility relies on the gas in the vicinity of the SN being compressed in a radiative shock, leading to a significant increase in density, and thus possibly to a reduction of the fragment mass, $M_{\rm BE}\propto
\rho^{-1/2}$, to values much lower than the one realized in primordial gas."
" During an intermediate epoch. therefore. low- and high-mass stars might form. almost simultaneously,"," During an intermediate epoch, therefore, low- and high-mass stars might form almost simultaneously."
 A similar idea had already been proposed in a prescient paper by (1986).. primarily to explain the so-called G-cwarl problem. i.e.. (he observed lack of metal- stars in our Galaxy (seealsoIkashlinskv&Rees1983:NakamuraUmenura2001).," A similar idea had already been proposed in a prescient paper by \citet{Cay86}, primarily to explain the so-called G-dwarf problem, i.e., the observed lack of metal-poor stars in our Galaxy \citep[see also][]{KasRee83,NakUme01}."
. It is to be expected that the mass range of Pop I5 stars extends down to <1M... and these stars at Che low-mass end of the Pop IL5 IAIF should still be around today.," It is to be expected that the mass range of Pop II.5 stars extends down to $\la 1\;\msun$, and these stars at the low-mass end of the Pop II.5 IMF should still be around today."
 Intriguingly and possibly providing a case in point. the recent discovery of (he extremely metal poor star CS 29498-0423 with |Fe/II]|2—3.7 and a significant overabundance of Mg and Si. has been tentatively interpreted as hinting at a new class of stars (Aokietal.2002).," Intriguingly and possibly providing a case in point, the recent discovery of the extremely metal poor star CS 29498-043 with $=-3.7$ and a significant overabundance of Mg and Si, has been tentatively interpreted as hinting at a new class of stars \citep{Aoketal02}."
. Ongoing surveys of extremely metal-poor halo stars. much improved in size and equality. should soon be able to test our prediction further.," Ongoing surveys of extremely metal-poor halo stars, much improved in size and quality, should soon be able to test our prediction further."
 Many of the uncertainties in our argument have to be addressed with detailed numerical simulations. and we plan to do so in future work.," Many of the uncertainties in our argument have to be addressed with detailed numerical simulations, and we plan to do so in future work."
" The framework presented in this paper does provide a coherent context for these numerical studies. and it highlights the important physical questions that define (he challenge of elucidating the end of the cosmic ""dark ages”."," The framework presented in this paper does provide a coherent context for these numerical studies, and it highlights the important physical questions that define the challenge of elucidating the end of the cosmic “dark ages”."
"al approximately 140.853 CIlz as the 942—845 transition of NaCN and the line at roughly 140.840 GHz as the 2,»—ly transition of formaldehyde.","at approximately $140.853\,$ GHz as the $9_{37}-8_{36}$ transition of NaCN and the line at roughly $140.840\,$ GHz as the $2_{12}-1_{11}\>$ transition of formaldehyde."
" The feature in the redshifted wing of the IbCO 245—ly line is actually (wo lines: it is the 7TTj2J=59/2—57/2 doublet of CH. The entire blend was originally identified by Cernicharoetal.(2000) as the 942—S46 transition of δαν,", The feature in the redshifted wing of the $_2$ CO $2_{12}-1_{11}\>$ line is actually two lines; it is the $^2\Pi_{1/2}\>J=59/2-57/2$ doublet of $_5$ H. The entire blend was originally identified by \citet{CGK00} as the $9_{37}-8_{36}$ transition of NaCN.
 Due to the line blending. it was not possible to obtain a reasonable lit for anv of the four lines if all fit parameters were allowed (to vary.," Due to the line blending, it was not possible to obtain a reasonable fit for any of the four lines if all fit parameters were allowed to vary."
 It was also not possible to fit any single line independently. so for this blend we fit all four lines simultaneously.," It was also not possible to fit any single line independently, so for this blend we fit all four lines simultaneously."
" In the simultaneous fit. we fixed the central frequencies and expansion velocities (6,=14.5kms1) ol the formaldehyde. C's] and δαν lines. while allowing the area and horn/center ratio (the line shape) parameters to vary. (hough we constrained the two CsI lines to have identical strengtlis and shapes."," In the simultaneous fit, we fixed the central frequencies and expansion velocities $v_{exp}=14.5{\rm \,km\,s^{-1}}$ ) of the formaldehyde, $_5$ H and NaCN lines, while allowing the area and horn/center ratio (the line shape) parameters to vary, though we constrained the two $_5$ H lines to have identical strengths and shapes."
 The four line simultaneous fit is displaved in the top panel of Figure L.., The four line simultaneous fit is displayed in the top panel of Figure \ref{cspec}.
 We note that NaCN and Cs1I are among the list of well studied molecules in IRC+10216 and (hat the line [Iuxes measured here are in line wil those measured previously [or transitions of similar excitation (e.g.Cernicharoetal.2000)., We note that NaCN and $_5$ H are among the list of well studied molecules in IRC+10216 and that the line fluxes measured here are in line with those measured previously for transitions of similar excitation \citep[e.g.][]{CGK00}.
.. As a sanitv check. we also examined a subtracted spectrum. removing (he fitted CID and NaCN profiles. presumably leaving only the 24»—ly transition of formaldehyde.," As a sanity check, we also examined a subtracted spectrum, removing the fitted $_5$ H and NaCN profiles, presumably leaving only the $2_{12}-1_{11}\>$ transition of formaldehyde."
 This spectrum is also plotted inFigure 1.. aud the formaldehyde line is readily apparent.," This spectrum is also plotted inFigure \ref{cspec}, and the formaldehyde line is readily apparent."
" The 140.340 GlIlz band also contains (wo lines between 140.9 and 140.95 GIlz: we identifv these lines as the 945;—S45 transition of NaCN and the 655—55, transition of SiCs based on the observations of Cernicharoetal.(2000).. who observed this region of the spectrum around IRC--10216 in a previous. shallower line survev."," The $140.840\,$ GHz band also contains two lines between $140.9$ and $140.95\,$ GHz; we identify these lines as the $9_{36}-8_{35}$ transition of NaCN and the $6_{25}-5_{24}$ transition of $_2$ based on the observations of \citet{CGK00}, who observed this region of the spectrum around IRC+10216 in a previous, shallower line survey."
 These lines are labelled in Figure 1.., These lines are labelled in Figure \ref{cspec}.
 The NaCN line in (his group also required a constrained fit due to the proximity of the strong SiC line. as noted in Table 3..," The NaCN line in this group also required a constrained fit due to the proximity of the strong $_2$ line, as noted in Table \ref{linefit}."
" Table 3 also contains fit parameters for a."" S1Cs line which we only observed at one local oscillator setting (due to band edge placement).", Table \ref{linefit} also contains fit parameters for a $^{30}$ $_2$ line which we only observed at one local oscillator setting (due to band edge placement).
 We do not plot Chis line in Figure 1.. but it would appear just bevond the blue edee of the 140.840 Giz band plot.," We do not plot this line in Figure \ref{cspec}, but it would appear just beyond the blue edge of the $140.840\,$ GHz band plot."
 In addition to observations of the central position. we also took data at 8 separate positions. using ILAMs mode.," In addition to observations of the central position, we also took data at $8$ separate off-center positions, using IRAM's mode."
" All of these observations were racially offset [rom the center position by 17"". which is a full beam offset at 140 GHz (i.e. the center of the central posiGion is separated from the center of the offset position bv one EWIIM οἱ the telescope beam)."," All of these observations were radially offset from the center position by $17\arcsec$, which is a full beam offset at $140\,$ GHz (i.e. the center of the central position is separated from the center of the offset position by one FWHM of the telescope beam)."
 We chose the 8 position angles such that we had (wo mapping grids or crosses. as drawn in Figure 3: one cross was oriented along the line al position angle 120 degrees. which has been suggested as the plane of a disk around IRC+10216 2000).. while the other cross was oriented al a position angle of 45 degrees to the first cross.," We chose the $8$ position angles such that we had two mapping grids or crosses, as drawn in Figure \ref{pattern}; one cross was oriented along the line at position angle $120\,$ degrees, which has been suggested as the plane of a disk around IRC+10216 \citep{tut00}, , while the other cross was oriented at a position angle of $45\,$ degrees to the first cross."
" The disk suggested by the observations of Tuthilletal.(2000) is only ~0.05""ru in", The disk suggested by the observations of \citet{tut00} is only $\sim 0.05\arcsec$ in
variety of Comptonized spectra may be generated. depending on the assumptions of the model.,"variety of Comptonized spectra may be generated, depending on the assumptions of the model."
 Further progress can only be made when the physical mechanism of energy dissipation is understood., Further progress can only be made when the physical mechanism of energy dissipation is understood.
 It could be thermal or nonthermal. and its radial dependence is important for the emerging spectrum.," It could be thermal or nonthermal, and its radial dependence is important for the emerging spectrum."
 Three heating mechanisms can operate in the subphotospheric region., Three heating mechanisms can operate in the subphotospheric region.
 Two of them (internal shocks and magnetic dissipation) rely of collisionless plasma processes. whose details remain uncertain because of their complexity.," Two of them (internal shocks and magnetic dissipation) rely of collisionless plasma processes, whose details remain uncertain because of their complexity."
 The third mechanism (collisional heating) is straightforward and can be modeled from first. principles., The third mechanism (collisional heating) is straightforward and can be modeled from first principles.
 It was recently proposed that collisional heating shapes the GRB spectra (B10)., It was recently proposed that collisional heating shapes the GRB spectra (B10).
 It will be the focus of the present paper., It will be the focus of the present paper.
 Collisional heating ts particularly strong (and inevitable) in jets that carry free neutrons (Derishevetal.1999;Bah-2000a:Rossietal.2006:Koers&Giannios 2007).," Collisional heating is particularly strong (and inevitable) in jets that carry free neutrons \citep{Derishev99,BahcallMeszaros2000,Fuller2000,MR00a,Rossi06,KG07}."
. Before the complete decoupling of neutrons and protons into two non-interacting components. the fading »-p collisions heat the protons to a mildly relativistic temperature.," Before the complete decoupling of neutrons and protons into two non-interacting components, the fading $n$ $p$ collisions heat the protons to a mildly relativistic temperature."
 Importantly. the #-p collisions also create ο plasma that efficiently radiates its energy.," Importantly, the $n$ $p$ collisions also create $e^\pm$ plasma that efficiently radiates its energy."
 The c plasma is cooled by radiative losses and heated by Coulomb collisions with hot protons., The $e^\pm$ plasma is cooled by radiative losses and heated by Coulomb collisions with hot protons.
 The rates of all these processes are well known and their effect on the jet and radiation can be calculated without additional assumptions (B10)., The rates of all these processes are well known and their effect on the jet and radiation can be calculated without additional assumptions (B10).
 In particular. one can show that the electron distribution function has two parts. thermal and nonthermal. and both Comptonize the radiation carried by the jet.," In particular, one can show that the electron distribution function has two parts, thermal and nonthermal, and both Comptonize the radiation carried by the jet."
" Three features of collisional heating make it à promising emission mechanism: (1) It peaks at radi r~0.1/,. not too early and not too late to produce a bright photospheric emission. ("," Three features of collisional heating make it a promising emission mechanism: (1) It peaks at radii $r\sim 0.1\Rph$, not too early and not too late to produce a bright photospheric emission. ("
2) It has a high radiative. efficiency. comparable to50%.. (,"2) It has a high radiative efficiency, comparable to. ("
3) The radiation spectrum produced by collisional heating can be accurately calculated from first principles.,3) The radiation spectrum produced by collisional heating can be accurately calculated from first principles.
" The striking result reported in BIO is that the spectrum has the Band-type shape with the high-energy photon index 2.5, which is consistent with observations."," The striking result reported in B10 is that the spectrum has the Band-type shape with the high-energy photon index $\beta\sim -2.5$ , which is consistent with observations."
 The formation of the radiation spectrum in à heated jet is a nonlinear problem that requires self-consistent calculation of the plasma and radiation behavior., The formation of the radiation spectrum in a heated jet is a nonlinear problem that requires self-consistent calculation of the plasma and radiation behavior.
 This problem was solved in BIO using a Monte-Carlo radiative transfer code in combination with an iterative technique., This problem was solved in B10 using a Monte-Carlo radiative transfer code in combination with an iterative technique.
 Our present work has two goals: First. to provide an independent check of the result of B10. we calculate the emitted spectrum using a different method: we solve the kinetic equations for the plasma and radiation in the jet.," Our present work has two goals: First, to provide an independent check of the result of B10, we calculate the emitted spectrum using a different method: we solve the kinetic equations for the plasma and radiation in the jet."
 We use the numerical code described in Vurm&Poutanen(2009.hereafter VPO9).., We use the numerical code described in \citet[][hereafter VP09]{VP09}.
 We develop a new version of the code to adapt it to GRB jets. calculate the emerging radiation. spectrum. and compare the results with those of BIO.," We develop a new version of the code to adapt it to GRB jets, calculate the emerging radiation spectrum, and compare the results with those of B10."
 Secondly. the numerical models in BIO were limited to weakly magnetized flows. p<1. and neglected synchrotron emission.," Secondly, the numerical models in B10 were limited to weakly magnetized flows, $\eB \ll 1$, and neglected synchrotron emission."
 In the present. paper. our models include synchrotron emission and self-absorption. and we explore the spectra produced by magnetized neutron-loaded jets for a range of0«sp<2.," In the present paper, our models include synchrotron emission and self-absorption, and we explore the spectra produced by magnetized neutron-loaded jets for a range of $0<\eB <2$."
 The paper is organized as follows., The paper is organized as follows.
 Section 2. describes the physical model and the setup of the simulations., Section \ref{GRB:Simulsetup} describes the physical model and the setup of the simulations.
 Section 3 describes our method of calculations., Section \ref{sec:method} describes our method of calculations.
 The results are presented in Sections + and 5 and discussed in Section 6.., The results are presented in Sections \ref{GRB:results1} and \ref{GRB:results2} and discussed in Section \ref{GRB:conclusions}.
 The central engines of GRBs must be compact and hot objects. which are almost certainly neutron rich.," The central engines of GRBs must be compact and hot objects, which are almost certainly neutron rich."
 The high density and temperature inevitably leads to .}-equilibrium (Imshenniketal.1967:Derishev1999;Beloborodov2003).," The high density and temperature inevitably leads to $\beta$ -equilibrium \citep{Imshennik67,Derishev99,Beloborodov03}."
 For all plausible parameters of the central engine. the equilibrium establishes a high fraction of free neutrons in the baryonie matter (this feature is seen. for example. in the accretion-disk model for GRBs. see Beloborodov2008 for a review).," For all plausible parameters of the central engine, the equilibrium establishes a high fraction of free neutrons in the baryonic matter (this feature is seen, for example, in the accretion-disk model for GRBs, see \citealt{Beloborodov08} for a review)."
 Therefore. neutrons are generally expected in GRB jets.," Therefore, neutrons are generally expected in GRB jets."
 The presence of neutrons creates perfect conditions for strong collisional heating., The presence of neutrons creates perfect conditions for strong collisional heating.
 In any variable baryonic jet. the inter-penetrating neutron and proton components inevitably develop relative motions (see BIO and references therein).," In any variable baryonic jet, the inter-penetrating neutron and proton components inevitably develop relative motions (see B10 and references therein)."
" The two components acquire different Lorentz factors T, and DL. and (rare) collisions between neutrons and protons dissipate enormous energy."," The two components acquire different Lorentz factors $\Gn$ and $\Gp$, and (rare) collisions between neutrons and protons dissipate enormous energy."
" The dissipation is strongest in those parts of the jet where D/T, is highest.", The dissipation is strongest in those parts of the jet where $\Gp/\Gn$ is highest.
 The variable jet consists of many causally disconnected shells., The variable jet consists of many causally disconnected shells.
 Different shells are accelerated to different Lorentz factors and emit different radiation. which is consistent with observed variability.," Different shells are accelerated to different Lorentz factors and emit different radiation, which is consistent with observed variability."
 In our simulations. we focus on the heating history and radiation of one shell.," In our simulations, we focus on the heating history and radiation of one shell."
 The shell is launched from the central source and accelerated at the expense of its internal energy until it enters the coasting matter-dominated phase of expansion., The shell is launched from the central source and accelerated at the expense of its internal energy until it enters the coasting matter-dominated phase of expansion.
" At radius {η the timescale for collisions. between the neutron and proton components exceeds the expansion timescale of the jet. and protons start to migrate relative to neutrons. forming a compound flow with T= L,."," At radius $R_{\rm n}$, the timescale for collisions between the neutron and proton components exceeds the expansion timescale of the jet, and protons start to migrate relative to neutrons, forming a compound flow with $\Gp\neq\Gn$."
 At this point. strong collisional heating begins and the radiation spectrum starts deviating from a blackbody.," At this point, strong collisional heating begins and the radiation spectrum starts deviating from a blackbody."
" Our simulation starts at A, and follows this evolution.", Our simulation starts at $R_{\rm n}$ and follows this evolution.
" We assume that DP and I4, remain approximately constant till the end of the simulation (see BIO for discussion of this approximation).", We assume that $\Gp$ and $\Gn$ remain approximately constant till the end of the simulation (see B10 for discussion of this approximation).
 Note that migration of neutrons between different shells of a strongly variable jet can create compound flows with DT., Note that migration of neutrons between different shells of a strongly variable jet can create compound flows with $\Gamma\gg\Gn$.
" δι may be defined as the radius where the neutron flow becomes “optically thin"" to collisions with protons.", $R_{\rm n}$ may be defined as the radius where the neutron flow becomes “optically thin” to collisions with protons.
" The corresponding ""optical depth"" for this process at radius + is given by m= . where i4, is the comovingπα number density of the neutron flow. £,,=bre is its kinetic luminosity (isotropic equivalent) and a,Pan3«10.7in? is the effective cross-section for nuclear collisions."," The corresponding “optical depth” for this process at radius $r$ is given by = = where $\nneutr$ is the comoving number density of the neutron flow, $\Ln = 4\pi \mneutr c^3 r^2 \Gn^2 \nneutr$ is its kinetic luminosity (isotropic equivalent) and $\sigman \sim 3 \times 10^{-26} \, {\rm cm}^2$ is the effective cross-section for nuclear collisions."
 Radius η is defined so that nG)—1. We aim to calculate the evolution of electron/positron and photon distribution functions in the heated jet.," Radius $R_{\rm n}$ is defined so that $\taun(R_{\rm n})=1$, = We aim to calculate the evolution of electron/positron and photon distribution functions in the heated jet."
 The photon spectrum is controlled by electrons (and. positrons) via Compton scattering and synchrotron emission., The photon spectrum is controlled by electrons (and positrons) via Compton scattering and synchrotron emission.
 Therefore. the key ingredient of any GRB model is how energy is injected into electrons.," Therefore, the key ingredient of any GRB model is how energy is injected into electrons."
 For the model studied in this paper. the electron heating is unambiguously determined by the collisional processes as described below (see Β10 for detailed discussion).," For the model studied in this paper, the electron heating is unambiguously determined by the collisional processes as described below (see B10 for detailed discussion)."
 The electron heating comes from baryons. in two forms:," The electron heating comes from baryons, in two forms:"
Equation (2)) below). then the radiation reaction-limitecl regime is reached while most of the orbit is outside the laver. where |D|=By.,"Equation \ref{eq-gamma_rad0}) ) below), then the radiation reaction-limited regime is reached while most of the orbit is outside the layer, where $|B|= B_0$."
 The evele-averaged radiation reaction force fq balances the electric force when , The cycle-averaged radiation reaction force $f_{\rm rad}$ balances the electric force when = .
"By comparing μας and με. we see that the radiation reaction limit is reached. before aa Only if the laver is sufficiently long. /45>0.5ETE7D,oE4."," By comparing $\gamma_{\rm max}$ and $\gamma_{\rm rad,*}$, we see that the radiation reaction limit is reached before $\gamma_{\rm max}$ only if the layer is sufficiently long, $l_{16} > 0.5\beta_{\rm rec}^{-1/2}\,B_{0,-3}^{-3/2}$."
 Then. the particle moves in a sinusoidal-like orbit with 52μα=const and efficiently converts all the energy. it gains from the electric field into svnchrotron 1photons with a characteristic enerevg hi4=(9/ imc = somewhat too low to explain the Crab Lares.," Then, the particle moves in a sinusoidal-like orbit with $\gamma\simeq \gamma_{\rm rad,0} ={\rm const}$ and efficiently converts all the energy it gains from the electric field into synchrotron photons with a characteristic energy = (3/2) _c = m_e c^2 =, somewhat too low to explain the Crab flares."
 llowever. despite the saturation of 5. the particle orbit continues to shrink towsrds the midplane with 6)exp(—Az/B8Peteacx)," However, despite the saturation of $\gamma$, the particle orbit continues to shrink towards the midplane with $\theta_0\sim \exp(-\beta_{\rm rec}z/3\rho_c \gamma_{\rm rad,*})$."
" Uf the laver is long enough. then gna, eventually shrinks below 9. reducing radiative losses and (hus making extremely. hieh particle energies possible."," If the layer is long enough, then $\ymax$ eventually shrinks below $\delta$, reducing radiative losses and thus making extremely high particle energies possible."
 Indeed. when the entire trajectory is contained deep within (he laver. (hemaximum perpendicular magnetic field sampled by the particle. Djs=Duas)vBogas/9. becomes smaller than Dy. thus reducing μμ for a fixed 5. whereas the accelerating force e£= remains unchanged.," Indeed, when the entire trajectory is contained deep within the layer, themaximum perpendicular magnetic field sampled by the particle, $B_{\rm max}=B_x(y_{\rm max}) \sim B_0\, y_{\rm max}/\delta$, becomes smaller than $B_0$, thus reducing $f_{\rm rad}$ for a fixed $\gamma$, whereas the accelerating force $eE_0 = e \beta_{\rm rec}\, B_0$ remains unchanged."
" Correspondingly. the radiation reaction limit 7,44 Hicreases in compensalion to maintain the balance between(he evele-averaged [μα and the constant electric acceleration: 54pad72sdasενας,O/{Ymax uas od Mav rise well above frac.μι a8 maxgas shrinks."," Correspondingly, the radiation reaction limit $\gamma_{\rm rad}$ increases in compensation to maintain the balance betweenthe cycle-averaged $f_{\rm rad}$ and the constant electric acceleration: $\gamma_{\rm rad} \approx \gamma_{\rm rad,*}\, \delta/\ymax$ , and may rise well above $\gamma_{\rm rad,*}$ as $y_{\rm max}$ shrinks."
 The, The
"photometric redshifts by extrapolating from an incomplete training set, if this is done with a survey such as DES it will create systematic errors on w of the order of ~10%..","photometric redshifts by extrapolating from an incomplete training set, if this is done with a survey such as DES it will create systematic errors on $w$ of the order of $\sim$."
 The possibility of using a template-fitting method to calculate photometric redshifts in regions where the training set is incomplete should be further investigated., The possibility of using a template-fitting method to calculate photometric redshifts in regions where the training set is incomplete should be further investigated.
 In this work we have shown the role of near infra-red photometry from the VISTA Hemisphere Survey in constraining photometric redshifts for the Dark Energy Survey., In this work we have shown the role of near infra-red photometry from the VISTA Hemisphere Survey in constraining photometric redshifts for the Dark Energy Survey.
 We have examined the effects of galaxy reddening and training sets on the photo-z estimate., We have examined the effects of galaxy reddening and training sets on the photo-z estimate.
" We have also studied the biases in the photometric redshift estimate when using training and testing sets with different sizes and levels of completeness and quantified the error in the dark energy equation of state parameter, w that arises from differences in these biases."," We have also studied the biases in the photometric redshift estimate when using training and testing sets with different sizes and levels of completeness and quantified the error in the dark energy equation of state parameter, $w$ that arises from differences in these biases."
 A method of clipping our galaxy catalogues by removing outliers based on the ANNz error estimate on the photometric redshift has been presented., A method of clipping our galaxy catalogues by removing outliers based on the ANNz error estimate on the photometric redshift has been presented.
" By applying a threshold error to our catalogues and rejecting all objects with errors bigger than this threshold error, we can effectively reduce the overall scatter on the photometric redshifts of our sample."," By applying a threshold error to our catalogues and rejecting all objects with errors bigger than this threshold error, we can effectively reduce the overall scatter on the photometric redshifts of our sample."
" Finally, we have conducted a full galaxy power spectrum analysis using our DES and DES+VHS catalogues and looked at how our clipping method can improve the uncertainties on our galaxy power spectrum measurements."," Finally, we have conducted a full galaxy power spectrum analysis using our DES and DES+VHS catalogues and looked at how our clipping method can improve the uncertainties on our galaxy power spectrum measurements."
 We find that there is an optimum threshold error at which we should clip our catalogues and this error depends on the catalogue being used and the redshift range in which we are evaluating the power spectrum., We find that there is an optimum threshold error at which we should clip our catalogues and this error depends on the catalogue being used and the redshift range in which we are evaluating the power spectrum.
" If we use a high value for the threshold error, the scatter on our photometric redshift estimate is high leading to large positional uncertainties and therefore large errors in the power spectrum due to cosmic variance."," If we use a high value for the threshold error, the scatter on our photometric redshift estimate is high leading to large positional uncertainties and therefore large errors in the power spectrum due to cosmic variance."
" However, if we adopt a very low value for our threshold error, we remove most of the galaxies from our sample before calculating the power spectrum and the resulting uncertainties in the power spectrum are dominated by shot noise."," However, if we adopt a very low value for our threshold error, we remove most of the galaxies from our sample before calculating the power spectrum and the resulting uncertainties in the power spectrum are dominated by shot noise."
 We find that the optimum threshold error is smaller for the DES+VHS catalogues compared to the DES only catalogues in the same redshift range and hence more outliers are removed from this sample before analysis., We find that the optimum threshold error is smaller for the DES+VHS catalogues compared to the DES only catalogues in the same redshift range and hence more outliers are removed from this sample before analysis.
 Adding the VHS NIR data thus helps us to compute the galaxy power spectrum more accurately out to higher redshifts than for the DES only case., Adding the VHS NIR data thus helps us to compute the galaxy power spectrum more accurately out to higher redshifts than for the DES only case.
" In summary, our main conclusions are:"," In summary, our main conclusions are:"
2008).,.
. While no unexpected. such effects are an unwelcome complication in the interpretation of measurements.," While not unexpected, such effects are an unwelcome complication in the interpretation of measurements."
 A further complication arises where radio sources show gross asymmetries and bends that must argely result from environments that are different on each side of he nucleus., A further complication arises where radio sources show gross asymmetries and bends that must largely result from environments that are different on each side of the nucleus.
 In order to alleviate these effects. we have chosen to study the nearby radio source270.. whose twin jets lie close o the plane of the skv and whose symmetry is largely intact.," In order to alleviate these effects, we have chosen to study the nearby radio source, whose twin jets lie close to the plane of the sky and whose symmetry is largely intact."
 In this paper we report a deep X-ray observation of its jet and counterjet., In this paper we report a deep X-ray observation of its jet and counterjet.
 2270 is hosted by the nearby elliptical galaxy =0.00746.][]trager..," 270 is hosted by the nearby elliptical galaxy \\citep[$z=
0.00746$,][]{trager}."
 The galaxy’s major-axis extent to 25 B mag 7 is 244412 aresec at position angle 160. and the ellipticity is Ο.Ε (de.Vaucouleursetαἱ.1991.takenfromhttp://nedwww.ipac.caltech.edu/}.," The galaxy's major-axis extent to 25 B mag $^{-2}$ is $244\pm 12$ arcsec at position angle $\sim 160^\circ$ , and the ellipticity is 0.11 \citep[][taken from
http://nedwww.ipac.caltech.edu/]{devauc}."
. The radio emission extends roughly 520 arcsec at a position angle ~SS” (Birkinshaw&Davies 1985)... and is symmetrical on the large scale. as seen in Figure |. where the black contours of radio emission show the brightest parts of the twin lobes.," The radio emission extends roughly 520 arcsec at a position angle $\sim 88^\circ$ \citep{birkd}, and is symmetrical on the large scale, as seen in Figure \ref{fig:xrbw}
 where the black contours of radio emission show the brightest parts of the twin lobes."
 The jet radio emission to the west is somewhat brighter than the counterjet to the east. but the overall morphology suggests that the radio emission is not highly affected by relativistic beaming.," The jet radio emission to the west is somewhat brighter than the counterjet to the east, but the overall morphology suggests that the radio emission is not highly affected by relativistic beaming."
 Indeed. Pineretal.(2001) have used the apparent speed and the jet-to-counterjet brightness ratio from VLBI images to deduce a jet speed of (0.46+0.02)e and an inclination angle of0=6333.," Indeed, \citet{piner} have used the apparent speed and the jet-to-counterjet brightness ratio from VLBI images to deduce a jet speed of $(0.46 \pm
0.02)c$ and an inclination angle of $\theta = 63^\circ \pm 3^\circ$."
 X-ray emission from wwas first detected using citepfberg.. but the improved resolution of wwas required for component separation.," X-ray emission from was first detected using \\citep{fberg}, but the improved resolution of was required for component separation."
 Worrall&Birkin-shaw(1994). used the PSPC to separate an unresolved nuclear component from resolved emission both spatially and spectrally. attributing the latter to the thermal atmosphere of the source (see also Fig. 10).," \citet{wb} used the PSPC to separate an unresolved nuclear component from resolved emission both spatially and spectrally, attributing the latter to the thermal atmosphere of the source (see also Fig. \ref{fig:xrbw}) )."
 Further investigation of the thermal component with aand citepdavis.osmond.finjones found that the gas temperature increased with distance from4261.. with a group atmosphere taking over from interstellar gas at the outer extremities of the radio source.," Further investigation of the thermal component with and \\citep{davis, osmond, finjones} found that the gas temperature increased with distance from, with a group atmosphere taking over from interstellar gas at the outer extremities of the radio source."
 Gas density and temperature profiles assuming spherical symmetry and using the ddata discussed here appear in Humphreyetal.(2009)... and see O'Sullivanetal.(2010). for further discussion of the galaxy and group gas properties.," Gas density and temperature profiles assuming spherical symmetry and using the data discussed here appear in \citet{hump}, and see \cite{newpaper} for further discussion of the galaxy and group gas properties."
 Improved measurements of the X-ray core using aand an earlier 32-ks oobservation (Sambrunaetal.2003:GliozziChiabergeetal.2003:Donato2004:Zezas2005). revealed that part of the nuclear emission suffers from high intrinsic absorption. with Ng several times 1077 7 (see also Section 3.45) contrary to the common trend towards low intrinsic absorption in radio galaxies whose bolometric luminosities are small compared with their Eddington limits (Evansetal.2006}.," Improved measurements of the X-ray core using and an earlier 32-ks observation \citep{sambruna, gliozzi, chiaberge,
donato, zezas} revealed that part of the nuclear emission suffers from high intrinsic absorption, with $N_{\rm H}$ several times $10^{22}$ $^{-2}$ (see also Section \ref{sec:xnucleus}) ) contrary to the common trend towards low intrinsic absorption in radio galaxies whose bolometric luminosities are small compared with their Eddington limits \citep{evans-cores}."
. Resolved X-ray emission from both the jet and counterjet were first seen in the earlier oobservation (Chiabergeetal.2003:Zezas2005).. showing to be the brightest known example of counterjet X-ray emission Tom a low-power radio galaxy.," Resolved X-ray emission from both the jet and counterjet were first seen in the earlier observation \citep{chiaberge, zezas}, showing to be the brightest known example of counterjet X-ray emission from a low-power radio galaxy."
 The radio jets close to the nucleus on either side are largely undisturbed. and iis thus an excellent source for addressing jet physics.," The radio jets close to the nucleus on either side are largely undisturbed, and is thus an excellent source for addressing jet physics."
 In this paper we concentrate on the resolved X-ray jet and counterjet. and their connection with the unresolved nucleus and their surroundings.," In this paper we concentrate on the resolved X-ray jet and counterjet, and their connection with the unresolved nucleus and their surroundings."
 Based on our results. we suggest that X-ray synchrotron emission in low-power jets may be used not only to address issues of particle acceleration but also to probe time-variable power released by an active-galaxy nucleus.," Based on our results, we suggest that X-ray synchrotron emission in low-power jets may be used not only to address issues of particle acceleration but also to probe time-variable power released by an active-galaxy nucleus."
 We adopt fy=TO km ! ο., We adopt $H_0 = 70$ km $^{-1}$ $^{-1}$.
 | aremin corresponds to 9.? kpe at the distance of NGC 4261., 1 arcmin corresponds to 9.2 kpc at the distance of NGC 4261.
 We observed ün FAINT dati mode with the Advanced CCD Imaging Spectrometer (ACTS) on board oon 2008 February 12 COBSID 9569. sequence 600710).," We observed in FAINT data mode with the Advanced CCD Imaging Spectrometer (ACIS) on board on 2008 February 12 (OBSID 9569, sequence 600710)."
 Details of the instrument and its modes of operation can be found in the PProposers? ObservatoryGuide?., Details of the instrument and its modes of operation can be found in the Proposers' Observatory.
. Results presented here use and the calibration database., Results presented here use and the calibration database.
" We re-calibrated the data. with random pixelization removed and bad pixels masked. following the software ""threads! from the XX-ray Center(CXC). to make a new level 2 events file."," We re-calibrated the data, with random pixelization removed and bad pixels masked, following the software `threads' from the X-ray Center, to make a new level 2 events file."
 Only events with grades 0.2.3.4.6 were used. as recommended.," Only events with grades 0,2,3,4,6 were used, as recommended."
 After removal of time intervals when the background deviated more than 37 above the average value. the calibrated dataset from 2008 has an observation durationof 100.538 Ks.," After removal of time intervals when the background deviated more than $3\sigma$ above the average value, the calibrated dataset from 2008 has an observation durationof 100.538 ks."
 We had earlier made a shorter observation of, We had earlier made a shorter observation of
The 5 Andromedae (0 And) planetary svstem is the first multiple planetary svstem discovered. bevond our own Solar System around. a solar-like star (Butler 1999).,The $\upsilon$ Andromedae $\upsilon$ And) planetary system is the first multiple planetary system discovered beyond our own Solar System around a solar-like star (Butler 1999).
 Not, Not
this method allows the simultaneous detection of multiple dust »opulations that have different extinction curves. something that is not possible when simply deter11iiniug the average galaxy. color in the sample.,"this method allows the simultaneous detection of multiple dust populations that have different extinction curves, something that is not possible when simply determining the average galaxy color in the sample."
" lu Section 5 we demoustrate this using clusters whicl have been ""artificially! endowed with dust.", In Section \ref{tests} we demonstrate this using clusters which have been `artificially' endowed with dust.
 Each set of galaxy clusters that we examine usiug our procecire contalus at least 25 clusters. with most samples containing more than 100. which are widely cis1jbuted over 1000 sq.," Each set of galaxy clusters that we examine using our procedure contains at least 25 clusters, with most samples containing more than 100, which are widely distributed over $\sim 1000$ sq."
 dee., deg.
 of high Galactic Latitude., of high Galactic Latitude.
 Assuming that the distribution of cirrus dist in our own Galaxy does uot inijc the distribution of APML clusters. the ellects of reddening aud obscuration due to Calactic dust ou galaxies in our Cluster aud Control Croups should average out.," Assuming that the distribution of cirrus dust in our own Galaxy does not mimic the distribution of APM clusters, the effects of reddening and obscuration due to Galactic dust on galaxies in our Cluster and Control Groups should average out."
 We generated lists of galaxies Iving in the Cluster aud Control Groups for apsroxtinately 110 APNI clusters., We generated lists of galaxies lying in the Cluster and Control Groups for approximately 140 APM clusters.
" The algorithia detailed iu the previous section was used ou tliree separate sets of clusters: (1) the coiiplete sample. (i) sample with clusters whose apparent raclii are 30"" (Le. i> 0.01). aud (iii) clusters with ruo>30! (2« Q.01)."," The algorithm detailed in the previous section was used on three separate sets of clusters: (i) the complete sample, (ii) sample with clusters whose apparent radii are $\le 30\arcmin$ (i.e. $z \geq 0.04$ ), and (iii) clusters with $r_{clust} > 30\arcmin$ $z < 0.04$ )."
 Individual clusters canιοί be analyzed uxing our method «die to the small number of galaxies per pixel iu the CM plane., Individual clusters cannot be analyzed using our method due to the small number of galaxies per pixel in the CM plane.
 I1 tlie case of the combined set. the average uumber ol galaxies per pixel is ou the order of 10s. whe‘eas the number per pixel for an average cluser would be on the order of 1 at most.," In the case of the combined set, the average number of galaxies per pixel is on the order of 10s, whereas the number per pixel for an average cluster would be on the order of 1 at most."
 Each of the three sets of clsters were divided even further iu order to search for dust ou various size scales by removing te ceutral. or the outer regious of the Cluster Groups.," Each of the three sets of clusters were divided even further in order to search for dust on various size scales by removing the central, or the outer regions of the Cluster Groups."
 We show the results of our analysis in Figures 3 aud 6.. where we plot coutour levels of £1; (eq. 2))," We show the results of our analysis in Figures \ref{trialsa} and \ref{trialsb}, , where we plot contour levels of $\xi_{k,l}$ (eq. \ref{statistic}) )"
" asa [uuction of the applied shift iu te CM plane. AR=&AVR. and ACB,—H)7Ap(ByHR)."," as a function of the applied shift in the CM plane, $\Delta R=k~\Delta_p R$, and $\Delta (B_J-R)=l~\Delta_p (B_J-R)$."
 Figure κ)+) shows tlie results rom selecting various plivsical scales from our combined data set of roughly 110 APM clusters while Fieure 6 shows various subsets of clusters with racii pκ)+)Q (upper panels) and >30 (lower panels)., Figure \ref{trialsa} shows the results from selecting various physical scales from our combined data set of roughly 140 APM clusters while Figure \ref{trialsb} shows various subsets of clusters with radii $ \le 30\arcmin$ (upper panels) and $> 30\arcmin$ (lower panels).
 Iu all our analysis of APM chsters. the minima in £j; are always at (0.0): in other wods. we detect no extinction or reddening «ue to dust in any of the cases.," In all our analysis of APM clusters, the minima in $\xi_{k,l}$ are always at $(0,0)$; in other words, we detect no extinction or reddening due to dust in any of the cases."
 Note that the lowest contoiw level in the plots is not at 0. but typically at 0.2—0.3. consistent. with what oue should expec giveu he definition of £54 (eq. 2))," Note that the lowest contour level in the plots is not at 0, but typically at $\sim 0.2-0.3$, consistent with what one should expect given the definition of $\xi_{k,l}$ (eq. \ref{statistic}) )"
 aud the level of Poissou noise in tle Llluber of galaxies per CM-plane »ixel., and the level of Poisson noise in the number of galaxies per CM-plane pixel.
 To estimate the robustuess of our results we implemened a bootstrapping algoritlin., To estimate the robustness of our results we implemented a bootstrapping algorithm.
" Me andoily selected 100 sets of galaxy clusters from our sauple (keepiig the total number of clusters in each set the same. aud equal to the original αμα])er). »erforued our analysis of each set separately. aud tabulated the AA aud ACB,—HR) valles οςTLSuxdliug to the ninuna of &7)."," We randomly selected 100 sets of galaxy clusters from our sample (keeping the total number of clusters in each set the same, and equal to the original number), performed our analysis of each set separately, and tabulated the $\Delta R$ and $\Delta (B_J-R)$ values corresponding to the minima of $\xi_{k,l}$."
 By selectingOm random combinations of clusters we are in ellec alugljentiug some and suppressiug other random subsets of the data., By selecting random combinations of clusters we are in effect augmenting some and suppressing other random subsets of the data.
 The degree to which tle varlous stbsets of the data agree serves to test the robustuess of our results., The degree to which the various subsets of the data agree serves to test the robustness of our results.
 We interpret the histogram of he uumber oftimes that each, We interpret the histogram of the number oftimes that each
aand modified Newtonian dynamics (MOND) contexts are divergent.,and modified Newtonian dynamics (MOND) contexts are divergent.
It is commonly believed that prompt emission by GRBs is due to synchrotron radiation by electrons accelerated in internal shocks associated with relativistic jets (with a bulk Lorentz boost I of 100-1000).,It is commonly believed that prompt emission by GRBs is due to synchrotron radiation by electrons accelerated in internal shocks associated with relativistic jets (with a bulk Lorentz boost $\Gamma$ of 100-1000).
 A review of the theoretical and observational status of GRBs is beyond the scope of this paper., A review of the theoretical and observational status of GRBs is beyond the scope of this paper.
 The reader is referred to Meszaros(2006).., The reader is referred to \citet{2006RPPh...69.2259M}.
 The average prompt photon emission is often described by fitting a Band function (Bandetal1993).., The average prompt photon emission is often described by fitting a Band function \citep{1993ApJ...413..281B}.
" High energy neutrino emission by GRBs in coincidence with the prompt y-ray photons has been proposed as a consequence of GRBs being a candidate to produce (UHE) ultra high energy (up to zz 10?"" eV) cosmic rays 1995).. Waxman&B"," High energy neutrino emission by GRBs in coincidence with the prompt $\gamma$ -ray photons has been proposed as a consequence of GRBs being a candidate to produce (UHE) ultra high energy (up to $\approx$ $^{20}$ eV) cosmic rays \citep{1995PhRvL..75..386W,1995ApJ...453..883V}."
ahcall(1997) have calculated the diffuse flux due to GRBs., \citet{1997PhRvL..78.2292W} have calculated the diffuse flux due to GRBs.
 Their work has been further improved by calculating neutrino emission for individual bursts (Guettaetal2004). and, Their work has been further improved by calculating neutrino emission for individual bursts \citep{2004APh....20..429G} and
where H is the scale height of the disk. jj2tanh(1). and the inclination of the field line (o=B-By at the disk surface z2 (Cao&Spruit2002)...,"where $H$ is the scale height of the disk, $\eta_{\rm i}=\tanh(1)$, and the inclination of the field line $\kappa_0=B_z/B_R^{\rm S}$ at the disk surface $z=H$ \citep{2002A&A...385..289C}."
 This expression can reproduce the basic features of the Kippenhahn-Schlütter model (Kippenhahn&Schlüter1957) well either for weak or strong field cases., This expression can reproduce the basic features of the Kippenhahn-Schlütter model \citep{1957ZA.....43...36K} well either for weak or strong field cases.
 The vertical distribution of the accretion disk with magnetic fields is not exactly the Gaussian distribution of an isothermal disk., The vertical distribution of the accretion disk with magnetic fields is not exactly the Gaussian distribution of an isothermal disk.
 In this case. we define the disk scale height H as p(H)=p(Ovexp(—1/2). the disk seale height H can then be evaluated numerically with the given magnetic field line shape (11))," In this case, we define the disk scale height $H$ as $\rho(H)=\rho(0)\exp(-1/2)$ , the disk scale height $H$ can then be evaluated numerically with the given magnetic field line shape \ref{b_shape_1}) )."
" As done by Cao&Spruit(2002).. we use a fitting formula to calculate the scale height of the disk in the rest of this work. where and &,, is defined as which can be calculated with the balance between the advection and diffusion of the fields in the accretion flow (see Section 2.2)."," As done by \citet{2002A&A...385..289C}, we use a fitting formula to calculate the scale height of the disk in the rest of this work, where and $\xi_{\rm bz}$ is defined as which can be calculated with the balance between the advection and diffusion of the fields in the accretion flow (see Section 2.2)."
 We find that this formula can reproduce the numerical results quite well., We find that this formula can reproduce the numerical results quite well.
" We note that equation (12)) reduces to H2c, /Ogin the absence of magnetic field."," We note that equation \ref{h_1}) ) reduces to $H=c_{\rm
s}/\Omega_{\rm K}$ in the absence of magnetic field."
" We define dimensionless quantities by where Bo is the strength of the putative external imposed homogeneous vertical magnetic fields. and Roy, is the outer radius of the accretion disk."," We define dimensionless quantities by where $B_0$ is the strength of the putative external imposed homogeneous vertical magnetic fields, and $R_{\rm out}$ is the outer radius of the accretion disk."
"We note, as WP found, that parameter errors scale linearly with for Α΄ points taken during the peak (defined as lasung [/j;)'..","We note, as WP found, that parameter errors scale linearly with for $N$ points taken during the peak (defined as lasting $4 \te$."
 Thus our results can be scaled for other numbers of observations with different values of σ., Thus our results can be scaled for other numbers of observations with different values of $\sigma$.
" Thus, we find that our results agree with those of WP if we assume the baseline magnitude is known and take a uniform sampling."," Thus, we find that our results agree with those of WP if we assume the baseline magnitude is known and take a uniform sampling."
"where fis the zone umber aloug the line of sight. dr), is its optical depth aud Πρες i8 the electron umuber density within the zone (47k).","where $k$ is the zone number along the line of sight, $d\tau_k$ is its optical depth and $n_{e,ijk}$ is the electron number density within the zone $(ijk)$."
" Refraction is important ouly below ~0.2 GIIz. where it accounts for the low-frequency turnover evident iu the spectra shown in Figure ὃ,"," Refraction is important only below $\sim 0.2$ GHz, where it accounts for the low-frequency turnover evident in the spectra shown in Figure 8."
 The large scale gas morphology through the ceutral dark cluster toward the end of the bydrodvuamical sinulation is shown in Figures | and 5., The large scale gas morphology through the central dark cluster toward the end of the hydrodynamical simulation is shown in Figures 4 and 5.
 Figure bis a plot of the column density viewed along the 2-axis while Figure 5 shows the emitted intensity G.c.. the integrated eiuissivitv alone the line-ofsieht).," Figure 4 is a plot of the column density viewed along the $z$ -axis while Figure 5 shows the emitted intensity (i.e., the integrated emissivity along the line-of-sight)."
" Of particular iutercst in these nuages is the appearance of streaks of high-velocity eas (""streamers) that are very renuüulsceut of features. such as the so-called “Bullet” seem near the Calactic ceuter (Yusef-Zacdel. et al."," Of particular interest in these images is the appearance of streaks of high-velocity gas (“streamers”) that are very reminiscent of features, such as the so-called “Bullet” seen near the Galactic center (Yusef-Zadeh, et al."
 1996)., 1996).
 Iu our simulation. these structures are produced predominantly within the wiud-wiud collision regions. and in the future. we shall consider in greater detail the possibility that the observed high-velocity eas colupouents near Ser Α are produced in this fashion.," In our simulation, these structures are produced predominantly within the wind-wind collision regions, and in the future, we shall consider in greater detail the possibility that the observed high-velocity gas components near Sgr A* are produced in this fashion."
 Note iu these figures the relative locations of the wind-producing stars (see also Fie., Note in these figures the relative locations of the wind-producing stars (see also Fig.
 1). aud the dominant role plaved by IRS I3E1 to the lower melt of the conpact radio source.," 1), and the dominant role played by IRS 13E1 to the lower right of the compact radio source."
 It should also be noted that the eas distribution iu a multiple-viud source cuviroument like that modeled here is distinctly ditfereut from that of a uuiform flow past à central accretor (Coker Melia 1997)., It should also be noted that the gas distribution in a multiple-wind source environment like that modeled here is distinctly different from that of a uniform flow past a central accretor (Coker Melia 1997).
 Uulike the latter. the former does not produce a large-scale bow shock. aud therefore the cuviroumental iupact of the eravitational focusing by the central dark mass las sguificautlv less order iu this case.," Unlike the latter, the former does not produce a large-scale bow shock, and therefore the environmental impact of the gravitational focusing by the central dark mass has significantly less order in this case."
 For exaniple. it is less likely iu a niultiple-wiud source environment that the interaction between a centralized mass and the Calactic ceuter wind cau craft the type of channeled flow required to produce the miui-cavity.," For example, it is less likely in a multiple-wind source environment that the interaction between a centralized mass and the Galactic center wind can craft the type of channeled flow required to produce the mini-cavity."
 We shall defer a iiore exteusive discussion of this poiut to a later publication iu which we report the results of a anultiple-wind source λαποια for the case in which the dark matter is iu the form of a massive black hole rather than a distributed dark cluster., We shall defer a more extensive discussion of this point to a later publication in which we report the results of a multiple-wind source simulation for the case in which the dark matter is in the form of a massive black hole rather than a distributed dark cluster.
 Iu Figure 6 we show the mass enclosed witlin a sphere of radius 0.1 Roy as a fiction of time. centered ou the midpoint of the dark cluster.," In Figure 6 we show the mass enclosed within a sphere of radius 0.1 $R_A$ as a function of time, centered on the midpoint of the dark cluster."
 Tn simular fashion. Figure 7 is a plot of the enclosed energy within this same volume.," In similar fashion, Figure 7 is a plot of the enclosed energy within this same volume."
 After reaching equilibrium several sound crossing times after the start of the simulation. the enclosed mass and cucrey beein to fluctuate aperiodicallv. reflecting the turbulent cell nature of the flow iu aud out of the central regiow.," After reaching equilibrium several sound crossing times after the start of the simulation, the enclosed mass and energy begin to fluctuate aperiodically, reflecting the turbulent cell nature of the flow in and out of the central region."
 Typically 2.7ος10.7M. of eas (the dashed line in this figure) is trapped within the cluster at any eiven time.," Typically $2.7\times 10^{-3}
\;M_\odot$ of gas (the dashed line in this figure) is trapped within the cluster at any given time."
 Note also that although the eas is highly supersonic. with most of," Note also that although the gas is highly supersonic, with most of"
but the PCA analysis indicates that it is possible to represent them in a space of N' dimensions (the number of chosen eigenvectors). with N«Ny.,"but the PCA analysis indicates that it is possible to represent them in a space of $N_\lambda'$ dimensions (the number of chosen eigenvectors), with $N_\lambda' \ll N_\lambda$."
 It is instructive to apply the same analysis based on principal components to a data set composed only of uncorrelated noise., It is instructive to apply the same analysis based on principal components to a data set composed only of uncorrelated noise.
" In the limit VN;--co and Nas,—c. the cross-product matrix is strictly equal to the identity."," In the limit $N_\lambda \to \infty$ and $N_\mathrm{obs} \to \infty$, the cross-product matrix is strictly equal to the identity."
 As a consequence. the eigenvectors are the canonical basis and the eigenvalues are all equal to 1.," As a consequence, the eigenvectors are the canonical basis and the eigenvalues are all equal to 1."
" However. since NV, is small. the cross-product matrix has non-diagonal elements and some spurious correlation may appear between different wavelength points."," However, since $N_\lambda$ is small, the cross-product matrix has non-diagonal elements and some spurious correlation may appear between different wavelength points."
 The bottom left and middle panels of Fig., The bottom left and middle panels of Fig.
 | show the first two eigenvectors of a matrix with the same size as the one of the observations but containing only gaussiannoise*., \ref{autovect} show the first two eigenvectors of a matrix with the same size as the one of the observations but containing only gaussian.
. The gaussian distribution of noise has a standard deviation equal to 107 I..., The gaussian distribution of noise has a standard deviation equal to $10^{-3}$ $_\mathrm{c}$.
 All the eigenvectors are noisy. similar to those shown in the figure.," All the eigenvectors are noisy, similar to those shown in the figure."
 The eigenvalues. represented in the bottom right panel of Fig.," The eigenvalues, represented in the bottom right panel of Fig."
 | are roughly the same for all the eigenvectors., \ref{autovect} are roughly the same for all the eigenvectors.
 Real spectro-polarimetric measurements can present some degree of correlation between different wavelength points but are also contaminated by uncorrelated noise., Real spectro-polarimetric measurements can present some degree of correlation between different wavelength points but are also contaminated by uncorrelated noise.
 Consequently. it is clear that the noise level will be an important issue that will restrain the denoising of the observations by means of principal components.," Consequently, it is clear that the noise level will be an important issue that will restrain the denoising of the observations by means of principal components."
 To simulate the real case. we add gaussian noise to each spectral line of our matrix of observations.," To simulate the real case, we add gaussian noise to each spectral line of our matrix of observations."
 We apply the denoising procedure to our simulated data set with different values of added noise in order to analyze how it behaves in different situations., We apply the denoising procedure to our simulated data set with different values of added noise in order to analyze how it behaves in different situations.
 In order to use only one reference signal to noise ratio we choose the ratio between the polarisation signal (Q. U or V) in the magnetically sensitive 630.2 nm line and the standard deviation of the added noise distribution.," In order to use only one reference signal to noise ratio we choose the ratio between the polarisation signal $Q$, $U$ or $V$ ) in the magnetically sensitive 630.2 nm line and the standard deviation of the added noise distribution."
 ote that this S/N σίνες an idea of the quality of the data for the most magnetically sensitive lines. while the median of the n/N distribution is located at a much lower value. typically Che order of magnitude smaller.," Note that this $S/N$ gives an idea of the quality of the data for the most magnetically sensitive lines, while the median of the $S/N$ distribution is located at a much lower value, typically one order of magnitude smaller."
 This is produced by the fact that most of the lines induce small polarization signals. while only few lines give conspicuous signals.," This is produced by the fact that most of the lines induce small polarization signals, while only few lines give conspicuous signals."
 Consequently. the distribution of amplitudes ts strongly shifted towards zero.," Consequently, the distribution of amplitudes is strongly shifted towards zero."
 For a given S/N. the eigenvectors of the cross-product matrix are computed as described in Section 2...," For a given $S/N$, the eigenvectors of the cross-product matrix are computed as described in Section \ref{sec:pca}."
 If any systematic Zeeman signature is present in the dataset. it will appear in the first few eigenvectors.," If any systematic Zeeman signature is present in the dataset, it will appear in the first few eigenvectors."
 Since uncorrelated noise is also present in the observations (perhaps even completely masking the line polarization signals). the rest of eigenvectors having smaller eigenvalues contain the contribution of this noise.," Since uncorrelated noise is also present in the observations (perhaps even completely masking the line polarization signals), the rest of eigenvectors having smaller eigenvalues contain the contribution of this noise."
 The filtering procedure consists on reconstructing the observed signal using only the first eigenvectors: where Ó' is the matrix of observations after the denoising procedure., The filtering procedure consists on reconstructing the observed signal using only the first eigenvectors: where $\hat{O'}$ is the matrix of observations after the denoising procedure.
 The matrix B’ contains the few first eigenvectors that have been retained as containing stellar magnetic signatures., The matrix $\hat{B'}$ contains the few first eigenvectors that have been retained as containing stellar magnetic signatures.
 The matrix C contains the coefficients of the, The matrix $\hat{C'}$ contains the coefficients of the
templates.,templates.
 Thus. we have no reason to exclude them when computing LEs.," Thus, we have no reason to exclude them when computing LFs."
 The inspection of these two figures indicates that the statistical properties of the photometric redshift s;auuples display the expected behaviour., The inspection of these two figures indicates that the statistical properties of the photometric redshift samples display the expected behaviour.
 According to our previous siauulations. aimuiug to reproduce the properties of the IIDFEs (Bolzouclla et citchyperz)). the τίdshitt distribution bevond the spectroscopic limits cau be cousidered as reliable.," According to our previous simulations, aiming to reproduce the properties of the HDFs (Bolzonella et \\cite{hyperz}) ), the redshift distribution beyond the spectroscopic limits can be considered as reliable."
 Iu Fig., In Fig.
 3 we plot the appareut colour ωνAv. versus he photometric redshift. as well as the colours computed roni the svuthetic SEDs of differcut types: from top to vottom. an Elliptical galaxy at fixed ages of 10 GCvr aud SCCar. evolving E. Sa. Bu. hu of fixed 1 Gyr aud αντ age.," \ref{z_ik} we plot the apparent colour $I_{814}-K_s$ versus the photometric redshift, as well as the colours computed from the synthetic SEDs of different types: from top to bottom, an Elliptical galaxy at fixed ages of $10$ Gyr and $5$ Gyr, evolving E, Sa, Im, Im of fixed $1$ Gyr and $0.1$ Gyr age."
 A few objects redder than old ellipticals seein to be present at 2<1 iu both fields., A few objects redder than old ellipticals seem to be present at $z<1$ in both fields.
" Also. a group of objects with £i,>26 and 1S2 is detected: hese objects are ΕΠΟΣ according to the usual criteria CChnatti et citecinatti: Scodegeiooo Silva 20001: Morioudo ot citemonm))."," Also, a group of objects with $I_{814}-K_s > 2.6$ and $1 \la z \la 2$ is detected: these objects are EROs according to the usual criteria Cimatti et \\cite{cimatti}; Scodeggio Silva \cite{scodeggio}; Moriondo et \\cite{mori}) )."
 Figure |. shows the redshift distributions. as obtained frou the photometric redshift computation. for the IIDEs catalogues.," Figure \ref{nzall} shows the redshift distributions, as obtained from the photometric redshift computation, for the HDFs catalogues."
 We have compared our result on the ΠΟΤΕ with the redshift distributions obtained bv Fernamucez-Soto et ((1999.. SUNY eroup) aud Foutana et ((2000.. Roma eroup). using the same catalog (sec Section 2)).," We have compared our result on the HDF-N with the redshift distributions obtained by Fern\'anndez-Soto et \cite{fsoto}, SUNY group) and Fontana et \cite{fonta1}, Roma group), using the same catalog (see Section \ref{catalogue}) )."
 In the IIDE-S. we compared our redshift distribution. obtained with the catalogue by Vauzella et al. (2001))," In the HDF-S, we compared our redshift distribution, obtained with the catalogue by Vanzella et al. \cite{vanzella}) )"
 with the Ας). computed using the SUNY catalogue by Feruáuudez-Soto et al. ((1999)), with the $N(z)$ computed using the SUNY catalogue by Fernánndez-Soto et al. \cite{fsoto}) )
 aud Foutana et ((2000))., and Fontana et \cite{fonta1}) ).
 We lave also computed a median redshift distribution. obtained with the Aloute Carlo method we applied to estimate the Iuninositv function. aud detailed in Section 12...," We have also computed a median redshift distribution, obtained with the Monte Carlo method we applied to estimate the luminosity function, and detailed in Section \ref{montecarlo}."
" This redshift distribution takes iuto accouut the probabilitv ""uctious of individual objects. in such a wav tha objects will be scattered in their allowed range of photometric redshifts. according to their Pu."," This redshift distribution takes into account the probability functions of individual objects, in such a way that objects will be scattered in their allowed range of photometric redshifts, according to their $P(z_{\rm phot})$."
 Objects with flat Pius) Will spread over a wide range of τμ]. Whereas objects with narrow Pss) will be always assigucc to a redshift close to the best fit onc.," Objects with flat $P(z_{\rm phot})$ will spread over a wide range of $z_{\rm phot}$, whereas objects with narrow $P(z_{\rm phot})$ will be always assigned to a redshift close to the best fit one."
 We show the result of 100 iterations. represeuted by the median aud the error bars. computed as and of the values distribution in cach redshitt bin.," We show the result of 100 iterations, represented by the median and the error bars, computed as and of the values distribution in each redshift bin."
 Considering the auter sample οἱ the IIDE-N presented in the left panel of Figure 1.. the KolmogorovSuiiurnov test shows that the hvpeiz aud SUNY distributions are compatible. whereas the πρ and Roma distributions are not.," Considering the fainter sample of the HDF-N presented in the left panel of Figure \ref{nzall}, the Kolmogorov-Smirnov test shows that the hyperz and SUNY distributions are compatible, whereas the hyperz and Roma distributions are not."
 We reject the lypothesis that two distributions were issued from tle same parent distributions when the significance level is lower than a conservative value of, We reject the hypothesis that two distributions were issued from the same parent distributions when the significance level is lower than a conservative value of.
 Using the same criteri. the mecian Ας} computed with the hvperz-Moute Carlo method is compatible with both the SUNY and the Roma distributions.," Using the same criterium, the median $N(z)$ computed with the hyperz-Monte Carlo method is compatible with both the SUNY and the Roma distributions."
" At Igy,<26. the KS test shows that cach N(2) is cousisteut with the others."," At $I_{814} \le 26$, the KS test shows that each $N(z)$ is consistent with the others."
" Iu the IIDE-S. we show the comparison of N(2) for the sale two limiting magnitudes. fs),<28.5 and fey)26."," In the HDF-S, we show the comparison of $N(z)$ for the same two limiting magnitudes, $I_{814} \le 28.5$ and $I_{814} \le 26$."
 Foutana et ((2000)) selected their stbsample imposing κι27.5 (dashed liue).," Fontana et \cite{fonta1}) ) selected their subsample imposing $I_{814} \le
27.5$ (dashed line)."
" The KS test for the subsiuple with Aa,<28.5 shows that the distributions are uot drawn from the same parent distribution. but cousideriug the bright subsample with Joy)<26 the distributions are fully cousisteut each other. even if they are obtaimect frou different photometric catalogues."," The KS test for the subsample with $I_{814} \le
28.5$ shows that the distributions are not drawn from the same parent distribution, but considering the bright subsample with $I_{814} \le
26$ the distributions are fully consistent each other, even if they are obtained from different photometric catalogues."
 Tn Figure 5 we show the same comparison carried out for the A;-baud selected subsample we used iu the following of this paper., In Figure \ref{nzksel} we show the same comparison carried out for the $K_s$ -band selected subsample we used in the following of this paper.
 We have analysed the distributions obtained when selecting objects with Ax25 aud A.<23.5., We have analysed the distributions obtained when selecting objects with $K_s \le 25$ and $K_s \le 23.5$.
" In the IIDEF-N he agreciment among the different Ας} distributions is remarkable: by means of the Ixolinogorov-Suirnuov test we obtain that we can reject the »ossibilitv that the distributions were drawn frou cliffercut went distributions with a confidence level ranging from 23% (ninunuuu)zaud 70% Guaxiunuu)at A,<25. auda xobabilitv ranging from [1 o 100 at A,x23.5."," In the HDF-N the agreement among the different $N(z)$ distributions is remarkable: by means of the Kolmogorov-Smirnov test we obtain that we can reject the possibility that the distributions were drawn from different parent distributions with a confidence level ranging from $23$ (minimum) and $70$ (maximum) at $K_s \le 25$, and a probability ranging from $41$ to $100$ at $K_s \le 23.5$."
 Tn the IIDE-S the same comparison shows that the ivperz and SUNY distributions are not compatible. eveu choosing the couservative value of 1% for the confidence evel. whereas the byperz and Boma results are mareiually consistent.," In the HDF-S the same comparison shows that the hyperz and SUNY distributions are not compatible, even choosing the conservative value of $1$ for the confidence level, whereas the hyperz and Roma results are marginally consistent."
 On the other haud. the redshift distributions of the bright subsample show a better agreement. with xobabilities iu each case larger than the chosen Πατ. even if the catalogues are different. im particular iu the NIR dataset.," On the other hand, the redshift distributions of the bright subsample show a better agreement, with probabilities in each case larger than the chosen limit, even if the catalogues are different, in particular in the NIR dataset."
 In the lower right panel we also show the redshift istribution obtained by Buduick ct ((20013) for a sample limited to Ax:23.5 and a catalogue built bx these authors using ISAAC-VLT images.," In the lower right panel we also show the redshift distribution obtained by Rudnick et \cite{rudnick}) ) for a sample limited to $K_s \le
23.5$ and a catalogue built by these authors using ISAAC-VLT images."
 Also in tliis case we found that the redshift distributions are compatible., Also in this case we found that the redshift distributions are compatible.
" Even considering the subsample with A,x2h. ie. the Iit chosen computing LEs. the redshift distributions are 1 general consistent."," Even considering the subsample with $K_s \le 24$, i.e. the limit chosen computing LFs, the redshift distributions are in general consistent."
 Iu conclusion. some small differences are observed iu the redshift distribution when using different methods. mostly affecting the faiutest samples. whereas the results are conrpatible for the bright oues.," In conclusion, some small differences are observed in the redshift distribution when using different methods, mostly affecting the faintest samples, whereas the results are compatible for the bright ones."
 At faint limits oei magnitude. the redshift diy.tribution obtained with the hvperz-Monte Carlo method. avoids iutrocducing spurious features. aud it makes the relevant N(:) distributions compatible.," At faint limits in magnitude, the redshift distribution obtained with the hyperz-Monte Carlo method avoids introducing spurious features, and it makes the relevant $N(z)$ distributions compatible."
 At brighter limits. by means of the comparison with the other photometric redshift estimates. we lave shown hat the results are cousistet aud characterized bv a Heh quality im the redshift determination.," At brighter limits, by means of the comparison with the other photometric redshift estimates, we have shown that the results are consistent and characterized by a high quality in the redshift determination."
 The differences among different authors become very smooth when the Monte Carlo approach is used: this procedure allow us to compute reliable LFs in Sect. 5., The differences among different authors become very smooth when the Monte Carlo approach is used; this procedure allow us to compute reliable LFs in Sect. \ref{resu}. .
Figure d shows that the nine NLAMDBO sources are not cistributed uuiornlv arotud the radio galaxy. although the sensitivity of our MAMDO map radially decreases from the radio οalaxy.,"Figure \ref{RMAMBOSN} shows that the nine MAMBO sources are not distributed uniformly around the radio galaxy, although the sensitivity of our MAMBO map radially decreases from the radio galaxy."
 The our brightest NLAMDBO sources (MOI. MO2. MYF and MU5) are all north of the radio ealaxy. while the deusest area of cenutters 1s Incated sottheast of the radio ealaxy (Venenuiuisetaεν2002:Mileval..2001).," The four brightest MAMBO sources (M01, M02, M07 and M05) are all north of the radio galaxy, while the densest area of emitters is located southeast of the radio galaxy \citep{ven02,mil04}."
.. This suggests that the centre of the proo-cluster is not necessarily to the southeast of the radio galaxy. as suggested by the cistribution of uitters.," This suggests that the centre of the proto-cluster is not necessarily to the southeast of the radio galaxy, as suggested by the distribution of emitters."
 None of the 11 spectroscopicaIx confirmed nitters in the NLAMDBO field were deected at 1.2 nuu Or 1.1 GIIz., None of the 14 spectroscopically confirmed emitters in the MAMBO field were detected at 1.2 mm or 1.4 GHz.
 Suuilulv. none of the LANBO sources show excess eenission Gf they are within the LOTS<:«1.123 coverage of the nurow-bauc ‘fhter).," Similarly, none of the MAMBO sources show excess emission (if they are within the $4.075<z<4.123$ coverage of the narrow-band filter)."
 A possible explanation would be that their large amounts of dust. as traced w the 1.2 nuu emission may rave quenched the eenission (6.g.wunhetal..1998).. puting them below he narrow-band excess cuoff y+ 215 ÀJ," A possible explanation would be that their large amounts of dust, as traced by the 1.2 mm emission may have quenched the emission \citep[\eg][]{kun98}, putting them below the narrow-band excess cutoff $_{\rm rest}>$ 15 )."
) This apparent absence of overla» betwee i]th. populations shows the nuJOLTAicc of using nuliple waveleust[um echuiques to obtain a lilore COPete picture of the xoto-cluster., This apparent absence of overlap between both populations shows the importance of using multiple wavelength techniques to obtain a more complete picture of the proto-cluster.
 €‘hapaetal.(2000) aud Webbetal.(2003). also re)ort a <20% overlap eween SAICs aud LBCs. with a few jotable exceptiois fey.Chapmanetal. 2002).," \citet{cha00} and \citet{webb03} also report a $<$ overlap between SMGs and LBGs, with a few notable exceptions \citep[\eg][]{cha02}."
. Iowever. Webbetal.(2003) do fiud a large cross-clusen1g correlation ampitude between both populations in the largest of their 3 fields.," However, \citet{webb03} do find a large cross-clustering correlation amplitude between both populations in the largest of their 3 fields."
" We do uot detect erough MÁAMDO sources in tιο TN 1912 ficd to perform such a cross-correlatic1n analysis. but our deection of asatistical ox""erdeusitv ¢ft MAMBO sources does suggest sone relation to the coined overdensity of eeniters (Venuehansetal.2002) aud LBCs 2001) in the sane field."," We do not detect enough MAMBO sources in the TN $-$ 1942 field to perform such a cross-correlation analysis, but our detection of a statistical overdensity of MAMBO sources does suggest some relation to the confirmed overdensity of emitters \citep{ven02} and LBGs \citep{mil04} in the same field."
 The fiüutuess of the in our objects also coutrasts the published spectra of SMCs. which ofen show bright Hines (Chapiiaιotal.2003).," The faintness of the in our objects also contrasts the published spectra of SMGs, which often show bright lines \citep{cha03}."
. Ilowever. these SAIC are οji average at lower redshifts (2)2.1 (Chapa-etal. 2003).. while our MAMDO sources are »otential members of the :—L1 proto-cluster.," However, these SMGs are on average at lower redshifts $\langle z \rangle \sim 2.4$ \citep{cha03}, while our MAMBO sources are potential members of the $z$ =4.1 proto-cluster."
 Altrough we on have A band information for 3) sources. which are all relatively bright class-II (2Ks hlW S21) or Il Nz5) sources (Ivisonetal.200koSuadl 2002a).. 3 of the 1 brightest NLAMDBO sources nav well be verv faint cass-ü (A= 21) sources which are much harder to obti rredshifts from (6.9.Danjierbaneretal.2002:Fraverctal.. 2001). either because they are more highly obscured. or becasο they are at higher redshifts.," Although we only have $K-$ band information for 3 sources, which are all relatively bright class-II $I-K\simlt 5$, $K\simlt $ 21) or class-I $I-K\simgt 5$ ) sources \citep{ivi00,sma02a}, 3 of the 4 brightest MAMBO sources may well be very faint class-0 $K\simgt 21$ ) sources which are much harder to obtain redshifts from \citep[\eg][]{dan02,fra04}, either because they are more highly obscured, or because they are at higher redshifts."
 Our success in findine higher redshift objects may have been helped by the use of a longer selection wavelength (1.2 mun instead of 850 pau)., Our success in finding higher redshift objects may have been helped by the use of a longer selection wavelength (1.2 mm instead of 850 $\mu$ m).
 Indeed. based on the low ratio of S50 qi nu to 1.2 uuu flux. several authors argue that a significant fraction of nia galaxies may be at 2>3 (οι.Ealesetal..2103:Aretxagaal.2003).," Indeed, based on the low ratio of 850 $\mu$ m to 1.2 mm flux, several authors argue that a significant fraction of mm galaxies may be at $z>3$ \citep[\eg][]{eal03,are03}."
. Most \LAMBO sources in the NTT deep field (Dannerbaucretal..2002) are also class) SOUILCCS. ST]yporting the trend that 1.2 nuu selected sources appear to )e fainter (and hence maybe at ueher redshift) than 850 ;nu selected sources.," Most MAMBO sources in the NTT deep field \citep{dan02} are also class-0 sources, supporting the trend that 1.2 mm selected sources appear to be fainter (and hence maybe at higher redshift) than 850 $\mu$ m selected sources."
 The detection of Xawoondsson idm several bright ΠΠ sources within fie proto-clusters sumroundius hree other WzRCs (Sualetal.2003b) suggests that AGN may be preseut witlin those sources., The detection of X-ray emission in several bright submm sources within the proto-clusters surrounding three other HzRGs \citep{sma03b} suggests that AGN may be present within those sources.
 Πα of the spectra presented by €!liauuanetal.(2003) also show ΥΡΟΠ AGN lines. bu we failed to detect these in our very deep VET. spectrosecoy of MOT awl MOD.," Half of the spectra presented by \citet{cha03} also show type-II AGN lines, but we failed to detect these in our very deep VLT spectroscopy of M01 and M05."
 IHencoe. we have no nedication fvat ACN are preseut là our ALAMIDO sources. but hei enussion lues may well have been obscure0l (eg.RalandetaL.2003).," Hence, we have no indication that AGN are present in our MAMBO sources, but their emission lines may well have been obscured \citep[\eg][]{reu03}."
" Ássuimius ALAMDO sources are at z— 1.1. aud ¢ust parameters =h0 Iv and j—1.5 (Beutordοal.1000).. we derive Lyn~--PERSLoeL,NoaSyoun in the rouge 10??L... Πρι star formationiiJxy raCS of 1000. to ο lo if the dust In ©itfively joatec| bv star quation (ey.OmoutctaL.2OL:DeDreck 3)."," Assuming the MAMBO sources are at $z$ =4.1, and dust parameters $T_{\rm d}=$ 50 K and $\beta$ =1.5 \citep{ben99}, we derive $L_{\rm FIR} \approx 2\times 10^{13}{\rm L_{\odot}} \times S_{\rm 1.2 mm}/{\rm mJy}$ in the range $\times 10^{13} {\rm L_{\odot}}$, implying star formation rates of 4000 to 12000 $_{\odot}$ $^{-1}$, if the dust is entirely heated by star formation \citep[\eg][]{omo01,deb03}."
". For MOS. which is possirly at 2 1.15. we derive ~--1.6«10PL, and a star fornatio1 rate of 1600 Moxy +."," For M08, which is possibly at $z$ =1.18, we derive $L_{\rm FIR} \approx 1.6\times 10^{13}{\rm L_{\odot}}$ and a star formation rate of 1600 $_{\odot}$ $^{-1}$."
 Such high stax fornation rates have been sorted for IEZRCis (Agchibaldeal..2001:Reulandetal. M. but our NAMDO imap incicates that they may also occur out to distances as far as 2N pe from the ceutral radio ealaxy.," Such high star formation rates have been reported for HzRGs \citep{arc01,reu04}, but our MAMBO map indicates that they may also occur out to distances as far as 2 Mpc from the central radio galaxy."
 If the MAMDO sources are really at += L1. this would sugeestOO that these proo-clusters have multiple. possibly aligned (Pentericcieta..2002) concentrations.," If the MAMBO sources are really at $z$ =4.1, this would suggest that these proto-clusters have multiple, possibly aligned \citep{pen02} concentrations."
 Iowever. the VET narrow-band image shows that none of the nuu sources seen in the TN 1912 proto-cluster have huge hhaloes. like those seen around the central Πρ (e.y.Venemausetal.2002:Reuland 2003).," However, the VLT narrow-band image shows that none of the mm sources seen in the TN $-$ 1942 proto-cluster have huge haloes, like those seen around the central HzRG \citep[\eg][]{ven02,reu03}."
. The ITZRC. therefore appears the best caudidates to evolve into the present-day eiaut elliptical., The HzRG therefore appears the best candidates to evolve into the present-day giant elliptical.
 We sununuuize the results from our imultiwavcleusthi observations of the D—ll proto-cluster surroundius TN 1912 as follows: e We detect 10 candidate nui sources with peak fluxes having 73 iu our MAMBO map., We summarize the results from our multi-wavelength observations of the $z$ =4.1 proto-cluster surrounding TN $-$ 1942 as follows: $\bullet$ We detect 10 candidate mm sources with peak fluxes having $>$ 3 in our MAMBO map.
" Of these. at least eieht sources with 5S4244,4 71.33 11. have possible radio and/or optical/near-IR. counterparts. aud 5 are coufirmed at S/NLl with pointed SCUBA subuuu photonetry."," Of these, at least eight sources with $S_{\rm 1.2 mm}>$ 1.3 mJy have possible radio and/or optical/near-IR counterparts, and 5 are confirmed at $>$ 4 with pointed SCUBA submm photometry."
" Three sources have Syoni,2 L0 mJy. while comrug with source counts from blank field surveys. we expect to find only 1 such source in the unassociatec field population."," Three sources have $S_{\rm 1.2 mm}>$ 4.0 mJy, while comparing with source counts from blank field surveys, we expect to find only 1 such source in the unassociated field population."
 e The radio-to-xubnuu aud ιαΚα. photonetric redshift estimates do not provide stroug constraiits Ol, $\bullet$ The radio-to-submm and mm-submm photometric redshift estimates do not provide strong constraints on
In general. the coalescence rate of a binary svstem containing an observable radio pulsar is defimed by where Αρα is the estimated number of pulsars in our Galaxy. with pulse profiles and orbital characteristics similar to those of the known systems. fi is a correction [actor for pulsar beanming and την. is the lifetime of the binary. svstem.,"In general, the coalescence rate of a binary system containing an observable radio pulsar is defined by where $N_{\rm {PSR}}$ is the estimated number of pulsars in our Galaxy with pulse profiles and orbital characteristics similar to those of the known systems, $f_{\rm b}$ is a correction factor for pulsar beaming and $\tau_{\rm life}$ is the lifetime of the binary system."
 In the following subsections. we caleulate the total lifetime of a pulsar binary aid derive the probability density hanction (PDF) of the Galactic coalescence rate. PR). for binaries.," In the following subsections, we calculate the total lifetime of a pulsar binary and derive the probability density function (PDF) of the Galactic coalescence rate, $P$ ), for binaries."
 In Table 1.. we summarize (he observational properties and relevant lifetimes for the 3 pulsar svstems considered in (his work.," In Table \ref{tab:obsprop}, we summarize the observational properties and relevant lifetimes for the 3 pulsar systems considered in this work."
 We define the lifetime of a coalescing pulsar binary τις lo be the sum of the current age of the observable pulsar and the remaining lifetime ol the svstem.," We define the lifetime of a coalescing pulsar binary $\tau_{\rm
life}$ to be the sum of the current age of the observable pulsar and the remaining lifetime of the system."
" Assiuning the pulsar spins down [rom an initial period 1) to the currently observed value J? (both in s) due to a non-decaving magnetic dipole radiation torque (see e.g. (1977)). ils current. ""spin-down age where P (in !) is the observed period derivative."," Assuming the pulsar spins down from an initial period $P_0$ to the currently observed value $P$ (both in s) due to a non-decaying magnetic dipole radiation torque (see e.g. ), its current “spin-down” age where $\dot{P}$ (in $^{-1}$ ) is the observed period derivative."
" For voung pulsars like J1141—6545. it is usually assumed that 2,«P so that τα reduces to the familiar characteristic age τιo—P/(2P)e1.5 Myr."," For young pulsars like $-$ 6545, it is usually assumed that $P_0 \ll P$ so that $\tau_{\rm sd}$ reduces to the familiar characteristic age $\tau_{\rm
c}=P/(2 \dot{P}) \simeq 1.5$ Myr."
 For the older recycled. pulsars. however. pointed out that this assumption is usually not. appropriate. since the weaker magnetic fields of these objects mean that their present spin periods are only moderately larger than the periods produced during accretion.," For the older recycled pulsars, however, pointed out that this assumption is usually not appropriate, since the weaker magnetic fields of these objects mean that their present spin periods are only moderately larger than the periods produced during accretion."
 Adopting the spin-up line [rom(1999). we may write Using the above (wo equations we caleulate τα for the two recycled pulsars JOT51+1807 and J1757—5322 to be 6.7 and 4.9 Gyr respectively (see Table 1)).," Adopting the spin-up line from, we may write Using the above two equations we calculate $\tau_{\rm sd}$ for the two recycled pulsars J0751+1807 and $-$ 5322 to be 6.7 and 4.9 Gyr respectively (see Table \ref{tab:obsprop}) )."
package (Mink2002).,package \citep{mi02}.
. The derived astvometry was compared with previously determined coordinates from. 2ALASS and showed no svstematie offset., The derived astrometry was compared with previously determined coordinates from 2MASS and showed no systematic offset.
 All coordinates presented in Table 3 are estimated to be accurate to < 1.5., All coordinates presented in Table 3 are estimated to be accurate to $\leq$ 1.5”.
 The spectroscopic observations were performed with the erism GI41. which is centered ab 1.4 jm. with a wavelength coverage betsveen 1.2 jm - 1.9 jm and a resolution of z 200 per pixel.," The spectroscopic observations were performed with the grism G141, which is centered at 1.4 $\micron$, with a wavelength coverage between 1.2 $\micron$ - 1.9 $\micron$ and a resolution of $\approx$ 200 per pixel."
 This covers the water-band feature at. 1.4: jm which is expected in the atmospheres of cool objects ancl is dillieult to observe from the ground. making IST ideal for these observations.," This covers the water-band feature at 1.4 $\micron$ which is expected in the atmospheres of cool objects and is difficult to observe from the ground, making HST ideal for these observations."
 To reduce these data. (he images were first run through the HST pipelines CalnicA and CalnieD. These pipelines perform the dark subtraction. using artilicial dark frames from the STScI website. as well as cosmic ray reduction.," To reduce these data, the images were first run through the HST pipelines CalnicA and CalnicB. These pipelines perform the dark subtraction, using artificial dark frames from the STScI website, as well as cosmic ray reduction."
 The pipelines were not used to co-add either (the dithers or the roll angles. because this did not improve the quality of the spectra over extracting (hem from individual images.," The pipelines were not used to co-add either the dithers or the roll angles, because this did not improve the quality of the spectra over extracting them from individual images."
 This means that for each object in the NGC! 1333 sample there can be a maximum of 12 extracted spectra., This means that for each object in the NGC 1333 sample there can be a maximum of 12 extracted spectra.
 The spectra were extracted using a custom IDL routine NICMOSIook (Pirzkaletal.1905)... which was designed to deal specifically with NICAIOS erism data.," The spectra were extracted using a custom IDL routine NICMOSlook \citep{pi98}, which was designed to deal specifically with NICMOS grism data."
 NICMOSlIook reduces spectra in a standard wav. by tracing the spectrum across the chip and summing the flix in the spatial dimension as well as subtracting a background region set bv (he user al each wavelength.," NICMOSlook reduces spectra in a standard way, by tracing the spectrum across the chip and summing the flux in the spatial dimension as well as subtracting a background region set by the user at each wavelength."
" To extract spectra, NICMOSIook needs a spectroscopic as well as a photometric image. {ο perform the wavelength calibration."," To extract spectra, NICMOSlook needs a spectroscopic as well as a photometric image, to perform the wavelength calibration."
 The photometric image gives (he (rue location of the objects position. which (hen gives (he zero point lor the wavelength of the spectrum.," The photometric image gives the true location of the object's position, which then gives the zero point for the wavelength of the spectrum."
 The photonmetric image used was the F160W image because ils central wavelength is closer to the central wavelength of G141., The photometric image used was the F160W image because its central wavelength is closer to the central wavelength of G141.
 The spectra were extracted with an extraction width of roughly 2 x FWIIM of the spectra. in general z 4 pixels.," The spectra were extracted with an extraction width of roughly 2 x FWHM of the spectra, in general $\approx$ 4 pixels."
 The background size was varied depending on the crowding of the field., The background size was varied depending on the crowding of the field.
 NICMOSIook also IHlatfields (he spectra during Cie extraction., NICMOSlook also flatfields the spectra during the extraction.
 This is necessary because the quantum efficiency of (he detector changes wilh both wavelength ancl position., This is necessary because the quantum efficiency of the detector changes with both wavelength and position.
 For this reason the spectra cannot be flatfielcled before extraction., For this reason the spectra cannot be flatfielded before extraction.
 NICMOSIook flatfields the spectra using a set of narrowband flatfield images. which are chosen according to the date of (he observations. from which the software constructs a 3-dimensional calibration datacube.," NICMOSlook flatfields the spectra using a set of narrowband flatfield images, which are chosen according to the date of the observations, from which the software constructs a 3-dimensional calibration datacube."
 signal-to-noise ratios (SNR) were estimated for all extracted spectra in two wavs., Signal-to-noise ratios (SNR) were estimated for all extracted spectra in two ways.
 First. a 3rd order polynomial was fitted to the combined average spectrum in the spectral region between 1.5 - 1.3 jm and the RAIS of the difference between the observed. continuum ancl the fit per pixel calculated.," First, a 3rd order polynomial was fitted to the combined average spectrum in the spectral region between 1.5 - 1.8 $\micron$ and the RMS of the difference between the observed continuum and the fit per pixel calculated."
 Second. a SNR was calculated for each individual pixel in the," Second, a SNR was calculated for each individual pixel in the"
The micro-amplification factor µ determined in component D depends on both the date and the wavelength (Table 2)).,The micro-amplification factor $\mu$ determined in component D depends on both the date and the wavelength (Table \ref{tab:ampli}) ).
" Considering theSiiv-Civ spectral region, the strongest variation of µ occurs between 1989 and 1993 (see also Figs."," Considering the spectral region, the strongest variation of $\mu$ occurs between 1989 and 1993 (see also Figs."
 2 and 6)), \ref{fig:specvis} and \ref{fig:sepavis}) ).
 It roughly corresponds to a relative photometric variation between A and D which can be observed in the V light curves of ppresented by Remy et al. (1996)), It roughly corresponds to a relative photometric variation between A and D which can be observed in the V light curves of presented by Remy et al. \cite{rem96}) )
" and @stensen et al. (1997)),"," and stensen et al. \cite{ost97}) ),"
 superimposed onto the common intrinsic variation of the 4 components., superimposed onto the common intrinsic variation of the 4 components.
" Between 1993 and 2000, the variation of µ is weaker, in agreement with the HST photometry in the F555W filter reported by Turnshek et al. (1997))"," Between 1993 and 2000, the variation of $\mu$ is weaker, in agreement with the HST photometry in the F555W filter reported by Turnshek et al. \cite{tur97}) )"
 and Chae et al. (2001))., and Chae et al. \cite{cha01}) ).
" At a given epoch, µ decreases with increasing wavelength, suggesting chromatic magnification of the continuum source."," At a given epoch, $\mu$ decreases with increasing wavelength, suggesting chromatic magnification of the continuum source."
 This is best seen in the 2005 data (obtained within a 2 month interval) which span the largest wavelength range., This is best seen in the 2005 data (obtained within a 2 month interval) which span the largest wavelength range.
" We emphasize that 1, when determined from the line profiles, is not contaminated by differential extinction (Sect. 4.1))."," We emphasize that $\mu$, when determined from the line profiles, is not contaminated by differential extinction (Sect. \ref{sec:method}) )."
" In Fig. 10,,"," In Fig. \ref{fig:chroml},"
 we plot the values of µ as a function of the wavelength of observation in the quasar rest-frame., we plot the values of $\mu$ as a function of the wavelength of observation in the quasar rest-frame.
" The magnification µ of an extended source close to a caustic can be written where Rs is the πεςsource radius, Rg the Einstein radius of the microlens projected onto the source plane, g is a constant on the order of unity, Z(d) a function which depends on the distance to the caustic and fo a constant “background” magnification (e.g. Schneider et al. 1992,,"," The magnification $\mu$ of an extended source close to a caustic can be written where $R_S$ is the source radius, $R_E$ the Einstein radius of the microlens projected onto the source plane, $g$ is a constant on the order of unity, $\zeta(d)$ a function which depends on the distance to the caustic and $\mu_0$ a constant “background” magnification (e.g. Schneider et al. \cite{sch92},"
 Witt et al. 1993))., Witt et al. \cite{wit93}) ).
" In the framework a simple model where the continuum is emitted by a Shakura-Sunyaev (1973)) thin accretion disk thermally radiating, Rs(A)ος1319 (e.g. Poindexter et al. 2008))"," In the framework a simple model where the continuum is emitted by a Shakura-Sunyaev \cite{sha73}) ) thin accretion disk thermally radiating, $R_S(\lambda) \propto \lambda ^{4/3}$ (e.g. Poindexter et al. \cite{poi08}) )"
" such that we may expect As seen in Fig. 10,,"," such that we may expect As seen in Fig. \ref{fig:chroml},"
" this model nicely reproduces the data using C = —0.55 and assuming for simplicity uo = 1 (i.e. no background (de- )magnification), supporting the idea of chromatic magnification of a continuum emitted by a Sunyaev accretion disk."," this model nicely reproduces the data using $C$ = $-$ 0.55 and assuming for simplicity $\mu_0$ = 1 (i.e. no background (de- )magnification), supporting the idea of chromatic magnification of a continuum emitted by a Shakura-Sunyaev accretion disk."
" In principle, we could have used the Lya+Nv and the Οπή 41909 emission lines, also present in the 2005 visible spectra, to measure µ at other wavelengths."," In principle, we could have used the $\alpha$ and the ] $\lambda$ 1909 emission lines, also present in the 2005 visible spectra, to measure $\mu$ at other wavelengths."
" Although tentative estimates do agree with the observed trend, the quality of the data is not sufficient to derive reliable values of µ at these wavelengths, due to the insufficient spectral resolution in the complex Lya+N region and to the fact that the m] line is truncated."," Although tentative estimates do agree with the observed trend, the quality of the data is not sufficient to derive reliable values of $\mu$ at these wavelengths, due to the insufficient spectral resolution in the complex $\alpha$ region and to the fact that the ] line is truncated."
" Clearly, with better quality data, it could be possible to separate M and µ at other wavelengths using additional line profiles and thus derive the temperature profile of the accretion disk."," Clearly, with better quality data, it could be possible to separate $M$ and $\mu$ at other wavelengths using additional line profiles and thus derive the temperature profile of the accretion disk."
" From the value of the constant C, we can derive a rough estimate of the size of the continuum source : Rs(A)/Re€4*510-36."," From the value of the constant $C$, we can derive a rough estimate of the size of the continuum source : $R_S(\lambda) / R_E \lesssim \,
\lambda^{4/3} \, 10^{-2C}$ ."
" The Einstein radius is computed to be ΚΕ~ 0.01 VM/Ms pc using z; = 1.0, a flat cosmology, Q,, = 0.27 and Ho = 70 km s! Mpc!."," The Einstein radius is computed to be $R_E
\simeq$ 0.01 $\sqrt{M/M_{\odot}}$ pc using $z_{l}$ = 1.0, a flat cosmology, $\Omega_m$ = 0.27 and $_{\rm 0}$ = 70 km $^{-1}$ $^{-1}$."
 M is the mass of the microlens., $M$ is the mass of the microlens.
" Then, Rs(A)€ 0.012 /M/Mg pc at the rest-frame UV wavelength A= 0.15 um. This limits tightens to Rs(A)x 0.007 VM/Mo pc with z = 1.88."," Then, $R_S(\lambda)
\lesssim$ 0.012 $\sqrt{M/M_{\odot}}$ pc at the rest-frame UV wavelength $\lambda$ = 0.15 $\mu$ m. This limits tightens to $R_S(\lambda) \lesssim$ 0.007 $\sqrt{M/M_{\odot}}$ pc with $z_{l}$ = 1.88."
" These values agree with the radii obtained for the lensed quasars HE1104—1805 and Q2237+0305 on the basis of their photometric variability due to microlensing (e.g. Poindexter et al. 2008,,"," These values agree with the radii obtained for the lensed quasars $-$ 1805 and $+$ 0305 on the basis of their photometric variability due to microlensing (e.g. Poindexter et al. \cite{poi08}, ,"
" Eigenbrod et al. 2008,,"," Eigenbrod et al. \cite{eig08},"
 Anguita et al. 2008b))., Anguita et al. \cite{ang08b}) ).
 Dilution of the quasar continuum -microlensed- by the host galaxy light -not microlensed- can affect the interpretation of the wavelength dependence of µ., Dilution of the quasar continuum –microlensed– by the host galaxy light –not microlensed– can affect the interpretation of the wavelength dependence of $\mu$.
" Denoting a=F.(host)/F.(qso) the ratio of the host and quasar continua at a given wavelength, we find that the micro-amplification factor of the quasar continuum pg is related to the measured u by if the host galaxy —u(1lto)-ais unresolved and macrolensed as the quasar (i.e. ΤΕΕ.=MμιF-(qso)+ MF.(host)), or by if the host galaxy is resolved and not macrolensed (i.e. if F.=MugF-(qso)+F-(host) assuming the host contained in the measurement aperture)."," Denoting $\alpha = F_{c}({\rm
host})/F_{c}({\rm qso})$ the ratio of the host and quasar continua at a given wavelength, we find that the micro-amplification factor of the quasar continuum $\mu_{q}$ is related to the measured $\mu$ by if the host galaxy is unresolved and macrolensed as the quasar (i.e. if $F'_{c} = M \, \mu_q \, F_{c}({\rm qso}) + M F_{c}({\rm
host})$ ), or by if the host galaxy is resolved and not macrolensed (i.e. if $F'_{c} = M
\, \mu_q \, F_{c}({\rm qso}) + F_{c}({\rm host})$ assuming the host contained in the measurement aperture)."
 Dilution by the host galaxy is expected to be stronger in the rest-frame visible-infrared than in the ultraviolet., Dilution by the host galaxy is expected to be stronger in the rest-frame visible-infrared than in the ultraviolet.
" In the rest-frame visible of a luminous high-redshift object, ax0.15 (e.g. Schramm et al. 2008))"," In the rest-frame visible of a luminous high-redshift object, $\alpha \lesssim 0.15$ (e.g. Schramm et al. \cite{sch08}) )"
 so that Hq does not differ from u by more than (even less if the host is partially resolved)., so that $\mu_q$ does not differ from $\mu$ by more than (even less if the host is partially resolved).
 This is negligible and within the uncertainties., This is negligible and within the uncertainties.
 Chae et al. (2001)), Chae et al. \cite{cha01}) )
 have obtained thefirst polarization mesurements of the four images of uusingthe HST., have obtained thefirst polarization mesurements of the four images of usingthe HST.
 The F555W filter was used., The F555W filter was used.
" They noted that, in March 1999, the polarization degree of component D might be higher that the polarization degree of the other components"," They noted that, in March 1999, the polarization degree of component D might be higher that the polarization degree of the other components"
bolometric flux. ie. integrated over all frequencies by the observer.,"bolometric flux, i.e. integrated over all frequencies by the observer."
 Its R95 formulation is And the redshift z rus as From the definition of E. Oogiven in the previous section. oue obtains. in Euclidean space-time. (see also Sxlos Labini et al..," Its R95 formulation is And the redshift $z$ runs as From the definition of $\Gamma$, given in the previous section, one obtains, in Euclidean space-time, (see also Sylos Labini et al.,"
 1998): where 5 ds the area of a spherical shell of radius r., 1998): where $S$ is the area of a spherical shell of radius $r$.
 Aveuine that the relevant distance for an analysis performed im a curved space-time is dj. Ribeiro defines Dol) as where 5 is presented as being the area of the observed spherical shell. of racius dj. aud is thus written It is here worth noting that the hunuinositv distance is the observable quantity relevant for measured distances.," Arguing that the relevant distance for an analysis performed in a curved space-time is $d_l$ , Ribeiro defines $\Gamma (d_l)$ as where $S$ is presented as being the area of the observed spherical shell, of radius $d_l$, and is thus written It is here worth noting that the luminosity distance is the observable quantity relevant for measured distances."
" It is thus correct to use this quantitv to evaluate an clementary change in the umber density iN, as measured by the observer on his lisht cone.", It is thus correct to use this quantity to evaluate an elementary change in the number density $N_c$ as measured by the observer on his light cone.
" But. when looking at a cross sectional area dS perpendicular to the helt rav aud subteuding a solid angle dQ. the observer must consider the area distance d,. also called ""observer area distance  bv Ellis (1971) aud. ""corrected Iuuinositv distance” by Ivistian aud Sachs (1966). defined by and related to the luminosity distance of an object nieasured with the same redshift by Therefore. the area of the observed spherical shell. at a eiven redshift z. used in Eq.(11)) ust be written From Eqs.(11)) and (12)). we eet an expression for d, as a function of the redshift :: Using Eqx.(9)). (LU). (18)) aud (19)). we then derive Dx As can be easily seen. by comparingS Eqs.(15)) and (Ls}}. this expression differs from the one that would proceed from R95 by a factor (1|Db "," But, when looking at a cross sectional area $dS$ perpendicular to the light ray and subtending a solid angle $d\Omega$, the observer must consider the area distance $d_a$, also called “observer area distance ” by Ellis (1971) and “corrected luminosity distance” by Kristian and Sachs (1966), defined by and related to the luminosity distance of an object measured with the same redshift by Therefore, the area of the observed spherical shell, at a given redshift $z$, used in \ref{gameds}) ) must be written From \ref{dleds}) ) and \ref{red}) ), we get an expression for $d_l$ as a function of the redshift $z$ : Using \ref{nc}) ), \ref{gameds}) ), \ref{sdla}) ) and \ref{delz}) ), we then derive $\Gamma(z)$: As can be easily seen, by comparing \ref{sdl}) ) and \ref{sdla}) ), this expression differs from the one that would proceed from R95 by a factor $(1+z)^4$ ."
Therefore. its sensitivity to the redshift value is wach lower.," Therefore, its sensitivity to the redshift value is much lower."
 The voluue clement ο. relevant for calculations ou the light cone. is Substituting Eq.(19)) iuto Eq.(21)) and iutegratiug. we obtain Substituting this expression iuto Eq.(7)) and using Eq.(20)). we eet which can be casily compared to Eqs.(19) aud (20) in R95.," The volume element $dV$, relevant for calculations on the light cone, is Substituting \ref{delz}) ) into \ref{dv}) ) and integrating, we obtain Substituting this expression into \ref{gamstarl}) ) and using \ref{gamzl}) ), we get which can be easily compared to Eqs.(19) and (20) in R95."
 We thus see that. contrary to Bibeiro's clau. D differs from the average nuniber deusitv «n> if the correct relativistic quantitics are retained.," We thus see that, contrary to Ribeiro's claim, $\Gamma^*$ differs from the average number density $<n>$ if the correct relativistic quantities are retained."
 It is teresting to note that. as d;(:). D Is an increasing function of τν hot a decreasing onc. as is Bibeiro's <4>(:) function.," It is interesting to note that, as $d_l(z)$ , $\Gamma^*$ is an increasing function of $z$, not a decreasing one, as is Ribeiro's $<n>(z)$ function."
 Furthermore. if we take temporarily ZZ=75 Wins + + to compare our results with those proposed iu R95. we see that jetween 2=0 and 2=1 the increase of our P is only (instead of a decrease for Ribeiro).," Furthermore, if we take temporarily $H_0=75$ Km $^{-1}$ $^{-1}$ to compare our results with those proposed in R95, we see that between $z=0$ and $z=1$ the increase of our $\Gamma^*$ is only (instead of a decrease for Ribeiro's)."
 A significant departure of E from its value in au Euclidean universe should therefore occur at redshifts lugher than claimed in R95., A significant departure of $\Gamma^*$ from its value in an Euclidean universe should therefore occur at redshifts higher than claimed in R95.
 Its custine conclusions thus do not agree with the analysis. which is confirmed by the results of our simulations preseuted below.," Its ensuing conclusions thus do not agree with the analysis, which is confirmed by the results of our simulations presented below."
 It is also interesting to note tha the Increase du ow at:= is markedly less than the Increase of the οἱ ratio. where rg=Πρ (n units e= 1) is the Euclidean distance.," It is also interesting to note that the increase in our $\Gamma^*$ at $z=1$ is markedly less than the increase of the $d_l/r_E$ ratio, where $r_E=z/H_0$ (in units $c=1$ ) is the Euclidean distance."
 Iu au Eiusteiu-de Sittermodel of the universe. the redshift-distauce relation therefore provides a large wpper-limit for the departure of D from it» Euclidean counterpart.," In an Einstein-de Sittermodel of the universe, the redshift-distance relation therefore provides a large upper-limit for the departure of $\Gamma^*$ from its Euclidean counterpart."
The simulations presented below show that this result can be generalized. to a wide class of standardhomogeueous models ofthe universe.,The simulations presented below show that this result can be generalized to a wide class of standardhomogeneous models ofthe universe.
galaxies.,galaxies.
 BECs are known to have positive color egracdicuts (i.e. center bluer than outsknt) on average and consist of both old stars and vouug stars (Menanteauetal.20012:Eliegreeu2006.2008. 2010a).," BEGs are known to have positive color gradients (i.e. center bluer than outskirt) on average and consist of both old stars and young stars \citep{men01a,elm05,fer05,lee06,lee08,lee10a}."
. It. las been also revealed that the internal color distributious of BECs are neither homologous nor svunmetric (Menauteauetal.2001:Lee 2006).. aud that some BECs are suspected to host active galactic nuclei (AGNs) (Menauteauetal.2005:Lee2006.Ww XN).," It has been also revealed that the internal color distributions of BEGs are neither homologous nor symmetric \citep{men04,lee06}, and that some BEGs are suspected to host active galactic nuclei (AGNs) \citep{men05,lee06,lee08}."
 Unlike typical τος early-type ealaxies (RECs) preferring high density euvironnieuts. BECs tend to reside i intermediate or low density enuvironnieuts (Ikaunappanetal.2009:Lee 2010b).," Unlike typical red early-type galaxies (REGs) preferring high density environments, BEGs tend to reside in intermediate or low density environments \citep{kan09,lee10b}."
. The origin aud future of DECs are still coutroversial., The origin and future of BEGs are still controversial.
 Observational evidence inclicates dat sole BECs may originate frou recent galaxy interaction or lucrecrs (Leeetal.2006).. but ie gas iufall is also a possible. iuechauism to »uld BECs (Moeuauteauοal.20015).," Observational evidence indicates that some BEGs may originate from recent galaxy interaction or mergers \citep{lee06}, but the gas infall is also a possible mechanism to build BEGs \citep{men01b}."
.. After sole tine passes; BECGs may evolve iuto typical RECs (Leeetal.2006.2007.2008).. but it is also »ossible that they evolve into bulges of ate-tyvpe ealaxies (Tameretal.2005:Kannappan 20093.," After some time passes, BEGs may evolve into typical REGs \citep{lee06,lee07,lee08}, but it is also possible that they evolve into bulges of late-type galaxies \citep{ham05,kan09}."
. In anv case. BECSs are xobablv iu the ornius/erowiug phase of carly-vpe galaxies or lees.," In any case, BEGs are probably in the forming/growing phase of early-type galaxies or bulges."
 To uuderstaud better the formation aud evolution of DECs. we need to inspect what happens in the inside of BECs.," To understand better the formation and evolution of BEGs, we need to inspect what happens in the inside of BEGs."
 Since more han a half of BEGs have evidence of their ποσοinteraction origiu (Leeetal.2006).. the iuteraction-iuduced. star formation (SE) is expected to make DEG blue.," Since more than a half of BEGs have evidence of their merger/interaction origin \citep{lee06}, the interaction-induced star formation (SF) is expected to make BEGs blue."
 Ou the other hand. galaxy mereger/iuteraction sonnetinies also causes ACN activity Usewley& 2009).. aud the power-law coutiima of AGNs may affect the color of BECs.," On the other hand, galaxy merger/interaction sometimes also causes AGN activity \citep{kew03,san05,com09}, and the power-law continua of AGNs may affect the color of BEGs."
" Leeeta.(2008) reported that nou-passive BEGs in the Sloan Digital Sky Survey (SDSS:Yorketal.2000). include both SE ealaxies and AGN host ealaxies. the umber ratio of which is about 3:5 for 22.9κ1M, 1207."," \citet{lee08} reported that non-passive BEGs in the Sloan Digital Sky Survey \citep[SDSS;][]{yor00} include both SF galaxies and AGN host galaxies, the number ratio of which is about 3:5 for $-22.9<{^{0.1}M_{r}}$ $<-20.7$."
 ILlowewer. optical spectra are casily obscured by dust. causing the missing of a sienificaut fraction of ACNs in optical surveys (Bestetal.2005).," However, optical spectra are easily obscured by dust, causing the missing of a significant fraction of AGNs in optical surveys \citep{bes05}."
. Moreover. eround-based spectroscopy the SDSS spectroscopy) of bright ealaxies in optical band often suffers the dilution of low-contrast emission lines due ο strong cussion lines (IIoctal.1997).. which also causes iuaccuracv du estimating ACN activity.," Moreover, ground-based spectroscopy the SDSS spectroscopy) of bright galaxies in the optical band often suffers the dilution of low-contrast emission lines due to strong emission lines \citep{ho97}, which also causes inaccuracy in estimating AGN activity."
 One of the most powerful tools to probe the SF/AGN activity is the spectroscopy in the infrared (IR) baud., One of the most powerful tools to probe the SF/AGN activity is the spectroscopy in the infrared (IR) band.
 Polveyeic aromatic hydrocarbons (PATI:Sclleren1981:Leger&Puget1981). have )ecn detected 3n various objects. such as region. protostellar clouds. dlanetary nebulae. aud SF galaxies (Peetersetal.2002:Tomy2002).," Polycyclic aromatic hydrocarbons \citep[PAHs;][]{sel84,leg84} have been detected in various objects, such as region, protostellar clouds, planetary nebulae, and SF galaxies \citep{pee02,hon02}."
.. Tt is shown that UV photons excite the vibrational and stretching modes of PAIL molecules. from he relaxation of which broad IR cussion lines (such as 3.3. 6.2e d.d. 8.6 and. 11.2 jan lines) are xoduced. (Cherchueff.&Barker19589:Shanctal.1991).," It is known that UV photons excite the vibrational and stretching modes of PAH molecules, from the relaxation of which broad IR emission lines (such as 3.3, 6.2, 7.7, 8.6 and 11.2 $\mu$ m lines) are produced \citep{che89,sha91}."
. Such ΡΑΠ emission features are good indicators of SE activity., Such PAH emission features are good indicators of SF activity.
 On the other haud.Voit(1992) argued. that PATIs in the ceutral region of an AGN cau o destructed by the stroug X-rav clussion from the ACN.," On the other hand,\citet{voi92} argued that PAHs in the central region of an AGN can be destructed by the strong X-ray emission from the AGN."
 Furthermore. ACN activity 1s known o suppress SF itself (AGNfertyetal. 2008)...," Furthermore, AGN activity is known to suppress SF itself \citep[AGN feedback;][]{ant08,raf08}."
 Thus. those properties of PATI Cluissionare believec to provide a reasonably clea diagnostic betweeu SE and ACN.," Thus, those properties of PAH emissionare believed to provide a reasonably clean diagnostic between SF and AGN."
 In adition. the shape of the IR coutiuunun is another indicator of SE/AQCN. because SE ealaxics aud ACNs have different soectral energy distribution (e.g.Tinanishietal.2008:Bisaliti2006. 2010)..," In addition, the shape of the IR continuum is another indicator of SF/AGN, because SF galaxies and AGNs have different spectral energy distribution \citep[e.g.][]{ima08,ris06,ris10}. ."
 Ilowever. no svscluatic IR spectroscopic surveys of DECs are sect rin the literature.," However, no systematic IR spectroscopic surveys of BEGs are seen in the literature."
 Iu this paper. we repor the result of the first DEG spectroscojc survey in the nem-IR (NIR) baud. using the ATARI (Muralzunietal.2007) TutraRed Camera (RC:Onakaetal.2007).," In this paper, we report the result of the first BEG spectroscopic survey in the near-IR (NIR) band, using the AKARI \citep{mur07} InfraRed Camera \citep[IRC;][]{ona07}."
. The outline of lis paper is as follows., The outline of this paper is as follows.
 Section 2 describes the target selection and the AIKARI spectroscopic observation., Section 2 describes the target selection and the AKARI spectroscopic observation.
 Section 3 explaius the reduction aud avalvsis of the AI&ARI data., Section 3 explains the reduction and analysis of the AKARI data.
 The extracted NIR 5]vectra of BEGs aud their features are shown in l., The extracted NIR spectra of BEGs and their features are shown in 4.
 The results are discussed iu 855.and the nuün results and thei miplicatiou aro sununanzed in 866.," The results are discussed in 5,and the main results and their implication are summarized in 6."
 Throughout this paper. we adopt the cosmological parameters: 5= 0.7. Q4— 0.7. and οap 03.," Throughout this paper, we adopt the cosmological parameters: $h=0.7$ $\Omega_{\Lambda}=0.7$ , and $\Omega_{M}=0.3$ ."
in Serenelletal.(2000).,in \citet{ssm:09}.
 The new models. in addition. adopt the uuclear reaction rates that were recently reconuuended in ὉFIL," The new models, in addition, adopt the nuclear reaction rates that were recently recommended in SFII."
 ΕΤΗ represents the first svstematic evaluation of hydrogen burning reactious since the SEI (Adclherecr and NACRE (Aneuloetal.1999) efforts., SFII represents the first systematic evaluation of hydrogen burning reactions since the SFI \citep{sfi} and NACRE \citep{nacre} efforts.
 SFIL provides a set of standard S-factors and uncertainties that reflect t16 progress made in laboratory aud theoretical nuc‘lear astroplivsics over the last decade., SFII provides a set of standard S-factors and uncertainties that reflect the progress made in laboratory and theoretical nuclear astrophysics over the last decade.
 The team of evaluators carefully considered which data sets were sufficiently complete aud well docimuenuted to be included: how t«» conibine data sets that span differeu enerev ranecs aud that nüeht © somewhat incousisteu in their error-bar clans: aud how to associate reasonable uncertainties to fittiug fncfious (when the data sets span the region of:wtrophlvsica interest) or to theoretica extrapolatious of data. (when LO data exist iu the solar Gamow peak)., The team of evaluators carefully considered which data sets were sufficiently complete and well documented to be included; how to combine data sets that span different energy ranges and that might be somewhat inconsistent in their error-bar claims; and how to associate reasonable uncertainties to fitting functions (when the data sets span the region of astrophysical interest) or to theoretical extrapolations of data (when no data exist in the solar Gamow peak).
 The treatnent of uncertanties was conservative. due to the iercasing relevance of the SSA as a laboratory for weak. nmeraction studies aud stellar evolution theory.," The treatment of uncertainties was conservative, due to the increasing relevance of the SSM as a laboratory for weak interaction studies and stellar evolution theory."
 The principal changes between SFII aud SFI are:, The principal changes between SFII and SFI are:
Secular aberration drift is an apparent alteration in the velocity of distant objects caused by the acceleration of the Solar System barycenter in the Milky Way.,Secular aberration drift is an apparent alteration in the velocity of distant objects caused by the acceleration of the Solar System barycenter in the Milky Way.
" The proper motion field of extragalactic sources affected by the secular aberration drift should present a dipolar structure with amplitude 4—6 microseconds of are per year (uas/yr) directed towards the Galactic center (see. e.g.. Fanselow 1983. Bastian 1995, Eubanks et al."," The proper motion field of extragalactic sources affected by the secular aberration drift should present a dipolar structure with amplitude 4--6 microseconds of arc per year $\mu$ as/yr) directed towards the Galactic center (see, e.g., Fanselow 1983, Bastian 1995, Eubanks et al."
" 1995, Sovers et al."," 1995, Sovers et al."
 1998. Migard 2002. Kovalevsky 2003. Kopeikin Makarov 2006).," 1998, Mignard 2002, Kovalevsky 2003, Kopeikin Makarov 2006)."
 Observations of extragalactic radio sources by very long baseline interferometry (VLBI) have been done since the late 1970s by various agencies. mainly the United States Navy and the National Aeronautic and Space Administration (NASA). and have been coordinated since 1998 by the International VLBI Service for Geodesy and Astrometry (IVS. Schlütter Behrend 2007).," Observations of extragalactic radio sources by very long baseline interferometry (VLBI) have been done since the late 1970s by various agencies, mainly the United States Navy and the National Aeronautic and Space Administration (NASA), and have been coordinated since 1998 by the International VLBI Service for Geodesy and Astrometry (IVS, Schlütter Behrend 2007)."
 The main purpose of these observations is to monitor the Earth orientation and realize terrestrial and celestial reference frames., The main purpose of these observations is to monitor the Earth orientation and realize terrestrial and celestial reference frames.
 One of the cornerstones of the VLBI technique ts the accurate realization of the celestial reference system using distant objects that supposedly have no detectable proper motion (Feissel Mignard 1998)., One of the cornerstones of the VLBI technique is the accurate realization of the celestial reference system using distant objects that supposedly have no detectable proper motion (Feissel Mignard 1998).
 The most recent realization of the International Celestial Reference Frame (ICRF2. Fey et al.," The most recent realization of the International Celestial Reference Frame (ICRF2, Fey et al."
 2009) provides absolute radio source coordinates for 3.414 sources.," 2009) provides absolute radio source coordinates for 3,414 sources."
 For 1.448 ICRFE2 sources observed in at least two sessions. the median formal error is 175 µας.," For 1,448 ICRF2 sources observed in at least two sessions, the median formal error is 175 $\mu$ as."
 For the most observed radio sources. the inflated position errors are 40 µας in each coordinate.," For the most observed radio sources, the inflated position errors are 40 $\mu$ as in each coordinate."
 This catalog did not include proper motions of sources because the dominant proper motion for any source is the somewhat temporary one associated with internal structure changes that produce apparent motions that are an order of magnitude greater than the aberrational drift (e.g.. Fey et al.," This catalog did not include proper motions of sources because the dominant proper motion for any source is the somewhat temporary one associated with internal structure changes that produce apparent motions that are an order of magnitude greater than the aberrational drift (e.g., Fey et al."
 1997)., 1997).
 Such proper motions are not constant with time and are about an order of magnitude higher than the 4—6 pas/yr proper motions expected from the secular aberration drift., Such proper motions are not constant with time and are about an order of magnitude higher than the 4–6 $\mu$ as/yr proper motions expected from the secular aberration drift.
 However. we believe that the secular aberration drift can be detected even with source position errors of ~ 100 µας and apparent motion formal errors of 10-100 jras/yr (e.g.. Feissel-Vernier 2003).," However, we believe that the secular aberration drift can be detected even with source position errors of $\sim$ 100 $\mu$ as and apparent motion formal errors of 10–100 $\mu$ as/yr (e.g., Feissel-Vernier 2003)."
 VLBI measurement of the secular aberration drift is important for both astrometry and astrophysics., VLBI measurement of the secular aberration drift is important for both astrometry and astrophysics.
 Indeed. it would give an independent estimate of the acceleration vector applied to the Solar System baryeenter without refering to objects within the Galaxy. so it can be used to constrain the mass of the central regions of the Milky Way.," Indeed, it would give an independent estimate of the acceleration vector applied to the Solar System barycenter without refering to objects within the Galaxy, so it can be used to constrain the mass of the central regions of the Milky Way."
 Moreover. the secular aberration drift produces a slow deformation of the celestial reference frame axes. and. if not corrected in geodetic VLBI sofware packages. could lead to contamination in estimates of other geodetic parameters. such as source coordinates and the Earth's nutation and precession (Titov 2010).," Moreover, the secular aberration drift produces a slow deformation of the celestial reference frame axes, and, if not corrected in geodetic VLBI sofware packages, could lead to contamination in estimates of other geodetic parameters, such as source coordinates and the Earth's nutation and precession (Titov 2010)."
 Gwinn et al. (, Gwinn et al. (
1997) have investigated the radio source velocity field obtained from the analysis of 16 years of VLBI data.,1997) have investigated the radio source velocity field obtained from the analysis of 16 years of VLBI data.
 Though the prime interest of their paper was to find an upper limit for the mass-energy of the cosmological gravitational wave background from the quadrupolar structure of the velocity field. their estimate of the Galactocentric acceleration is higher than expected by a factor of two with comparable standard errors.," Though the prime interest of their paper was to find an upper limit for the mass-energy of the cosmological gravitational wave background from the quadrupolar structure of the velocity field, their estimate of the Galactocentric acceleration is higher than expected by a factor of two with comparable standard errors."
 MacMillan (2005) processed another seven years of data but did not detect this effect., MacMillan (2005) processed another seven years of data but did not detect this effect.
 Using the OCCAM geodetic VLBI analysis software package. Titov (2009) reports a statistically significant dipole harmonic. but its magnitude was far from the theoretical value.," Using the OCCAM geodetic VLBI analysis software package, Titov (2009) reports a statistically significant dipole harmonic, but its magnitude was far from the theoretical value."
 The approach used by MaeMillan (2005) and Titov (2009) consisted of a direct estimation of the Galactocentric acceleration without intermediate estimation of the proper motions., The approach used by MacMillan (2005) and Titov (2009) consisted of a direct estimation of the Galactocentric acceleration without intermediate estimation of the proper motions.
 Our approach differs from these two., Our approach differs from these two.
 We used a three-step procedure that includes (1) a production of source coordinate time series from analysis of VLBI delays. (1) a least-square fit of proper motions to source coordinate time series. including editing large anomalous proper motions. and (iit) the fit of the Galactocentric acceleration to proper motions.," We used a three-step procedure that includes (i) a production of source coordinate time series from analysis of VLBI delays, (ii) a least-square fit of proper motions to source coordinate time series, including editing large anomalous proper motions, and (iii) the fit of the Galactocentric acceleration to proper motions."
 This three-step analysis helped to remove outliers that significantly biased the results of the previous analyses., This three-step analysis helped to remove outliers that significantly biased the results of the previous analyses.
lenses hence encode information of the cosmology. the galaxy mass profiles ancl (he evolution history of galaxies.,"lenses hence encode information of the cosmology, the galaxy mass profiles and the evolution history of galaxies."
 At present. the lens sample is too small to constrain all the ingredients simultaneously.," At present, the lens sample is too small to constrain all the ingredients simultaneously."
 Most previous studies concentrated on constraining the cosmological constant assumine non-evolving populations of lenses (e.g.. Fukugita et 11992: Ixochanek. 1996: ]lelbig et 11999).," Most previous studies concentrated on constraining the cosmological constant assuming non-evolving populations of lenses (e.g., Fukugita et 1992; Kochanek 1996; Helbig et 1999)."
 Under (his assumption. (the most recent lens statistics study of Chae et ((2002) finds that the lens statistics are best-fitted by a present-day matter density of Ong&0.3 and a cosmological constant of O44)20.7.," Under this assumption, the most recent lens statistics study of Chae et (2002) finds that the lens statistics are best-fitted by a present-day matter density of $\omega0 \approx 0.3$ and a cosmological constant of $\lambda0 \approx 0.7$."
 This result is consistent with results from a variety of other studies. including the cosmic microwave background radiation (e.g.. de Dernardis et 22000: Spergel et al.," This result is consistent with results from a variety of other studies, including the cosmic microwave background radiation (e.g., de Bernardis et 2000; Spergel et al."
 2003). Type Ia supernovae at cosmological distances (e.g.. Riess et al.," 2003), Type Ia supernovae at cosmological distances (e.g., Riess et al."
 1993: Perlmutter et al., 1998; Perlmutter et al.
 1999). and the large-scale structures in (he universe (e.g.. Peacock et al.," 1999), and the large-scale structures in the universe (e.g., Peacock et al."
 2001)., 2001).
 In light of the convergence of the cosmological model. it becomes important (o use eravitational lenses for a different purpose: to study the evolution of galaxies (their number density and cdenanmücal properüies) in the Quy=0.3.04450.7 cosmologv. (hereafter ACDAI).," In light of the convergence of the cosmological model, it becomes important to use gravitational lenses for a different purpose: to study the evolution of galaxies (their number density and dynamical properties) in the $\omega0=0.3, \lambda0=0.7$ cosmology (hereafter $\Lambda$ CDM)."
 Lensing limits on galaxw evolution have been explored by Mao (1991). Mao ]xochanek (1994). Rix et ((1994). and. Jain et ((2000). all of which used. lenses.," Lensing limits on galaxy evolution have been explored by Mao (1991), Mao Kochanek (1994), Rix et (1994), and Jain et (2000), all of which used optically-selected lenses."
 Lensing is sensitive to the evolution of galaxy properties as the lensine probability is xnol. while separations are xa7: here 5» is the number density and σ is the velocity dispersion of typical lenses.," Lensing is sensitive to the evolution of galaxy properties as the lensing probability is $\propto n\sigma^4$, while separations are $\propto \sigma^2$; here $n$ is the number density and $\sigma$ is the velocity dispersion of typical lenses."
 For example. (for a fixed. cosmological model) a decreasing number densitv of galaxies with redshift(z) lowers the lensing rate and the mean redshift of lenses while a decreasing velocity dispersion with 2 lowers the lensing rate. ihe mean lens redshift and the mean angular size of image separations.," For example, (for a fixed cosmological model) a decreasing number density of galaxies with $z$ ) lowers the lensing rate and the mean redshift of lenses while a decreasing velocity dispersion with $z$ lowers the lensing rate, the mean lens redshift and the mean angular size of image separations."
 Therelore. through a careful analvsis of the lens redshifts. image separations and lensing probabilitv we can constrain (he evolution of (he number density. aud dvnamical properties of galaxies.," Therefore, through a careful analysis of the lens redshifts, image separations and lensing probability we can constrain the evolution of the number density and dynamical properties of galaxies."
 Most lensing galaxies are massive earlv-0vpe galaxies as (μον dominate the lensing due to their larger central mass concentrations., Most lensing galaxies are massive early-type galaxies as they dominate the lensing cross-sections due to their larger central mass concentrations.
 Gravitational lenses therefore provide a unique mass-selected sample to study the evolution of earlv-tvpe galaxies. independent of and complementary (o (he traditional redshift survevs of galaxies (e.g.. Fried et al.," Gravitational lenses therefore provide a unique mass-selected sample to study the evolution of early-type galaxies, independent of and complementary to the traditional redshift surveys of galaxies (e.g., Fried et al."
 2001: Tin et al., 2001; Im et al.
 2002)., 2002).
 This is a much debated research area., This is a much debated research area.
 There exist (wo different. views on the formation and evolution of earlv-tvpe galaxies. namely a monolithic collapse model (Eegen. Lvnden-Dell. Sandage 1962) aud a merger hypothesis (Toomre Toomre 1972).," There exist two different views on the formation and evolution of early-type galaxies, namely a monolithic collapse model (Eggen, Lynden-Bell, Sandage 1962) and a merger hypothesis (Toomre Toomre 1972)."
 In the monolithic collapse model. early-(wpe galaxies are thought to have formed rapidly at hieh redshift and then evolve passively to the present-day.," In the monolithic collapse model, early-type galaxies are thought to have formed rapidly at high redshift and then evolve passively to the present-day."
 The merger model is a natural consequence of the hierarchical structure formation theory., The merger model is a natural consequence of the hierarchical structure formation theory.
 Semi-analvtic implementations of this model predict a continuous formation of ellipticals and hence a certain fraction of massive early-type galaxies must have formed since z~1 IxIxauffinann 1996: Baneh et 11996: IxaufBinann et 11999): the fraction depends on the assumed cosmology and other, Semi-analytic implementations of this model predict a continuous formation of ellipticals and hence a certain fraction of massive early-type galaxies must have formed since $z \sim 1$ Kauffmann 1996; Baugh et 1996; Kauffmann et 1999); the fraction depends on the assumed cosmology and other
for the fainter part. flatter than Hasingeretal.(1998).,"for the fainter part, flatter than \citet{has98}."
. The difference probably results from the fact that we are sampling different parts of the Log(N)-Log(S)., The difference probably results from the fact that we are sampling different parts of the Log(N)-Log(S).
 A slight slope flattening of the Log(N)-Log(S) is also observed in the 2-10 keV band. around fluxes of 2«107! ere em s7!. although the data are consistent with a single power-law with a cumulative slope index à=1.344))).," A slight slope flattening of the Log(N)-Log(S) is also observed in the 2-10 keV band, around fluxes of $2\times10^{-14}$ erg $^{-2}$ $^{-1}$, although the data are consistent with a single power-law with a cumulative slope index $\alpha=1.34^{+0.11}_{-0.10}$."
 A single power-law fit is tenable also for the 5-10 keV Log(N)-Log(S) and gives a slope a=1.54433. An extensive analysis of the X-ray broad-band properties of the sources and the optical follow-up of a hard X-ray selected sample will be the subjects of forthcoming papers (Baldi et al., A single power-law fit is tenable also for the 5-10 keV Log(N)-Log(S) and gives a slope $\alpha=1.54^{+0.25}_{-0.19}$ An extensive analysis of the X-ray broad-band properties of the sources and the optical follow-up of a hard X-ray selected sample will be the subjects of forthcoming papers (Baldi et al.
 in prep.)., in prep.).
 We thank A. De Luca for developing the hot pixel cleaning algorithm., We thank A. De Luca for developing the hot pixel cleaning algorithm.
 We are also grateful to G. Zamorani. G. C. Perola and all members of the team for useful discussions.," We are also grateful to G. Zamorani, G. C. Perola and all members of the team for useful discussions."
 We also thank the referee for useful suggestions which improved the presentation. of the results., We also thank the referee for useful suggestions which improved the presentation of the results.
 AB and SM acknowledge partial financial support by ASI L/R/190/00 contract., AB and SM acknowledge partial financial support by ASI I/R/190/00 contract.
"Following traditional practice. it is asstuned that the ME does not evolve with time. so that the umber density of stays ever created per unif mass and per unit time. C(NL£) (teriied the Creation Function by Miller Scalo 1979) can be separated into mass- and time-depeudent functions: (Miller&Scalo 1979).. where W(AD is the ME. b(f) is the birth rate (απο density of stars born per unit time}. aud Z;, is the age of the Galaxy, assumed lere to be LO Car.","Following traditional practice, it is assumed that the MF does not evolve with time, so that the number density of stars ever created per unit mass and per unit time, $C({\rm M},t)$ (termed the Creation Function by Miller Scalo 1979) can be separated into mass- and time-dependent functions: \citep{mil79}, , where $\Psi({\rm M})$ is the MF, $b(t)$ is the birth rate (number density of stars born per unit time), and $T_o$ is the age of the Galaxy, assumed here to be 10 Gyr."
 The separation of C(ALf£) enables the exanunation of the MIF aud birth rates separately.," The separation of $C({\rm M},t)$ enables the examination of the MF and birth rates separately."
 The metallicity CZ) and mass ratio (q) distributions are also assumed to be independent so that they may be treated separately as well., The metallicity $Z$ ) and mass ratio $q$ ) distributions are also assumed to be independent so that they may be treated separately as well.
 Clearly. these assuniptious av not accurately reflect the detailed stellar formation history of the Calasxy (ce... chemical evolution is ignored: see Edvardsson et 11993). but are adequate for current substellar ALF deteruiuations.," Clearly, these assumptions may not accurately reflect the detailed stellar formation history of the Galaxy (e.g., chemical evolution is ignored; see Edvardsson et 1993), but are adequate for current substellar MF determinations."
 Six MES were examined. includiug five power-law distributions: with a = 0.0. 0.5. 1.0. 1.5. and 2.0: aud the lognormal ME. from Chabrier(2001 ):: (sce also Miller Scalo 1979). where AL. = 0.1 AL. and o0 = 0.627.," Six MFs were examined, including five power-law distributions: with $\alpha$ = 0.0, 0.5, 1.0, 1.5, and 2.0; and the lognormal MF from \citet{cha01}: : (see also Miller Scalo 1979), where $_c$ = 0.1 $_{\sun}$ and $\sigma$ = 0.627."
 Note that V(log M) = Indo M (AT)., Note that $\Psi$ ${\rm \log{M}}$ ) = $\ln{10}$ M $\Psi$ (M).
 The baseline simulations iucorporate a lass rauge 0.01 < ALD < 0.1 ML... with the upper limit set by the evolutionary models and the lower limit set to provide enough objects in the higher mass bius. particularly for the steeper power Luvs.," The baseline simulations incorporate a mass range 0.01 $\leq$ M $\leq$ 0.1 $_{\sun}$, with the upper limit set by the evolutionary models and the lower limit set to provide enough objects in the higher mass bins, particularly for the steeper power laws."
" Lower mass liuits ranging from 0.001 to 0.015 AL. were also examined iu order to measure the iuflueuce of a nimium “cut-off” mass (M,,;,) in brown dwarf formation (6 1.2.1).", Lower mass limits ranging from 0.001 to 0.015 $_{\sun}$ were also examined in order to measure the influence of a minimum “cut-off” mass $_{min}$ ) in brown dwarf formation $\S$ 4.2.4).
" The age distribution is related to the birth by P(t)=Το,f).", The age distribution is related to the birth by $P(t) = b(T_o-t)$.
 The birth rate does not iuflueuce the shape of the MIF if the latter is assumed not to evolve., The birth rate does not influence the shape of the MF if the latter is assumed not to evolve.
 However. asbrown dwarfs themsclyes evolve thermally. the age distribution caninfluence the LE.," However, asbrown dwarfs themselves evolve thermally, the age distribution caninfluence the LF."
 Five birth rates were examined. as illustrated in Figure 1:," Five birth rates were examined, as illustrated in Figure 1:"
sample's selection function plays a large part in determining the behaviour of the derived. distribution. with samples selected in the UV and Hi probing very dillerent populations with very different underlving star formation behaviour and dust content (e.g. 2... 2)).,"sample's selection function plays a large part in determining the behaviour of the derived distribution, with samples selected in the UV and IR probing very different populations with very different underlying star formation behaviour and dust content (e.g. \citealt{2005ApJ...619L..51B}, , \citealt{2006ApJS..164...38I}) )."
 The behaviour of dust is also of ercal interest. as extinction. correlating stronely with star formation. plavs a major role in the derived star formation properties of a sample.," The behaviour of dust is also of great interest, as extinction, correlating strongly with star formation, plays a major role in the derived star formation properties of a sample."
" Aleasuring the luminosity [function and the SER distribution function has been an important goal for many vears. with much work being dedicated towards successfully estimating the local luminosity function (?.. 2.. 2.. 2)). the star formation rate distribution (e.g. ?.. ὃν, 7)). and the closely related total star formation rate volume density (?.. 7.. 2))."," Measuring the luminosity function and the SFR distribution function has been an important goal for many years, with much work being dedicated towards successfully estimating the local luminosity function \citealt{2003ApJ...587L..89T}, , \citealt{2005ApJ...619L..47S}, \citealt{2005ApJ...619L..15W}, \citealt{2010MNRAS.401...35W}) ), the star formation rate distribution (e.g. \citealt{2002NewA....7..337B}, \citealt{2005ApJ...619L..59M}, \citealt{2007ApJS..173..415M}) ), and the closely related total star formation rate volume density \citealt{1995ApJ...455L...1G}, \citealt{1998ApJ...495..691T}, \citealt{2004ApJ...615..209H}) )."
 These studies. often based on the large homogenous IAS and. GALEN survey datasets. have provided. good constraints at the bright. hiehly star forming end. of the distribution.," These studies, often based on the large homogenous IRAS and GALEX survey datasets, have provided good constraints at the bright, highly star forming end of the distribution."
 However. the behaviour of the faint. ened slope is still somewhat unexplored. with most studies sullering significant incompleteness at low luminosities ο... and SERs (2104M;ve Dh ," However, the behaviour of the faint end slope is still somewhat unexplored, with most studies suffering significant incompleteness at low luminosities $\lta10^{8-9} \;\mathrm{L}_{\sun}$ ) and SFRs $\lta10^{-1} \;\mathrm{M}_{\sun}\; \mathrm{yr}^{-1}$ )."
These less wellexploree populations at the faint end. being numerically dominant. orovide a vital insight into the hierarchical galaxy formation μα...," These less well-explored populations at the faint end, being numerically dominant, provide a vital insight into the hierarchical galaxy formation process."
" Lt is a well established result that the low-mass slope of he predicted. Dark Matter (DM) mass function is steeper han the observed faint-end slope of the luminosity function (LE): this ciserepaney is often framed as the well known ""missing satellite problem (2: 2)).", It is a well established result that the low-mass slope of the predicted Dark Matter (DM) mass function is steeper than the observed faint-end slope of the luminosity function (LF): this discrepancy is often framed as the well known `missing satellite problem' \citealt{1999ApJ...522...82K}; \citealt{Moore:1999aa}) ).
 Phe faint-end slope of the uminosityv function is a wavy of quantifving the number of aunt galaxies. and while models adopting a Press-Schechter ormatism predict a steep faint-end slope of a~LS (?:: 7)). observational results suggest a significantly shallower value (2:5: 2)).," The faint-end slope of the luminosity function is a way of quantifying the number of faint galaxies, and while models adopting a Press-Schechter formalism predict a steep faint-end slope of $\alpha \sim -1.8$ \citealt{1974ApJ...187..425P}; \citealt{2009MNRAS.398.1150B}) ), observational results suggest a significantly shallower value \citealt{1999ApJ...521...50M}; \citealt{2003MNRAS.344..307L}) )."
 Llowever. the form of the observed LE is several steps removed from the DM mass function. being the end result of many non-linear barvonic processes.," However, the form of the observed LF is several steps removed from the DM mass function, being the end result of many non-linear baryonic processes."
 A multitude of physical processes dictate how DM haloes collect and retain their eas. setting the form of the mass function which. through star formation. produces the luminosity function.," A multitude of physical processes dictate how DM haloes collect and retain their gas, setting the form of the mass function which, through star formation, produces the luminosity function."
" Measuring and comparing the distribution functions pertaining to these barvonie processes (i.e. the the hydrogen mass function. the star formation rate function. and the luminosity Function) has the potential to shed light on the possible reasons that barvons are under-abuncdant in low massclark matter haloes. causing the shallow slope and ""missing satellites."," Measuring and comparing the distribution functions pertaining to these baryonic processes (i.e. the the hydrogen mass function, the star formation rate function, and the luminosity function) has the potential to shed light on the possible reasons that baryons are under-abundant in low mass dark matter haloes, causing the shallow slope and `missing satellites'."
 To attempt to obtain good constraints on the extreme faint end. we use data from the Local Volume Legacy survey. a volume-limited dataset of galaxies within 11 Alpe with deep multi-band Spitzer. optical. and GALEN imaging.," To attempt to obtain good constraints on the extreme faint end, we use data from the Local Volume Legacy survey, a volume-limited dataset of galaxies within 11 Mpc with deep multi-band Spitzer, optical, and GALEX imaging."
 Being volume limited. this survey provides an excellent. ancl unbiased. database from which to calculate IR and UV luminosities ancl star formation rates. which can be combined with the same parameters from the larger datasets.," Being volume limited, this survey provides an excellent and unbiased database from which to calculate IR and UV luminosities and star formation rates, which can be combined with the same parameters from the larger datasets."
 For the first time. therefore. we are able to augment our large survey datasets with deep imaging of this statistically complete sample of local galaxies. and. extend the local luminosity ancl star formation rate distributions functions 1-2 orders of magnitude deeper.," For the first time, therefore, we are able to augment our large survey datasets with deep imaging of this statistically complete sample of local galaxies, and extend the local luminosity and star formation rate distributions functions 1-2 orders of magnitude deeper."
 In this paper. we present statistically complete SER distribution functions and. luminosity functions. which extend 1-2. orders of magnitude. deeper. than previous studies.," In this paper, we present statistically complete SFR distribution functions and luminosity functions which extend 1-2 orders of magnitude deeper than previous studies."
 The data provide the first ever SER. distribution study complete at both the bright and extreme faint. ends. and give an unbiased view of the distribution of star ormation and structure at z0.," The data provide the first ever SFR distribution study complete at both the bright and extreme faint ends, and give an unbiased view of the distribution of star formation and structure at $z\sim 0$."
 We use our samples o accurately constrain the true value - ancl breakdown - of the star formation rate volume density in the z=0 Universe., We use our samples to accurately constrain the true value - and breakdown - of the star formation rate volume density in the $z=0$ Universe.
" 1n order to investigate the role of extinction (which correlates strongly with SER). we also use HX and UV photometry to analyse the behaviour of dust obscured star formation (measured via the ratio Linων). and esent ""extinction distribution functions. analogous to the uminosity. function."," In order to investigate the role of extinction (which correlates strongly with SFR), we also use IR and UV photometry to analyse the behaviour of dust obscured star formation (measured via the ratio $\mathrm{L}_{\mathrm{IR}} / \mathrm{L}_{\mathrm{FUV}}$ ), and present `extinction distribution functions', analogous to the luminosity function."
 Throughout this paper we assume a concordance .\- CDAL cosmology with parameters (lly.OQ...Qa)=(72kms!Alpe 10.27. 0.73).," Throughout this paper we assume a concordance $\Lambda$ -CDM cosmology with parameters $_0, \; \Omega_m, \; \Omega_{\Lambda}) = (72 \;\mathrm{km s}^{-1}\; \mathrm{Mpc}^{-1}$ , 0.27, 0.73)."
" The Local Volume sample. or which peculiar motions overwhelm the Llubble How. has distances based on a number of direct distance estimators wherever possible (see 8:2,2,:2))."," The Local Volume sample, for which peculiar motions overwhelm the Hubble flow, has distances based on a number of direct distance estimators wherever possible (see )."
 As mentioned in £11 above. for any observationallv-defined sample the nature of the selection function. will eo a long wav to determining the physical properties of the sample.," As mentioned in 1 above, for any observationally-defined sample the nature of the selection function will go a long way to determining the physical properties of the sample."
 Any non-volume-limitecl sample usecl to calculate a luminosity function must. be strictly flux limited. at some wavelength. so that when weightings are applied. (for example volume weightings. as per the LfVins method) the resulting function. &(L). becomes statistically representative of the underlying population as a whole (an inconsistent Hux limit would. lead to bias through inconsistent volume COLLECLLONS)," Any non-volume-limited sample used to calculate a luminosity function must be strictly flux limited at some wavelength, so that when weightings are applied (for example volume weightings, as per the $1/\mathrm{V}_{\mathrm{max}}$ method) the resulting function, $\Phi(\mathrm{L})$, becomes statistically representative of the underlying population as a whole (an inconsistent flux limit would lead to bias through inconsistent volume corrections)."
 The wavelength at which this limit is imposed. will to a large extent determine the properties of the sample., The wavelength at which this limit is imposed will to a large extent determine the properties of the sample.
 An IR-sclectecl sample (selected at 607 for example) which is then crossed. with a UV database will contain more clusty. obscured systems. than an initially FUV-selected: sample which is subsequently cross matched with an LI database.," An IR-selected sample (selected at $60\mu$ m for example) which is then crossed with a UV database will contain more dusty, obscured systems than an initially FUV-selected sample which is subsequently cross matched with an IR database."
" ? demonstrated that the bolometric luminosity fanctions [or Ilt. ancl UV. selected: samples. divereccd for. Iuminous systems (Li,25«101 Le). with a significant. proportion of the massive. luminous svstems detected: strongly at 60j/m being non-detected in the EUM due to their extreme dust attenuation."," \cite{2007ApJS..173..404B} demonstrated that the bolometric luminosity functions for IR and UV selected samples diverged for luminous systems $\mathrm{L}_{\mathrm{bol}} > 5 \times 10^{10}\; \mathrm{L}_{\sun}$ ), with a significant proportion of the massive, luminous systems detected strongly at $\mu$ m being non-detected in the FUV due to their extreme dust attenuation."
 At the low-Iuminosity end. samples suller [rom the opposite bias: late-type. metal poor. star forming galaxies may lack sulliciont dust. to. re-process the UV photons into LR emission. and as a result can be weak- or non-cletections at 607m despite having measurable levels of ongoing star formation (?:: ?7)).," At the low-luminosity end, samples suffer from the opposite bias; late-type, metal poor, star forming galaxies may lack sufficient dust to re-process the UV photons into IR emission, and as a result can be weak- or non-detections at $\mu$ m despite having measurable levels of ongoing star formation \citealt{2009ApJ...703..517D}; \citealt{Calzetti:2010aa}) )."
 Clearly. defining a sample in terms of selection at a single wavelength: anc extrapolating global properties. wil lead to an inaccurate view of the behaviour of star formation and extinction in the overall population.," Clearly, defining a sample in terms of selection at a single wavelength and extrapolating global properties will lead to an inaccurate view of the behaviour of star formation and extinction in the overall population."
 In this paper. we therefore use separate datasets selected. at. 607. ane GALIZN FUY - hereafter. these will be referred to as the “LR- and “UV-sclectecl samples respectively.," In this paper, we therefore use separate datasets selected at $\mu$ m and GALEX FUV - hereafter, these will be referred to as the `IR-selected' and `UV-selected' samples respectively."
 Phese duasamples can then becombined. and the resultant. datase encapsulates the full range. of star. forming. behaviour. ensuring that both obscured and unobscured star formation are represented.," These dualsamples can then becombined, and the resultant dataset encapsulates the full range of star forming behaviour, ensuring that both obscured and unobscured star formation are represented."
We found that these results were generally. consistcnt with others in the literature (e.g.Spolaoretal. 2009:: Rawleetal. 2010)) once dillerent galaxy selection ellects were taken into account.,We found that these results were generally consistent with others in the literature (e.g.\citealt{Spolaor+09}; ; \citealt{Rawle+10}) ) once different galaxy selection effects were taken into account.
 The overall observational picture can be framed within the predictions. of. hvdrodynamical. and chenioclynamical simulations of ewaxy formation (Bekki&Shiova 1999:: lxawata 2001: Ἱναναία&Gibson2003: Ixobavashi 2O04:: llopkinsetal. 2009: Pipinoetal. 2010: Tortoraetal. 2011)) whieh include various physical. processes that can influence the gradients: supernovae and GN. feedback. ealaxy mergers. interactions with environment. etc.," The overall observational picture can be framed within the predictions of hydrodynamical and chemodynamical simulations of galaxy formation \citealt{BS99}; ; \citealt{Kawata01}; ; \citealt{KG03}; \citealt{Ko04}; \citealt{Hopkins+09a}; \citealt{Pipino+10}; \citealt{Tortora+11CGsim}) ) which include various physical processes that can influence the gradients: supernovae and AGN feedback, galaxy mergers, interactions with environment, etc."
 A fresh and qualitative overview of this physical scenario is given by Fig. L.," A fresh and qualitative overview of this physical scenario is given by Fig. \ref{fig:fig1},"
 where we show the isodensity contours for the colour-mass diagram (which we have already. discussed in T|10) and superimpose somoe lines (in the inset) and arrows which indicate the ellicieney of the highlighted physical processes., where we show the isodensity contours for the colour-mass diagram (which we have already discussed in T+10) and superimpose some lines (in the inset) and arrows which indicate the efficiency of the highlighted physical processes.
 In particular. the observations for high-mass IZECis seem to &enerallv suggest the inlluence of mergers and AGN feedback (Ixobavashi2004:: Hopkinsetal. 2009:: Sijackietal. 20073). with an efficiency. possibly increasing with mass (see red arrows and lines in the inset panel).," In particular, the observations for high-mass ETGs seem to generally suggest the influence of mergers and AGN feedback \citealt{Ko04}; \citealt{Hopkins+09a}; \citealt{Sijacki+07}) ), with an efficiency possibly increasing with mass (see red arrows and lines in the inset panel)."
 On the other hand. the population gradients of LCs and. less massive LGBA may be mainty driven by infall ancl supernovac [cecdhack (Larson1974.1975: Ixawata 2001: Ixawata&Gibson 2003:: Ixobavashi2004: Pipinoetal.2010: Tortoraetal.20113). which are more efficient at very low masses (see blue lines andarrows).," On the other hand, the population gradients of LTGs and less massive ETGs may be mainly driven by infall and supernovae feedback \citealt{Larson74, Larson75}; \citealt{Kawata01}; \citealt{KG03}; \citealt{Ko04}; \citealt{Pipino+10}; \citealt{Tortora+11CGsim}) ), which are more efficient at very low masses (see blue lines andarrows)."
 In this paper we consider the implications of the observed. colour. metallicity and age gradients from. T|10 [or variations of stellarralio AlfL with radius.," In this paper we consider the implications of the observed colour, metallicity and age gradients from T+10 for variations of stellar $M/L$ with radius."
 This issue is relevant for understanding the detailed distribution of mass within galaxies: both the stellar. mass and the dark matter. whose distribution is inferred by substracting the stellar mass from the total cynamical mass.," This issue is relevant for understanding the detailed distribution of mass within galaxies: both the stellar mass and the dark matter, whose distribution is inferred by substracting the stellar mass from the total dynamical mass."
" ""here has been scarcely any literature to date that focussed on this topic. except for a study in spiral galaxies by lortinari&Salucci(2010).."," There has been scarcely any literature to date that focussed on this topic, except for a study in spiral galaxies by \cite{PS10}."
 As a first approximation. one could. use single-colour semi-enmpirical models such as those of Bell&deJong(2001) and Bellctal.(2008.BOShereafter). to &o directly. [rom colour &radients to eeradients.," As a first approximation, one could use single-colour semi-empirical models such as those of \cite{BdJ01} and \citet[B03
hereafter]{Bell+03} to go directly from colour gradients to gradients."
 Llowever. we wish to make use of multiple colours. to allow for departures of individual galaxies from. mean trends. to explore the systematies thoroughly. and to trace the dependence on stellar parameters. like age and metallicity.," However, we wish to make use of multiple colours, to allow for departures of individual galaxies from mean trends, to explore the systematics thoroughly, and to trace the dependence on stellar parameters, like age and metallicity."
 This paper is organized as follows., This paper is organized as follows.
 The sample ancl the spectral analysis ave discussed in 82?.. in where we also show the ability of our fitting procedure to recover the correct eeracients. and we examine possible svstematies.," The sample and the spectral analysis are discussed in \ref{sec:sample}, , in where we also show the ability of our fitting procedure to recover the correct gradients, and we examine possible systematics."
 Ehe results ave shown in $??.. while §?? is devoted to a discussion of the physical implications and. conclusions.," The results are shown in \ref{sec:results}, while \ref{sec:conclusions} is devoted to a discussion of the physical implications and conclusions."
 Our database consists of50000 Low redshift. (0.0033= 0.05) ealaxies in the NYU Value-Acdded Galaxy. Catalog extracted from SDSS DIA (Blantonetal.2005.hereafter D05)).. recently analvzed in|107.," Our database consists of$50\,000$ low redshift $0.0033 \leq z
\leq 0.05$ ) galaxies in the NYU Value-Added Galaxy Catalog extracted from SDSS DR4 \citealt[hereafter
B05]{Blanton05}) , recently analyzed in."
. Following T|10. we have identified 1Γς as those svstems with a r-band Sérsic index satisfving the condition 2.5x:px5.5 and with a concentration index C'72.6.," Following T+10, we have identified ETGs as those systems with a r-band $\rm
\acute{e}$ rsic index satisfying the condition $2.5 \leq n \leq
5.5$ and with a concentration index $C > 2.6$."
 Ehe final EPC sample consists of 10508 galaxies.," The final ETG sample consists of $10\,508$ galaxies."
 Of the remaining entries of our database. ~27800 are ΓΣ. defined as objects with Cost2.6. ηx2.5. and oy5150km/s.," Of the remaining entries of our database, $\sim 27\,800$ are LTGs, defined as objects with $C
\leq 2.6$, $n \leq 2.5$, and $\sigc\leq 150~\rm km/s$."
 We have used the multi-band: structural parameters eiven by BOS to derive the colour profile (CY336GR) of each ealaxy as the dillerences between the (logarithmic) surface brightness measurements in the two bands. NV and Y.," We have used the multi-band structural parameters given by B05 to derive the colour profile $(X-Y)(R)$ of each galaxy as the differences between the (logarithmic) surface brightness measurements in the two bands, $X$ and $Y$."
 We have defined the colour gradient as the angular coellicient of the relation VYY vs logFf Rar. Vxy=Ad)log(RiBa) measured in mag + (omitted in the Following). where lis the r-band. ellective radius.," We have defined the colour gradient as the angular coefficient of the relation $X-Y$ vs $\log R/\Re$ , $\displaystyle \nabla_{X-Y} = \frac{\delta
(X-Y)}{\delta \log (R/R_{\rm eff})}$ , measured in mag $^{-1}$ (omitted in the following), where is the r-band effective radius."
 By definition. a positive epacdient. Vxy2O. means that a galaxy isredder as £& increases. while it is bluer outward for a negative eraclicnt.," By definition, a positive gradient, $\nabla_{X-Y}>0$, means that a galaxy isredder as $R$ increases, while it is bluer outward for a negative gradient."
 Vhe fit of svnthetiecolours isperformed on the colours at, The fit of syntheticcolours isperformed on the colours at
The surface reactions used in the gas-grain chemical network normally involve the diffusion of one or more reaction partners across surface sites via thermal hopping. a process known as the Langmui-Hinshelwood mechanism.,"The surface reactions used in the gas-grain chemical network normally involve the diffusion of one or more reaction partners across surface sites via thermal hopping, a process known as the Langmuir-Hinshelwood mechanism."
 At low temperatures. H atoms are a most important surface reactant. and are responsible for the hydrogenation of heavy atoms landing on grains into their saturated forms (e.g. O into HO. C into CH).," At low temperatures, H atoms are a most important surface reactant, and are responsible for the hydrogenation of heavy atoms landing on grains into their saturated forms (e.g. O into $_{2}$ O, C into $_{4}$ )."
 Of the major ice species in our model (H*O. CO. 0Ο). CO: has proven to be the most difficult to produce on surfaces via such processes (e.g.Ruffle&Herbst2001)..," Of the major ice species in our model $_2$ O, CO, $_{2}$ ), $_{2}$ has proven to be the most difficult to produce on surfaces via such processes \citep[e.g.][]{RH01III}."
 We have adopted an activation energy E4 = 290 K. as measured by Roseretal.(2001).. for the important CO2(s) formation reaction between CO(s) and O(s) (Nummelinetal.2001).," We have adopted an activation energy $E_{\rm A}$ = 290 K, as measured by \citet{RVMP01}, for the important $_{2}$ (s) formation reaction between CO(s) and O(s) \citep{Nummelin01}."
. Following the approach of Ruffle&Herbst (2001).. we also consider a decreased activation energy E4 of 130 K (see.forTielens&Hagen 1982)..," Following the approach of \citet{RH01III}, , we also consider a decreased activation energy $E_{\rm A}$ of 130 K \citep[see, for example,][]{gdh86, FEA79, TH82}. ."
 Gas-phase species can also react directly with surface species via the Eley-Rideal mechanism (Ruffle&Herbst 2001)., Gas-phase species can also react directly with surface species via the Eley-Rideal mechanism \citep{RH01III}.
. To be efficient. such a mechanism requires a reactant with reasonably high surface coverage. and CO(s) meets this criterion once CO(g) starts to acerete.," To be efficient, such a mechanism requires a reactant with reasonably high surface coverage, and CO(s) meets this criterion once CO(g) starts to accrete."
 To possibly enhance the rate of CO» formation. the Eley-Rideal reaction O(g)+CO(s)- was added to the network. with a rate coefficient kup given by where Kyeeo is the aceretion rate (s7!) of O(g) onto the grain surface. Geo.) is the average fractional surface coverage of CO(s). and the Boltzmann factor depends on a reaction barrier. Eypep.," To possibly enhance the rate of $_{2}$ formation, the Eley-Rideal reaction $\rm{O(g)}+\rm{CO(s)} \rightarrow \rm{CO_{2}(s)}$ was added to the network, with a rate coefficient $k_{\rm ER}$ given by where $k_{\rm acc,O}$ is the accretion rate $^{-1}$ ) of O(g) onto the grain surface, $\theta_{\rm CO(s)}$ is the average fractional surface coverage of CO(s), and the Boltzmann factor depends on a reaction barrier, $E_{\rm A, E-R}$."
 A value of Eypep= OK is assumed in order to investigate the full capability of this reaction.," A value of $E_{\rm A,E-R}= 0$ K is assumed in order to investigate the full capability of this reaction."
 To consider a reasonable portion of parameter space. we have run a considerable number of models.," To consider a reasonable portion of parameter space, we have run a considerable number of models."
 The ones chosen for discussion comprise all four shock models listed in Table I. each with two different evolutions of granular temperature and two different barriers for the diffusive O + CO reaction.," The ones chosen for discussion comprise all four shock models listed in Table \ref{tbl-shockmods}, each with two different evolutions of granular temperature and two different barriers for the diffusive O $+$ CO reaction."
 Models that contain 7450=15 K and Ey = 290 K are labeled the A set. while models that contain Tyg=20 K and E4 = 130K are labeled the B set.," Models that contain $T_{\rm d,0} = 15$ K and $E_{\rm A}$ = 290 K are labeled the A set, while models that contain $T_{\rm d,0} = 20$ K and $E_{\rm A}$ = 130 K are labeled the B set."
 With the shock models 1-4. we label the chemical models I-A. I-B. ete.," With the shock models 1-4, we label the chemical models 1-A, 1-B, etc."
 Models with other sets of varied parameters were also considered. but our choices for A and B parameters span the range of variability of results.," Models with other sets of varied parameters were also considered, but our choices for A and B parameters span the range of variability of results."
 We report our principal results as column densities. rather than fractional abundances to facilitate comparison with the available observations (see next section)., We report our principal results as column densities rather than fractional abundances to facilitate comparison with the available observations (see next section).
 The column densities are obtained from the calculated fractional abundances at a given visual extinetion by using a relation between the visual extinction and the total column density for hydrogen. Ny=1.9κ107Ac em? mag?! (Whittet2003).. although there is some variation in the adopted value of this conversion factor in the literature.," The column densities are obtained from the calculated fractional abundances at a given visual extinction by using a relation between the visual extinction and the total column density for hydrogen, $N_{\rm H}=1.9\times10^{21} A_V$ $^{-2}$ $^{-1}$ \citep{dustbook03}, although there is some variation in the adopted value of this conversion factor in the literature."
" As in BO4. the gaseous spectes H» and CO form after an enhanced density ng, at low temperatures even though the extinction is very low."," As in B04, the gaseous species $_{2}$ and CO form after an enhanced density $n_{\rm H,disc}$ at low temperatures even though the extinction is very low."
 A variety of other species start to form. both in the ice phase and in the gas.," A variety of other species start to form, both in the ice phase and in the gas."
" After H» and CO. the most abundant gas-phase species by Ay=2 include No. CN. H?O. and CH,. although the last three possess columns of =10! em? or less. nearly three orders of magnitude smaller than the first three."," After $_{2}$ and CO, the most abundant gas-phase species by $A_V=2$ include $_{2}$, CN, $_{2}$ O, and $_4$, although the last three possess columns of $\approx 10^{14}$ $^{-2}$ or less, nearly three orders of magnitude smaller than the first three."
 The high-metal elemental abundances maintain a high electron fractional abundance exceeding 107°. which diminishes the power of the ton-molecule chemistry to produce larger molecules.," The high-metal elemental abundances maintain a high electron fractional abundance exceeding $^{-6}$, which diminishes the power of the ion-molecule chemistry to produce larger molecules."
 There are no molecular ions among the top 50 species of any model considered at this time. nor does the largest abundance of a molecular ton come within six orders of magnitude of the CO(g) abundance.," There are no molecular ions among the top 50 species of any model considered at this time, nor does the largest abundance of a molecular ion come within six orders of magnitude of the CO(g) abundance."
 Our principal results are contained in two figures. 1 which we plot assorted column densities vs Ay and N[H;]. respectively: the latter is more commonly used in studies of CO(g).," Our principal results are contained in two figures, in which we plot assorted column densities vs $A_V$ and $N[{\rm H_{2}}]$, respectively; the latter is more commonly used in studies of CO(g)."
 The column densities of CO(g). N[CO(g)]. and ar assortment of ices are plotted against Ay in Figure 2. for the A and B parameter sets and all four shock models.," The column densities of CO(g), $N$ [CO(g)], and an assortment of ices are plotted against $A_V$ in Figure \ref{fig-2NH} for the A and B parameter sets and all four shock models."
 The columi densities are shown for visual extinction through Ay=3.3. as further extensions of the evolution are definitely beyond the scope of the shock hydrodynamies (see discussion in BO4 regarding the importance of gravity).," The column densities are shown for visual extinction through $A_{V} = 3.3$, as further extensions of the evolution are definitely beyond the scope of the shock hydrodynamics (see discussion in B04 regarding the importance of gravity)."
" The extinction. used i Figure 2. is considered to be ""edge-to-center"" for a line of sight into dense or translucent material.", The extinction used in Figure \ref{fig-2NH} is considered to be “edge-to-center” for a line of sight into dense or translucent material.
 Thus an observed extinetior of 2 mag would correspond to 1 mag in the edge-to-center model frame., Thus an observed extinction of 2 mag would correspond to 1 mag in the edge-to-center model frame.
 The column density for CO(g) is plotted against N[H>] in Figure 3. along with column densities for the other major forms of carbon in the gas - C and C. In both figures. solid lines represent models with the A parameter set. while dashed lines representthe Bset.," The column density for CO(g) is plotted against $N$ $_2$ ] in Figure \ref{fig-3CO} along with column densities for the other major forms of carbon in the gas - $^+$ and C. In both figures, solid lines represent models with the A parameter set, while dashed lines representthe Bset."
and will have no meaning in terms of the probability that the model corresponds to what is observed.,and will have no meaning in terms of the probability that the model corresponds to what is observed.
 We will however refer to the obtained values of the reduced X? (cf., We will however refer to the obtained values of the reduced $\chi^2$ (cf.
 Sect. 3.3)).," Sect. \ref{results1}) ),"
 because it is the usual way to express the quality of a fit., because it is the usual way to express the quality of a fit.
" At a given frequency v. we model the light curve 5,(f) of a single outburst of amplitude -1(77). starting at time t=fy and peaking att=ty|fu(i) by if fySf|tas) and by iff>ty|ἐπ)."," At a given frequency $\nu$, we model the light curve $S_{\nu}(t)$ of a single outburst of amplitude $A(\nu)$, starting at time $t\!=\!t_0$ and peaking at $t=t_0+t\dmrm{rise}(\nu)$ by if $t_0\leq t < t_0 + t\dmrm{rise}(\nu)$ and by if $t\geq t_0 + t\dmrm{rise}(\nu)$."
 The exponents p(v) and ο] define the shape of the light curve at frequeney 7 and te.y(77) Is the o-folding decay time of the flareat frequency v., The exponents $\rho(\nu)$ and $\phi(\nu)$ define the shape of the light curve at frequency $\nu$ and $t\dmrm{fall}(\nu)$ is the $\mrm{e}$ -folding decay time of the flareat frequency $\nu$.
 Different time profiles of an outburst defined byEqs. (1)), Different time profiles of an outburst defined byEqs. \ref{rise}) )
 and (2)) are shown in Fig. 1., and \ref{fall}) ) are shown in Fig. \ref{func}.
 Rather than constraining the outburst parameters (A(7/). fasQ0). μι). pv) and ο) at each of the twelve light curve's frequencies. we describe their logarithm by a cubic spline which we parameterize at only four frequencies spaced by 0.75ddex and covering the GGHz range (see Fig. 3)).," Rather than constraining the outburst parameters $A(\nu)$, $t\dmrm{rise}(\nu)$, $t\dmrm{fall}(\nu)$, $\rho(\nu)$ and $\phi(\nu)$ ) at each of the twelve light curve's frequencies, we describe their logarithm by a cubic spline which we parameterize at only four frequencies spaced by dex and covering the GHz range (see Fig. \ref{param}) )."
 This reduces the number of free parameters by a factor three. while keeping the parameterization completely model-independent.," This reduces the number of free parameters by a factor three, while keeping the parameterization completely model-independent."
 We thus need a total of 5«| parameters to fully characterize the spectral and temporal evolution of an outburst. 1.e. à surface in the three dimensional (55.7.t)-space (cf.," We thus need a total of $5\!\times\!4$ parameters to fully characterize the spectral and temporal evolution of an outburst, i.e. a surface in the three dimensional $(S,\nu,t)$ -space (cf."
 Fig. 4))., Fig. \ref{evo1}) ).
 We impose that all individual outbursts are self-similar. in the sense that they all have the same evolution pattern. ie. the same shape of the surface in the (55.ij.f)-space.," We impose that all individual outbursts are self-similar, in the sense that they all have the same evolution pattern, i.e. the same shape of the surface in the $(S,\nu,t)$ -space."
 What we allow to change from one outburst to the other is the normalization in flux 5. frequency v and time f. which changes. respectively. the amplitude of the outburst (strong or weak). the frequency at which the emission. peaks (high- or frequency peaking) and the time scale of the evolution (long-lived or short-lived).," What we allow to change from one outburst to the other is the normalization in flux $S$, frequency $\nu$ and time $t$, which changes, respectively, the amplitude of the outburst (strong or weak), the frequency at which the emission peaks (high- or low-frequency peaking) and the time scale of the evolution (long-lived or short-lived)."
 A change in normalization corresponds to a shift of the position of the outburst's characteristic surface in the (logS.logv.f)-space.," A change in normalization corresponds to a shift of the position of the outburst's characteristic surface in the $(\log{S},\log{\nu},\log{t})$ -space."
 To define this position. we take the point of maximum flux as an arbitrary reference point on the surface.," To define this position, we take the point of maximum flux as an arbitrary reference point on the surface."
 On average among all individual outbursts. this point 1s located at (log55.ij.fj) and this average normalization defines what we call the of 3C 273.," On average among all individual outbursts, this point is located at $(\langle\log{S}\rangle,\langle\log{\nu}\rangle,\langle\log{t}\rangle)$ and this average normalization defines what we call the of 3C 273."
" We denote by AlogS. Alogr» and Alocsf the logarithmic shifts of this point with respect to the average position, i.e. Alogk=loghlogkj.¥k—S.v.f."," We denote by $\Delta\log{S}$, $\Delta\log{\nu}$ and $\Delta\log{t}$ the logarithmic shifts of this point with respect to the average position, i.e. $\Delta\log{k}\!=\!\log{k}-\langle\log{k}\rangle,\,\forall\,k\!=\!S,\,\nu,\,t$."
 These 123 logarithmic shifts plus the 12 different start times fy of the flares give a total of 15 parameters used to define the specificity of all outbursts., These $12\!\times\!3$ logarithmic shifts plus the $12$ different start times $t_0$ of the flares give a total of $48$ parameters used to define the specificity of all outbursts.
 The superimposed decays of the outbursts that started before 1979 are simply modelled by an hypothetical event of amplitude AyQ7} at time £.=ty|fined)1979.0 and decaying with the c-folding time £g(7) of the typical outburst at frequency 1’., The superimposed decays of the outbursts that started before 1979 are simply modelled by an hypothetical event of amplitude $A_0(\nu)$ at time $t=t_0+t\dmrm{rise}(\nu)=1979.0$ and decaying with the $\mrm{e}$ -folding time $t\dmrm{fall}(\nu)$ of the typical outburst at frequency $\nu$.
 The variation of the amplitude ;19(7/) with frequency is modelled by a cubic spline as for the five other variables. but parameterized at four slightly lower frequencies (log(v/GHz)= 0.5. 1.0. 1.5 and 2.0). due to the fact that AjQ7) is only well constrained for the radio light curves.," The variation of the amplitude $A_0(\nu)$ with frequency is modelled by a cubic spline as for the five other variables, but parameterized at four slightly lower frequencies $\log{(\nu/\mbox{GHz})}=$ 0.5, 1.0, 1.5 and 2.0), due to the fact that $A_0(\nu)$ is only well constrained for the radio light curves."
" Finally. we assume a constant contribution to the light curves due to the quiescent emission of the jet’s hot spot 3C 273A. This emission ts modelled with a power law spectrum as given in Türrler et al. (1999a)),"," Finally, we assume a constant contribution to the light curves due to the quiescent emission of the jet's hot spot 3C 273A. This emission is modelled with a power law spectrum as given in Türrler et al. \cite{TPC99}) )."
 To summarize. this first parameterization uses a total of 72 (20|Ls D) parameters to adjust the 4352 observational points in the twelve light curves.," To summarize, this first parameterization uses a total of $72$ $20+48+4$ ) parameters to adjust the 4352 observational points in the twelve light curves."
 The great number of free parameters still leaves more than four thousand degrees of freedom (d.o.f.), The great number of free parameters still leaves more than four thousand degrees of freedom (d.o.f.)
 to the fit., to the fit.
 The simultaneous fitting of the twelve light curves is performed by many iterative fits of small subsets of the 72 parameters., The simultaneous fitting of the twelve light curves is performed by many iterative fits of small subsets of the $72$ parameters.
 Fig., Fig.
 2 shows three representative light. curves among the twelve fitted simultaneously with the outbursts parameterized as described in Sect. 3.2.., \ref{fit1} shows three representative light curves among the twelve fitted simultaneously with the outbursts parameterized as described in Sect. \ref{model1}. .
 The major features of the light curves are reproduced by the model with only about one outburst every 1.5 year starting simultaneously at all frequencies., The major features of the light curves are reproduced by the model with only about one outburst every 1.5 year starting simultaneously at all frequencies.
 The overall fit has a reduced 4? value of 42.—4?/d.o.f. 16.1., The overall fit has a reduced $\chi^2$ value of $\chi\dmrm{red}^2\!\equiv\!\chi^2/\mbox{d.o.f.}\!=\!16.1$ .
 The main discrepancy between the model and the observations arises, The main discrepancy between the model and the observations arises
in (he Abell catalog.,in the Abell catalog.
 We suggest that the mean non-thermal luminosity associated with rich clusters is L4.o1xI0! eres | (30eV-1TeV) in the local Universe., We suggest that the mean non-thermal luminosity associated with rich clusters is $\bar{L_{\gamma c}}\sim 1\times 10^{44}$ erg s $^{-1}$ (30eV-1TeV) in the local Universe.
 Recent theoretical predictions and modeling have suggested that. of the EGRB arises from upscattering of CAIB photons by relativistic electrons in the large-scale gravitational shocks of the IGM. including elusters.," Recent theoretical predictions and modeling have suggested that of the EGRB arises from upscattering of CMB photons by relativistic electrons in the large-scale gravitational shocks of the IGM, including clusters."
 If such models are valid our results suggest. that e]—10% of this component of the EGRB originates [rom structures associated with rich clusters wilh 0<z«I., If such models are valid our results suggest that $\sim 1-10$ of this component of the EGRB originates from structures associated with rich clusters with $0<z<1$.
 Specilic predictions for the population of clusters alone have estimated a contribution to the EGRD of 0.5—2*4... dominated by svstems with z<0.2 (Colafrancesco&Blasi1993).," Specific predictions for the population of clusters alone have estimated a contribution to the EGRB of $0.5-2$ , dominated by systems with $z\leq0.2$ \citep{col98}."
. several caveats are worth noting., Several caveats are worth noting.
 First. diffuse Galactic gamma-ray emission is very hard (0 remove. even with the current best model for its distribution.," First, diffuse Galactic gamma-ray emission is very hard to remove, even with the current best model for its distribution."
 We believe we have accounted for the residual effects of this emission by assessing significance via our Monte Carlo simulations. but it may still introduce svstematic errors to our results.," We believe we have accounted for the residual effects of this emission by assessing significance via our Monte Carlo simulations, but it may still introduce systematic errors to our results."
 Second. we cannot. wilh these data. prove that the excess emission associated wilh clusters is not due to some population of discrete sources. themselves positively correlated with clusters. although there is no a priori reason (o believe this to be the case.," Second, we cannot, with these data, prove that the excess emission associated with clusters is not due to some population of discrete sources, themselves positively correlated with clusters, although there is no a priori reason to believe this to be the case."
 Third. estimating (he cluster luminosity by measuring the intensitv in the innermost 1° bin is likely to be an underestimate. given both the width of the EGRET PSF and the evidence for more extended. correlated. enussion.," Third, estimating the cluster luminosity by measuring the intensity in the innermost $^{\circ}$ bin is likely to be an underestimate, given both the width of the EGRET PSF and the evidence for more extended, correlated, emission."
 Finally. several potential sources of random error exist serving {ο increase (he uncertainty in quoted numbers - which. therefore. should be considered as broad estimates.," Finally, several potential sources of random error exist serving to increase the uncertainty in quoted numbers - which, therefore, should be considered as broad estimates."
 It should also be noted that another possible mechanism [or producing non-thermal emission in clusters is svuehrotron emission in radio galaxies contained in the clusters. that produce relativistic electrons which then up-scatter to produce (he high energv ganuna-ravs (e.g. Giovanninietal.(1993):Petrosian (2001))).," It should also be noted that another possible mechanism for producing non-thermal emission in clusters is synchrotron emission in radio galaxies contained in the clusters, that produce relativistic electrons which then up-scatter to produce the high energy gamma-rays (e.g. \citet{gio93,pet01}) )."
 In this case the particles are accelerated in the nuclei of radio galaxies. and this could account lor the observed gamnma-ray Inminosities.," In this case the particles are accelerated in the nucleii of radio galaxies, and this could account for the observed gamma-ray luminosities."
 In (his case. one would expect to see a stronger correlation signal with rich clusters containing radio galaxies.," In this case, one would expect to see a stronger correlation signal with rich clusters containing radio galaxies."
 Ledlow&Owen(1995) find that some of &=2 Abell clusters have evidence of radio emitting member galaxies (in a statistical subsample of 393 clusters of all richnesses)., \citet{led95} find that some of $R=2$ Abell clusters have evidence of radio emitting member galaxies (in a statistical subsample of 393 clusters of all richnesses).
 However. more sophisticated analvsis is required owing to the small munber of racio surveved clusters and we will investigate (his in an extension to the present work.," However, more sophisticated analysis is required owing to the small number of radio surveyed clusters and we will investigate this in an extension to the present work."
 As discussed by Colafrancesco&Blasi(1993).. Loeb&Waxman (2000).. WaxmanLoeb (2000).. and lxeshetetal.(2002).. the appearance of non-thermal emission in ICM/cluster structures is transient. with a timescale of ~ 10?vrs.," As discussed by \citet{col98}, \citet{loe00}, \citet{wax00}, and \citet{kes02}, the appearance of non-thermal emission in IGM/cluster structures is transient, with a timescale of $\sim 10^{9}$ yrs."
 Therefore the emission disappears almost entirely as soon as there is no strong shocking of gas., Therefore the emission disappears almost entirely as soon as there is no strong shocking of gas.
 If. we consider our sample of rich. Abell clusters then we know from other wavelengths that the majority of these low-z systems are probablyin a state closely approximating hyedrostatie equilibrium, If we consider our sample of rich Abell clusters then we know from other wavelengths that the majority of these $z$ systems are probablyin a state closely approximating hydrostatic equilibrium
The existence of non-axvsimnnetrie conrponeuts of the ealactic potential has been frequeutlv imwoked as an cficient way το transport eas from the ealaxy scale down to the nucleus to fue the ACN.,The existence of non-axysimmetric components of the galactic potential has been frequently invoked as an efficient way to transport gas from the galaxy scale down to the nucleus to fuel the AGN.
 In particular. the shocks and eravitational torques induced by a galactic bar can make the C»eas loose aneularOo momentu auk therefore facilitate the fuelling mechanisi (Shlosman et al.," In particular, the shocks and gravitational torques induced by a galactic bar can make the gas loose angular momentum and therefore facilitate the fuelling mechanism (Shlosman et al."
 L989)., 1989).
 Nevertheless. it has been sugeestedOO that there is not a preference for Sevtert nuclei to occur in barre ealaxies (ITeckiuian 1980: Simi et al.," Nevertheless, it has been suggested that there is not a preference for Seyfert nuclei to occur in barred galaxies (Heckman 1980; Simkim et al."
 1980)., 1980).
 This result is confirmed by statistical analyses on the catalogue morphology of Sevfert ealaxies (Moles et al., This result is confirmed by statistical analyses on the catalogued morphology of Seyfert galaxies (Moles et al.
 1995) iux through near-intrared inagiug studdies (MeLeod Ricke 1995: Alonso-Ierrero et al., 1995) and through near-infrared imaging dies (McLeod Rieke 1995; Alonso-Herrero et al.
 1996: Mulchaev Reean 1997)., 1996; Mulchaey Regan 1997).
 We lave started a large observational program aimed at studying the coucditious for the onset of activity iu ealactic nuclei., We have started a large observational program aimed at studying the conditions for the onset of activity in galactic nuclei.
 Based on previous work (Moles et al., Based on previous work (Moles et al.
 1995). we search for detailed morphological aud kinematical differences between active aud non active galaxies of simular global morphology.," 1995), we search for detailed morphological and kinematical differences between active and non active galaxies of similar global morphology."
 We particularly pay attention to those that could facilita| tlie transport of gas towards the very central regions and he nucleus (for an example of the complete analysis of one of our sample galaxies see Pérez ct al., We particularly pay attention to those that could facilitate the transport of gas towards the very central regions and the nucleus (for an example of the complete analysis of one of our sample galaxies see Pérrez et al.
 1999)., 1999).
 For this purpose. we are obtaining optical aud near-infrared images aud long slit spectroscopy WwoH.h the best possible spatial resolution.," For this purpose, we are obtaining optical and near-infrared images and long slit spectroscopy with the best possible spatial resolution."
 We selected a sample of ealaxics. out of which LS host active nuclei and the rest formu: a control sample of nou active galaxies (sce Sect. 1)).," We selected a sample of galaxies, out of which 18 host active nuclei and the rest form a control sample of non active galaxies (see Sect. \ref{sample}) )."
The overall redshift completeness of the PLASLL survey is O.64. with redshifts for of the galaxies in the catalogue.,"The overall redshift completeness of the FLASH survey is 0.64, with redshifts for of the galaxies in the catalogue."
 We now consider the dependence of this completeness on observing conditions (field). on apparent magnitude. anc on redshift itself.," We now consider the dependence of this completeness on observing conditions (field), on apparent magnitude, and on redshift itself."
 Because we observe galaxies with PLALR one Schmidt plate at a time. there is a field-to-field. variation in completeness that depends on the quality of the observing conditions for each field.," Because we observe galaxies with FLAIR one Schmidt plate at a time, there is a field-to-field variation in completeness that depends on the quality of the observing conditions for each field."
 Figure 6 shows the Schmidt fields that tile the FLASL survey region. with the field number shown in the top left and the field redshift completeness centred on each Ποια.," Figure \ref{fig:fieldcomp}
 shows the Schmidt fields that tile the FLASH survey region, with the field number shown in the top left and the field redshift completeness centred on each field."
 The redshift’ completeness across the survey region 1s reasonably uniform given the variation in galaxv number density across the skv in the region (seeFigure 12)., The redshift completeness across the survey region is reasonably uniform given the variation in galaxy number density across the sky in the region (seeFigure \ref{fig:hcc}) ).
 The median completeness over all 42 fields is 0.63., The median completeness over all 42 fields is 0.63.
 The reason his figure is lower than the overall completeness is that he lower-completeness fields tend to be near the edee of he survey region. where the number of galaxies within he survey is too small to be able to use the FLAIR system ellicientlv. so that galaxy redshifts in these fields are edominantly Crom the literature.," The reason this figure is lower than the overall completeness is that the lower-completeness fields tend to be near the edge of the survey region, where the number of galaxies within the survey is too small to be able to use the FLAIR system efficiently, so that galaxy redshifts in these fields are predominantly from the literature."
 ‘Table 1. lists the statistics of each Schmidt field. in he survey region., Table \ref{tab:fields} lists the statistics of each Schmidt field in the survey region.
 For each field the table contains the ollowing information: field. LD number. IX. and. oof he field centre (D1950). number of target galaxies. number of redshifts obtained. and field completeness.," For each field the table contains the following information: field ID number, R.A. and of the field centre (B1950), number of target galaxies, number of redshifts obtained, and field completeness."
 The signal-to-noise ratio in the observed. spectra decreases as the target objects’ apparent magnitudes become fainter., The signal-to-noise ratio in the observed spectra decreases as the target objects' apparent magnitudes become fainter.
 This results in a decreasing redshift completeness at. fainter apparent magnitudes which. if not corrected for. will result in an incorrect selection function ancl skew the results obtained from magnitude-dependent analyses tthe luminosity function).," This results in a decreasing redshift completeness at fainter apparent magnitudes which, if not corrected for, will result in an incorrect selection function and skew the results obtained from magnitude-dependent analyses the luminosity function)."
 The completeness as a function of apparent magnitucle for the whole survey sample is shown in Figure 4.. and for," The completeness as a function of apparent magnitude for the whole survey sample is shown in Figure \ref{fig:mag_comp_all}, , and for"
"However, when the central object is a rotating black hole. the situation will be quite different.","However, when the central object is a rotating black hole, the situation will be quite different."
 As before. when By grows to a fraction of Buy. a significant MIID outflow will develop. exerting a braking torque on the inner portion of the disk.," As before, when $B_{p0}$ grows to a fraction of $B_{d0}$, a significant MHD outflow will develop, exerting a braking torque on the inner portion of the disk."
 Because of the cdiffercutial fine drageine. however. even ifali the IKepleriau augular momentum were removed by the outflowing MIID power. the disk aaterial still will be sheareditself. coutinuing to increase the local magnetic field streugth and contiuung to produce MIID power.," Because of the differential frame dragging, however, even if the Keplerian angular momentum were removed by the outflowing MHD power, the disk material still will be sheared, continuing to increase the local magnetic field strength and continuing to produce MHD power."
 The Appendix shows quantitatively how this occurs iu the herr uetric. producing feld cuhancement evervwhere in the creospheric region.," The Appendix shows quantitatively how this occurs in the Kerr metric, producing field enhancement everywhere in the ergospheric region."
 Tension in this erowine. working field hen produces a back-reaction ou the disk plasina falling oward the hole. accelerating it relative to the rotating yane in a direction to the black hole spin.," Tension in this growing, working field then produces a back-reaction on the disk plasma falling toward the hole, accelerating it relative to the rotating frame in a direction to the black hole spin."
 Dv the tine the magnetized accretion flow enters the jack hole. it could have a significantmeyative angular uonmentun which. when added to the hole. would decrease he spin.," By the time the magnetized accretion flow enters the black hole, it could have a significant angular momentum which, when added to the hole, would decrease the spin."
 Therefore. because of frame drageiug. power flows from hole to disk to field and finally to the outflow. exerting a braking torque on the hole itself.," Therefore, because of frame dragging, power flows from hole to disk to field and finally to the outflow, exerting a braking torque on the hole itself."
 No lounger sclregulated bv the accretion flow. [δρυ i the creospleric reeion above the disk now can continue to grow at the expense of the hole rotational enerey to order Byy. given by the disk solutions with a I&err aetric (Novikov&Thorne 1973)).," No longer self-regulated by the accretion flow, $B_{p0}$ in the ergospheric region above the disk now can continue to grow at the expense of the hole rotational energy to order $B_{d0}$, given by the disk solutions with a Kerr metric \cite{nt73}) )."
 For the gas pressure/clectrou-scattering accretion model asstuned here. this value is Equation (8)) then becomes (This differs from equation (16) of Moderski1996 bv a factor of a. from equation (9)). aud a factor of 0.P from the slightly smaller creospheric rotation rato.)," For the gas pressure/electron-scattering accretion model assumed here, this value is Equation \ref{bh_mhd_power_eq}) ) then becomes (This differs from equation (16) of \cite{ms96} by a factor of $\alpha$, from equation \ref{ss_eq}) ), and a factor of $\sim 0.4^{2}$ from the slightly smaller ergospheric rotation rate.)"
" While £3515,0 appears to scale with i9 aud i» i a ΠΛΑΤΟ simular to an accretion Iuminosity. it is portant to note that it does so only because of the scaling ou Bay."," While $L_{MHD}$ appears to scale with $m_{9}$ and $\dot{m}$ in a manner similar to an accretion luminosity, it is important to note that it does so only because of the scaling on $B_{d0}$."
 The energv liberated is derived mainly from the rotation of the hole itself and potentially mach larger than the accretion huuinosity (Lu%LOMeresbagi3 a).," The energy liberated is derived mainly from the rotation of the hole itself and potentially much larger than the accretion luminosity $L_{acc} \approx 10^{43} {\rm erg s^{-1}} \, m_{9} \, \dot{m}_{\sm 3}$ )."
 To turn Layyp ito an observed radio power. we note that. over long periods of tine. the central cugine may generate a jet oulv a fraction ¢ of the time.," To turn $L_{MHD}$ into an observed radio power, we note that, over long periods of time, the central engine may generate a jet only a fraction $\zeta$ of the time."
 Furthermore. oulv a fraction € of this average power is radiated at CGIIz frequencies in à 1 GIIz bandwidth aud observed with standard radio telescopes.," Furthermore, only a fraction $\epsilon$ of this average power is radiated at GHz frequencies in a 1 GHz bandwidth and observed with standard radio telescopes."
" From millimeter aud optical observations of repeated outbursts of radio galaxies aud quasars (sec.οι, Robson 1992)). we estimate typical duty cvele values to be ¢~OL0.1."," From millimeter and optical observations of repeated outbursts of radio galaxies and quasars (see, \cite{rob92}) ), we estimate typical duty cycle values to be $\zeta \sim 0.01 - 0.1$."
 The radio efficiency. is also very uncertain. but we note that for the nicroquasars GRO J1655-10 and GRS 1915|105. estimates ofe are of order 10? or so. just for the instautancous moving jet (Meier 1996)).," The radio efficiency is also very uncertain, but we note that for the microquasars GRO J1655-40 and GRS 1915+105, estimates of$\epsilon$ are of order $10^{-5}$ or so, just for the instantaneous moving jet \cite{m96}) )."
 A canonical value of €—10.7 waschosen for these calculations by comparing equation (12)) with the maxima observed radio source power (102*WIIz 1!) for a radio galaxy that should have a 10°AL. hole iu its nucleus (Ap21.0).," A canonical value of $\epsilon \sim 10^{-3}$ waschosen for these calculations by comparing equation \ref{lmhd_eq}) ) with the maximum observed radio source power $10^{27} \, {\rm W \, Hz^{-1}}$ ) for a radio galaxy that should have a $10^{9} \, M_{\sun}$ hole in its nucleus $M_B \approx -21.0$ )."
 This value for e is consistent with total conversion rates of ~(0.010.1 (see.cy. 1991)}. with oulv a fraction of that racdiated iu a 1 CGIIz bandwidth.," This value for $\epsilon$ is consistent with total conversion rates of $\sim 0.01-0.1$ (see, \cite{leahy91}) ), with only a fraction of that radiated in a 1 GHz bandwidth."
 Scaling to these values. the calculated radio power for differcut black hole spins is where iis still the average accretion rate ouly. uot that averaged over the eutire duty cvele.," Scaling to these values, the calculated radio power for different black hole spins is where $\dot{m}$ is still the average accretion rate only, not that averaged over the entire duty cycle."
 The final step in the model is to identify the maguetic switch as the cause of the FR ΤΠ break., The final step in the model is to identify the magnetic switch as the cause of the FR I/II break.
 That is. one should observe a distinct chanee iu the character of the jet produced when Layyp=Lei. or when the radio power attains the critical value To compute an observed break power we need an expression for the corona density p.y at Ry.," That is, one should observe a distinct change in the character of the jet produced when $L_{MHD} = L_{crit}$, or when the radio power attains the critical value To compute an observed break power we need an expression for the corona density $\rho_{c0}$ at $R_{0}$."
 Of course. coronal physics is poorly understood. even for the sun.," Of course, coronal physics is poorly understood, even for the sun."
 Nevertheless. we cau obtain a rough estimate by assuniug that the corona is produced locally by the accretion disk. so the disk and corona. deusities should scale together where pagg is the internal disk deusitv at Ry aud ip<< Lois a parameter which. while not necessarily a true constant. is asstmed to vary sufficiently slowly with the other parameters that it docs not affect significautlv the scaling of our results with ag. j. aud viv.," Nevertheless, we can obtain a rough estimate by assuming that the corona is produced locally by the accretion disk, so the disk and corona densities should scale together where $\rho_{d0}$ is the internal disk density at $R_{0}$ and $\eta << 1$ is a parameter which, while not necessarily a true constant, is assumed to vary sufficiently slowly with the other parameters that it does not affect significantly the scaling of our results with $m_{9}$ , $j$, and $\dot{m}$."
 The parameter is expected to be quite siuall: in the solar corona. for example. 4.=p.fp.(r—1)10.P im the region between the photosphere and ον.," The parameter $\eta$ is expected to be quite small; in the solar corona, for example, $\eta_{\sun} \equiv \rho_{c\sun} /
\rho_{\sun}(\tau=1) \sim 10^{-8} - 10^{-11}$ in the region between the photosphere and $2 R_{\sun}$."
 Tf the corona is a pair plasma. y (which is a mass ratio) could be very small indeed.," If the corona is a pair plasma, $\eta$ (which is a mass ratio) could be very small indeed."
 Combining equations (5)). (15)). aud (16)). and again using the eas pressure/electrou-scatteriug di: one obtains This is the maguetic switch model prediction for the locus of the FR I/II break as a fuuctiou of black hole mass.," Combining equations \ref{lcrit_eq}) ), \ref{pfr_eq}) ), and \ref{rhoc_eq}) ), and again using the gas pressure/electron-scattering disk, one obtains This is the magnetic switch model prediction for the locus of the FR I/II break as a function of black hole mass."
 This section applies the model to the FR 1Ἡ breals in the radio-optical plane. shows that the rotationally magnetic switch consistent. with long source morphology lifetimes. and addresses the more general questions of evolution and radio loud radioquiet sources.," This section applies the model to the FR I/II break in the radio-optical plane, shows that the rotationally-triggered magnetic switch consistent with long source morphology lifetimes, and addresses the more general questions of evolution and radio loud radioquiet sources."
 The important features of the model aud its applications are shown iu figure L.., The important features of the model and its applications are shown in figure \ref{roplane_fig}. .
"parameter 7,0.01. consistent wilh a capture al (he vortex periphery and a slow cit towarc the core in a lime scale Tí72¢/(8O7,): (2) the parameter measuring the turbulence inside the vortex is a,270.01.","parameter $\taus\simeq 0.01$, consistent with a capture at the vortex periphery and a slow drift toward the core in a time scale $T_{\rm capt} \simeq 2q/(3\Omega \taus)$; (2) the parameter measuring the turbulence inside the vortex is $\alpha_{\rm v} \simeq 0.01$."
 As a result. the correponding size of particles confined in the outer regions is of the order of a centimeter.," As a result, the correponding size of particles confined in the outer regions is of the order of a centimeter."
 The proposed model can account for the main characteristics of 115Ds observations., The proposed model can account for the main characteristics of 15D's observations.
 It is able to reproduce (he lisht-curve and permits to fit pretty well the data. (, It is able to reproduce the light-curve and permits to fit pretty well the data. (
1) The period and duration of (he eclipses may be reproduced by (the rotation around the star of a giunt swarm of solid particles. trapped in a persistent. gaseous vortex. (,"1) The period and duration of the eclipses may be reproduced by the rotation around the star of a giant swarm of solid particles, trapped in a persistent gaseous vortex. ("
2) During totality. the rise up of the flix can result from a lower opacity of the central regions. possibly coming from a particle size segregation (the larger the particle. (he deeper and the faster the confinement inside the vortex). (,"2) During totality, the rise up of the flux can result from a lower opacity of the central regions, possibly coming from a particle size segregation (the larger the particle, the deeper and the faster the confinement inside the vortex). ("
3) The estimated size of the (rapped particles ranges from 1 to 10 centimeters. consistent with the erev absorption observed in IXKIL115D. The presence and growth of bigger particles is verv likely due to the high densities and (he low relative velocities.,"3) The estimated size of the trapped particles ranges from 1 to 10 centimeters, consistent with the grey absorption observed in 15D. The presence and growth of bigger particles is very likely due to the high densities and the low relative velocities."
 This is in agreement with the expected predominance of large grains in PMS disks which seems also required to explain sub-millimeter ancl millimeter disk observations (eg.Nattaetal.1999). (, This is in agreement with the expected predominance of large grains in PMS disks which seems also required to explain sub-millimeter and millimeter disk observations \citep{Nal99}. (
4) The shift of the reverse peak observed in successive (ransils could be explained by eas rotation inside (he vortex. the inner zone rotating at half the orbital frequency similarly to what happens in smaller vortices (Chnavanis2000).,"4) The shift of the reverse peak observed in successive transits could be explained by gas rotation inside the vortex, the inner zone rotating at half the orbital frequency similarly to what happens in smaller vortices \citep{Cha00}."
". We can also speculate that: (1) ""secular. changes of the light-curve could result. from a dynamical evolution of the swarm of particles: Gi) the wavy behaviour al egress could correspond (o a turbulent wake of the vortex."," We can also speculate that: (i) $``$ secular"" changes of the light-curve could result from a dynamical evolution of the swarm of particles; (ii) the wavy behaviour at egress could correspond to a turbulent wake of the vortex."
 One of (he main assuniption in our model is that the disk is seen nearly edge-on. al an angle /~2—37. and wilh an optical depth Ay~0.6—0.7 magnitude during the out-of-eclipse periods.," One of the main assumption in our model is that the disk is seen nearly edge-on, at an angle $i\sim 2-3^\circ$, and with an optical depth $A_0\sim 0.6-0.7$ magnitude during the out-of-eclipse periods."
 This is a possible situation if the dust particles have grown and settled down to the mid-plane forming a flea! sub-disk dominated by laree grains., This is a possible situation if the dust particles have grown and settled down to the mid-plane forming a $\sl flat$ sub-disk dominated by large grains.
 The predominance of large particles in IXKIL115D is. indeed. completely consistent with our results.," The predominance of large particles in 15D is, indeed, completely consistent with our results."
 However. we wal to stress that no numerical simulation is. nowadays. able to reproduce (he huge azimutal extent of the vortex required to fit the observations.," However, we want to stress that no numerical simulation is, nowadays, able to reproduce the huge azimutal extent of the vortex required to fit the observations."
" Further hydrodynamical simulations and theoretical modeling of the particle dvnamies are necessary (o confirm and improve our model,", Further hydrodynamical simulations and theoretical modeling of the particle dynamics are necessary to confirm and improve our model.
"lenses and the lensing cross section: For a constant comoving lens clensity (πίς=mp(1+z)') and using the relation between proper distance and redshift dr,,,/4dz2'=c/[(1+29 L2(2)). Equation 5. can be rewritten as Note that although this derivation assumes that all lenses have the same mass M. the result is unchanged by instead integrating over some mass distribution function (?)..","lenses and the lensing cross section: For a constant comoving lens density $n(z') = n_0\,(1+z')^3$ ), and using the relation between proper distance and redshift $dr_{prop}/dz' =
c/[(1+z')H(z')]$ , Equation \ref{eq:pz-first} can be rewritten as Note that although this derivation assumes that all lenses have the same mass $M$, the result is unchanged by instead integrating over some mass distribution function \citep{lens:press-gunn}."
" Integrating Equation 6 over the lens redshift z vields the optical depth: Figure 1. shows the optical depth plotted for various values of Q, in the given cosmological model.", Integrating Equation \ref{eq:pz} over the lens redshift $z'$ yields the optical depth: Figure \ref{fig:opt-depth} shows the optical depth plotted for various values of $\Omega_c$ in the given cosmological model.
" From the optical depth. one can calculate the Poisson probability of a line of sight encountering a given number of lenses: Now. the probability ofa point source at z being maenilied by j5, can be written formally where p,(4) is the probability [orn lenses to produce a total magnification of j4,."," From the optical depth, one can calculate the Poisson probability of a line of sight encountering a given number of lenses: Now, the probability of a point source at $z$ being magnified by $\mu_u$ can be written formally as where $p_n(\mu_u)$ is the probability for $n$ lenses to produce a total magnification of $\mu_u$."
" Given that the optical cleptl is small for any realistic choice of ο) at low to moderate redshift (see Figure 1)).multiple-lensing events will be quite rare. so Equation 9. is well-approximated by where py(f,) isgiven in Equation 4.. and"," Given that the optical depth is small for any realistic choice of $\Omega_c$ at low to moderate redshift (see Figure \ref{fig:opt-depth}) ),multiple-lensing events will be quite rare, so Equation \ref{eq:pu-formal} is well-approximated by where $p_1(\mu_u)$ isgiven in Equation \ref{eq:p1}, and"
compression of the interstellar magnetic field bevond (he heliopause may cause ions {ο align prelerentially perpendicular to the interstellar magnetic field through conservation of the first adiabatic invariant and conservation of energy.,compression of the interstellar magnetic field beyond the heliopause may cause ions to align preferentially perpendicular to the interstellar magnetic field through conservation of the first adiabatic invariant and conservation of energy.
 This will also lead (o a special orientation of peaked ENA emissions perpendicular to the interstellar magnetic field. and max help (o explain both how the ribbon is formed aud (the even more surprising fine structure observed in it (MeComas et al..," This will also lead to a special orientation of peaked ENA emissions perpendicular to the interstellar magnetic field, and may help to explain both how the ribbon is formed and the even more surprising fine structure observed in it (McComas et al.,"
 2009)., 2009).
 The basic idea of secondary. ENA generation of the IBEX ribbon was futher examined by Heerikluisenetal.(2010)., The basic idea of secondary ENA generation of the IBEX ribbon was further examined by \cite{heerick10}.
. These authors assumed (hat pickup protons in the outer heliosheath have and retain a partial shell distribution and that their re-neutralization is effectively instantaneous., These authors assumed that pickup protons in the outer heliosheath have and retain a partial shell distribution and that their re-neutralization is effectively instantaneous.
 This approach is significantly different from ours in this study since we solve consistently for the motion of pickup protons along magnetic [ield lines in (he scatter-Iree limit. ancl (hus include the motion of PUIs along the field line between their pickup and reneutralization.," This approach is significantly different from ours in this study since we solve consistently for the motion of pickup protons along magnetic field lines in the scatter-free limit, and thus include the motion of PUIs along the field line between their pickup and reneutralization."
 The IBEX observations of an unexpected narrow ribbon of enhanced ENA [Iuxes raise fundamental euestions about the origin of these particles., The IBEX observations of an unexpected narrow ribbon of enhanced ENA fluxes raise fundamental questions about the origin of these particles.
 MeConmasetal.(2009). considered 6 possible sources of the ribbon including5 both sources inside and bevond the heliopause., \cite{mccomas09} considered 6 possible sources of the ribbon including both sources inside and beyond the heliopause.
 The idea that the source lies in (he outer heliosheath has at least two arguments., The idea that the source lies in the outer heliosheath has at least two arguments.
 First. il was recentlv demonstrated in Izmodenovetal.(2009) that high-enerey charged (pickup) protons can arise in the outer heliosheath due to charge exchange between interstellar protons and ENAs originated inside the heliopause.," First, it was recently demonstrated in \citet{izmod09} that high-energy charged (pickup) protons can arise in the outer heliosheath due to charge exchange between interstellar protons and ENAs originated inside the heliopause."
 The energy distribution of the pickup protons has maximum near 1l keV and extends (o energies of about LO keV. Secondly. it was noted in MeComasetal.(2009) and Schwadronetal.(2009). Chat the ribbon position. as seen [rom the Earth. coincides closely with the likely magnetic field direction located just bevond the heliopause. where. according to the recent MIID models. the interplanetary," The energy distribution of the pickup protons has maximum near 1 keV and extends to energies of about 10 keV. Secondly, it was noted in \citet{mccomas09} and \citet{schwadr09} that the ribbon position, as seen from the Earth, coincides closely with the likely magnetic field direction located just beyond the heliopause, where, according to the recent MHD models, the interplanetary"
Iu this letter we have used the high quality IIR diagram of the IIvades to place lanits on the scatter in the metallicities of this cluster.,In this letter we have used the high quality HR diagram of the Hyades to place limits on the scatter in the metallicities of this cluster.
 We find that the scatter in |Fo/TI| has standard deviation approximately 0.03 dex., We find that the scatter in [Fe/H] has standard deviation approximately 0.03 dex.
 This estimate is an inproveinieut upon previous spectroscopic estimates and suggests that open clusters are formed from gas that is homogenous im mectallicity., This estimate is an improvement upon previous spectroscopic estimates and suggests that open clusters are formed from gas that is homogenous in metallicity.
 We have introduced the possibility that the extremely low scatter from a predicted isochrone can be used to constrain stellar pollution scenarios which involve the addition of hieh-Z material to the star from processes associated with planet formation., We have introduced the possibility that the extremely low scatter from a predicted isochrone can be used to constrain stellar pollution scenarios which involve the addition of high-Z material to the star from processes associated with planet formation.
 This can be done if stellar evolution tracks are calculated which include the addition of high-Z metals after formation to see where stars lic ou the IIR diagram., This can be done if stellar evolution tracks are calculated which include the addition of high-Z metals after formation to see where stars lie on the HR diagram.
 If polluted stars lie off the isochrone for unpolluted stars then the scatter of stars in the IIR diagram can be used to constrain the extent of stellar pollution., If polluted stars lie off the isochrone for unpolluted stars then the scatter of stars in the HR diagram can be used to constrain the extent of stellar pollution.
 Otherwise comparisons between spectroscopic measured metallicities aud those estimated from the mean properties of the cluster can be used to identify stars which are likely to have been polluted by ligh-Z material., Otherwise comparisons between spectroscopic measured metallicities and those estimated from the mean properties of the cluster can be used to identify stars which are likely to have been polluted by high-Z material.
 Since few stars in voung clusters lave had colmprehcusive inetallicity analyses. subsequeut studies could be expanded anc verified. and including more stars with a larecr ranee of mass.," Since few stars in young clusters have had comprehensive metallicity analyses, subsequent studies could be expanded and verified, and including more stars with a larger range of mass."
" Care should be taken that stars with discrepant ietallicities are not thrown out of the sample aud assumed to be non-cluster ποος,", Care should be taken that stars with discrepant metallicities are not thrown out of the sample and assumed to be non-cluster members.
 Currently there is no observational evidence that immetallicities of different mass stars dn a giveu cluster differ., Currently there is no observational evidence that metallicities of different mass stars in a given cluster differ.
 Comparison of spectroscopically measured metallicities in the Uvades. Pleiades aud NCC 2261 between Co aud Fo stars find that the moetallicities are the same within the observational uncertainties (approximately 0.1 dex) (sineetal.2000:Fricleaard1990a).," Comparison of spectroscopically measured metallicities in the Hyades, Pleiades and NGC 2264 between G and F stars find that the metallicities are the same within the observational uncertainties (approximately 0.1 dex) \citep{king,friel90a}."
. Despite the fact that the ISM (along Uspe line of sight absorption lues) appears to have a uniform imetallicity (Meyer.Jura&Cardelli1998).. the metallicitics of moving eroups differ significautl (Fricl&Bocseaard1990a.b).. and the vouugest ones seem to have increasingly subsolar values.," Despite the fact that the ISM (along 1kpc line of sight absorption lines) appears to have a uniform metallicity \citep{meyer}, the metallicities of moving groups differ significantly \citep{friel90a,friel90b}, and the youngest ones seem to have increasingly subsolar values."
 There also are extremely metal rich stars some of which are not vouug., There also are extremely metal rich stars some of which are not young.
 The high metallicity stars are kinematically related. they tend to be iu moving eroups or streams (Soubiran1999).," The high metallicity stars are kinematically related, they tend to be in moving groups or streams \citep{soubiran}."
 This sugeests that high metallicity stars are foriued iu clusters that are uuiforui in metalliitv just Like lower metallicity stars., This suggests that high metallicity stars are formed in clusters that are uniform in metallicity just like lower metallicity stars.
 This may rule out a scenario where even a few percent of stars have their imetallicity. increased significantly by the affects of planetary processes and would instead support a model where higher metallicity stars are more likely to harbor short period planets., This may rule out a scenario where even a few percent of stars have their metallicity increased significantly by the affects of planetary processes and would instead support a model where higher metallicity stars are more likely to harbor short period planets.
 However there is good evidence that oue parent star of an extra solar plauet has eugulfed anuetarv imaterial: the detection of Lithium 6 in ID N29[3 (Israclianetal.2001)., However there is good evidence that one parent star of an extra solar planet has engulfed planetary material; the detection of Lithium 6 in HD 82943 \citep{israelian}.
. Further study of uuiforui »pulatious such as are found i vouug clusters may be able to deteriuue the relative statistical importance of otl scenarios., Further study of uniform populations such as are found in young clusters may be able to determine the relative statistical importance of both scenarios.
 This work could not have been carried out without iclpful discussions with Don Garnett. Caty Pilachowski. Jonathan Lunine. Ari Laor. Patrick Young. Dave Aructt. Adam Burrows. Nadva Corlova. aud Mike \lever.," This work could not have been carried out without helpful discussions with Don Garnett, Caty Pilachowski, Jonathan Lunine, Ari Laor, Patrick Young, Dave Arnett, Adam Burrows, Nadya Gorlova, and Mike Meyer."
 I thauk Joss de Bruijue for providing nie with information on the hieh fidelity Ibvades. sample., I thank Joss de Bruijne for providing me with information on the high fidelity Hyades sample.
" A. ον geratefully acknowledges hospitality aud support as a Visitor at the ουπολι,", A. C. gratefully acknowledges hospitality and support as a Visitor at the Technion.
"location where e, changes The solution region above Σι, (between 2),c]300 and the sonic surface at ὃς2 4.7) represents the iwind.","location where $v_r$ changes The solution region above $\tilde z_{\rm b}$ (between $\tilde
z_{\rm b} \approx 3.0$ and the sonic surface at $\tilde z_{\rm s}
\approx 4.7$ ) represents the ."
". 1n this region e, is 2O. rm, is super-Ixeplerian and the magnetic field transfers angular momentuni back to the matter. causing the uppermost lavers of the cise to be driven out centrifueally (with the mass Hux in the wind ellectively fixed by the reeularity condition at the sonic critical point of the cise solution)."," In this region $\vr$ is $>0$, $\vf$ is super-Keplerian and the magnetic field transfers angular momentum back to the matter, causing the uppermost layers of the disc to be driven out centrifugally (with the mass flux in the wind effectively fixed by the regularity condition at the sonic critical point of the disc solution)."
" The distinguishing properties of the four Hall. sub-regimes. as inferred in the context of the hyclrostatic approximation. are listed in ""Table A2.."," The distinguishing properties of the four Hall sub-regimes, as inferred in the context of the hydrostatic approximation, are listed in Table \ref{table:constraints1}."
" To see how the expressions derived. in this approximation for |db,/db..|. bh—hjfhy (where h is the magnetically reduced: density scaleheight. defined as the location where the density drops to poffe) and. off compare with the properties of the full solutions listed. in “Table 1.. we need to relate the parameters Cj. Jo) and Yo used in the two-component plasma formulation of Paper L to the tensor conductivity components emploved in this work."," To see how the expressions derived in this approximation for $|db_r/db_\phi|_0$, $\tilde h \equiv h/h_{\rm T}$ (where $h$ is the magnetically reduced density scaleheight, defined as the location where the density drops to $\rho_{\rm 0}/\sqrt{e}$ ) and $\tilde z_{\rm
b}/\tilde h$ compare with the properties of the full solutions listed in Table \ref{table:boundary}, we need to relate the parameters $\betio$, $\beteo$ and $\upo$ used in the two-component plasma formulation of Paper I to the tensor conductivity components employed in this work."
 Phe parameter Yo is the mid-plane value of the neutralion coupling function given by equation A2 and is equal to the Elsasser number in the ambipolar diffusivity limit., The parameter $\upo$ is the mid-plane value of the neutral–ion coupling function given by equation \ref{eq:eta} and is equal to the Elsasser number in the ambipolar diffusivity limit.
 The relationships between these variables and the conductivity ratios can be obtained [rom equations 195L.97 and reduce. in the limit gy—0/3.«xm1. lo and The conductivity ratio σησι serves to distinguish. the ambipolar cilfusivity domain (60Z9σιX ση) from the Hall cülfusivitv regime (σι.ση] &o): theambipolar dilfusivitv limit corresponds to this ratio tending to 0. whereas the Hall limit corresponds to this ratio tending to," The relationships between these variables and the conductivity ratios can be obtained from equations I.95–I.97 and reduce, in the limit $q\equiv \beta_{\rm i}/\beta_{\rm e} \ll 1$, to and The conductivity ratio $\sigH/\tilde \sigma_\perp$ serves to distinguish the ambipolar diffusivity domain $\sigpar \gg \sigP \gg |\sigH|$ ) from the Hall diffusivity regime $\sigP \ll |\sigH| \ll \sigpar$ ): theambipolar diffusivity limit corresponds to this ratio tending to 0, whereas the Hall limit corresponds to this ratio tending to"
2001).,2001).
" The resulting source positions and fluxes. calculated using all the data from the projected baselines longer than 1.5 KA shown in reftable:RTJonrees plin,ags."," The resulting source positions and fluxes, calculated using all the data from the projected baselines longer than 1.5 $\lambda$ shown in \\ref{table:RT_sources_opt_ir_mags}."
 Lhesesourceswerethensiublracted from all the visibilites., These sources were then subtracted from all the visibilites.
" To check that no residual flux from any of the sources was present. source-subtracted maps were made using baselines of length 1.5 KA and greater (shown in reftig:long,aselinesourceubl racted)). and2.0 EA and greater."," To check that no residual flux from any of the sources was present, source-subtracted maps were made using baselines of length 1.5 $\lambda$ and greater (shown in \\ref{fig:long_baselines_source_subtracted}) ), and 2.0 $\lambda$ and greater."
 These maps are both consistent with noise., These maps are both consistent with noise.
 Finally. a map was made using source-subtracted visibilities with projected baselines in the range 0.65—1.0 kA. corresponding to the single shortest physical baseline.," Finally, a map was made using source-subtracted visibilities with projected baselines in the range 0.65–1.0 $\lambda$, corresponding to the single shortest physical baseline."
" This map showed a negative feature centred at 02 33"" 20.8"". 30° 21 58.7"" (2000). and with a minimum flux of in the dirty map of 6755iy."," This map showed a negative feature centred at $^h$ $^m$ $^s$, $^\circ$ $^\prime$ $^{\prime \prime}$ (J2000), and with a minimum flux of in the dirty map of $-675
\mu$ Jy."
 The synthesized beam of this single-baseline map has strong sidelobes. so the rms level on the map even well outside the primary beam is stil somewhat affected by the presence of the negative feature at the centre of the map.," The synthesized beam of this single-baseline map has strong sidelobes, so the rms level on the map even well outside the primary beam is still somewhat affected by the presence of the negative feature at the centre of the map."
" We therefore made a very large map (2048 2048 5"".square pixels) and measured the map rms in. 128-pixe square regions where the map apppeared by eye to be unaffected by sidelobes from the central source.", We therefore made a very large map (2048 $\times$ 2048 $^{\prime \prime}$ -square pixels) and measured the map rms in 128-pixel square regions where the map apppeared by eye to be unaffected by sidelobes from the central source.
 The rms levels agreed wth the expected noise. due to the telescope system temperature. of 95 py.," The rms levels agreed wth the expected noise, due to the telescope system temperature, of 95 $\mu$ Jy."
" The short-baseline map was to a flux limit of 285pJy. and this map is presented in reacted. reftig:short,asetines;ourcesub"," The short-baseline map was to a flux limit of $- 285
\mu$ Jy, and this map is presented in \\ref{fig:short_baselines_source_subtracted}."
"lOninspeclionoflheazimulhselend visibilities reftig:azimuthal,,cerage,fol) ). dhe flirontheshorlestbasclineis- pulv. with an rms of 170 jdy (note that since the shortest baseline has an elliptical track in the aperture plane. this rms includes the variation in the amplitude of the SZ signal as a function of uv distance)."," On inspection of the azimuthally-averaged visibilities \\ref{fig:azimuthal_average_plot}) ), the flux on the shortest baseline is $-695 \mu$ Jy, with an rms of 170 $ \mu$ Jy (note that since the shortest baseline has an elliptical track in the aperture plane, this rms includes the variation in the amplitude of the SZ signal as a function of $uv$ distance)."
 We obtained /i-band imaging of the TOC 70233.34-3021. field using the Imaging Grism Instrument on the McDonald Observatory 2.7-m Smith Reflector in 1998 October. 1999 October. and 2000 September. and ./-band imaging of the field at the Calar Alto 3.5- telescope in 1999 August.," We obtained $R$ -band imaging of the TOC J0233.3+3021 field using the Imaging Grism Instrument on the McDonald Observatory 2.7-m Smith Reflector in 1998 October, 1999 October, and 2000 September, and $J$ -band imaging of the field at the Calar Alto 3.5-m telescope in 1999 August."
 Additionally. we obtained several ÁA-band images of the TOC J0233.343021 field with UFTI at UKIRT in 2000 July.," Additionally, we obtained several $K$ -band images of the TOC J0233.3+3021 field with UFTI at UKIRT in 2000 July."
 Full details. and extensive discussion of other candidate cluster galaxies. will be presented in C2001.," Full details, and extensive discussion of other candidate cluster galaxies, will be presented in C2001."
 Using these images. in combination with high-resolution VLA imaging (C2001). we have identitied the host galaxies of the central five NVSS radiosources shown in Fig. ??..," Using these images, in combination with high-resolution VLA imaging (C2001), we have identified the host galaxies of the central five NVSS radiosources shown in Fig. \ref{fig:nvss_plus_long_baselines}."
" For four of these galaxies. we have obtained 3.5"" -diameter aperture magnitudes in all filters: for the fifth galaxy. only /? and ./ magnitudes are available."," For four of these galaxies, we have obtained $^{\prime
\prime}$ -diameter aperture magnitudes in all filters; for the fifth galaxy, only $R$ and $J$ magnitudes are available."
 The measured magnitudes of the galaxies are presented in Table ?? along with the positions of the host galaxy identifications., The measured magnitudes of the galaxies are presented in Table \ref{table:RT_sources_opt_ir_mags} along with the positions of the host galaxy identifications.
 The positions were measured from the /i-band image. with an astrometric solution using the positions of stars measured in the Cambridge APM survey.," The positions were measured from the $R$ -band image, with an astrometric solution using the positions of stars measured in the Cambridge APM survey."
 The magnitudes of the galaxies are all roughly ἐνzz18. z20. F0zz24.," The magnitudes of the galaxies are all roughly $K \approx 18$, $J \approx 20$, $R \approx 24$."
 In 2000 January. we attempted spectroscopy of the radiogalaxies with ISIS on the WHT.," In 2000 January, we attempted spectroscopy of the radiogalaxies with ISIS on the WHT."
 We used a 2-aresee slit. with sensitivity on the ISIS red arm up to 9340Α.," We used a 2-arcsec slit, with sensitivity on the ISIS red arm up to 9340."
. We took spectra of A (at PAZIIS. the axis of its double structure: see C2001). galaxies B C (as a pair with PA=113). and galaxies D E tas a pair with PA=77).," We took 1800-second spectra of A (at PA=118, the axis of its double structure; see C2001), galaxies B C (as a pair with PA=113), and galaxies D E (as a pair with PA=77)."
 Paint red continuum was detected spectroscopicallyin all the cases expected (i.e. the + brightest). but no emission lines were evident.," Faint red continuum was detected spectroscopicallyin all the cases expected (i.e. the 4 brightest), but no emission lines were evident."
"drop-out technique; at the lower redshifts, the Y—J colour might be too blue to enter our selection, and at the higher redshift end of our range the effect of the a forest means that an increasingly large fraction of the J-band filter is absorbed, so only the most UV-luminous galaxies will appear in our apparent-magnitude-limited sample.","drop-out technique; at the lower redshifts, the $Y-J$ colour might be too blue to enter our selection, and at the higher redshift end of our range the effect of the $\alpha$ forest means that an increasingly large fraction of the $J$ -band filter is absorbed, so only the most UV-luminous galaxies will appear in our apparent-magnitude-limited sample."
" We quantify this effect, and hence constrain the luminosity function through our observed number counts."," We quantify this effect, and hence constrain the luminosity function through our observed number counts."
 The probability of recovering a high-redshift galaxy as a function of redshift and rest frame UV luminosity can be found with simulations., The probability of recovering a high-redshift galaxy as a function of redshift and rest frame UV luminosity can be found with simulations.
" To perform these simulations we add into the images a large number of fake galaxies, with properties similar to those of the observed high-redshift population (ie. compact with half-light radii rn,z01, large Lyman-a forest decrement of DA~0.99 and blue rest-frame UV colours)."," To perform these simulations we add into the images a large number of fake galaxies, with properties similar to those of the observed high-redshift population (i.e. compact with half-light radii $r_{hl}\approx 0\farcs1$, large $\alpha$ forest decrement of $D_A\approx 0.99$ and blue rest-frame UV colours)."
 We then run our selection procedure and infer the probability of recovering such galaxies as a function of redshift and magnitude (see Figure 5))., We then run our selection procedure and infer the probability of recovering such galaxies as a function of redshift and magnitude (see Figure \ref{fig:selection}) ).
 We adopt the effective volume approach as described in Steidel et al. (, We adopt the effective volume approach as described in Steidel et al. (
"1999) and Stanway, Bunker McMahon (2003), such that the probability of recovering a galaxy in our survey depends on the redshift and absolute UV magnitudes, p(Muv,z), and from this the effective survey volume can be calculated (νε).","1999) and Stanway, Bunker McMahon (2003), such that the probability of recovering a galaxy in our survey depends on the redshift and absolute UV magnitudes, $p(M_{UV},z)$, and from this the effective survey volume can be calculated $V_{eff}$)."
" We use a Gaussian distribution of spectral slopes, with (8)=—2.2 and σ(β)= 0.5, reflecting the generally"," We use a Gaussian distribution of spectral slopes, with $\langle \beta \rangle = -2.2$ and $\sigma(\beta)=0.5$ , reflecting the generally"
such inflow were possible bv some unknown mechanisin. the accretion speed should be so ligh that am accretion shock would be unavoidable. transforming ram pressure into thermal pressure aud probably ecucrating a hydraulic juup in the edge of the disk aud other observable features such as Πα cmission.,"such inflow were possible by some unknown mechanism, the accretion speed should be so high that an accretion shock would be unavoidable, transforming ram pressure into thermal pressure and probably generating a hydraulic jump in the edge of the disk and other observable features such as $\alpha$ emission."
 The blydraulic jump will divert the flow around the disk decreasing dramatically its confinement ability., The hydraulic jump will divert the flow around the disk decreasing dramatically its confinement ability.
" A colerent external (intergalactic) magnetic field could in principle contribute to the confinement πο certain cireunstauces,", A coherent external (intergalactic) magnetic field could in principle contribute to the confinement under certain circunstances.
 It is important to distinguish here between the intracluster and the wider intergalactic field., It is important to distinguish here between the intracluster and the wider intergalactic field.
 As we are dealing with ealaxics in the field. we are interested in the amplitude of the latter one.," As we are dealing with galaxies in the field, we are interested in the amplitude of the latter one."
 At present. setting observational constraints to the streneth and topology of the intergalactic magnetic fields is a challeging task (see Ryu. Ikaug Bicriuaun 1988: Irouberg et al.," At present, setting observational constraints to the strength and topology of the intergalactic magnetic fields is a challeging task (see Ryu, Kang Biermann 1988; Kronberg et al."
 2001: aud Widrow 2002 for a review) aud. therefore. they are very uncertain.," 2001; and Widrow 2002 for a review) and, therefore, they are very uncertain."
 Coherent maeguetie fields on scales ich ercater than the typical sizes of galaxies could be as large as ~OL pC in the optimistic situation that extra auplification through batteries. dvuamos or other routes of magnetization were operative.," Coherent magnetic fields on scales much greater than the typical sizes of galaxies could be as large as $\sim 0.1$ $\mu$ G in the optimistic situation that extra amplification through batteries, dynamos or other routes of magnetization were operative."
 We assune the ideal situation that initially the preealactic magnetic field is uniform and fills the protogalaxy making a finite anele a with the rotation axis of the galaxy., We assume the ideal situation that initially the pregalactic magnetic field is uniform and fills the protogalaxy making a finite angle $\alpha$ with the rotation axis of the galaxy.
 As the disk is forming by accretion of eas. tle magnetic field lines are advected with the flow.," As the disk is forming by accretion of gas, the magnetic field lines are advected with the flow."
 It is wellknown in studies of the collapse ofinterstellar clouds that. under fux-freezing conditions. this compression produces that the total magnetic coutributiou to the virial theorem Is positive (see. ce. Chapt.," It is well-known in studies of the collapse of interstellar clouds that, under flux-freezing conditions, this compression produces that the total magnetic contribution to the virial theorem is positive (see, e.g., Chapt."
 21 of Shu 1992 aud Chapt., 24 of Shu 1992 and Chapt.
 11 of Mestel 1999) aud. therefore. the confinement by the external field is mipeded.," 11 of Mestel 1999) and, therefore, the confinement by the external field is impeded."
 Iu order to have coufineient. he disk must be able to eet rid of the magnetic flux.," In order to have confinement, the disk must be able to get rid of the magnetic flux."
 Iu the strict evliudrical sviuuietzx. the external magnetic field nmst be vertical aud then Eq. (11))," In the strict cylindrical symmetry, the external magnetic field must be vertical and then Eq. \ref{eq:magtermtxt}) )"
 with Bog=0 cls us that magnetic confinement is ouly possible if the voluuc-averaged vertical component of the maeuetic field within V is strictly simaller than the magnetic field at the )oundary., with $B_{\phi {\rm S}}=0$ tells us that magnetic confinement is only possible if the volume-averaged vertical component of the magnetic field within $V$ is strictly smaller than the magnetic field at the boundary.
 Since eg is directed inward aud Qeg/oR«0 in the AICP. |D. increases at every point of tle disk due o the ieard advection (see also Reves-BRuiz Stepiusli 1996: Moss Shukurov 2001): under those conditious confinement is not viable.," Since $v_{R}$ is directed inward and $\partial v_{R}/\partial R<0$ in the MCP, $|B_{z}|$ increases at every point of the disk due to the inward advection (see also Reyes-Ruiz Stepinski 1996; Moss Shukurov 2001); under those conditions confinement is not viable."
 This problem could be alleviated if magnetic field lines were expelled from the disk by aiubipolar and turbulent diffusion., This problem could be alleviated if magnetic field lines were expelled from the disk by ambipolar and turbulent diffusion.
 However. a thorough numerical study by Toward&I&ulsrud(1997) shows that |B.] is also increased by ambipol shearing ternis to a value considerably larger than its initial value.," However, a thorough numerical study by \citet{how97}
 shows that $|B_{z}|$ is also increased by ambipolar shearing terms to a value considerably larger than its initial value."
 Iu the lack of a mechanism able to reduce |B.) within disks. pregalactic magnetic fields do not help to coufine ealactic disks.," In the lack of a mechanism able to reduce $|B_{z}|$ within disks, pregalactic magnetic fields do not help to confine galactic disks."
 This example clearly illustrates why confinement is so rare in nature., This example clearly illustrates why self-confinement is so rare in nature.
 Models based ou the coufinement by anchored fields to the galactic nucleus will present a similar problem. as [D| is expected to decline outwards.," Models based on the confinement by anchored fields to the galactic nucleus will present a similar problem, as $|B_{z}|$ is expected to decline outwards."
 Battaner Florido (1995. 2000) aud Battaner. Florido Jnueénuez-Viceute (2002) poiuted out that: (1) the magnetic alternative to dark halos does not suffer frou the problem of excessive flaring of the disk and (2) the existence of truncated. stellar disks is explained satistactorily iu the maenuetic alternative.," Battaner Florido (1995, 2000) and Battaner, Florido Jiménnez-Vicente (2002) pointed out that: (1) the magnetic alternative to dark halos does not suffer from the problem of excessive flaring of the disk and (2) the existence of truncated stellar disks is explained satisfactorily in the magnetic alternative."
 Putting aside the problemi of radial confinement discussed in the previous sections. if ds stimulating to see whether altcruative scenarios to conventional ideas have the poteutial of explaining other phenomena in a natural wav.," Putting aside the problem of radial confinement discussed in the previous sections, it is stimulating to see whether alternative scenarios to conventional ideas have the potential of explaining other phenomena in a natural way."
 If this occurs. alternative models could attrack much more attention.," If this occurs, alternative models could attrack much more attention."
 For iustance. some curpirical correlations of galactic warps have led to Castro-Rocdrigenez et al. (," For instance, some empirical correlations of galactic warps have led to guez et al. ("
2002) to sugecst teutatively that outer rotation curves couk not be caused bv the preseuce of a iassive halo.,2002) to suggest tentatively that outer rotation curves could not be caused by the presence of a massive halo.
 Iu the next sections we will show that the flaring problem relmaius unsolved aud that the existence of truucatec stellar disks cannot be considered a prediction of the magnetic alternative., In the next sections we will show that the flaring problem remains unsolved and that the existence of truncated stellar disks cannot be considered a prediction of the magnetic alternative.
 If galaxies contained ouly stars aud the observed eas. the initial protodisk of mass A; should follow approximately a Iseplerian rotation. e?—GA/R. at large R. and evolves to the preseut-day exponential disk of iiass Mg aud radial scale leugth Ry. due to magnetic contraction.," If galaxies contained only stars and the observed gas, the initial protodisk of mass $M_{\rm i}$ should follow approximately a Keplerian rotation, $v_{\phi}^{2}=GM_{\rm i}/R$, at large $R$, and evolves to the present-day exponential disk of mass $M_{\rm f}$ and radial scale length $R_{\rm d}$ , due to magnetic contraction."
" The peak velocity of the final disk is Cha,=0.62απRa (e.g. Binney Tremaine 198T)"," The peak velocity of the final disk is $v_{\rm max}=0.62 \sqrt{GM_{\rm f}/R_{\rm d}}$ (e.g., Binney Tremaine 1987)."
 For simplicity. we assume that bevoud 2.27R4 the rotation curve of a present ealaxy is flat with cieular veocitY Cas.," For simplicity, we assume that beyond $2.27 R_{\rm d}$ the rotation curve of a present galaxy is flat with circular velocity $v_{\rm max}$."
 During the MCP. a ving of gas located initially at a radius Ry must squeeze p to a final radius Re. such as cua=VGALG roni angular momentum conservation.," During the MCP, a ring of gas located initially at a radius $R_{\rm i}$ must squeeze up to a final radius $R_{\rm f}$, such as $v_{\rm max} R_{\rm f}=\sqrt{GM_{\rm i}R_{\rm i}}$ from angular momentum conservation."
 Therefore. the oeutial radius is eiven by the simple expression Ry=δεΑΙRe RM.," Therefore, the initial radius is given by the simple expression $R_{\rm i}=0.38 (M_{\rm f}/M_{\rm i})(R_{\rm f}/R_{\rm d})R_{\rm f}$ ."
 For a normal spiral galaxy like ιο Milkv Wa. with a ucarlv fat rotation curve up to je last measured poiut. typically at ~30 kpc. aud with Rax Lkpe. the AICP should be able to squeeze the disk at least ἅοι ~85 kpc to ~30 kpe iu radius. provided 16 niass loss in outflows is negligible. corresponding to a vpieal radial velocity of 3 kan |! for oue Hubble time.," For a normal spiral galaxy like the Milky Way, with a nearly flat rotation curve up to the last measured point, typically at $\sim 30$ kpc, and with $R_{\rm d}\approx 4$ kpc, the MCP should be able to squeeze the disk at least from $\sim 85$ kpc to $\sim 30$ kpc in radius, provided the mass loss in outflows is negligible, corresponding to a typical radial velocity of $3$ km $^{-1}$ for one Hubble time."
 According to 83.. this contraction is an unufeasible process due to the global expansive nature of magnetic fields.," According to \ref{sec:core}, this contraction is an unfeasible process due to the global expansive nature of magnetic fields."
 The level of contraction may be significantly reduced if a substantial fraction of the mass is lost in winds following carly epoch of star bursts. ic. MgMi.," The level of contraction may be significantly reduced if a substantial fraction of the mass is lost in winds following early epoch of star bursts, i.e. $M_{\rm f}<M_{\rm i}$."
 Receutly. Zhao(2002) has estimated that less than X of the total mass of the Galaxy has been lost during its lifetime.," Recently, \citet{zha02} has estimated that less than $4\%$ of the total mass of the Galaxy has been lost during its lifetime."
 Even assuming an amount of mass loss 3 times larger than this value. the maeguitude of the contraction of the disk is still huge. from 12 kpc to 30 kpc. again for Ry= Lkpe.," Even assuming an amount of mass loss $3$ times larger than this value, the magnitude of the contraction of the disk is still huge, from $72$ kpc to $30$ kpc, again for $R_{\rm d}=4$ kpc."
 The AICP represeuts a serious threat to the sugeestion by Battaner.Florido.Jimeénez-Viceute(2002). that. since stars do not fecl the galactic magnetic field. the escape of stars formed bevoud a certain galactocentric radius could explain the observed truncation of stellar disks.," The MCP represents a serious threat to the suggestion by \citet{bat02} that, since stars do not feel the galactic magnetic field, the escape of stars formed beyond a certain galactocentric radius could explain the observed truncation of stellar disks."
 The calculations by Battaner et al. (, The calculations by Battaner et al. (
2002) are rooted iu a scenario that docs not consider the AICP. with the gas disks being iu anomalous rotation from the very beginning.,"2002) are rooted in a scenario that does not consider the MCP, with the gas disks being in anomalous rotation from the very beginning."
 If oue tries to iuclude the MCP iu themodel. the stellar distribution function depends on different assuniptions that concern the details of the AICP. such as the radial velocity of the contraction. or the relationship between," If one tries to include the MCP in themodel, the stellar distribution function depends on different assumptions that concern the details of the MCP, such as the radial velocity of the contraction, or the relationship between"
Msolar<m<4 Msolarasstudiedhere. feqimpliesMace =0.23£0.03Msolar.,"$<m< 4$ as studied here, \\ref{eq9} implies $M_{\rm acc} = 0.23 \pm 0.03$."
". Taken at face value. such a high M,.. figure would imply that a considerable amount of mass is accreted during the PMS phase. If this is the case. the PMS evolution of moderate-mass stars. say <2Msolar.. should be reconsidered and recaleulated taking into account the high M, values."," Taken at face value, such a high $M_{\rm acc}$ figure would imply that a considerable amount of mass is accreted during the PMS phase, If this is the case, the PMS evolution of moderate-mass stars, say $<
2$, should be reconsidered and recalculated taking into account the high $\dot M_{\rm acc}$ values."
 One of the effects of 99 is that the appropriate fz3yrs for a star that reaches the MS with à given mass will be longer than the value estimated by evolutionary models that assume Mj=0 after the first few 10 vyr (e.g. Baraffe et al., One of the effects of 9 is that the appropriate $t_{\rm ZAMS}$ for a star that reaches the MS with a given mass will be longer than the value estimated by evolutionary models that assume $\dot M_{\rm acc} \equiv 0$ after the first few $10^4$ yr (e.g. Baraffe et al.
 2009)., 2009).
 Since the star had an appreciably lower mass for a very long time (several Myr) at the early epochs. it was necessarily evolving more slowly.," Since the star had an appreciably lower mass for a very long time (several Myr) at the early epochs, it was necessarily evolving more slowly."
 Extrapolating these findings to lower masses than those that we have reached here (1.6. <0.4 Msolar)) would carry the implication that no star could form with a mass below a suitable minimum value., Extrapolating these findings to lower masses than those that we have reached here (i.e. $< 0.4$ ) would carry the implication that no star could form with a mass below a suitable minimum value.
 For 3346 this would correspond to a lower cut-off of the initial mass function at ~0.2Msolar.," For 346 this would correspond to a lower cut-off of the initial mass function at $\sim
0.2$."
". In the Milky Way. where M,∙ 1s about five times lower but the metallicity is about 8 times higher so that the PMS evolutionary times are a factor of about 2.8 longer. the cut-off would occur at ~0.06Msolar."," In the Milky Way, where $\dot M_{\rm acc}$ is about five times lower but the metallicity is about 8 times higher so that the PMS evolutionary times are a factor of about $2.8$ longer, the cut-off would occur at $\sim
0.06$."
. On the other hand. it is also possible that for low-mass stars circumstellar dises are eroded in a finite time that is appreciably shorter than their tzaqis.," On the other hand, it is also possible that for low-mass stars circumstellar discs are eroded in a finite time that is appreciably shorter than their $t_{\rm ZAMS}$."
 For example. if for stars below we were to limit the integrals in 99 between 0 and the smaller of fyayjg and MMyr. the accreted mass would decrease proportionally to the mass. t.e. In practice this would still require some refinements to models of stellar evolution in the PMS phase. but the increase in mass would never be higher than 20MMPvr as the upper limit for the integration stems from the fact that near the MS we still find stars that have strong Ho excess. thus indicating the presence of considerable accretion.," For example, if for stars below we were to limit the integrals in 9 between 0 and the smaller of $t_{\rm ZAMS}$ and Myr, the accreted mass would decrease proportionally to the mass, i.e. In practice this would still require some refinements to models of stellar evolution in the PMS phase, but the increase in mass would never be higher than Myr as the upper limit for the integration stems from the fact that near the MS we still find stars that have strong $\alpha$ excess, thus indicating the presence of considerable accretion."
 Deeper observations of 3346 reaching well below on the MS can easily show whether the dises of lower- objects are indeed eroded before they reach the ZAMS., Deeper observations of 346 reaching well below on the MS can easily show whether the discs of lower-mass objects are indeed eroded before they reach the ZAMS.
 It would be sufficient to verify whether strong Ha excess emission is still shown by stars smaller than near the MS. or. alternatively. whether for objects of these lower masses strong Ho emission is only seen at younger ages. when they are still well above the MS.," It would be sufficient to verify whether strong $\alpha$ excess emission is still shown by stars smaller than near the MS, or, alternatively, whether for objects of these lower masses strong $\alpha$ emission is only seen at younger ages, when they are still well above the MS."
 If at these masses the dises are indeed eroded before reaching the ZAMS. Ho emission should systematically taper off for objects older than a cut-off age.," If at these masses the discs are indeed eroded before reaching the ZAMS, $\alpha$ emission should systematically taper off for objects older than a cut-off age."
 In any case. these new results can have important implications for some of the assumptions and boundary conditions used in theories of star formation and early evolution.," In any case, these new results can have important implications for some of the assumptions and boundary conditions used in theories of star formation and early evolution."
 We have studied the properties of the stellar populations in the field of the 3346 cluster in the SMC as observed with the ACS camera on board the Hubble Space Telescope (Sabbi et al., We have studied the properties of the stellar populations in the field of the 346 cluster in the SMC as observed with the ACS camera on board the Hubble Space Telescope (Sabbi et al.
 2007)., 2007).
 To this aim. we have employed a novel. self-consistent method that allows us to reliably identify PMS stars undergoing active mass accretion. regardless of their age.," To this aim, we have employed a novel, self-consistent method that allows us to reliably identify PMS stars undergoing active mass accretion, regardless of their age."
 The method. fully described in II. does not require spectroscopy and combines broad-band V and / photometry with narrow-band Ho imaging to identify all stars with excess Hai emission and allows us to derive the accretion luminosity Lace and mass accretion rate M. for all of them.," The method, fully described in I, does not require spectroscopy and combines broad-band $V$ and $I$ photometry with narrow-band $H\alpha$ imaging to identify all stars with excess $\alpha$ emission and allows us to derive the accretion luminosity $L_{\rm acc}$ and mass accretion rate $\dot M_{\rm acc}$ for all of them."
 The main results of this work can be summarised as follows., The main results of this work can be summarised as follows.
the top x-axis gives the turnofÉ masses corresponding to the progenitor ages in the lower x-axis.,the top x-axis gives the turnoff masses corresponding to the progenitor ages in the lower x-axis.
 The results clearly indicates that the magnitude of the gradient. slopes of both oxvgen and neon increase going [rom the older population PNe toward the voung ones. indicating a steepening of the a-element eradient. wilh Gime in the disk of the Milkv. Wav.," The results clearly indicates that the magnitude of the gradient slopes of both oxygen and neon increase going from the older population PNe toward the young ones, indicating a steepening of the $\alpha$ -element gradient with time in the disk of the Milky Way."
 The age intervals are not sharply defined. given the large uncertainties of data and models defining (he different PN Types. especially for the voung PNe. as discussed earlier in the paper.," The age intervals are not sharply defined, given the large uncertainties of data and models defining the different PN Types, especially for the young PNe, as discussed earlier in the paper."
 Nonetheless. given the scaling between ages and masses. choosing a lower limiting mass Of 3 or 4 even M. instead of 2 M. lor Type I PNe would not change the implications of Fig.," Nonetheless, given the scaling between ages and masses, choosing a lower limiting mass of 3 or 4 even $_{\odot}$ instead of 2 $_{\odot}$ for Type I PNe would not change the implications of Fig."
 8 al all., 8 at all.
 It is also worth noting that the dispersion in the log(O/II) along the Galactic disk varies with: PN tvpe. being smaller for Tvpe II PNe and larger both lor Types I and II] (see Table 3).," It is also worth noting that the dispersion in the log(O/H) along the Galactic disk varies with PN type, being smaller for Type II PNe and larger both for Types I and III (see Table 3)."
 The gradient evolution with (ime relies on the results shown in Figures + through 6., The gradient evolution with time relies on the results shown in Figures 4 through 6.
 While the Type II PN sample is well represented ancl unambiguouslv leads to the calculated gradient slope. the other PN Types should be discussed further.," While the Type II PN sample is well represented and unambiguously leads to the calculated gradient slope, the other PN Types should be discussed further."
 First. the Type ] PN gradient indeed seems to rely on a very flew points. especially at hieh galactocentric distances. with a single PN at 16 kpc skewing the slope.," First, the Type I PN gradient indeed seems to rely on a very few points, especially at high galactocentric distances, with a single PN at 16 kpc skewing the slope."
 On (he other haad. even excluding this Tvpe I PN from the gradient. calculation we still obtain a steep (negative) slope (although the dispersion is higher. \log(O/ID)/.NRe;=—0.046£0.041 dex +) confirming the steepening with time of gradients from Tvpe H (to Tvpe I PNe.," On the other hand, even excluding this Type I PN from the gradient calculation we still obtain a steep (negative) slope (although the dispersion is higher, $ {\rm \Delta~log(O/H) / \Delta~R_G=-0.046\pm0.041}$ dex $^{-1}$ ) confirming the steepening with time of gradients from Type II to Type I PNe."
 second. Ivpe III PNe represent the very old disk population. thus the question of eradient blurring due to radial migration occurs.," Second, Type III PNe represent the very old disk population, thus the question of gradient blurring due to radial migration occurs."
 However. in Figure 6 we see that no PN seenis to stand out as having a particularly high or low abundance for its location. which could be the case if migration had erased or flattened a significant gracient.," However, in Figure 6 we see that no PN seems to stand out as having a particularly high or low abundance for its location, which could be the case if migration had erased or flattened a significant gradient."
 Moreover. comparison of Figures 4. 5 and 6 shows that most Tvpe III PNe at Re< Tkpc have on average lower abundances (han those seen at the same radii in the Type I and Type II," Moreover, comparison of Figures 4, 5 and 6 shows that most Type III PNe at $_{\rm G}<$ 7kpc have on average lower abundances than those seen at the same radii in the Type I and Type II"
The spectrum of the H 14134117 lens system is presented in Fig.,The spectrum of the H 1413+117 lens system is presented in Fig.
 and shows no evidence of water maser emission., \ref{all-spectra} and shows no evidence of water maser emission.
 There was significant|. local RFT detected with the Effelsberg radio telescope between 6.219 and 6.231 GHz. which resulted in part of the spectrum between 600 and 1200 + from the systemic velocity of H 1413117 having to be removed.," There was significant local RFI detected with the Effelsberg radio telescope between 6.219 and 6.231 GHz, which resulted in part of the spectrum between 600 and 1200 $^{-1}$ from the systemic velocity of H 1413+117 having to be removed."
 The rms noise of he spectrum within a 10 kmss7 wide channel is 0.9 mJy., The rms noise of the spectrum within a 10 $^{-1}$ wide channel is 0.9 mJy.
 This corresponds to an upper limit of « 77700L.. (30) for an unlensed Water maser (assuming a lens magnification of ~ 110)., This corresponds to an upper limit of $<$ $L_\odot$ $\sigma$ ) for an unlensed water maser (assuming a lens magnification of $\sim$ 10).
 We note hat our observations are about half as sensitive as those presented by Wilneretal.(1999). but probe a much larger velocity range. out to around +22000kkmss+.," We note that our observations are about half as sensitive as those presented by \citet{wilner99}, but probe a much larger velocity range, out to around $\pm$ $^{-1}$."
 The region of the Effelsberg spectrum affected by RFI was also not investigated by Wilneretal.(1999) because their spectrum was confined to velocities between 633 and | . relative to our of definitionthe systemic velocity of the quasar.," The region of the Effelsberg spectrum affected by RFI was also not investigated by \citet{wilner99} because their spectrum was confined to velocities between $-$ 633 and $+$ $^{-1}$, relative to our definition of the systemic velocity of the quasar."
 Therefore. we also find no evidence of emission from this system.," Therefore, we also find no evidence of time-variable emission from this system."
 The spectrum of MS 1512-eB58 is presented in Fig., The spectrum of MS 1512–cB58 is presented in Fig.
 | and shows no evidence of clear water maser emission within + from the systemic velocity., \ref{all-spectra} and shows no evidence of clear water maser emission within $\pm$ $^{-1}$ from the systemic velocity.
 The rms noise of the spectrum is 1.2 mJy for a 10 ! wide channel., The rms noise of the spectrum is 1.2 mJy for a 10 $^{-1}$ wide channel.
 This corresponds to an upper limit of < 22300L.. (3) for the luminosity of any water maser emission within a 10 kmss+ wide channel., This corresponds to an upper limit of $<$ $L_\odot$ $\sigma$ ) for the luminosity of any water maser emission within a 10 $^{-1}$ wide channel.
 This calculation assumes a lensing magnification of 50., This calculation assumes a lensing magnification of 50.
 The spectrum of our final candidate water maser galaxy is also shown in Fig. |.., The spectrum of our final candidate water maser galaxy is also shown in Fig. \ref{all-spectra}.
 We find no water maser emission in the spectrum. which has an rms noise of O.9mmly within a wide channel.," We find no water maser emission in the spectrum, which has an rms noise of mJy within a $^{-1}$ wide channel."
 The upper-limit to the luminosity of water maser emission is therefore « 1]200L. (30) after correcting for the magnification from the cluster lens (~445)., The upper-limit to the luminosity of water maser emission is therefore $<$ $L_\odot$ $\sigma$ ) after correcting for the magnification from the cluster lens $\sim$ 45).
 Comparing with the recent observations of SMM 71635946612 by Edmondsetal.(009) and Wage&Momjian(2009).. we find that our spectrum is about half as sensitive as those taken with the EVLA.," Comparing with the recent observations of SMM J16359+6612 by \citet{edmonds09} and \cite{wagg09}, we find that our spectrum is about half as sensitive as those taken with the EVLA."
 The Effelsberg observation confirms that there is no time- emission from this system and probes a much larger range in velocity., The Effelsberg observation confirms that there is no time-variable emission from this system and probes a much larger range in velocity.
 The EVLA observations were processed with, The EVLA observations were processed with
quadrant of NGCdue 253 was performed by 2 (2s: s also 7. who detected | lain sequence stars in tle halo near the niünor axis of the galaxy. most likely associated with the outflow from the central starburst}.,"quadrant of NGC 253 was performed by \citeauthor{davidge10-n253} \citeyear{davidge10-n253}; see also \citealp{comeron-etal01}, who detected blue main sequence stars in the halo near the minor axis of the galaxy, most likely associated with the outflow from the central starburst)."
 He detected an extended fattened distribution of asvauptotic eiaut brauch (AGB) stars extending out to 13 kpe from the isk plane., He detected an extended flattened distribution of asymptotic giant branch (AGB) stars extending out to $13$ kpc from the disk plane.
" Based ou the similarity of the disk aud halo A-baud huinosity fictions. which implies au exteuded tar formation history. he coucluded that these stars were ay,eposited iuto the halo via the disruption of the galactic eSisk iu an interaction."," Based on the similarity of the disk and halo $K$ -band luminosity functions, which implies an extended star formation history, he concluded that these stars were deposited into the halo via the disruption of the galactic disk in an interaction."
 The data were not of sufficient epth to detect ROB stars. which would be far more munerous aud trace stellar populations of all ages.," The data were not of sufficient depth to detect RGB stars, which would be far more numerous and trace stellar populations of all ages."
" ""Theoretical models of galaxy formation iu a AC'DAL universe have been used to predict the properties of stellar halos. usually with the assumption that the stars hat are today found in the halo were accreted in he form of satellite ealaxics that were subsequently idally shredded (227?).."," Theoretical models of galaxy formation in a $\Lambda$ CDM universe have been used to predict the properties of stellar halos, usually with the assumption that the stars that are today found in the halo were accreted in the form of satellite galaxies that were subsequently tidally shredded \citep{bkw01,bj05,purcell-etal07,cooper-etal10}."
 These models predict that halos contain several percent of the total stellar mass of the ealaxy. and have very chuupy structure. particularly iu he outer reeious.," These models predict that halos contain several percent of the total stellar mass of the galaxy, and have very clumpy structure, particularly in the outer regions."
 They also predict significant galaxy-o-galaxy scatter in the total size ancl morphology of the halo structure (?2??)..," They also predict significant galaxy-to-galaxy scatter in the total size and morphology of the halo structure \citep{bj05,bell-etal08,johnston-etal08}."
 While these propertics qualitatively match observed galactic halos. quantitative ueasurelents of a larecr suuple of halos is required iu order to adequately test these models and distinenish )etween them.," While these properties qualitatively match observed galactic halos, quantitative measurements of a larger sample of halos is required in order to adequately test these models and distinguish between them."
 Qur goal is to test these models by obtaining resolved data of stellar halos in a larger. aud therefore more representativo. sample of galaxies.," Our goal is to test these models by obtaining resolved data of stellar halos in a larger, and therefore more representative, sample of galaxies."
 This will allow us to aceuratelv deteriune stellar halo structural properties. and to investigate their correlation with the elobal parameters of the associated ealaxics.," This will allow us to accurately determine stellar halo structural properties, and to investigate their correlation with the global parameters of the associated galaxies."
 The larger sample will help determine any scatter between svstenis due to a stochastic formation history., The larger sample will help determine any scatter between systems due to a stochastic formation history.
 Because of the large deeree of predicted aud observed substructure. obtaining an unbiased view of the halo requires a large field of view that is only possible usine grounud-based: cata.," Because of the large degree of predicted and observed substructure, obtaining an unbiased view of the halo requires a large field of view that is only possible using ground-based data."
 However. identification aud accurate photometry of stars in galaxies bevoud the Local Group is far superior with UST.," However, identification and accurate photometry of stars in galaxies beyond the Local Group is far superior with HST."
 We are therefore using a combination of TST and erouncd-based wide-field imagine., We are therefore using a combination of HST and ground-based wide-field imaging.
 Iu this paper. we present the first results of this progranunie. a study of NGC 253. along with a detailed deseriptiou of the analvsis required to turn the data iuto quantitative elobal measures of halo structure.," In this paper, we present the first results of this programme, a study of NGC 253, along with a detailed description of the analysis required to turn the data into quantitative global measures of halo structure."
 We adopt a distance modulus to NGC 253 of i)AL27.71 (D=3.418 Apc) throughout. as determined usine GIIOSTS (?).. aud consistent with but towards the upper end of distance estimates based ou a variety of inethods catalogued in the NASA Extragalactic (NED).," We adopt a distance modulus to NGC 253 of $m-M=27.71$ $D=3.48$ Mpc) throughout, as determined using GHOSTS \citep{ghosts11}, and consistent with but towards the upper end of distance estimates based on a variety of methods catalogued in the NASA Extragalactic (NED)."
" Note that we use the term ""halo"" to refer to the stellar halo of the galaxy. not the dark matter halo. uuless specifically stated otherwise."," Note that we use the term “halo” to refer to the stellar halo of the galaxy, not the dark matter halo, unless specifically stated otherwise."
 The ground-based data were obtained at the Magellan Ll (Baade} 6.511 telescope usiug the Inamori-\laecllau Areal Camera SSpectroeraph (INLACS). a mosaic of 8 2k x Uk CCDs. in f/2 configuration.," The ground-based data were obtained at the Magellan 1 (Baade) 6.5m telescope using the Inamori-Magellan Areal Camera Spectrograph (IMACS), a mosaic of $8$ 2k x 4k CCDs, in f/2 configuration."
 This results in a (072/pixel plate scale., This results in a $0\farcs2$ /pixel plate scale.
 The field was centered. 11 south of the ceuter of NGC 253. aud therefore iuchided a significant fraction of the galactic disk within the 21 muivignetted field of view.," The field was centered $14\arcmin$ south of the center of NGC 253, and therefore included a significant fraction of the galactic disk within the $24\arcmin$ unvignetted field of view."
 At the assumed clistance of NGC 253. the field extends 30 kpc from the galaxy center.," At the assumed distance of NGC 253, the field extends $30$ kpc from the galaxy center."
 Ou the nights of 17 Ls October 2009. we obtained 1800 seconds in CTIO {ας aud 2650 seconds in Bessel W-band.," On the nights of 17 -- 18 October 2009, we obtained $1800$ seconds in CTIO $I$ -band and $2650$ seconds in Bessel $V$ -band."
 A further 1000 seconds were obtained in V-baud on the night of 20 May 2010., A further $1000$ seconds were obtained in $V$ -band on the night of 20 May 2010.
 The seeciug full width half max (FWOAD was »proxinatelv 170 1”? diving the exposures., The seeing full width half max (FWHM) was approximately $1\farcs0$ – $1\farcs2$ during the exposures.
 Laudolt photometric standards were taken on these aud other welts during the same rus., Landolt photometric standards were taken on these and other nights during the same runs.
 Data were reduced using standard IRAF assisted dy scripts developed for reducing A\Larvlancd-Magellan Tunable Filterdata?.," Data were reduced using standard IRAF, assisted by scripts developed for reducing Maryland-Magellan Tunable Filter."
. Note that although the IMACS doctunentation suggests usiug the coluunu and row overscan regions distead of bias frames. we found that in the case of broadband iniaenmeg. saturated stars near the edees of the chips cau bleed iuto the overscan regious. making them unusable.," Note that although the IMACS documentation suggests using the column and row overscan regions instead of bias frames, we found that in the case of broadband imaging, saturated stars near the edges of the chips can bleed into the overscan regions, making them unusable."
 We therefore used bias frames that were obtained during the runs., We therefore used bias frames that were obtained during the runs.
 Sky subtraction was complicated by the presence of the ealactic disk iu the feld. which made it impossible to fit an accurate low-order surface to the sky.," Sky subtraction was complicated by the presence of the galactic disk in the field, which made it impossible to fit an accurate low-order surface to the sky."
 We therefore determined the sky level using a two-dimensional median filter of width 101 pixels: this width was chosen to be aree chough that photometry tests on individual stars ound no bias. but small enough that the skv was flat ο within 1. 6i large scales.," We therefore determined the sky level using a two-dimensional median filter of width $101$ pixels; this width was chosen to be large enough that photometry tests on individual stars found no bias, but small enough that the sky was flat to within $1\%$ on large scales."
 This results iu eood image quality over most of the frame. but clearly iuvalidates the xortiou of the field that iucludes the disk of NGC 253 which coutains smooth features on scales much larger hau the filter width.," This results in good image quality over most of the frame, but clearly invalidates the portion of the field that includes the disk of NGC 253, which contains smooth features on scales much larger than the filter width."
 However. as the goal of this work is to study the halo of the ealaxy. this does uot affect our results;," However, as the goal of this work is to study the halo of the galaxy, this does not affect our results."
 An astrometric solution over the ficld was computed using the USNO A2.0 catalog (?).., An astrometric solution over the field was computed using the USNO A2.0 catalog \citep{usno-a2}.
 A third order polvionial solution was required in order that the astrometric residuals did not vary systematically across the field: the resulting solution achieved an riis scatter of 072L. comparable to the pixel size.," A third order polynomial solution was required in order that the astrometric residuals did not vary systematically across the field; the resulting solution achieved an rms scatter of $0\farcs24$ , comparable to the pixel size."
 Two sample reeious of the final stacked aud sky-ubtracted L-band image are shown in Figure 1: one located ~9/5 (10 kpe} frou the center of the ealaxyv auc oue located at ~22 (22 kpe). both approximately along the minor axis.," Two sample regions of the final stacked and sky-subtracted $I$ -band image are shown in Figure \ref{fig:field}: one located $\sim 9\farcm5$ $10$ kpc) from the center of the galaxy and one located at $\sim 22\arcmin$ $22$ kpc), both approximately along the minor axis."
 The faint poiut sources that are abundant iu the inner field but almost completely absent iu the outer field are RGB stars in the halo of NGC 253., The faint point sources that are abundant in the inner field but almost completely absent in the outer field are RGB stars in the halo of NGC 253.
 We ignore the very outermost part of the mosaic. which is conrpromused by the shadow of the guide star camera. whose location varied between exposures.," We ignore the very outermost part of the mosaic, which is compromised by the shadow of the guide star camera, whose location varied between exposures."
 We use data from the IIST/CGIIOSTS survev (27) overlapping with our INLAC'S pointing for the purposes of carrvineout our artificial star tests. and to test our juterred stellar deusities Gu the high stellar deusity," We use data from the HST/GHOSTS survey \citep{dejong-etal07,ghosts11}
 overlapping with our IMACS pointing for the purposes of carryingout our artificial star tests, and to test our inferred stellar densities (in the high stellar density"
pair potential fitted to the gas gun (27.47) data (Ross. Ree Young 1933).,"pair potential fitted to the gas gun $(P,V)$ data (Ross, Ree Young 1983)."
 All four EOS presented here (hus produce the same molecular Hugoniot., All four EOS presented here thus produce the same molecular Hugoniot.
 The One Component Plasma (OCP) forms the basis of the metallic {hid [ree energy. with a zero-temperature electron EOS including exchange and correlation (Ilolmesetal.1995).," The One Component Plasma (OCP) forms the basis of the metallic fluid free energy, with a zero-temperature electron EOS including exchange and correlation \citep{holmes95}."
". An entropy term (Fi=0.42. where / is (he Boltzmann constant) is introduced in the metallic [ree energy. and acljusted to fit the reshock temperature data (0,——2.7)."," An entropy term $F_{\rm fit}=\delta_e kT$, where $k$ is the Boltzmann constant) is introduced in the metallic free energy and adjusted to fit the reshock temperature data $\delta_e=-2.7$ )."
" This model was found to agree well with the subsequent NOVA IHHuegoniot (Collinsetal.1998:DaSilva1997) with the same value οἱ ὁ,."," This model was found to agree well with the subsequent NOVA Hugoniot \citep{collins98,dasilva97} with the same value of $\delta_e$."
 The LM model is particularly useful for the present study as it can be easily modified to fit various data sets., The LM model is particularly useful for the present study as it can be easily modified to fit various data sets.
 We found (hat the original LM model (Ross1998.1999) predicts an anomalous acliabat for Jupiter.," We found that the original LM model \citep{ross98,ross99} predicts an anomalous adiabat for Jupiter."
 This can be seen in a sequence of adiabats with decreasing specific entropy (Fig. 3))., This can be seen in a sequence of adiabats with decreasing specific entropy (Fig. \ref{fig:J_ad_Ross}) ).
 The turnover of the acliabat around Mbar has been reported. before (Nellis. Ross llolmes 1995) and is a direct consequence of fitting the reshock temperature measurements.," The turnover of the adiabat around $\,$ Mbar has been reported before (Nellis, Ross Holmes 1995) and is a direct consequence of fitting the reshock temperature measurements."
 At higher pressures. a cusp in the adiabat develops at low entropies.," At higher pressures, a cusp in the adiabat develops at low entropies."
 While this behavior is not strictly prohibited (hermmodvnamically. it is very suspicious.," While this behavior is not strictly prohibited thermodynamically, it is very suspicious."
" Ht arises from (he constant o,k shift in entropy of (he metallie adiabats in the model.", It arises from the constant $\delta_e k$ shift in entropy of the metallic adiabats in the model.
 This causes a mismatch between (he molecular and metallic acliabats that is significant when the entropy becomes small enough. as in Jupiter and Saturn.," This causes a mismatch between the molecular and metallic adiabats that is significant when the entropy becomes small enough, as in Jupiter and Saturn."
 This means that while the original form of the LM model reproduces ihe NOVA IIugoniot quite well. it is inadecquate for EOS calculations at temperatures well below that of (he principal IHugoniot.," This means that while the original form of the LM model reproduces the NOVA Hugoniot quite well, it is inadequate for EOS calculations at temperatures well below that of the principal Hugoniot."
 We computed four LM EOS that produce better-behaved. acliabats:, We computed four LM EOS that produce better-behaved adiabats:
hose is the week field approximation for the magnetism. aa Zeeman split of less then the intrinsic line therma and pressure broadening. which is. however. safely applicable to a field of onwt a few Kilogauss in a main sequence B star.,"those is the week field approximation for the magnetism, a Zeeman split of less then the intrinsic line thermal and pressure broadening, which is, however, safely applicable to a field of only a few kilogauss in a main sequence B star."
" A concern prompting ""ur more caution when interpreting the derived numbers is that the spectra are. in fact. an average over the stellar disk longitudinal field component. without considering tthe limb darkening (See?.orafulldiscussionofthelimitationsandcaveats method).."," A concern prompting far more caution when interpreting the derived numbers is that the spectra are, in fact, an average over the stellar disk longitudinal field component, without considering the limb darkening \citep[See][for
 a full discussion of the limitations and caveats of the
 method]{2002A&A...389..191B}."
 We used a 47 minimization approach to fit the linear model Vl=1X+ to the observational data: this model is based on eqn., We used a $\chi^2$ minimization approach to fit the linear model $V/I=A+\langle B_z \rangle \times x$ to the observational data; this model is based on eqn.
 1. with the addition of a constant term st.ῇ)., \ref{eqn_reg} with the addition of a constant term $A$ .
. The formal errors of 1//. which are derived from the wel known photon- and detector noise characteristics. were used to assign weights for the 47 fit.," The formal errors of $V/I$, which are derived from the well known photon- and detector noise characteristics, were used to assign weights for the $\chi^2$ fit."
 The derived (2.5 are given in Table I.., The derived $\langle B_z \rangle$ are given in Table \ref{tab_result}.
 Given recent concerns expressed overthe error bars associated with FORSI magnetic field measurements (see.e.g.2).. we have exercised particular care in deriving the errors listed in Table |..," Given recent concerns expressed over the error bars associated with FORS1 magnetic field measurements \citep[see, e.g.][]{2009MNRAS.398.1505S}, we have exercised particular care in deriving the errors listed in Table \ref{tab_result}."
 Applying the bootstrap Monte Carlo approach described by ?.. we used the observational data from each epoch to construct one million corresponding synthetic datasets.," Applying the bootstrap Monte Carlo approach described by \citet{1992nrfa.book.....P}, we used the observational data from each epoch to construct one million corresponding synthetic datasets."
 We then applied the same fitting approach as above to obtain et and £D.) values for every synthetic dataset., We then applied the same fitting approach as above to obtain $A$ and $\langle B_z \rangle$ values for every synthetic dataset.
 These values exhibit a distribution about those derived from the actual observations: Fig., These values exhibit a distribution about those derived from the actual observations; Fig.
 2. shows this distribution or the 5546669.330 observations. both as a 2-dimensional orobability density map. and as a |-dimensional probability density ‘unction in (D2.5.," \ref{fig_FORSB} shows this distribution for the 669.330 observations, both as a 2-dimensional probability density map, and as a 1-dimensional probability density function in $\langle B_z \rangle$."
 Plotted over the map are contours of constant X7 hat enclose68.3%... and of the synthetic datasets: ikewise. plotted over the probability density function are the symmetric bounds enclosing68.3%... and of the synthetic datasets.," Plotted over the map are contours of constant $\chi^2$ that enclose, and of the synthetic datasets; likewise, plotted over the probability density function are the symmetric bounds enclosing, and of the synthetic datasets."
 The bounds for the case are assigned as he error bars quoted in Table 1: because the probability density unction in Fig., The bounds for the case are assigned as the error bars quoted in Table \ref{tab_result}; because the probability density function in Fig.
 2. is close to Gaussian (and likewise for the data Tom other epochs). these error bars can be regarded as approximate 1-7 contidence intervals.," \ref{fig_FORSB} is close to Gaussian (and likewise for the data from other epochs), these error bars can be regarded as approximate $\sigma$ confidence intervals."
 Note. however. that we have not made any specific assumptions regarding the propagation of errors through our modeling process: the Monte Carlo approach naturally results in error estimates that directly reflect the combined characteristics of the observations and the modeling.," Note, however, that we have not made any specific assumptions regarding the propagation of errors through our modeling process; the Monte Carlo approach naturally results in error estimates that directly reflect the combined characteristics of the observations and the modeling."
 Next to the Stokes V. spectra. we examined the intensity spectra observed by FORSI.," Next to the Stokes $V$ spectra, we examined the intensity spectra observed by FORS1."
 Due to the very high S/N and the large rotational velocity. even the relatively low dispersion of FORSI allows high quality equivalent width CEW) measurements of the spectral lines.," Due to the very high $S/N$ and the large rotational velocity, even the relatively low dispersion of FORS1 allows high quality equivalent width (EW) measurements of the spectral lines."
 The lines. with the exception of H7. show little variability. except in the very core.," The lines, with the exception of $\beta$ , show little variability, except in the very core."
 The spectra taken at MJD—54656.146 and 54669.186 show enhanced core absorption., The spectra taken at MJD=54656.146 and 54669.186 show enhanced core absorption.
 At these epochs the D.) had null-detections. which is in agreement with the crossing of the magnetic equator through the line of sight. because this is the time at which a corotating cloud ofcircumstellar matter. located at the crossing of the magnetic and rotationalequators. is expected to pass in front of the star.," At these epochs the $\langle B_z \rangle$ had null-detections, which is in agreement with the crossing of the magnetic equator through the line of sight, because this is the time at which a corotating cloud ofcircumstellar matter, located at the crossing of the magnetic and rotationalequators, is expected to pass in front of the star."
declining. trend until about LJ) 3540. followed. bv. two shorter duration Lares similar to those seen in the 2000 and 2002 outbursts (Wachter et al.,declining trend until about HJD 3540 followed by two shorter duration flares similar to those seen in the 2000 and 2002 outbursts (Wachter et al.
 2000: Wijnanels et al., 2000; Wijnands et al.
 2001. Wijnands. 2004).," 2001, Wijnands, 2004)."
 There was also a large excess (~0.45 mag) inthe band but not in the other colours on HJD 3527., There was also a large excess $\sim 0.45$ mag) in the band but not in the other colours on HJD 3527.
 |ecay time-scale during the first outburst (LTD 3522 to 3540) was ~Y davs for both the optical anc| the X-ray enmussion., The decay time-scale during the first outburst (HJD 3522 to 3540) was $\sim 7$ days for both the optical and the X-ray emission.
 This is very similar to that during the discovery outburst in 1998 (Ciles et al., This is very similar to that during the discovery outburst in 1998 (Giles et al.
 1999) but. much shorter than the ~22 cay time-scale for N'TE 0929.314 (Giles et al., 1999) but much shorter than the $\sim 22 $ day time-scale for XTE J0929–314 (Giles et al.
 2005)), 2005).
 Visual inspection of the light. curves suggests that there is à weak correlation between the optical and X-ray Hluxes with the optical emission. preceding the X-ray by a [Kew days as in the 1998 outburst (Giles et al., Visual inspection of the light curves suggests that there is a weak correlation between the optical and X-ray fluxes with the optical emission preceding the X-ray by a few days as in the 1998 outburst (Giles et al.
 1999)., 1999).
 To test this hypothesis we converted the optical magnitudes to fluxes normalised to the peak ASAL count rate and calculated the sum of the squares of the dillerences between the the two Iluxes for delays varving between -10 and |10 days., To test this hypothesis we converted the optical magnitudes to fluxes normalised to the peak ASM count rate and calculated the sum of the squares of the differences between the the two fluxes for delays varying between -10 and +10 days.
 There is a broad. minimum. for the whole dataset. corresponding to the optical emission leading the rav by -1 to 13 cays.," There is a broad minimum, for the whole dataset, corresponding to the optical emission leading the X-ray by -1 to +3 days."
 When this analysis is restricted to the data during the secondary outbursts commencing at I1JD 3542 there is no significant minimum., When this analysis is restricted to the data during the secondary outbursts commencing at HJD 3542 there is no significant minimum.
 This lack of correlation may well be a consequence of inadequate sampling and the low statistical significance of the ASAI cata at these times., This lack of correlation may well be a consequence of inadequate sampling and the low statistical significance of the ASM data at these times.
 In Fig., In Fig.
 3 we plot the and colour index changes during the observations., 3 we plot the and colour index changes during the observations.
 Where two measurements in the same colour were mace on the same night we use the mean of the two., Where two measurements in the same colour were made on the same night we use the mean of the two.
 Phe solid line represents a linear fit to the data during the first steady decline in intensity ending at. 111) 3540., The solid line represents a linear fit to the data during the first steady decline in intensity ending at HJD 3540.
 The anomalous point at ILJD. 3527 is discussed. in section 3.3., The anomalous point at HJD 3527 is discussed in section 3.3.
 Neither it nor the points measured after 1191) 3540 are included in the linear fits., Neither it nor the points measured after HJD 3540 are included in the linear fits.
 Many oof our nightly: measurements were mace consecutively ancl therefore cluster within a binary phase duration of «0.25 but this was not always the case., Many of our nightly measurements were made consecutively and therefore cluster within a binary phase duration of $< 0.25$ but this was not always the case.
 In the absence of a continuous set of measurements we are unable to confirm the continued existence of the 0.12 mag binary modulation in observed during the 1998 outburst by Ciles et al. (, In the absence of a continuous set of measurements we are unable to confirm the continued existence of the 0.12 mag binary modulation in observed during the 1998 outburst by Giles et al. (
1999).,1999).
 Larger. modulations are reported during the quiescent state (Campana ct al., Larger modulations are reported during the quiescent state (Campana et al.
 1998) but no modulation value is known for either state., 1998) but no modulation value is known for either state.
 We have taken the Campana ct al., We have taken the Campana et al.
 values (scaled down to the 1998 outburst. modulation). assumed a modulation a little smaller than atV. and then used the precise ephemoeris of Chakrabarty Morgan (1998) to examine the and shifts expected if any such mocdulations were actually esent.," values (scaled down to the 1998 outburst modulation), assumed a modulation a little smaller than at, and then used the precise ephemeris of Chakrabarty Morgan (1998) to examine the and shifts expected if any such modulations were actually present."
 We find that for the introduced: simulated mocdulation ‘errors’ are «0.05 mag., We find that for the introduced simulated modulation 'errors' are $< 0.05$ mag.
 The scatter in the lower xuel of fig 3 indicates that a modulation of this order may » present although we have no positive indication for it., The scatter in the lower panel of fig 3 indicates that a modulation of this order may be present although we have no positive indication for it.
 For he data the modulation “errors” are perhaps twice the size for and of course rather dependent on the assumed modulation., For the data the modulation 'errors' are perhaps twice the size for and of course rather dependent on the assumed modulation.
 However. the small scatter about the inear fit in the upper panel of fig 3 suggests that minimal modulation was actually present.," However, the small scatter about the linear fit in the upper panel of fig 3 suggests that minimal modulation was actually present."
 The most obvious features in both plots is the large increase due to the [lare in on ILJID 3527 :uxl two smaller but significant increases on LLJD 3624 and y3546., The most obvious features in both plots is the large increase due to the flare in on HJD 3527 and two smaller but significant increases on HJD 3624 and 3546.
 We discuss these anomalies below., We discuss these anomalies below.
 Overall increased linearly (became redder) until the intensity had decreased by about two magnitudes at the end of the primary outburst at LLJD 3540., Overall increased linearly (became redder) until the intensity had decreased by about two magnitudes at the end of the primary outburst at HJD 3540.
 Εις is illustrated by the solid lines in Fig., This is illustrated by the solid lines in Fig.
 3., 3.
" The fact that the decline in emission is greatest in the blue end. of the spectrum is suggestive of a viscosity driven. ""inside out” transition leading to a decline in inner disc emission.", The fact that the decline in emission is greatest in the blue end of the spectrum is suggestive of a viscosity driven “inside out” transition leading to a decline in inner disc emission.
The presence of soft N-ray cinission (either thermal or photoionized) on a 22 kpe scale is interesting.,The presence of soft X-ray emission (either thermal or photoionized) on a $\gtrsim$ 2 kpc scale is interesting.
 If the X- cluission originates from hot eas heated by the AGN outflow. it would be strong evidence for AGN feedback to the ISM on galaxy-sceales. reseiibliug the larger scale AGN feedback from radio bubbles interacting with the intracluster media secu in the Perseus cluster2003).," If the X-ray emission originates from hot gas heated by the AGN outflow, it would be strong evidence for AGN feedback to the ISM on galaxy-scales, resembling the larger scale AGN feedback from radio bubbles interacting with the intracluster medium seen in the Perseus cluster."
. Iu the photoiouized scenario. cousidering the lieht travel time from the nucleus to the ΤΗ eas lanes (10! yy) and the photoionization recombination timescale tc.220007/102)πο s1997).. we cam constraiu the timescale when NGC L151 was experiencing such au Eddinetou-limit outburst pliase to be teay<2.5«10! vr ago.," In the photoionized scenario, considering the light travel time from the nucleus to the HI gas lanes $10^4$ yr) and the photoionization recombination timescale $t_{rec}\approx 200 (n/10^9)^{-1}
(T/10^5)^{0.7}$ s, we can constrain the timescale when NGC 4151 was experiencing such an Eddington-limit outburst phase to be $t_{Edd}< 2.5\times 10^4$ yr ago."
 Otherwise the N-rav eas would no longer cuit recombination lines., Otherwise the X-ray gas would no longer emit recombination lines.
 If the soft diffuse cussion is due to shock heating associated with a nuclear outflow instead. the current Lag—42xLoh (Kaspi et al.," If the soft diffuse emission is due to shock heating associated with a nuclear outflow instead, the current $L_{bol}=7.3\times 10^{43}$ (Kaspi et al."
 2005) indicates LuaífLggc0.01., 2005) indicates $L_{bol}/L_{Edd}\sim 0.01$.
" Assunmüng that ~5% of the power iis gone to create the IIT cavity2005)... ~[10 vear of such ACN heating is needed to produce the thermal cucrey content of the rot eas (Ey,c3s10°! erg)."," Assuming that $\sim 5\%$ of the power has gone to create the HI cavity, $\sim 4\times 10^4$ year of such AGN heating is needed to produce the thermal energy content of the hot gas $E_{th}\sim 3\times 10^{54}$ erg)."
 This is much shorter than he current cooling time of the lot gas (7.~105 xr)., This is much shorter than the current cooling time of the hot gas $\tau_c\sim 10^8$ yr).
" The efficiency at which the ACN outflow deposits its ducmatic energv in the ISM cau be as low as 103, xovided that the hot eas is confined aud the duration of he strong nuclear outflow Is 7,4:promEΤε."," The efficiency at which the AGN outflow deposits its kinematic energy in the ISM can be as low as $10^{-4}$, provided that the hot gas is confined and the duration of the strong nuclear outflow is $\tau_{outflow}\ll \tau_c$."
 Uowever. perpendicular to the plane. the gas could expand aud cool efficiently via adiabatic expansion uuless some other confinement exists.," However, perpendicular to the plane, the gas could expand and cool efficiently via adiabatic expansion unless some other confinement exists."
 Asstmineg an outflow velocity of Voutflow~105 lan + for the free expuusiou perpendicular to the disk (c£., Assuming an outflow velocity of $_{outflow}\sim 10^3$ km $^{-1}$ for the free expansion perpendicular to the disk (c.f.
" thermal velocity dispersion Ve,=Vieti,~LOO bins 1) the relevant timescale is the time span in which the hot gas expands to the scale height above disk plane AR/VontflomcLO? ντ,"," thermal velocity dispersion $_{th}=\sqrt{k_BT/m_p}\sim 100$ km $^{-1}$ ), the relevant timescale is the time span in which the hot gas expands to the scale height above disk plane $\Delta R/{\rm v}_{outflow}\sim 10^5$ yr."
 This timescale again tuples that the timescale to the last AGN outflow heating must be <LO? vr. in order to replenish the N-rav gas expanding out of the plane that is cooling rapidly.," This timescale again implies that the timescale to the last AGN outflow heating must be $<10^5$ yr, in order to replenish the X-ray gas expanding out of the plane that is cooling rapidly."
 Compared to the estimated AGN lifetime (10*.—105 vis). the fact that we observe the ACN-host interaction is intriguing eiven the short timescale to the last episode of activity.," Compared to the estimated AGN lifetime $10^7 - 10^8$ yrs), the fact that we observe the AGN-host interaction is intriguing given the short timescale to the last episode of activity."
 Such episodic outbursts must occur frequently2010).. and our finding implies they may occupy =21% of the ACN lifetime.," Such episodic outbursts must occur frequently, and our finding implies they may occupy $\gtrsim$ of the AGN lifetime."
 To sununarize. we have discovered soft (0.31 keV) diffuse N-rav cunission (LosopecLO ere 13 in the ceutral —2 kpe of NGC. 1151. which cannot be attributed ο PSF scattering or clectron scattering of the nuclear Cluission. or the inteerated ciuission from unresolved zdut point sources.," To summarize, we have discovered soft (0.3–1 keV) diffuse X-ray emission $L_{0.5-2 keV}\sim 10^{39}$ erg $^{-1}$ ) in the central $\sim$ 2 kpc of NGC 4151, which cannot be attributed to PSF scattering or electron scattering of the nuclear emission, or the integrated emission from unresolved faint point sources."
 If the eas is of photoionized origin and represcuts he relic of past ACN activities. the nucleus of NGC 1151 aust have experienced an Eddinetou-limited biel uninositv phase.," If the gas is of photoionized origin and represents the relic of past AGN activities, the nucleus of NGC 4151 must have experienced an Eddington-limited high luminosity phase."
" Alternatively, ACN outflows may rave inechanically heated the eas to N-rav αμτις cluperature (AD=0.25 keV)."," Alternatively, AGN outflows may have mechanically heated the gas to X-ray emitting temperature $kT=0.25$ keV)."
 This hoto eas could be confined in the ΤΠ cavity by the thermal pressure of the III eas and dynamic pressure of iufalliug gas along the arge-scale stellar bar., This hot gas could be confined in the HI cavity by the thermal pressure of the HI gas and dynamic pressure of infalling gas along the large-scale stellar bar.
 For both scenarios. the ΑΝμον interaction iust rave occured relatively recently.," For both scenarios, the AGN-host interaction must have occured relatively recently."
 For the ACN outflow reating to work. the deposit of mechanical energy. niust rave happened Z10? xy ago to replenish the lot eas. which is expanding out of the plane uuless prevented x other confining mechanism.," For the AGN outflow heating to work, the deposit of mechanical energy must have happened $\lesssim 10^5$ yr ago to replenish the hot gas, which is expanding out of the plane unless prevented by other confining mechanism."
 Whereas for photoionized eas from a past AGN outburst. the timescale to the lightly active pliase inst be <2.5«10! y ago.," Whereas for photoionized gas from a past AGN outburst, the timescale to the highly active phase must be $\lesssim 2.5\times 10^4$ yr ago."
 This short iuescale to the last episode of high activity phase may imply such outbursts occupy Z of the ACN lifetime., This short timescale to the last episode of high activity phase may imply such outbursts occupy $\gtrsim$ of the AGN lifetime.
 We thank the anouvious referee for helpful comunucuts., We thank the anonymous referee for helpful comments.
 This work is partially supported from NASA eraut GOs- and NASA Contract. NASS-39073 (CNC)., This work is partially supported from NASA grant GO8-9101X and NASA Contract NAS8-39073 (CXC).
 CD acknowledges support from DFC erauts. SCIT 536/1-1 and SCID 536/1-2 as part of the SPP 1177., GD acknowledges support from DFG grants SCH 536/4-1 and SCH 536/4-2 as part of the SPP 1177.
"Simple radial surlace-brightness plots for the 16 objects chosen for modelling reveals that there is sullicient. signal-to-noise to allow fitting to a tvpical radius of c4"".",Simple radial surface-brightness plots for the 16 objects chosen for modelling reveals that there is sufficient signal-to-noise to allow fitting to a typical radius of $\simeq4\asec$.
 It is therefore necessary to have a point spread function (PSE) to convolve our model galaxies with which has good signal-to-noise out to at least this radius., It is therefore necessary to have a point spread function (PSF) to convolve our model galaxies with which has good signal-to-noise out to at least this radius.
" Due to the synthetic PSE's produced. by the software package (Ixrist. 1995) being unable to reproduce the WEPC2’s scattered light halo outside a radius of 21.5"". it is clear that an empirical PSE is required for the modelling of these data."," Due to the synthetic PSF's produced by the software package (Krist 1995) being unable to reproduce the WFPC2's scattered light halo outside a radius of $\simeq 1.5\asec$, it is clear that an empirical PSF is required for the modelling of these data."
 A glance at Table l reveals that all of the exposures utilised here were image through either the E785LP or ESIAW filters., A glance at Table \ref{bigtab} reveals that all of the exposures utilised here were imaged through either the F785LP or F814W filters.
 Fhis require the acquisition of two relatively deep PSE's imaged with the correct filter/chip combination. which were also locatec close to the average chip position of the 3C1 galaxies. in order to avoid the noticeable positional variation of the WEPC2 DSP.," This required the acquisition of two relatively deep PSF's imaged with the correct filter/chip combination, which were also located close to the average chip position of the 3CR galaxies, in order to avoid the noticeable positional variation of the WFPC2 PSF."
 Xn interrogation of the LIST PSE search too ooduced. disappointing results. with all available PSL's either too faint or too small in angular extent.," An interrogation of the HST PSF search tool produced disappointing results, with all available PSF's either too faint or too small in angular extent."
 Fortunately. two suitable stars were present. on. the exposures of 968 (FTSSLP) and 3€239 (ESIAW).," Fortunately, two suitable stars were present on the exposures of 3C41 (F785LP) and 3C239 (F814W)."
 Due to the need. for sullicient. depth in the PSE wings. both of these stars had. saturated: cores in even the shortes exposures.," Due to the need for sufficient depth in the PSF wings, both of these stars had saturated cores in even the shortest exposures."
 To overcome this problem a mocified: version of the PSE re-sampling technique described in. MeLurea£ 1999b was implemented., To overcome this problem a modified version of the PSF re-sampling technique described in McLure 1999b was implemented.
" Both PSIswere only. saturate within a radius of <0.3"" of their core. well inside the radius where can accurately reproduce the empirical PSP."," Both PSF'swere only saturated within a radius of $\le0.3\asec$ of their core, well inside the radius where can accurately reproduce the empirical PSF."
" Alaking use of this. the two PSE's had their core replace: with the equivalent. model. the relative scaling xung determined by matching the [lux in an annulus AnotherH""0.7"""," Making use of this, the two PSF's had their core replaced with the equivalent model, the relative scaling being determined by matching the flux in an annulus between $0.4\asec<r<0.7\asec$."
" advantage of this ⋅⋠ approach is that due to ability to produce model SES at up to 50 times oversampling. sub-pixel centring of the PSF can be matched to that of the galaxies to an accuracy of <0.005""."," Another advantage of this approach is that due to ability to produce model PSF's at up to 50 times oversampling, sub-pixel centring of the PSF can be matched to that of the galaxies to an accuracy of $\leq 0.005\asec$."
 Although not as crucial as with the quasar host work it is still important to overcome the severe unclersampline of WEPC2. especially as several of the 3CR galaxies modelled proved to have a substantial point-source contribution.," Although not as crucial as with the quasar host work it is still important to overcome the severe undersampling of WFPC2, especially as several of the 3CR galaxies modelled proved to have a substantial point-source contribution."
 The modelling of these objects proceeded. in. an identical fashion to that of the z20.2 AGN sample studied. by AleLure (1999a) and Dunlop (1999)., The modelling of these objects proceeded in an identical fashion to that of the $z\simeq0.2$ AGN sample studied by McLure (1999a) and Dunlop (1999).
 Lach object was modellecl with both a standard. Freeman disc template and de Vaucouleurs £F7 prolile. with no priori assumptions being made about position angle. axial ratio or," Each object was modelled with both a standard Freeman disc template and de Vaucouleurs $r^{1/4}$ profile, with no assumptions being made about position angle, axial ratio or"
photosphere and heat and “evaporate” the chromospheric plasma. providing the hot material that emits the X-ray flare observed between 0.1 and 10 keV. Indeed. the flares for which long flaring loops were derived in Favataetal.(2005a) appear to take place in CTTs with no significant ongoing accretion (hence their “accretion” tubes are not mass-loaded).,"photosphere and heat and “evaporate” the chromospheric plasma, providing the hot material that emits the X-ray flare observed between 0.1 and 10 keV. Indeed, the flares for which long flaring loops were derived in \citet{ffr+2005} appear to take place in CTTs with no significant ongoing accretion (hence their “accretion” tubes are not mass-loaded)."
 In YSOs. which are still strongly accreting. the density of material confined in the accretion tubes 1s expected to be high.," In YSOs, which are still strongly accreting, the density of material confined in the accretion tubes is expected to be high."
" For a fiducial case of mass accretion rate of 1077M. yr! Muzerolleetal.(2001) estimate densities of my=10—10"" cm at 2R.. that is. column densities of the order of ΔΝ~1074 em"" assuming a stellar radius of 10!! em."," For a fiducial case of mass accretion rate of $10^{-7} M_{\odot}$ $^{-1}$, \citet{mch01} estimate densities of $n_{\rm H} = 10^{12} - 10^{13}$ $^{-3}$ at $\sim 2 {\rm R_{*}}$, that is, column densities of the order of $N \sim 10^{23} - 10^{24}$ $^{-2}$ assuming a stellar radius of $^{11}$ $^{-2}$."
 In Elias 29 with an accretion luminosity of 28.8 Lo. the accretion rate is estimated to be one order of magnitude higher (Nattaetal.. 2006)). so the column density within the accretion tubes will be correspondingly higher.," In Elias 29 with an accretion luminosity of 28.8 $L_{\sun}$ , the accretion rate is estimated to be one order of magnitude higher \citealp{nts06}) ), so the column density within the accretion tubes will be correspondingly higher."
 With such a high density. the accelareted electrons produced by a reconnection event in the accretion tube cannot reach the stellar photosphere. but will be decelerated in situ. producing non-thermal (hard) X-ray emission and collisional ionisation of elements such as Fe. which will be largely neutral at a typical temperature of ~7000 K (Muzerolleetαἰ.. 2001)).," With such a high density, the accelareted electrons produced by a reconnection event in the accretion tube cannot reach the stellar photosphere, but will be decelerated in situ, producing non-thermal (hard) X-ray emission and collisional ionisation of elements such as Fe, which will be largely neutral at a typical temperature of $\sim 7000$ K \citealp{mch01}) )."
 A sketch of the Ellas 29 star-disk system with the hypothesised accretion tube where the electron acceleration is taking place is shown in reffig:el29.., A sketch of the Elias 29 star-disk system with the hypothesised accretion tube where the electron acceleration is taking place is shown in \\ref{fig:el29}.
 To explain the observed long-lasting fluorescent emission. the electron acceleration mechanism should be sustained over several days.," To explain the observed long-lasting fluorescent emission, the electron acceleration mechanism should be sustained over several days."
 While we have no detailed physical mechanism to propose for this. we note that many of the long-lasting flares observed in YSOs by Favataetal.(2005a) require sustained heating of the flaring material to explain the observed duration of the events. which in some cases lasted a few days.," While we have no detailed physical mechanism to propose for this, we note that many of the long-lasting flares observed in YSOs by \citet{ffr+2005} require sustained heating of the flaring material to explain the observed duration of the events, which in some cases lasted a few days."
 The thermodynamic cooling time of the flaring plasma would be much shorter: therefore. to explain the observed long flaring events in YSOs. magnetic reconnection must be ongoing. on time scales of days. similar to what is observed for the Fe 6.4 keV fluorescent emission from Elias 29.," The thermodynamic cooling time of the flaring plasma would be much shorter; therefore, to explain the observed long flaring events in YSOs, magnetic reconnection must be ongoing, on time scales of days, similar to what is observed for the Fe 6.4 keV fluorescent emission from Elias 29."
 The same electron acceleration mechanism could therefore result in long-lasting flares. if it takes place in systems with little ongoing accretion. or in strong Fe 6.4 keV emission. if it takes place in systems with strong ongoing accretion.," The same electron acceleration mechanism could therefore result in long-lasting flares, if it takes place in systems with little ongoing accretion, or in strong Fe 6.4 keV emission, if it takes place in systems with strong ongoing accretion."
 In their study. Emslieetal.(1986) provide a relation (Fig.," In their study, \citet{epd86} provide a relation (Fig."
 | in their paper) between the flux emitted in the 6.4 keV line in the case of electron ionisation. as a function of the column density of material between the electrons’ injection point and the “cold” Fe (where Fe 6.4 keV emission can take place).," 1 in their paper) between the flux emitted in the 6.4 keV line in the case of electron ionisation, as a function of the column density of material between the electrons' injection point and the “cold” Fe (where Fe 6.4 keV emission can take place)."
 This column density is zero in our case. since. as mentioned above. the material in the accretion tubes is mostly neutral at temperatures of ~7000 K. In the diagram. different curves are given for different values of the spectral index. I. of the non-thermal X-ray emission. and the line flux is normalised to an amplitude of non-thermal (power-law) emission of ¢=1 photons em? s! keV! (at 1 keV).," This column density is zero in our case, since, as mentioned above, the material in the accretion tubes is mostly neutral at temperatures of $\sim 7000$ K. In the diagram, different curves are given for different values of the spectral index, $\Gamma$, of the non-thermal X-ray emission, and the line flux is normalised to an amplitude of non-thermal (power-law) emission of $a =
1$ photons $^{-2}$ $^{-1}$ $^{-1}$ (at 1 keV)."
 The flux observed for Elias 29 in the Fe 6.4 keV emission line in the integrated spectrum is ~107° photons em s' and from the plot by Emslieetal..1986 one can see that such a flux can be obtained for a very wide choice of parameter combinations. with F ranging from 2.1 to 5.0. fore =107—I photons em7* s7! keV-!.," The flux observed for Elias 29 in the Fe 6.4 keV emission line in the integrated spectrum is $\sim 10^{-6}$ photons $^{-2}$ $^{-1}$, and from the plot by \citealp{epd86} one can see that such a flux can be obtained for a very wide choice of parameter combinations, with $\Gamma$ ranging from 2.1 to 5.0, for $a = 10^{-4} - 1$ photons $^{-2}$ $^{-1}$ $^{-1}$."
" The parameters F and e refer to the hard X-ray emission. typically at E220 keV. As discussed by Holman (2003).. however. the X-ray spectrum of the non-thermal X-ray emission flattens below the low-energy cutoff (E,.) of the suprathermal electron responsible for the emission,"," The parameters $\Gamma$ and $a$ refer to the hard X-ray emission, typically at $E \ga 20$ keV. As discussed by \cite{holman03}, , however, the X-ray spectrum of the non-thermal X-ray emission flattens below the low-energy cutoff $E_c$ ) of the suprathermal electron responsible for the emission."
" For example. the bremsstrahlung emission from a beam of electrons. with a power-law spectrum with spectral index 6=3 and low-energy cutoff E,=50 keV. has a spectral index of [D -2at E>E... but flattens tol~1.3 between | and 10 keV. For the electron beam model in solar flares. electrons are assumed to have energy higher than 20 keV (e.g. Brownetal.. 1990)) and recent RHESSI data indicates a low-energy cutoff that is typically E.=20—40 keV and can be as high 70 keV (Holmanetal.. 2003))."," For example, the bremsstrahlung emission from a beam of electrons, with a power-law spectrum with spectral index $\delta = 3$ and low-energy cutoff $E_c = 50$ keV, has a spectral index of $\Gamma \sim 2$ at $E >
E_c$, but flattens to $\Gamma \sim 1.3$ between 1 and 10 keV. For the electron beam model in solar flares, electrons are assumed to have energy higher than 20 keV (e.g. \citealp{bkm+90}) ) and recent RHESSI data indicates a low-energy cutoff that is typically $E_c = 20
- 40$ keV and can be as high 70 keV \citealp{hss+03}) )."
" Using the code by G. D. Holman. we verified that a non thermal hard X-ray emission with «=107—1 photons em s! keV7! and=2—5. from a population of electrons with E,240 keV. between | and 10 keV. would be buried in the star's thermal emission and thus cannot be constrained without hard-X ray observations."," Using the code by G. D. Holman, we verified that a non thermal hard X-ray emission with $a = 10^{-4} - 1$ photons $^{-2}$ $^{-1}$ $^{-1}$ and $\Gamma =
2-5$, from a population of electrons with $E_{c} \ga 40$ keV, between 1 and 10 keV, would be buried in the star's thermal emission and thus cannot be constrained without hard-X ray observations."
 Note. also. that a valueof a=1077—| photons em s! Κεντ! would not require an unusually high power input in accelerated electrons.," Note, also, that a valueof $a = 10^{-4} - 1$ photons $^{-2}$ $^{-1}$ $^{-1}$ would not require an unusually high power input in accelerated electrons."
 If our proposed mechanism of in situ. electron bombardment of neutral or nearly-neutralmaterial is correct.," If our proposed mechanism of in situ electron bombardment of neutral or nearly-neutralmaterial is correct,"
with the IFMR.,with the IFMR.
" Various theoretical question marks in SPS modeling can be answered, or at least constrained, by the IFMR."," Various theoretical question marks in SPS modeling can be answered, or at least constrained, by the IFMR."
 The degree of convective overshooting in stellar evolutionary models has an appreciable impact on implied progenitor masses and the ir» rrelationship;hence. forapredictedlightout putintheT P—AGB phase. therewillbeastatisticallysigni ficantdif fer üss," The degree of convective overshooting in stellar evolutionary models has an appreciable impact on implied progenitor masses and the $-$ relationship; hence, for a predicted light output in the TP-AGB phase, there will be a statistically significant difference between the remnant mass functions predicted with different overshooting parameters."
uchthatthecoregrowssigni ficantlyduringtheT P—AGB phase. wecan placemst ructive boundSONCOnVveCTIVeOVer:," If we demand that is such that the core grows significantly during the TP-AGB phase, we can place instructive bounds on convective overshoot prescriptions for some clusters."
 populationsre presentama joruncertaintvin$PSmodels(Conrovet al.2009).," Additionally, the luminosity of TP-AGB populations represent a major uncertainty in SPS models \citep{Conroy09}."
.Ourresultsshowingcoremassgro phaseasa functiono f initialmassrepresentafloortoT P—AGBstarluminosityandcanalread vexcludesomelin, Our results showing core mass growth in the TP-AGB phase as a function of initial mass represent a floor to TP-AGB star luminosity and can already exclude some limiting case models.
i AGBstarsdonotcontributetoa population sintegrated lightcannolongerbe physicallymotivated," Obviously, models assuming that TP-AGB stars do not contribute to a population's integrated light can no longer be physically motivated."
 Asmentionedabove AGbstarsissufficieuttoconstrainT P—AGBlightout put," As mentioned above, this work and assumptions as to the chemical yields of TP-AGB stars is sufficient to constrain TP-AGB light output."
 Thesenewconstraintsarevaluableto futureSPSmodel sam The most prominent sources of error in this analysis are the WD and cluster observations., These new constraints are valuable to future SPS models and will greatly reduce the uncertainties associated with the one of the most theoretically dubious stellar evolutionary phases.. The most prominent sources of error in this analysis are the WD and cluster observations.
 Measurements of WD surface gravities and temperatures provide constraints on remnant masses and cooling times while cluster observations yield the cluster age and composition., Measurements of WD surface gravities and temperatures provide constraints on remnant masses and cooling times while cluster observations yield the cluster age and composition.
 Uncertainties in theory linking these observations and the parameters of interest are typically small compared to the observational errors (S09)., Uncertainties in theory linking these observations and the parameters of interest are typically small compared to the observational errors (S09).
" By assuming either perfect WD or cluster measurements and repeating our analysis, we find that uncertainties in these two types of observables contribute similarly to the resulting error in core mass growth during the TP-AGB phase."," By assuming either perfect WD or cluster measurements and repeating our analysis, we find that uncertainties in these two types of observables contribute similarly to the resulting error in core mass growth during the TP-AGB phase."
 More precise WD or cluster measurements will provide stricter bounds on the energy output of TP-AGB stars., More precise WD or cluster measurements will provide stricter bounds on the energy output of TP-AGB stars.
 Another potential method to improve the precision of the IFMR (and our results) would be à survey specifically designed to find young white dwarts in open clusters., Another potential method to improve the precision of the IFMR (and our results) would be a survey specifically designed to find young white dwarfs in open clusters.
" In any given cluster, the youngest white dwarfs would correspond to the oldest progenitor lifetimes, reducing the impact of uncertainties in the WD and cluster observations on the implied initial mass andΙτρ."," In any given cluster, the youngest white dwarfs would correspond to the oldest progenitor lifetimes, reducing the impact of uncertainties in the WD and cluster observations on the implied initial mass and."
. More precise measurements of the IFMR would directly lead to more precise calculations in our procedure., More precise measurements of the IFMR would directly lead to more precise calculations in our procedure.
 Determining the metallicity dependence of TP-AGB star energy output is an intriguing next step., Determining the metallicity dependence of TP-AGB star energy output is an intriguing next step.
" However, there are not enough open clusters at extreme metallicities to empirically calibrate such a dependence."," However, there are not enough open clusters at extreme metallicities to empirically calibrate such a dependence."
" We would need to increase the number of clusters in our sample, even if the overall spread in metallicity does not change."," We would need to increase the number of clusters in our sample, even if the overall spread in metallicity does not change."
" With many more clusters, one could reproduce our results in bins of metallicity."," With many more clusters, one could reproduce our results in bins of metallicity."
" After empirically characterizing how the energy output of the phase scales with composition (albeit over a limited range), theory may be able to provide a physical model for TP-AGB star luminosity as a function of initial mass and stellar density."," After empirically characterizing how the energy output of the TP-AGB phase scales with composition (albeit over a limited range), theory may be able to provide a physical model for TP-AGB star luminosity as a function of initial mass and stellar density."
" If said model could be extrapolated to low metallicity, it would provide a theoretical constraint on the contribution of TP-AGB stars to the spectral energy distributions of high redshift galaxies."," If said model could be extrapolated to low metallicity, it would provide a theoretical constraint on the contribution of TP-AGB stars to the spectral energy distributions of high redshift galaxies."
" At high metallicity, the lowest mass stars may go directly to the white dwarf stage, missing the core He burning stages and beyond (Kilicetal.2007).."," At high metallicity, the lowest mass stars may go directly to the white dwarf stage, missing the core He burning stages and beyond \citep{Kilic07}."
 This effect is not currently included in population synthesis models and might be important for the low redshift properties of giant elliptical galaxies., This effect is not currently included in population synthesis models and might be important for the low redshift properties of giant elliptical galaxies.
 Physically motivated priors on TP-AGB star light output would be a substantial advancement of semi-analytic galaxy evolution models as SPS codes still have widely varied characterizations of these stars., Physically motivated priors on TP-AGB star light output would be a substantial advancement of semi-analytic galaxy evolution models as SPS codes still have widely varied characterizations of these stars.
(he radio spectral index (0.4-0.6) between 1.4 and 4.87 GlIz.,the radio spectral index (0.4-0.6) between 1.4 and 4.87 GHz.
 Unless the spectral index of the unseen electrons is dramatically steeper (han that of the electrons producing the observed radio emission. itis not possible for a signilicant fraction of the X-ray emission to be due to (his mechanism because of these spectral differences.," Unless the spectral index of the unseen electrons is dramatically steeper than that of the electrons producing the observed radio emission, it is not possible for a significant fraction of the X-ray emission to be due to this mechanism because of these spectral differences."
 On the basis of the morphological similarities between the radio and X-ray. emission and the limitations of competing models. it has previously been concluded. (hat both the X-ray and radio flux from (he jet originate in svuchrotvon emission from a population of ultrarelativistic electrons.," On the basis of the morphological similarities between the radio and X-ray emission and the limitations of competing models, it has previously been concluded that both the X-ray and radio flux from the jet originate in synchrotron emission from a population of ultrarelativistic electrons."
 The knots are presumably sites of shocks where particles are accelerated to 5~ 10* to produce the observed X-ray emission. assuming an equipartition magnetic field (Feigelsone£af.1981:Burns.Feigelson.andSchreier1983).," The knots are presumably sites of shocks where particles are accelerated to $\gamma\sim$ $^{7-8}$ to produce the observed X-ray emission, assuming an equipartition magnetic field \citep{fei81,bur83}."
. The main argument against a svnchrotron hypothesis has been that the lifetime due to svuchrotron emission of the ullrvarelativistic particles is considerablv less (han the travel time down the jet. thus requiring that the particles be reaccelerated (Feigelsone£a£.1951).. perhaps multiple times. as (hey travel along the jet to the radio lobe.," The main argument against a synchrotron hypothesis has been that the lifetime due to synchrotron emission of the ultra-relativistic particles is considerably less than the travel time down the jet, thus requiring that the particles be reaccelerated \citep{fei81}, perhaps multiple times, as they travel along the jet to the radio lobe."
 In fact. in the inner jet region. the particles mist undergo local accelerations multiple times within a knot.," In fact, in the inner jet region, the particles must undergo local accelerations multiple times within a knot."
 The lifetime of the X-ray emitting electrons in (he inner jet is ~20 vears. an order of magnitude less than the travel time from the nucleus to knot D. The lifetime of the particles bevond knot D is longer (760 vears) due to the lower magnetic field. but still consiclerably less than the travel time into the radio lobe.," The lifetime of the X-ray emitting electrons in the inner jet is $\sim$ 20 years, an order of magnitude less than the travel time from the nucleus to knot B. The lifetime of the particles beyond knot B is longer $\sim$ 60 years) due to the lower magnetic field, but still considerably less than the travel time into the radio lobe."
 Based on estimates of the knot sizes from. observations (hundreds of parsecs). ib was argued that the large volumes of the knots pushed acceleration models to their limits (Burns.Feigelson.andSchreier1933).," Based on estimates of the knot sizes from observations (hundreds of parsecs), it was argued that the large volumes of the knots pushed acceleration models to their limits \citep{bur83}."
. This criticism is considerably weakened based on our observations because each previously observed. X-ray. knot is actually composed of several smaller knots., This criticism is considerably weakened based on our observations because each previously observed X-ray knot is actually composed of several smaller knots.
 Thus the knot volume ancl the shock sizes are considerably sialler (tens of pe rather than hundreds of pe) than previously supposed., Thus the knot volume and the shock sizes are considerably smaller (tens of pc rather than hundreds of pc) than previously supposed.
 Local acceleration of the electrons within the knots is still a necessary component of this model. however.," Local acceleration of the electrons within the knots is still a necessary component of this model, however."
 The morphological dillerences between the X-ray. and radio emission complicate this V.imple model for (he emission., The morphological differences between the X-ray and radio emission complicate this simple model for the emission.
 These differences argue strongly against the simple hypothesis iab a single. spatially coincident. population of relativistic electrons is responsible for both 1ο X-ray. and radio emission (Feigelsonefal.1931:Burns.Feigelson.andSchreier1933). in je nuclear ancl inner jet regions. except perhaps at knots NI/NXI and AL1/AXI.," These differences argue strongly against the simple hypothesis that a single, spatially coincident, population of relativistic electrons is responsible for both the X-ray and radio emission \citep{fei81,bur83}
 in the nuclear and inner jet regions, except perhaps at knots N1/NX1 and A1/AX1."
 Our data support a model in which the X-ray emission is indeed svncehrotvon radiation from relativistic parücles. but the radio emission originates. at least in part. from a spatially distinct. less energetic population of electrons.," Our data support a model in which the X-ray emission is indeed synchrotron radiation from relativistic particles, but the radio emission originates, at least in part, from a spatially distinct, less energetic population of electrons."
 The only other extragalactic jet that has been as well studied as the Cen A jet is that in AIST. and there are some remarkable morphological similarities between (he two jets.," The only other extragalactic jet that has been as well studied as the Cen A jet is that in M87, and there are some remarkable morphological similarities between the two jets."
 sinall differences between the X-ray. optical. and radio morphology. of the jet in M8T on," Small differences between the X-ray, optical, and radio morphology of the jet in M87 on"
than in model 01.,than in model 01.
 The loops supporting the immer bar in model OL also undergo a lot more pulsation than iu model 01 (Figure 1)., The loops supporting the inner bar in model 04 also undergo a lot more pulsation than in model 01 (Figure 4).
 The pattern speed of the inner bar for model 03 is 90iL., The pattern speed of the inner bar for model 03 is 90.
 There is a gap between loops 5 and 9. which corresponds in the rine-width diagram (Fieure lc) to the region of the inner arch. where its dark spine is apparently uvissine.," There is a gap between loops 8 and 9, which corresponds in the ring-width diagram (Figure 1c) to the region of the inner arch, where its dark spine is apparently missing."
 We could not fiud good approximations to the loops in this region., We could not find good approximations to the loops in this region.
 This means that even if à more detailed search succeeds in findiug the loops there. regular orbits are uot trapped well there. hence the support for the iuner bar is strongly reduced.," This means that even if a more detailed search succeeds in finding the loops there, regular orbits are not trapped well there, hence the support for the inner bar is strongly reduced."
 Nevertheless. re extent of the immer bar. in terius of support from the we. orbits. ds still ereater than in the models above.," Nevertheless, the extent of the inner bar, in terms of support from the $x_2$ orbits, is still greater than in the models above."
 The last loop supporting the iuner bar is the sth loop. and because of the gap past it. where we could not fiud loops.it has the longest seni-niajor axis among these loops.," The last loop supporting the inner bar is the 8th loop, and because of the gap past it, where we could not find loops, it has the longest semi-major axis among these loops."
 ILoxcever. the extent of the bar may still be slightly larger if a loop supporting this bar were fouud past the 8th loop.," However, the extent of the bar may still be slightly larger if a loop supporting this bar were found past the 8th loop."
 Contrarv to the models analysed so far. the loops supporting the immer bar iu model 03 do not behave coherently throughout rotation: when the augle between the bars in the nuposed potential is iu Figure 2 there is a twist in loops supporting the immer bar. reflected by a change in the PA in Fiewre 32e.," Contrary to the models analysed so far, the loops supporting the inner bar in model 03 do not behave coherently throughout rotation: when the angle between the bars in the imposed potential is, in Figure 2 there is a twist in loops supporting the inner bar, reflected by a change in the PA in Figure 3e."
 One can see that the deeper inside the bar the loop is located. the more it accelerates.," One can see that the deeper inside the bar the loop is located, the more it accelerates."
 Au upturn iu the PA for the iuncr-most loop appears now to be more consistent with the trend exhibitedbv the other loops. but for the sake of consistency with the other models we still exclude the iuncr-most loop from the loops supporting the immer bar.," An upturn in the PA for the inner-most loop appears now to be more consistent with the trend exhibited by the other loops, but for the sake of consistency with the other models we still exclude the inner-most loop from the loops supporting the inner bar."
 Iu model 03. (loopsthe mean PA of the loops supporting the oemer bar 25) is larger thanin the models above (Table1). hence the immer bar in this model undergoes more acceleration.," In model 03, the mean PA of the loops supporting the inner bar (loops 2–8) is larger than in the models above (Table 1), hence the inner bar in this model undergoes more acceleration."
 The loops supporting the immer bar iu model 03 are more eccentric compared to the models above aud they imply greater pulsation of the iuner bar than in the models above. but this pulsation is no longer consistent among the loops (Figure 1).," The loops supporting the inner bar in model 03 are more eccentric compared to the models above and they imply greater pulsation of the inner bar than in the models above, but this pulsation is no longer consistent among the loops (Figure 4)."
 The pattern speed of the iuner bar in model 02 is 80iL., The pattern speed of the inner bar in model 02 is 80.
 The last among the loops supporting the inner bar is the 9th loop. but among them the sth one has the longest scuu-imajor axis aud therefore serves as the bar length estimate.," The last among the loops supporting the inner bar is the 9th loop, but among them the 8th one has the longest semi-major axis and therefore serves as the bar length estimate."
 In the rine-width ciaerau in Figure b. the initial conditions for loops 8-10 iu model 02 are taken from a ridec-like feature iu the right leg of the ner arch.," In the ring-width diagram in Figure 1b, the initial conditions for loops 8-10 in model 02 are taken from a ridge-like feature in the right leg of the inner arch."
 This ridge departs upwards from the course of the dark spine of the immer arch on the lower right., This ridge departs upwards from the course of the dark spine of the inner arch on the lower right.
 On the vieht of this ridge. there are light erev areas indicating poorly trapped trajectories. and ou the left there is white area indicating chaotic trajectories with initial conditions iuuediatelv next to those of the loops.," On the right of this ridge, there are light grey areas indicating poorly trapped trajectories, and on the left there is white area indicating chaotic trajectories with initial conditions immediately next to those of the loops."
 Finding loops ou this ridge is unexpected since stable double frequency orbits should be well surrounded by trapped trajectories., Finding loops on this ridge is unexpected since stable double frequency orbits should be well surrounded by trapped trajectories.
 Iu fact in model 03. for corresponding iuitial conditions (hetween loops 8 aud 9 there) we were not able to fiud loops at all: trapping of trajectories was increasing when moving to the left iu the rine-width diagram in this region. but then the trajectories were abruptly becoming chaotic before they couverged ou loops.," In fact in model 03, for corresponding initial conditions (between loops 8 and 9 there) we were not able to find loops at all: trapping of trajectories was increasing when moving to the left in the ring-width diagram in this region, but then the trajectories were abruptly becoming chaotic before they converged on loops."
 We checked he appearance of orbits around the ridge iu model Figure 5 shows the ath loop. aud points sampled roni the trajectories whose initial positions are just sanaller aud larger than that of the loop (1.6. come from adjacent pixels iu the rine-width diagram).," We checked the appearance of orbits around the ridge in model 02: Figure 5 shows the 8th loop, and points sampled from the trajectories whose initial positions are just smaller and larger than that of the loop (i.e. come from adjacent pixels in the ring-width diagram)."
 As expected. he trajectory with initial conditions from the lett of the ridge (Figure 5. left-hand panel) looks chaotic. while the rajectorv whose initial conditions are to the right of the ridee (Figure 5. right-hand panel). although reeular. docs iot. align with the loop. and hence it does not support he bar.," As expected, the trajectory with initial conditions from the left of the ridge (Figure 5, left-hand panel) looks chaotic, while the trajectory whose initial conditions are to the right of the ridge (Figure 5, right-hand panel), although regular, does not align with the loop, and hence it does not support the bar."
 Therefore the outermost loops of the ΠΙΟ bar (loops 7-9) in 1nodel 02 provide esseutially no support for lis bar in teris of trapped trajectories., Therefore the outermost loops of the inner bar (loops 7-9) in model 02 provide essentially no support for this bar in terms of trapped trajectories.
 The resolution of the riug-width diagram is barely sufficient to locate hese loops. hence in Figure 2 they appear as a collection of noticeably scattered points.," The resolution of the ring-width diagram is barely sufficient to locate these loops, hence in Figure 2 they appear as a collection of noticeably scattered points."
 The twisting of the loops in the inner bar is even more dramatic iu model 02 than iu model 03., The twisting of the loops in the inner bar is even more dramatic in model 02 than in model 03.
" In Figure 2. these loops formu, a spiral rather than a bar. when the anele between the bars in the imposed poteutial is"," In Figure 2, these loops form a spiral rather than a bar, when the angle between the bars in the imposed potential is."
 As in model 03. the iunerauost orbit is excluded from loops supporting the innerbar.," As in model 03, the inner-most orbit is excluded from loops supporting the inner bar."
 Loops supporting the inner bar iu model 02 strougly pulsatebut the change im eccentrieitv Is even more inconsistent among them thaw in the models analysed above. as shown in Figure 1.," Loops supporting the inner bar in model 02 strongly pulsate but the change in eccentricity is even more inconsistent among them than in the models analysed above, as shown in Figure 4."
 The pattern speed for the immer bar m model 02E is T0|., The pattern speed for the inner bar in model 02E is 70.
 This is a new model. not preseuted iu ALAOS. and it extends the rauge of models frou: NLÀAOS to lower pattern speeds.," This is a new model, not presented in MA08, and it extends the range of models from MA08 to lower pattern speeds."
" The rine-width diagram for this model, preseuted in Figure 1a. looks very differcut from those for all the other models analysed in this paper."," The ring-width diagram for this model, presented in Figure 1a, looks very different from those for all the other models analysed in this paper."
 The loops. presented in the bottom row of Figure 2. are also significantly different.," The loops, presented in the bottom row of Figure 2, are also significantly different."
 The inuermost two loops have the initial couditious taken from the W-shaped feature branching from the outer arch in Figure la., The innermost two loops have the initial conditions taken from the $V$ -shaped feature branching from the outer arch in Figure 1a.
 These loops see to rotate in the opposite direction to the two bars., These loops seem to rotate in the opposite direction to the two bars.
 There is a region of siall rine-widtls between the left, There is a region of small ring-widths between the left
than that of (classical) bulges.,than that of (classical) bulges.
" In this context, the scenario entertained by ?,, in which the spheroid forms after the dissolution of a bar, is highly improbable, given that the typical bar is about two orders of magnitude less massive than the spheroid in Sombrero ?)."," In this context, the scenario entertained by , in which the spheroid forms after the dissolution of a bar, is highly improbable, given that the typical bar is about two orders of magnitude less massive than the spheroid in Sombrero ."
". On the other hand, the bulge obtained in modelBDH, although somewhat small for its mass, falls near the relation drawn for classical bulges, although the locus it occupies in this relation is also not unusual for a disc-like, pseudo-bulge."," On the other hand, the bulge obtained in model, although somewhat small for its mass, falls near the relation drawn for classical bulges, although the locus it occupies in this relation is also not unusual for a disc-like, pseudo-bulge."
" In summary, our structural analysis shows that the most appropriate model for Sombrero is modelBDH, with a spheroid comprising an inner bulge and an outer halo."," In summary, our structural analysis shows that the most appropriate model for Sombrero is model, with a spheroid comprising an inner bulge and an outer halo."
" Nevertheless, we cannot rule out the possibility that the spheroid is in fact an elliptical galaxy (as suggested in Fig. 8)),"," Nevertheless, we cannot rule out the possibility that the spheroid is in fact an elliptical galaxy (as suggested in Fig. \ref{fig:scale}) ),"
" that, unlike most elliptical galaxies, displays intricate structural properties, due to the presence of a massive embedded disc, and the possibility that this disc have formed a bar and thus experienced effects from secular evolution."," that, unlike most elliptical galaxies, displays intricate structural properties, due to the presence of a massive embedded disc, and the possibility that this disc have formed a bar and thus experienced effects from secular evolution."
" This opens up the possibility that the current view on the Sombrero spheroid as a single entity, a large classical merger-built bulge, is mistaken."," This opens up the possibility that the current view on the Sombrero spheroid as a single entity, a large classical merger-built bulge, is mistaken."
 These possibilities will be discussed in more detail in the next section., These possibilities will be discussed in more detail in the next section.
" Another important way to address the nature of the spheroid in Sombrero is by looking at minor axis profiles, where the influence of the inner sub-structures, disc and rings is minimized."," Another important way to address the nature of the spheroid in Sombrero is by looking at minor axis profiles, where the influence of the inner sub-structures, disc and rings is minimized."
 In Fig., In Fig.
" 9 we plot the minor axis profiles of the galaxy and modelsBD, andVEBD."," \ref{fig:minor} we plot the minor axis profiles of the galaxy and models, and."
" One sees that modelBDH provides a better fit, although it does not fit well the outer parts."," One sees that model provides a better fit, although it does not fit well the outer parts."
" This is very likely because the halo has a fixed ellipticity of 0.28, somewhat larger than the ellipticity of the outermost parts of the galaxy, which is around 0.1."," This is very likely because the halo has a fixed ellipticity of 0.28, somewhat larger than the ellipticity of the outermost parts of the galaxy, which is around 0.1."
" Model provides a better fit than modelBD, attesting the importance of having an ellipticity varying with radius."," Model provides a better fit than model, attesting the importance of having an ellipticity varying with radius."
" At the inner parts, models and provide equally good fits, whereas model is slightly off."," At the inner parts, models and provide equally good fits, whereas model is slightly off."
" In this context, we produced further fits by excluding all sub-structures apart from the spheroid."," In this context, we produced further fits by excluding all sub-structures apart from the spheroid."
" This was done by masking a rectangular area centred at the galaxy center, and with dimensions 36” x 430”, along the disc major axis."," This was done by masking a rectangular area centred at the galaxy center, and with dimensions 36” $\times$ 430”, along the disc major axis."
" The minor axis dimension is twice the disc scale height (as in model BD), and the major axis dimension was chosen to allow only for the part dominated by the spheroid, as seen in the galaxy profile (see e.g. Fig. 2))."," The minor axis dimension is twice the disc scale height (as in model ), and the major axis dimension was chosen to allow only for the part dominated by the spheroid, as seen in the galaxy profile (see e.g. Fig. \ref{fig:profs_bd}) )."
" The new fits include a Sérrsic spheroid only, with the Sérrsic index left free."," The new fits include a Sérrsic spheroid only, with the Sérrsic index left free."
" Model has free ellipticity, model has a fixed ellipticity of 0.1, i.e. reproducing the outermost parts of the galaxy ellipticity radial profile (see Fig. 6)),"," Model has free ellipticity, model has a fixed ellipticity of 0.1, i.e. reproducing the outermost parts of the galaxy ellipticity radial profile (see Fig. \ref{fig:geo}) ),"
" and finally, model has varying ellipticity, as in modelVEBD."," and finally, model has varying ellipticity, as in model."
 The resulting parameters are shown in Table 1.., The resulting parameters are shown in Table \ref{tab:result}.
 Minor axis profiles from these fits are shown in Fig. 10.., Minor axis profiles from these fits are shown in Fig. \ref{fig:minor_s}.
" These new models generally provide better fits to the outer parts of the galaxy than the models in the previous sub-section, but fail in the inner parts."," These new models generally provide better fits to the outer parts of the galaxy than the models in the previous sub-section, but fail in the inner parts."
 This evidently reflects the design of this exercise and the influence the disc component has in the global fits of the previous sub-section., This evidently reflects the design of this exercise and the influence the disc component has in the global fits of the previous sub-section.
 Model is the one that fails the most in the outermost parts., Model is the one that fails the most in the outermost parts.
" Model is overall quite good but fails at intermediate radii, i.e. from about 50” to 100""."," Model is overall quite good but fails at intermediate radii, i.e. from about 50” to 100”."
" Finally, model can account for the change in ellipticity in the outermost parts, but generally does not result in a very good fit."," Finally, model can account for the change in ellipticity in the outermost parts, but generally does not result in a very good fit."
 This exercise illustrates the fact that even excluding from the fit all, This exercise illustrates the fact that even excluding from the fit all
The emerging magnetic flux of active regions on the solar surface appears in the form of separate loops: they then coalesce to form opposite polarity poles that separate and increase m size (?)..,The emerging magnetic flux of active regions on the solar surface appears in the form of separate loops; they then coalesce to form opposite polarity poles that separate and increase in size \citep{str99}.
 Earlier studies of velocities of plasma motions were carried out in the region of emerging magnetic fields at different height levels of the Sun., Earlier studies of velocities of plasma motions were carried out in the region of emerging magnetic fields at different height levels of the Sun.
 Investigations into the chromospheric level have been made by ?? and others (seereviewby?)..," Investigations into the chromospheric level have been made by \citet{bru67,bru69} and others \citep[see review
by][]{cho93}."
 Velocities at the photospheric level were measured with different methods (discussed in more detail below)., Velocities at the photospheric level were measured with different methods (discussed in more detail below).
" In the last decade. helioseismological methods have been applied to study the subphotospheric level (??????).,"," In the last decade, helioseismological methods have been applied to study the subphotospheric level \citep{cha99,kos00,zha08,kos09,kom08,kom09}."
 Direct measurements of photospherie velocities revealed negative values (blue Doppler shift) on the polarity inversion line of emerging magnetic fields., Direct measurements of photospheric velocities revealed negative values (blue Doppler shift) on the polarity inversion line of emerging magnetic fields.
 The events under consideration. were located in the central part of the solar disk. and. therefore. vertical motions were measured.," The events under consideration were located in the central part of the solar disk, and, therefore, vertical motions were measured."
 Negative velocities or plasma upflows up to 1 km s'∣ were observed in separate magnetic loops inside emerging active regions (??222)..," Negative velocities or plasma upflows up to 1 km $^{-1}$ were observed in separate magnetic loops inside emerging active regions \citep{tar90,lit98,str99,kub03,gri09}."
 ?. revealed high velocities of about 1.7 km s! at the beginning of the powerful active region emergence NOAA 10488 at heliographie coordinates NO8 E31 (By = + 4.9)., \citet{gri07} revealed high velocities of about 1.7 km $^{-1}$ at the beginning of the powerful active region emergence NOAA 10488 at heliographic coordinates N08 E31 $B_{0}$ = + 4.9).
 Indirect measurements of the horizontal velocities of the photospheric flow that accompanies the emergence of active regions were made by tracking the displacement of individual magnetic elements., Indirect measurements of the horizontal velocities of the photospheric flow that accompanies the emergence of active regions were made by tracking the displacement of individual magnetic elements.
 The velocities obtained ranged from 0.1 to 14 km s! (222).," The velocities obtained ranged from 0.1 to 1.4 km $^{-1}$ \citep{fra72,sch73,str99}."
 2) studied 45 bipolar pairs in the emerging active region., \citet{bar90} studied 45 bipolar pairs in the emerging active region.
 The opposite polarity poles were separating from each other with velocities of 0.5-3.5 km s*'. decreasing with time (consequently. the drift with regard to the polarity inversion line is only half that value).," The opposite polarity poles were separating from each other with velocities of 0.5–3.5 km $^{-1}$, decreasing with time (consequently, the drift with regard to the polarity inversion line is only half that value)."
 ? calculated the separation velocities of the external boundaries of the photospheric magnetic flux in the active region NOAA 10488., \citet{gri09} calculated the separation velocities of the external boundaries of the photospheric magnetic flux in the active region NOAA 10488.
 The velocities decreased as the magnetic fields emerged: they were 2—2.5 kmss! by the end of the first hour and 0.3 km s! in two hours and a half., The velocities decreased as the magnetic fields emerged: they were 2–2.5 $^{-1}$ by the end of the first hour and 0.3 km $^{-1}$ in two hours and a half.
 It should be noted that there are several papers dealing with the study of vertical (2) and horizontal (2???) velocities at photospheric level during the appearance of ephemeral active regions.," It should be noted that there are several papers dealing with the study of vertical \citep{gug06} and horizontal \citep{har73,cho87,hag01}
 velocities at photospheric level during the appearance of ephemeral active regions."
 We do not discuss these works in detail. since they concern another spatial scale of emerging magnetic fields.," We do not discuss these works in detail, since they concern another spatial scale of emerging magnetic fields."
 The papers above are the researches into photospheric velocities that accompany the emergence of magnetic fields in active regions located only in the central zone of the solar disk., The papers above are the researches into photospheric velocities that accompany the emergence of magnetic fields in active regions located only in the central zone of the solar disk.
 Plasma motions in the emerging magnetic fields distant from the disk center have neverbeen studied before., Plasma motions in the emerging magnetic fields distant from the disk center have neverbeen studied before.
 We used full solar disk magnetograms and Dopplergrams in the photospheric line Ni 1 6768 A and continuum images obtained on board the space observatory SOHO/MDI (?).., We used full solar disk magnetograms and Dopplergrams in the photospheric line Ni I 6768 $\mbox{\AA}$ and continuum images obtained on board the space observatory SOHO/MDI \citep{sch95}.
 The temporal resolution of the magnetograms and Dopplergrams is| ninute. that of the continuum. 96 minutes.," The temporal resolution of the magnetograms and Dopplergrams is 1 minute, that of the continuum, 96 minutes."
" The spatial resolution of the data is 4"". the pixel size is approximately 2."," The spatial resolution of the data is $\arcsec$, the pixel size is approximately $\arcsec$."
 Magnetograms with a 1.8.2 calibration level were used (?).., Magnetograms with a 1.8.2 calibration level were used \citep{ulr09}.
 Besides photospheric velocities of solar plasma motions. the Dopplergrams include the contribution of 1) the velocity of the differential solar rotation: 2) the velocity of the SOHO satellite with regard to the Sun: 3) the instrumental distortions caused by nonuniform transmission of MDUSs filter systems on the field of view.," Besides photospheric velocities of solar plasma motions, the Dopplergrams include the contribution of 1) the velocity of the differential solar rotation; 2) the velocity of the SOHO satellite with regard to the Sun; 3) the instrumental distortions caused by nonuniform transmission of MDI's filter systems on the field of view."
 The technique described in ? was used to separate photospheric velocities., The technique described in \citet{gri07} was used to separate photospheric velocities.
 The negative velocity on SOHO/MDI Dopplergrams corresponds to the blue Doppler shift (motion of matter toward the observer): the positive one. to the red Doppler shift (motion of matter away from the observer).," The negative velocity on SOHO/MDI Dopplergrams corresponds to the blue Doppler shift (motion of matter toward the observer); the positive one, to the red Doppler shift (motion of matter away from the observer)."
 There are regular errors in velocity measurements at different distances from the disk center., There are regular errors in velocity measurements at different distances from the disk center.
 For example. Doppler velocities near the limb are underestimated by up to 40 m s! because of limb darkening (?)..," For example, Doppler velocities near the limb are underestimated by up to 40 m $^{-1}$ because of limb darkening \citep{alb85}."
 Additionally. we deal with different height levels above the solar surface in measurements of velocities at the disk center and near the limb because of the different optical thickness of the atmosphere.," Additionally, we deal with different height levels above the solar surface in measurements of velocities at the disk center and near the limb because of the different optical thickness of the atmosphere."
 We study the velocities that accompany the emerging magnetic fields., We study the velocities that accompany the emerging magnetic fields.
 Therefore. the precise spatial superimposition of the used data is vary important.," Therefore, the precise spatial superimposition of the used data is vary important."
 For this purpose. the region of the emerging magnetic. fields was taken from the time sequences of magnetograms and Dopplergrams. taking into account solar rotation.," For this purpose, the region of the emerging magnetic fields was taken from the time sequences of magnetograms and Dopplergrams, taking into account solar rotation."
 An approximate value of the displacement of the region was calculated with the differential rotation law for photospheric magnetic fields (2).., An approximate value of the displacement of the region was calculated with the differential rotation law for photospheric magnetic fields \citep{sno83}. .
 The exact tracking of the region under investigation. was performed by applying two magnetograms adjacent in time. with the," The exact tracking of the region under investigation was performed by applying two magnetograms adjacent in time, with the"
where jjο,where $\eta\sim\mathcal{O}(\delta^2)$.
 The horizontal shear rate of this base flow is: Keplerian shear is anticvelonic (as indicated by the negative sign). and is comparable in magnitude to the rotation rate itself.," The horizontal shear rate of this base flow is: Keplerian shear is anticyclonic (as indicated by the negative sign), and is comparable in magnitude to the rotation rate itself."
 As we will see. this has important consequences for the properties of vortices in protoplanetary disks.," As we will see, this has important consequences for the properties of vortices in protoplanetary disks."
 Shear will stvetel ancl tear a vortex apart unless: (1) the vortex rotates in the sume sense as (he ambient shear. and (2) the vortex is al least as strong as the shear (Moore1990.1993).," Shear will stretch and tear a vortex apart unless: (1) the vortex rotates in the same sense as the ambient shear, and (2) the vortex is at least as strong as the shear \citep{ms71,kida81,marcus90a,marcus93}."
. The first condition implies that all vortices in a protoplanetary disk will be antievelones: that is. (he vortices will rotate in a sense opposite of that of the overall rotation of the disk.," The first condition implies that all long-lived vortices in a protoplanetary disk will be anticyclones; that is, the vortices will rotate in a sense opposite of that of the overall rotation of the disk."
" Let ὃ be the characteristic horizontal speed of gasaround a vortex wilh aspect ratio y=A./A,>1. where A, (along the radial direction in the disk) is the minor axis and A. (along the azimuthal direction in the disk) is the major axis."," Let $\tilde{v}_{\bot}$ be the characteristic horizontal speed of gasaround a vortex with aspect ratio $\chi\equiv\Lambda_{\phi}/\Lambda_r>1$, where $\Lambda_r$ (along the radial direction in the disk) is the minor axis and $\Lambda_{\phi}$ (along the azimuthal direction in the disk) is the major axis."
 The characteristic z-component of in (he core of the vortex is d.ονὃΑμ., The characteristic $z$ -component of in the core of the vortex is $\tilde{\omega}_z\sim\tilde{v}_{\bot}/\Lambda_r$.
 The second condition above implies that the horizontal speed of gas around the vortex must be of order the dillerential velocity across the vortex due to the ambient shear. or equivalently. that the characteristic vorticity is of order the shear rate: d~6{Αρoy. Moore&Sallman(1971). and Ixida(1981). analytically determined (he exact steady-state solution for a 2D elliptical vortex of uniform. vorticity embecdcded in a uniform shear and showed that: for a vortex that rotates in the same sense as the shear.," The second condition above implies that the horizontal speed of gas around the vortex must be of order the differential velocity across the vortex due to the ambient shear, or equivalently, that the characteristic vorticity is of order the shear rate: $\tilde{\omega}_z\sim\tilde{v}_{\bot}/\Lambda_r\sim\sigma_K.$ \citet{ms71} and \citet{kida81} analytically determined the exact steady-state solution for a 2D elliptical vortex of uniform vorticity embedded in a uniform shear and showed that: for a vortex that rotates in the same sense as the shear."
 A vortex (hat is weak relative to the shear will be stretched oul aud. have a large aspect ratio. whereas a vortex (hat is comparable in strength to the shear will be more compact and have an aspect ratio closer to unity.," A vortex that is weak relative to the shear will be stretched out and have a large aspect ratio, whereas a vortex that is comparable in strength to the shear will be more compact and have an aspect ratio closer to unity."
" Because the Keplerian shear is comparable in magnitude to the Ixeplerian rotation rate. the rotational period of gas around the core of the vortex. 7,,,=4a/w.. is of order the orbital period of the vortex around the protostar. 7,= 27/Q."," Because the Keplerian shear is comparable in magnitude to the Keplerian rotation rate, the rotational period of gas around the core of the vortex, $\tau_{vor}\equiv 4\pi/\tilde{\omega}_z$, is of order the orbital period of the vortex around the protostar, $\tau_{orb}\equiv 2\pi/\Omega$ ."
 The Rossby number is defined to be the ratio of these two timescales:, The Rossby number is defined to be the ratio of these two timescales:
Collapse supernovae {ype IL as well as type Ib.c).,"collapse supernovae (type II as well as type Ib,c)."
 The GOODS cdaa Lor core collapse supernovae have been placed in two bins at z=0.340.2 and z—0.7+0.2., The GOODS data for core collapse supernovae have been placed in two bins at $z = 0.3 \pm 0.2$ and $z = 0.7 \pm 0.2$.
 Their results Qvhich have been corrected for the effects of extinction) show SN rates which are signiicautly higher than tle local rate (at 5Q) which is taken Grom ?.., Their results (which have been corrected for the effects of extinction) show SN rates which are significantly higher than the local rate (at $z = 0$ ) which is taken from \citet{cappellaro:99}.
 Since the timescale between star formation aud the core collapse explosion is very short (particularly [or very massive stars). there will be no contribution of Population HL stars to type IL rates at PEE ," Since the timescale between star formation and the core collapse explosion is very short (particularly for very massive stars), there will be no contribution of Population III stars to type II rates at $z \le 1$."
"As retmarked above. the SN tyye IHE rate is directly related to the overall SFR and therefor to the assumed astration rate vy, ane our )est fit model does an excellent job at describing the existing cala."," As remarked above, the SN type II rate is directly related to the overall SFR and therefore to the assumed astration rate $\nu_1$ and our best fit model does an excellent job at describing the existing data."
 Note that. had. we chosen a Sclunidt law (V.xGeas). LL woud net be possible simultaneosly fit bot1i the local (2= 0) daaas well as the GOODS data.," Note that, had we chosen a Schmidt law $\Psi \propto \sigma_{\mathrm gas})$, it would not be possible to simultaneously fit both the local $z = 0$ ) data as well as the GOODS data."
 I1 general. inodels wih a Sclunidt aw Call Dot account [ο τα laree eiliaucement in the SN rate.," In general, models with a Schmidt law can not account for a large enhancement in the SN rate."
 Therefore ounce the peak of the SFR has |een fixed (at zzz 1). the local SN rate is too large.," Therefore once the peak of the SFR has been fixed (at $z \approx 1$ ), the local SN rate is too large."
 The slope o‘the type HE rate (wih respect to 'ecshift) is also dependent ou the IMF slope. wy.," The slope of the type II rate (with respect to redshift) is also dependent on the IMF slope, $x_1$."
 In Figure 6.. we show this dependeuce displaying the preclicted type IL rates for ey=1.1.1.3 aud 1.5.," In Figure \ref{fig:SNII}, we show this dependence displaying the predicted type II rates for $x_1 = 1.1, 1.3$ and 1.5."
 Of course tie choice of vy and wry also has a stroug e[ect ou the metallicity and detailed chemical history as cliscussed it more detail iu the next section when Population ILI stars are included iu the analysis., Of course the choice of $\nu_1$ and $x_1$ also has a strong effect on the metallicity and detailed chemical history as discussed in more detail in the next section when Population III stars are included in the analysis.
" It is Importau to note that while the data at 2x1 is well described by the mocel. the same uodel predicts much higher SN ""utes αἱ higher redshift."," It is important to note that while the data at $z \le 1$ is well described by the model, the same model predicts much higher SN rates at higher redshift."
 BIudeec. he SN type II rate peaks at a redshift z22 αἱ a ‘ate which is uea‘ly 8 times the observed rate at z2Q.T.," Indeed, the SN type II rate peaks at a redshift $z \approx 3$ at a rate which is nearly 8 times the observed rate at $z \approx 0.7$."
 Ht will bei ierestiug to see whether [uture data will be able to probe the SN rate at higher redshitt., It will be interesting to see whether future data will be able to probe the SN rate at higher redshift.
 These elevated rates uay be detectable iu spite of the expected iicrease in dust extinction due to the early production of netals., These elevated rates may be detectable in spite of the expected increase in dust extinction due to the early production of metals.
 Au alternaive to the direct detection of SN at high redshilt is the indirect detection of SN hrough neutrinos., An alternative to the direct detection of SN at high redshift is the indirect detection of SN through neutrinos.
 The predicted neutrino background is very sensitive to the assumed. chemical evolutiou moclel (???)..," The predicted neutrino background is very sensitive to the assumed chemical evolution model \citep{ando:04,iocco:05,dosv:05}."
 As it the case for type HE supernovae. the data (also taken from GOODS (2))) is binued iuto our redshift |uus at z— 0.1. 0.8. 1.2. and 1.6. each with a spread of 0.2.," As in the case for type II supernovae, the data (also taken from GOODS \citep{dahlen:04}) ) is binned into four redshift bins at $z = 0.4$ , 0.8, 1.2, and 1.6, each with a spread of $\pm 0.2$ ."
 These are shown in the ower panel o [Figuree 5. alonge with data from ?.. 2.. 2.. 2.. 2.. ?. . aud ?..," These are shown in the lower panel of Figure \ref{fig:model0SN} along with data from \citet{reiss:00}, \citet{hardin:00}, \citet{pain:02}, \citet{madgwick:03}, \citet{tonry:03}, \citet{strolger:04}, and \citet{blanc:04}."
 Ouce again. the local (c Q) «lata is taken from ?..," Once again, the local $z = 0$ ) data is taken from \citet{cappellaro:99}."
 Unlike the type II data. there appears to be a definie peak iu the ate of tyye Ia supernovae around z~1.," Unlike the type II data, there appears to be a definite peak in the rate of type Ia supernovae around $z \sim 1$."
 As tl ep'ogenitors of type Ia supernovae are stars with intermediate masses beween 2 aud 5 L.. the rate for SN type Ia is controlled entirely by parameters within the normal mode of star ormation.," As the progenitors of type Ia supernovae are stars with intermediate masses between 2 and 8 $_\odot$, the rate for SN type Ia is controlled entirely by parameters within the normal mode of star formation."
 Tlere are however two additional parameters that must be adjusted to fi the observed data. namely he fraction of intermediate mass stars that explode as SN type Ia. aud he time delay )etween the cleath of the progenitor aud the SN explosion.," There are however two additional parameters that must be adjusted to fit the observed data, namely the fraction of intermediate mass stars that explode as SN type Ia, and the time delay between the death of the progenitor and the SN explosion."
 In most previous models of Galactic chemical evoltion. this delay was takento be of order 1 Cyr.," In most previous models of Galactic chemical evolution, this delay was takento be of order 1 Gyr."
 Recently. inany stujes havebeen," Recently, many studies havebeen"
A quite different proposal is that some of the irregular süiruceture arises [from unstable ballistic transport (Lissauer1984:Durisenetal.1989.1992:1995).,"A quite different proposal is that some of the irregular structure arises from unstable ballistic transport \citep{lis84,dur89,dur92,dur95}."
 hupacts by interplanetary particles erode the ring particles and redistribute (heir mass to adjacent annuli., Impacts by interplanetary particles erode the ring particles and redistribute their mass to adjacent annuli.
 This process is unstable because high-densitwv regions of the ring tend to absorb more of the ejecta than neighboring regions., This process is unstable because high-density regions of the ring tend to absorb more of the ejecta than neighboring regions.
 However. the growth rate of the instability is negligible for optical depths 21.5.," However, the growth rate of the instability is negligible for optical depths $\gtrsim
1.5$."
 Thus. ballistic transport is probably unable to explain ihe irregular structure in the bulk of the D ring.," Thus, ballistic transport is probably unable to explain the irregular structure in the bulk of the B ring."
 In (his section we investigate (he dynamics. stresses. and stability of solid. ancl liquid rine phases.," In this section we investigate the dynamics, stresses, and stability of solid and liquid ring phases."
 We examine a ring orbiting a point mass MM., We examine a ring orbiting a point mass $M$.
 Test particles on circular orbils travel at the Keplerian angular speed O&(r)=(GM/r*)2., Test particles on circular orbits travel at the Keplerian angular speed $\Omega_K(r)=(GM/ r^3)^{1/2}$.
" We employ a rotating Cartesiancoordinate system in which e, points radially outward. e, points in the direction of rotation. and the origin. at radius J. rotates around (he central mass al angular speed ο=O&(R)."," We employ a rotating Cartesiancoordinate system in which $\hat\bfe_x$ points radially outward, $\hat\bfe_y$ points in the direction of rotation, and the origin, at radius $R$, rotates around the central mass at angular speed $\Omega\equiv\Omega_K(R)$."
 We work al distances [from the origi (hat are small compared to the orbital radius £2 (Hills approximation)., We work at distances from the origin that are small compared to the orbital radius $R$ (Hill's approximation).
 We ignore the vertical structure of the disk. treating it as a razor-(hin sheet with surface density pi(x./).," We ignore the vertical structure of the disk, treating it as a razor-thin sheet with surface density $\mu(\bfx,t)$."
 We restrict ourselves to axisvmmetrie disturbances. so thal O/Oy— 0.," We restrict ourselves to axisymmetric disturbances, so that $\p/\p y=0$ ."
 The Euler and continnity equations then read, The Euler and continuity equations then read
X-ray Hovac are low-mass binary systems in which a compact object undergoes unstable aceretion from a late-tvpe companion star. releasing energy in the form of N-ravs.,"X-ray novae are low-mass binary systems in which a compact object undergoes unstable accretion from a late-type companion star, releasing energy in the form of X-rays."
 In à number of cases there is substantial evidence that the compact object is a black hole (see Tanaka Shibazaki 1996 and van Paradijs MeClintock. 1995 for reviews)., In a number of cases there is substantial evidence that the compact object is a black hole (see Tanaka Shibazaki 1996 and van Paradijs McClintock 1995 for reviews).
 Some of the best evidence for the presence of a black hole is obtained [rom the measurement of the orbital velocity of the companion star. leading to a determination of the mass function of the svstem and hence the minimum mass of the compact object.," Some of the best evidence for the presence of a black hole is obtained from the measurement of the orbital velocity of the companion star, leading to a determination of the mass function of the system and hence the minimum mass of the compact object."
 When this mass exceeds the maximum mass for a neutron star (~3.2 M.: Rhoacles luulini 1974). a black hole seems the only remaining possibility.," When this mass exceeds the maximum mass for a neutron star $\sim 3.2$ $\rm{M}_\odot$; Rhoades Ruffini 1974), a black hole seems the only remaining possibility."
 The X-ray nova CRO 11655.40 is one such system. and was discovered on July 27 1994 with BATSE on board the Compton Gamma Rav Observatory (Zhang et 11904).," The X-ray nova GRO J1655–40 is one such system, and was discovered on July 27 1994 with BATSE on board the Compton Gamma Ray Observatory (Zhang et 1994)."
 ]t has been studied extensively during the past two vears in X-ravs and at optical and radio wavelengths (Dailvn et 11995a and b. Zhang ct 11995. van der Llooft et 11998).," It has been studied extensively during the past two years in X-rays and at optical and radio wavelengths (Bailyn et 1995a and b, Zhang et 1995, van der Hooft et 1998)."
 Strong evidence that he compact object in. X-Ray Nova Sco is a black hole was presented. by DBailvn et ((1995b) who initially established a spectroscopic period of 2.601d:0.027 davs. and suggested a mass function. f(AM)=3.163:0.15 AL.," Strong evidence that the compact object in X-Ray Nova Sco is a black hole was presented by Bailyn et (1995b) who initially established a spectroscopic period of $2.601\pm0.027$ days, and suggested a mass function $f(M)=3.16\pm0.15$ ."
.. An improved. value of f(AM)=3.24x0.00 wwas presented bv Orosz ανα (1997). derived from a radial velocity semi-amplitude A»=228.2+2.2," An improved value of $f(M)=3.24\pm0.09$ was presented by Orosz Bailyn (1997), derived from a radial velocity semi-amplitude $K_{2}=228.2\pm2.2$."
 Vheir fitted: values of inclination and the mass ratio then implied a black hole mass of Ad;=7.01£0.22M., Their fitted values of inclination and the mass ratio then implied a black hole mass of $M_1=7.01\pm0.22$.
.. llowever. in calculating the racial velocity semi-amplitude. from which the mass function is derived. Orosz Dailvn (1997) used both quiescent. data. (taken in. 1996 February 2425) outburst data (taken in 995 April 30Alay 4). while Bailvn ct al. (," However, in calculating the radial velocity semi-amplitude, from which the mass function is derived, Orosz Bailyn (1997) used both quiescent data (taken in 1996 February 24–25) outburst data (taken in 1995 April 30--May 4), while Bailyn et al. ("
1996b) used just outburst data. ancl in both cases a sinusoidal fit was performed.,"1996b) used just outburst data, and in both cases a sinusoidal fit was performed."
 We suggest that using outburst data in this way may lead to an incorrect result., We suggest that using outburst data in this way may lead to an incorrect result.
 The effect. of substantial heating of the secondary can shift the celfective centre’ of the secondary. weighted by the strength of the absorption lines. from the centre of mass of the star. às described. in section 2.," The effect of substantial heating of the secondary can shift the `effective centre' of the secondary, weighted by the strength of the absorption lines, from the centre of mass of the star, as described in section 2."
 This results in a significant distortion of the radial velocity curve and renders a sinusoidal fit to be clearly inadequate. leading to a spuriouslv high racial velocity semi-amplitucde.," This results in a significant distortion of the radial velocity curve and renders a sinusoidal fit to be clearly inadequate, leading to a spuriously high radial velocity semi-amplitude."
 Phe masses of the binary components derived from this will therefore be incorrect., The masses of the binary components derived from this will therefore be incorrect.
 In this paper. we intend to consider only the outburst radial velocity data (from 1995 April 30.May 4).," In this paper, we intend to consider only the outburst radial velocity data (from 1995 April 30–May 4)."
 We use a model which incorporates the basic ellects of X-ray heating of the secondary. and. attempt to derive a more realistic range for the radial velocity. semi-amplitude.," We use a model which incorporates the basic effects of X-ray heating of the secondary, and attempt to derive a more realistic range for the radial velocity semi-amplitude."
. Using the orbital period and the range of inclinations obtained by van der Hooft et CC1998). we obtain (νο) solutions.," Using the orbital period and the range of inclinations obtained by van der Hooft et (1998), we obtain $K_{2},q$ ) solutions."
 Using, Using
Astrophysical jets and disks (Livio2009) seem to be inter-related. notably in the case of young stellar objects (YSOs). wherein jet signatures are well correlated with the infrared excess and accretion rate of the circumstellar disk (Cabritetal.1990:Hartiganetal. 2004).,"Astrophysical jets and disks \citep{Liv09} seem to be inter-related, notably in the case of young stellar objects (YSOs), wherein jet signatures are well correlated with the infrared excess and accretion rate of the circumstellar disk \citep{CES90, HEP04}."
. Disks provide the plasma which is outflowing in the jets. while Jets in turn provide the disk with the needed angular momentum removal in order that aceretion in the protostellar object takes place (Hartmann2009).," Disks provide the plasma which is outflowing in the jets, while jets in turn provide the disk with the needed angular momentum removal in order that accretion in the protostellar object takes place \citep{Har09}."
. On the theoretical front. the most widely accepted description of this accretion-ejection phenomenon (Ferreira2007) 1s based on the interaction of a large scale magnetic field with an accretion disk around the central object.," On the theoretical front, the most widely accepted description of this accretion-ejection phenomenon \citep{Fer07} is based on the interaction of a large scale magnetic field with an accretion disk around the central object."
 Then. plasma is channeled and magneto-centrifugally accelerated along the open magnetic field lines threading the accretion disk. as firstly described in Blandford&Payne(1982).," Then, plasma is channeled and magneto-centrifugally accelerated along the open magnetic field lines threading the accretion disk, as firstly described in \citet{BlP82}."
. Several works have extended this study either by semi-analytic models using radially self-similar solutions of the full magnetohydrodynamies (MHD) equations with the disk treated as a boundary condition (Vlahakis& 1998).. or. by selfconsistently treating numerically the disk-jet system (e.g.Zannietal.2007).," Several works have extended this study either by semi-analytic models using radially self-similar solutions of the full magnetohydrodynamics (MHD) equations with the disk treated as a boundary condition \citep{VlT98}, or, by selfconsistently treating numerically the disk-jet system \citep[e.g.][]{ZFR07}."
 The original Blandford&Payne(1982) model. however. has serious limitations for a heeded meaningful comparison of its predictions with observations.," The original \citet{BlP82} model, however, has serious limitations for a needed meaningful comparison of its predictions with observations."
 Singularities exist at the jet axis. the outflow is not asymptotically superfast. and most importantly. the disk has no intrinsic scale with the result that the jet formally extends to radial infinity. to mention just a few.," Singularities exist at the jet axis, the outflow is not asymptotically superfast, and most importantly, the disk has no intrinsic scale with the result that the jet formally extends to radial infinity, to mention just a few."
 The singularity at the axis can be easily taken care of by numerical simulations extending the analytical solutions close to this symmetry axis (Graciaetal.2006.GVTO6hereafter)., The singularity at the axis can be easily taken care of by numerical simulations extending the analytical solutions close to this symmetry axis \citep[][GVT06 hereafter]{GVT06}.
 The outflow speed at laree distances may be tuned to cross the corresponding limiting characteristic. with the result that the terminal wind solution is causally disconnected from the disk and hence perturbations downstream of the superfast transition (as modified by self-similarity) cannot affect the whole structure ofthe steady disk-wind outflow which has also been shown to be structurally stable (Matsakosetal.2008.MOShereafter).," The outflow speed at large distances may be tuned to cross the corresponding limiting characteristic, with the result that the terminal wind solution is causally disconnected from the disk and hence perturbations downstream of the superfast transition (as modified by self-similarity) cannot affect the whole structure ofthe steady disk-wind outflow which has also been shown to be structurally stable \citep[][M08 hereafter]{MTV08}."
 The next step of introducing a seale in the disk has been undertaken in à previous paper (Stuteetal.2008).. wherein we presented numerical simulations of truncated flows whose initial conditions are based on analytical self-similar models.," The next step of introducing a scale in the disk has been undertaken in a previous paper \citep{STV08}, wherein we presented numerical simulations of truncated flows whose initial conditions are based on analytical self-similar models."
 In order to test our truncated models. we ther applied our simulations to observations (Stuteetal.2010.paperIhereafter)...," In order to test our truncated models, we then applied our simulations to observations \citep[][paper I hereafter]{SGT10}."
 In recent years. many NIR and optical data have become available exploring the morphology and kinematics of the jet launching region (e.g.Dougadosetal.2000:Rayetal.2007:Dougados2008.andreferences therein)..," In recent years, many NIR and optical data have become available exploring the morphology and kinematics of the jet launching region \citep[e.g.][and references therein]{DCL00,RDB07,Dou08}."
 HST and adaptive optics observations give access to the innermost regions of the wind. where the acceleration and collimation occurs (Rayetal.1996;Dougados2000:Woitas2002:Hartiganetal. 2004).," HST and adaptive optics observations give access to the innermost regions of the wind, where the acceleration and collimation occurs \citep{RMD96,DCL00,WRB02,HEP04}."
. Since YSO jets emit in a number of atomic (and molecular) lines. we use a synthetic emission code to create emission maps in different forbidden lines which were used by other authors to extract the jet width from images.," Since YSO jets emit in a number of atomic (and molecular) lines, we use a synthetic emission code to create emission maps in different forbidden lines which were used by other authors to extract the jet width from images."
 —| paper 1. we compared the observed jet widths with those extracted from our synthetic images.," In paper I, we compared the observed jet widths with those extracted from our synthetic images."
 A similar study has been done by Ferreira(1997);Casse&(2000):Garciaetal.(2001);Dougadosetal. (2004).. however. using untruncated disk wind models.," A similar study has been done by \citet{Fer97,CaF00,GCF01,DCF04}, however, using untruncated disk wind models."
 In paper L. however. we always assumed the inclination of the jet to be 907.," In paper I, however, we always assumed the inclination of the jet to be $^\circ$."
" The influence of varying inclinations on the ""Serived synthetic jet widths will be examined now in this paper.", The influence of varying inclinations on the derived synthetic jet widths will be examined now in this paper.
 The remainder of the paper is organized as follows: we briefly review the initial set-up of the numerical simulations and describe our procedure for the comparisor with observations in Sec. 2.., The remainder of the paper is organized as follows: we briefly review the initial set-up of the numerical simulations and describe our procedure for the comparison with observations in Sec. \ref{sec_models}.
 The results of our studies are presented in Sec. 3., The results of our studies are presented in Sec. \ref{sec_results}.
 In Sec., In Sec.
 + we describe our best-fit models for each object in the sample and compare derived inclinations with literature values., \ref{sec_bestfit} we describe our best-fit models for each object in the sample and compare derived inclinations with literature values.
 Finally. we conclude with the implications of the results in terms of the structure of the disk and the respective launching radi of the jets m YSOs.," Finally, we conclude with the implications of the results in terms of the structure of the disk and the respective launching radii of the jets in YSOs."
 In appendix Appendix A:.. we present the extracted line widths for all models. runs and inclinations for the sake," In appendix \ref{sec_all_lw}, , we present the extracted line widths for all models, runs and inclinations for the sake"
" O=2 // D) = 0.025 + 0.005 mas derived from the flux fitting method (see Section ??)) to the stellar radius, yielding a distance D = 360 + 75 pe.","$\Theta$ = $2$ / $D$ = 0.025 $\pm$ 0.005 mas derived from the flux fitting method (see Section \ref{check}) ) to the stellar radius, yielding a distance $D$ = 360 $\pm$ 75 pc."
" It is reassuring that these values are similar, given that both methods rely on parameters [rom the stellar models and are therefore not independent."," It is reassuring that these values are similar, given that both methods rely on parameters from the stellar models and are therefore not independent."
" We take the distance to be the weighted mean of the two values, J) = 343 + 36 pc."," We take the distance to be the weighted mean of the two values, $D$ = 343 $\pm$ 36 pc."
 The stellar mass can also be determined through an empirical calibration betweenΤομ. παπά |Fe/H] found from eclipsing binaries (22)..," The stellar mass can also be determined through an empirical calibration between, and [Fe/H] found from eclipsing binaries \citep{Torres10, Enoch10}."
" This yields a larger mass, == 1.002 + 0.034 (than that [rom the evolutionary model fit, GT∶∶↻⋅≋↼∔⊖⋮⋮↕⋮⋮↓∎∣∣⋯∙⋅⋅"," This yields a larger mass, = 1.002 $\pm$ 0.034 than that from the evolutionary model fit, = $^{+0.067}_{-0.040}$."
"⋀↿⊃∪⋝⋝⋔∣∁∁≳⋯⋝∁↕⋝⊔⋯⇂∨⋁⋀≦↕⊃−∃∇↕⋝∪∏∐↧∁∁∟∣⇇↔⊺∁∪⇂⊔↧∁∟⋮≳⇂∣⋔⊺≳⋯∪∏⋁∖⇁↥↧∁↕⊲∁⊔↧∁↕⊲∁≳∐⊲∁∣⋯ ⋅ ⋅ ""AG ⋅ . ⋅ ⋝⇂≳∐⋪⋝∪⇂↸⋝↥⋯∐≳∐⋪⊺≺⋅∐⊳ ∐≳", A possible cause is that WASP-37 is on the edge of the calibration where there are no stars of similar and [Fe/H] and.
⋯∟∣∣∊∁∕∐∣≳⋯∟⊔∩∑∸⋅↙∕⋅⋅⋀⋝↥⋯∐≳∐⋪⋝∏∐∐⊔∪∏∖↖⇁≳⇂⋝∣↸∪⋃⋯∣⇂↸⋯⋪∨∨⊽⋀⊟↕⊃−⊋∣∼⋝∁," A similar situation was found for WASP-21, see \citet{Bouchy09}."
∁⇀∙≱⋅⋅⇂∏⋯⋪∟∣∁↕⋪⋯∪↻⇂≳∏∏≳⊔⋪∪↻∐⋝⊓⋪∁⋝⋃∥∙ ∖↖⇁∁↥↧≳⇂∖⇁∁⋯∣∖⋮∁∏⋔∁⋝⇂∁∥≳∐⊲⋯∐⋝⋝⋯↻∁⊔↕∁≳⇂∖⇁∁↕⊲≳⇂⇣↔⊺∁∪∣↸⊔↧∁⋯⇁∪∖⇁≳⇂∣⋃∁⋝⋀⊔⋆∶⋃⋅⊖⊋≤∶∶↻⋅∣⊋↻↓∖∉∙⋅," In order to obtain a robust result, we have taken the stellar mass to be the average of the two values = 0.925 $\pm$ 0.120."
"⋅ We used archival magnitudes from 2MASS, DENIS. NOMAD. TASS and CMCI4 to derive an independent value of the stellar temperature using the flux-fitüing method (?).."," We used archival magnitudes from 2MASS, DENIS, NOMAD, TASS and CMC14 to derive an independent value of the stellar temperature using the flux-fitting method \citep{Morossi85}."
" We find Togas = 5840 + 140 K, in excellent agreement with the temperature derived [rom spectral analysis."," We find $_{\rm flux}$ = 5840 $\pm$ 140 K, in excellent agreement with the temperature derived from spectral analysis."
 Independent measures of aand |Fe/H]| can be derived from the shape of the spectral cross correlation function (CCF)., Independent measures of and [Fe/H] can be derived from the shape of the spectral cross correlation function (CCF).
" Using the empirical calibrations of ?. for the SOPHIE HE mode, and inputing Contrast = 4.4 + and FWHM = 9.48 4- 0.06 Ifrom the SOPHIE CCF and B—V = 0.60 + 0.05 (calculatedusing2)., we find |Fe/H]ecj = —0.32 + 0.17 and esinvcop 2 3.2 + 10°."," Using the empirical calibrations of \citet{Boisse10} for the SOPHIE HE mode, and inputing Contrast = 4.4 $\pm$ and FWHM = 9.48 $\pm$ 0.06 from the SOPHIE CCF and $B-V$ = 0.60 $\pm$ 0.05 \citep[calculated using][]{Casagrande10}, we find $_{\rm CCF}$ = $-$ 0.32 $\pm$ 0.17 and $_{\rm CCF}$ = 3.2 $\pm$ 1.0."
. These are also in excellent agreement with the values found [rom spectral analysis., These are also in excellent agreement with the values found from spectral analysis.
 The stellar evolution modelling provides colour indices which can be compared to observational values., The stellar evolution modelling provides colour indices which can be compared to observational values.
 We corrected the2MASS magnitudes [or exünction using A(V) 23.1 xE(B—V). appropriately scaled for each passband following ?..," We corrected the2MASS magnitudes for extinction using $A(V)$ = 3.1 $\times E(B-V)$, appropriately scaled for each passband following \citet*{Cardelli89}."
 The model colours were transformed [rom the ESO tothe 2MASS system following ?.., The model colours were transformed from the ESO tothe 2MASS system following \citet{Carpenter01}.
" The 2MASS and model colours agree well, as shown in Table 3.."," The 2MASS and model colours agree well, as shown in Table \ref{colors}."
" The large relative error on pprovides Πο constraint on the maximum stellar rotation period Po,«2zA/esin?;<21i d. or gyrochronological age; 1—28 Gyr (?).."," The large relative error on provides little constraint on the maximum stellar rotation period $P_{\rm rot} < 2 \pi R_{*} / v \sin i < 21^{+44}_{-9}$ d, or gyrochronological age; 1–28 Gyr \citep{Barnes07}."
" Equally, the lithium abundance only implies a lower limit on the age of 1-2 Gyr (2)."," Equally, the lithium abundance only implies a lower limit on the age of 1–2 Gyr \citep*{SestitoRandich05}."
" Given the low metallicity of WASP-37, we investigated whether it could be a member of the thick disk using the method outlined in ?.."," Given the low metallicity of WASP-37, we investigated whether it could be a member of the thick disk using the method outlined in \citet{Bouchy10}. ."
" We calculated the space velocities according to ?., using the proper motions,"," We calculated the space velocities according to \citet{Johnson87}, , using the proper motions,"
resppectivelv.,pectively.
In the past. this criterion. was used by Luvten to search for wide svstems by visual inspection of photographic plates: his prospect was performed over more than 40 vears. from Luvten(1940). to Luvten(1987) and led to the discovery of 6121. systems.,"In the past, this criterion was used by Luyten to search for wide systems by visual inspection of photographic plates; his prospect was performed over more than 40 years, from \citet{Luyten40}
 to \citet{Luyten87} and led to the discovery of 6121 systems."
 However. despite their. large number. these systenis were not used to derive the properties of wide binaries: the main reason is that the selection was based on subjective criteria. and. as à Consequence. it would be hard to estimate its inconipleteness. ancl also the contamination by optical pairs.," However, despite their large number, these systems were not used to derive the properties of wide binaries; the main reason is that the selection was based on subjective criteria, and, as a consequence, it would be hard to estimate its incompleteness, and also the contamination by optical pairs."
 Moreover. since the components of Luxten's double stars are rather faint (around 16th magnitude). it would be cillicult. to. improve the selection with complementary data.," Moreover, since the components of Luyten's double stars are rather faint (around 16th magnitude), it would be difficult to improve the selection with complementary data."
 Another selection of CPAL pairs was performed. by Lalbwachs(1986).. on the basis of the ACIX2/3 catalogue (Lacroute&Valbousquet.1974).. Four hundred and. thirty-nine CPM double stars were selected. and. the number of optical pairs was estimated to be around 40.," Another selection of CPM pairs was performed by \citet{Halbwachs86}, on the basis of the AGK2/3 catalogue \citep{Lacroute74}.. Four hundred and thirty-nine CPM double stars were selected, and the number of optical pairs was estimated to be around 40."
 Since the ACH stars are brighter than 12th. magnitude. it was decided to select the physical binaries on the basis of the RY of the components.," Since the AGK3 stars are brighter than 12th magnitude, it was decided to select the physical binaries on the basis of the RV of the components."
 As is the case for the proper motions. the RY are mareinally allected by the orbital motion. and the components of wide binaries are therefore also common IV stars.," As is the case for the proper motions, the RV are marginally affected by the orbital motion, and the components of wide binaries are therefore also common RV stars."
 Two hundred. and. sixtv-six stars were selected. and they were measured with the spectrovelocimeterCORAVEL.," Two hundred and sixty-six stars were selected, and they were measured with the spectrovelocimeter."
 This programme was initiated in 1986. and it was initiallv supposed. to end after a few vears.," This programme was initiated in 1986, and it was initially supposed to end after a few years."
 However. several stars had. variable velocities. since these wide pair components were themselves spectroscopic binarics (SB).," However, several stars had variable velocities, since these wide pair components were themselves spectroscopic binaries (SB)."
 It was then necessary to extend the observations over about 20 vears. in order to derive the SB orbital elements. ancl thus the systemic velocity of these stars.," It was then necessary to extend the observations over about 20 years, in order to derive the SB orbital elements, and thus the systemic velocity of these stars."
 The present paper is devoted to these variable velocity stars. and to the derivation of their SB elements.," The present paper is devoted to these variable velocity stars, and to the derivation of their SB elements."
 The selection of the wide binaries and. the derivation of their statistical properties will be treated. in a second. paper., The selection of the wide binaries and the derivation of their statistical properties will be treated in a second paper.
 The paper is organized as follows: theRV observations are presented in section 2.. and the catalogue of the RV nieasurements is in section 3...," The paper is organized as follows: theRV observations are presented in section \ref{sec:RV}, and the catalogue of the RV measurements is in section \ref{sec:RVcat}."
 The calculation of the SB orbits is in section 4: for the doublelined SB (8D2). we present a method to take into account the measurements related to blends of the two components.," The calculation of the SB orbits is in section \ref{sec:SB}; for the double–lined SB (SB2), we present a method to take into account the measurements related to blends of the two components."
 Some interesting points relating to the new SB that we have discovered. are discussed in the conclusion. section 5..," Some interesting points relating to the new SB that we have discovered are discussed in the conclusion, section \ref{sec:conclusion}."
 The observational programme concerned a Iarge part of the CPM stars listed in Halbwachs(1986)., The observational programme concerned a large part of the CPM stars listed in \citet{Halbwachs86}.
. In that paper. the CPM stars were presented in two tables. according to the time Lo—6/p. where @ is the apparent separation and sr is the proper motionof the pair of stars.," In that paper, the CPM stars were presented in two tables, according to the time $T=\theta/\mu$, where $\theta$ is the apparent separation and $\mu$ is the proper motionof the pair of stars."
 Three hundred and twenty-six pairs of CPM stars with 2«1000 wears are Listed in the first table. and the estimated frequency of optical pairs is only 1.3.," Three hundred and twenty-six pairs of CPM stars with $T < 1000$ years are listed in the first table, and the estimated frequency of optical pairs is only 1.3."
%.. Phe second table contains 119 CPAL pairs with J between 1000 ancl 3500 vears: this range corresponcls to an expected frequency. of optical pairs of 40%., The second table contains 113 CPM pairs with $T$ between 1000 and 3500 years; this range corresponds to an expected frequency of optical pairs of 40.
. Since the proportion of optical pairs is rather important in the second table. all these 113 CPAL pairs were selected in a first step.," Since the proportion of optical pairs is rather important in the second table, all these 113 CPM pairs were selected in a first step."
 For the first table. it was decided to observe only the components of pairs with a primary brighter than 9.15 mag.," For the first table, it was decided to observe only the components of pairs with a primary brighter than 9.15 mag."
 This limit was chosen since it appeared that the Cil 2/3 is not complete bevond that magnitude., This limit was chosen since it appeared that the AGK 2/3 is not complete beyond that magnitude.
 The speetrovelocimeter (Baranneetal.1979) is installed on the Swiss Lena telescope. at the Observatory of Llaute-Provence (OLLP).," The spectrovelocimeter \citep{Baranne79} is installed on the Swiss 1-m telescope, at the Observatory of Haute-Provence (OHP)."
 For slow rotators with spectral νρος around WO. it may. provide RY measurements with a srecision around 0.8 Kms. Llowever. it becomes inellicient or stars earlier than F5. except when their spectra contain metallic tines.," For slow rotators with spectral types around K0, it may provide RV measurements with a precision around 0.3 km/s. However, it becomes inefficient for stars earlier than F5, except when their spectra contain metallic lines."
 Therefore. it is possible to get BN. with a precision near 1 km/s for the Amtype stars.," Therefore, it is possible to get RV with a precision near 1 km/s for the Am–type stars."
 Due to his restriction. our final programme contains 266 stars observable withCORAVEL: 00 are extracted [rom the first able. anc 176 [rom the second onc.," Due to this restriction, our final programme contains 266 stars observable with: 90 are extracted from the first table, and 176 from the second one."
 The programme started. in. 1986. but several stars were already measured. with at this time.," The programme started in 1986, but several stars were already measured with at this time."
 The observations of the stars with a constant velocity were stopped. after 1 or 2 vears. but several SB were detected. and i was necessary to prolong the programme in order to cover their periods and to derive their orbital elements.," The observations of the stars with a constant velocity were stopped after 1 or 2 years, but several SB were detected, and it was necessary to prolong the programme in order to cover their periods and to derive their orbital elements."
 In practice. when a SB was observed. its wide companion was usually observed too.," In practice, when a SB was observed, its wide companion was usually observed too."
 Lhe observations were performed until the decomamissioning of that instrument in 2000., The observations were performed until the decommissioning of that instrument in 2000.
 Thev led to the detection of 66 stars with variable radial velocity. according to the (47) test at the 1% threshold.," They led to the detection of 66 stars with variable radial velocity, according to the $P(\chi^2)$ test at the 1 threshold."
 In 2007. we had the opportunity to obtain very accurate RV measurements with the spectrograph installed on the 1.93 m telescope at OLLP (Perruchotetal.2008:3ouchvetal. 2009).," In 2007, we had the opportunity to obtain very accurate RV measurements with the spectrograph installed on the 1.93 m telescope at OHP \citep{Perruchot08, Bouchy09}."
. Ten stars with variable RV. but. for which it was still not. possible to derive an orbit. were observed in during 2 semesters.," Ten stars with variable RV, but for which it was still not possible to derive an orbit, were observed in during 2 semesters."
 The spectra were registered. in high.resolution mode with a signal-to-noise ratio between 20 ancl 200. depending on the brightness of the star.," The spectra were registered in high–resolution mode with a signal-to-noise ratio between 20 and 200, depending on the brightness of the star."
 Fhanks to the automatic pipeline. hey were crosscorrelated with a mask close to the actual spectral type of the star.," Thanks to the automatic pipeline, they were cross–correlated with a mask close to the actual spectral type of the star."
 For the single-lined binaries (SB1). he RV was directly obtained. [rom the pipeline.," For the single-lined binaries (SB1), the RV was directly obtained from the pipeline."
 For two SB2. the crosscorrelation function (CCE) was fitted with wo normal distributions in order to derive the velocities of th components.," For two SB2, the cross–correlation function (CCF) was fitted with two normal distributions in order to derive the velocities of both components."
 The automatic pipeline provides also RV uncertainties. xut they are obviously underestimated.," The automatic pipeline provides also RV uncertainties, but they are obviously underestimated."
 Boisseetal.(2010) estimate that 3 terms must be quadraticeally: added: the instrumental drift is around 3 m the guiding error is. around 4 m 1 and they cerivecl. from⋅ their. own observations an additional error of S m ," \citet{Boisse10} estimate that 3 terms must be quadratically added: the instrumental drift is around 3 m $^{-1}$, the guiding error is around 4 m $^{-1}$, and they derived from their own observations an additional error of 8 m $^{-1}$."
"We finally obtain an additional error of 9.4 m s.+. that is rounded upto lO mis.ο,"," We finally obtain an additional error of 9.4 m $^{-1}$, that is rounded up to 10 m $^{-1}$."
 Ht is worth noticing that this error does not apply to the radial velocity with respect to the Sun. but to variationsof RV due to an exoplanetary companion. je to RY Iluctuations over a range not larger than around 100 ni," It is worth noticing that this error does not apply to the radial velocity with respect to the Sun, but to variationsof RV due to an exoplanetary companion, ie to RV fluctuations over a range not larger than around 100 m $^{-1}$."
 Since we are observing SD with stellar components. the amplitude of the RV is as large as several km s. +. and the error could be larger than 10 m 1 .," Since we are observing SB with stellar components, the semi-amplitude of the RV is as large as several km $^{-1}$ , and the error could be larger than 10 m $^{-1}$ ."
 Moreover. we must introduce an olfset between the measurements and the measurements.," Moreover, we must introduce an offset between the measurements and the measurements."
 TFhis olfset depends on the star.," This offset depends on the star,"
"individual credible intervals, are of course very prior-dependent which again indicates that the data are not very discriminatory with respect to the prior.","individual credible intervals, are of course very prior-dependent which again indicates that the data are not very discriminatory with respect to the prior."
" On the other hand, rather drastic priors or small sub-samples must be used in order to significantly affect the credible intervals of the joint posterior, especially for the hyper-parameters method."," On the other hand, rather drastic priors or small sub-samples must be used in order to significantly affect the credible intervals of the joint posterior, especially for the hyper-parameters method."
" Figure 4 shows the individual clusters’ constraints on a, y, and n."," Figure \ref{fi:zdep} shows the individual clusters' constraints on $\alpha$, $\gamma$, and $n$."
" As could be expected given the varying nature of our results, there is perhaps the slightest of hints of a redshift-dependence in the constraints but our sample size does not allow us to probe such an issue in detail."," As could be expected given the varying nature of our results, there is perhaps the slightest of hints of a redshift–dependence in the constraints but our sample size does not allow us to probe such an issue in detail."
" It should again be noted that any hint of a redshift-dependence could actually be caused by a mass—dependence instead, since the two lowest redshift clusters in our sample are also the least massive."," It should again be noted that any hint of a redshift–dependence could actually be caused by a mass--dependence instead, since the two lowest redshift clusters in our sample are also the least massive."
 A different picture emerges when we consider the overlap of the individual clusters’ credible intervals for the slopeNFW model., A different picture emerges when we consider the overlap of the individual clusters' credible intervals for the slopeNFW model.
" For example, no value of a is contained in all 11 credible intervals, and only the very short range (1.08,1.10) is contained in all but two intervals."," For example, no value of $\alpha$ is contained in all 11 credible intervals, and only the very short range $(1.08,1.10)$ is contained in all but two intervals."
 Likewise the NFW a=1 case is excluded from four of the eleven intervals., Likewise the NFW $\alpha=1$ case is excluded from four of the eleven intervals.
" These results, as well as a visual inspection of Figure 3,, puts strong doubts about the concept of a universal shape parameter."," These results, as well as a visual inspection of Figure \ref{fi:third}, puts strong doubts about the concept of a universal shape parameter."
 The situation is not quite as compelling for the transNFW and Sérrsic models which is likely the reason that they do not stand out from the NFW in the model selection., The situation is not quite as compelling for the transNFW and Sérrsic models which is likely the reason that they do not stand out from the NFW in the model selection.
" In fact, we believe it is a reasonable statement that the success of the slopeNFW model is precisely due to the very different preferred values of o from cluster to cluster."," In fact, we believe it is a reasonable statement that the success of the slopeNFW model is precisely due to the very different preferred values of $\alpha$ from cluster to cluster."
 This puts a strong question mark against the idea of universal third parameter., This puts a strong question mark against the idea of universal third parameter.
" We conclude that there is moderate evidence for the slopeNFW model to be preferred over the simple NFW, while the transNFW and Sérrsic models do not stand out against the two-parameter NFW profile."," We conclude that there is moderate evidence for the slopeNFW model to be preferred over the simple NFW, while the transNFW and Sérrsic models do not stand out against the two-parameter NFW profile."
" If the inner slope of the slopeNFW model is universal, we constrain it to be close to —1 but preferably slightly steeper."," If the inner slope of the slopeNFW model is universal, we constrain it to be close to $-1$ but preferably slightly steeper."
" However, the spread of the individual clusters' preferred ranges suggests that the inner slope is not universal."," However, the spread of the individual clusters' preferred ranges suggests that the inner slope is not universal."
 We also comment that the method of hyper-parameters method could in principle be extended to the model selection analysis., We also comment that the method of hyper-parameters method could in principle be extended to the model selection analysis.
" As a matter of fact, since the slopeNFW ‘contains’ both the NFW and the Moore models as subsets, we can derive the corresponding Bayes factor for the Moore profile which is only 10."," As a matter of fact, since the slopeNFW `contains' both the NFW and the Moore models as subsets, we can derive the corresponding Bayes factor for the Moore profile which is only 10."
 This is of course a drastic reduction numerically but it does not alter the conclusion and anyway corresponds to rather convincing odds of about 20000:1.," This is of course a drastic reduction numerically but it does not alter the conclusion and anyway corresponds to rather convincing odds of about $20\,000:1$."
 So far we have discussed the interpretation of our results with respect to the statistical evidence., So far we have discussed the interpretation of our results with respect to the statistical evidence.
" However, a number of biases, or systematic uncertainties, can be thought of that may affect our results."," However, a number of biases, or systematic uncertainties, can be thought of that may affect our results."
" Loosely, these can be grouped into biases that affect both the individual cluster mass modeling and the combined analysis, and selection effects that only influence the latter."," Loosely, these can be grouped into biases that affect both the individual cluster mass modeling and the combined analysis, and selection effects that only influence the latter."
 'The analysis rests on the ability to produce deprojected temperature and density profiles with uncertainties that correctly mirror the uncertainties in the spectral analysis of the X-ray data., The analysis rests on the ability to produce deprojected temperature and density profiles with uncertainties that correctly mirror the uncertainties in the spectral analysis of the X-ray data.
 This has been discussed extensively in Kaastraetal., This has been discussed extensively in \citet{2004A&A...413..415K}.
" The basic assumption in determining the mass(2004).. distribution of a galaxy cluster is that the cluster is relaxed, and obeys the equation of hydrostatic equilibrium."," The basic assumption in determining the mass distribution of a galaxy cluster is that the cluster is relaxed, and obeys the equation of hydrostatic equilibrium."
" Numerical simulations indicate that the additional pressure associated with turbulence and bulk motion in the ICM yields an underestimate of the mass in the region of 5—2096 with the larger values corresponding to large radii, Τ50ο and greater (Nagaietal.2007;Piffaretti&Valdarnini2008;Lau 2009).."," Numerical simulations indicate that the additional pressure associated with turbulence and bulk motion in the ICM yields an underestimate of the mass in the region of $5-20$ with the larger values corresponding to large radii, $r_{500}$ and greater \citep{2007ApJ...655...98N,2008A&A...491...71P,2009arXiv0903.4895L}. ."
 We do not expect this bias to exceed in the present case since we do not model further out than to ~2500., We do not expect this bias to exceed in the present case since we do not model further out than to $\sim r_{2500}$.
" On the other hand, the same numerical simulations indicate that if the turbulent pressure is accounted for, an accurate mass reconstruction is possible."," On the other hand, the same numerical simulations indicate that if the turbulent pressure is accounted for, an accurate mass reconstruction is possible."
 This point demonstrates that deviations from spherical symmetry are not a major concern in the error budget., This point demonstrates that deviations from spherical symmetry are not a major concern in the error budget.
 A related question is whether the parameterized profiles should be tested against the total mass distribution or the dark matter mass profile only., A related question is whether the parameterized profiles should be tested against the total mass distribution or the dark matter mass profile only.
" While the predictions of numerical simulations are founded in dark matter-only simulations, it is not clear how much a simulated dark matter-only mass profile is modified by the presence of baryons."," While the predictions of numerical simulations are founded in dark matter-only simulations, it is not clear how much a simulated dark matter-only mass profile is modified by the presence of baryons."
" Observationally, the ICM contributes about 5—15% of the total density in a cluster, again increasing with radius in the range of interest here, so formally there is a difference between the total and the dark matter profile's radial dependence."," Observationally, the ICM contributes about $5-15$ of the total density in a cluster, again increasing with radius in the range of interest here, so formally there is a difference between the total and the dark matter profile's radial dependence."
" To test the impact of this, we have rerun the statistical analyses described above without subtracting the ICMmass from the mass estimate of Equation (3).."," To test the impact of this, we have rerun the statistical analyses described above without subtracting the ICMmass from the mass estimate of Equation ."
 We find only minor, We find only minor
catalog from these nuaging data are detailed by (2008).,catalog from these imaging data are detailed by .
. Accurate redshift information is fundiunentallv imuportant for clustering studies., Accurate redshift information is fundamentally important for clustering studies.
 Photometric redshifts were calculated using the EAZY software //www., Photometric redshifts were calculated using the EAZY software .
astro.vale.edufeazy/.. EAZY fts lear combinations of galaxy templates to the observed photometry using a Bayesian prior for the distribution of redshifts as a function of apparent magnitude., EAZY fits linear combinations of galaxy templates to the observed photometry using a Bayesian prior for the distribution of redshifts as a function of apparent magnitude.
 The teiiplate set was carefully chosen to provide high-quality photometric redshifts for A-selected galaxies iu several current deep surveys., The template set was carefully chosen to provide high-quality photometric redshifts for $K$ -selected galaxies in several current deep surveys.
 To assess the aceuracy of our photometric redshifts we colupare to a sape 119 spectroscopic redshifts available in this field2008)., To assess the accuracy of our photometric redshifts we compare to a sample 119 spectroscopic redshifts available in this field.
. The nornalized median absolute deviation of A:(1|i) Is ONALp=0.033., The normalized median absolute deviation of $\Delta z/(1+z)$ is $\sigma_{\rm NMAD} = 0.033$.
 As nearly all spectroscopic redshifts iu this fieId are at 2<1.5. and none are for DRGs. we use pubic data on the well-studied Chandra Deep Field-South {CDE-S) to obtain an estimate of the expected plotometric redshift accuracy for the ealaxics currently under consideration.," As nearly all spectroscopic redshifts in this field are at $z<1.5$, and none are for DRGs, we use public data on the well-studied Chandra Deep Field-South (CDF-S) to obtain an estimate of the expected photometric redshift accuracy for the galaxies currently under consideration."
 We use the photometric catalog of (2008).. and exclude the CV7T filters in order to match the filter set available in the UDS.," We use the photometric catalog of , and exclude the $UVH$ filters in order to match the filter set available in the UDS."
 Restricting the sample to the DRCs in CDF-S with spectroscopic redshifts. we fud txapyp=0.065. with no appareut dependence on DRC redshift aud uceleible systematic offsets.," Restricting the sample to the DRGs in CDF-S with spectroscopic redshifts, we find $\sigma_{\rm NMAD} = 0.065$, with no apparent dependence on DRG redshift and negligible systematic offsets."
 This sugeests that the DRG photometric redshifts in our study also lave an accuracy of OAD~0.060.07. although this can ouly be taken as a rough indication.," This suggests that the DRG photometric redshifts in our study also have an accuracy of $\sigma_{\rm NMAD} \sim 0.06-0.07$, although this can only be taken as a rough indication."
 The overall redshift distibutiou of the sunuple is used to deproject the observed angular correlation fiction in order to inter the spatial correlation function., The overall redshift distribution of the sample is used to deproject the observed angular correlation function in order to infer the spatial correlation function.
 This distribution is often estimated simply using a histogram of the photometric redshifts. without taking iuto account the redshift uucertaimties.," This distribution is often estimated simply using a histogram of the photometric redshifts, without taking into account the redshift uncertainties."
 Because we select DRGs in a specific photometric redshift range. if is expected +rat randon errors will cause galaxies to scatter into the sample from both lower aud higher redshifts.," Because we select DRGs in a specific photometric redshift range, it is expected that random errors will cause galaxies to scatter into the sample from both lower and higher redshifts."
 It is also expected that galaxies within the redshift sclection window will scatter out. however this will not affect the clusteriug so long as such galaxies are drawn randomly roni the sample (svstematic errors which cause galaxies o scatter into or out of the selection window are much nore problematic).," It is also expected that galaxies within the redshift selection window will scatter out, however this will not affect the clustering so long as such galaxies are drawn randomly from the sample (systematic errors which cause galaxies to scatter into or out of the selection window are much more problematic)."
 To estimate the redshift distribution of the galaxies in our sample. we use Monte Carlo simulations iu which the observed photometry is serturbed according the the photometric uncertainties2007a).," To estimate the redshift distribution of the galaxies in our sample, we use Monte Carlo simulations in which the observed photometry is perturbed according the the photometric uncertainties."
. This procedure broadens he distzxibution relative to what would be inferred from a ustoeraim of photometric redshifts. resulting iu a higher correlation leueth: we return to this point in 6..," This procedure broadens the distribution relative to what would be inferred from a histogram of photometric redshifts, resulting in a higher correlation length; we return to this point in \ref{sec:discussion}."
 For a detailed description of the techuiques used to perform correlation function measurements. sce(2007a).," For a detailed description of the techniques used to perform correlation function measurements, see."
. Tere we brieflv describe the method. and point out differences between the method used here aud in the previous paper.," Here we briefly describe the method, and point out differences between the method used here and in the previous paper."
 The aneular two-point correlation function is calculated using the estimator., The angular two-point correlation function is calculated using the estimator.
 Uucertaiuties are estimated using bootstrap resampling. which vields uncertainties that are sienificautly lavecr than those expected frou Poisson statistics.," Uncertainties are estimated using bootstrap resampling, which yields uncertainties that are significantly larger than those expected from Poisson statistics."
 We follow the method of to correct for the integral coustraimt. although the correction is πα] eiven the large size of the field.," We follow the method of to correct for the integral constraint, although the correction is small given the large size of the field."
 A potentially serious problem for measurements of the auto-correlation function arises from variable scusitivity or calibration errors across tlie nuages., A potentially serious problem for measurements of the auto-correlation function arises from variable sensitivity or calibration errors across the images.
 We ieasured he correlation function of stars in order to test whether such variations induce artificial clustering. obtaining the expected result that stars are uuclustered.," We measured the correlation function of stars in order to test whether such variations induce artificial clustering, obtaining the expected result that stars are unclustered."
 We have also neasured the clusteringof DRGs selected from catalogs in which the J and A zeropoiuts are varied within their uncertainties Independently for each of the poiutiugs that nake up the UDS nosaic., We have also measured the clusteringof DRGs selected from catalogs in which the $J$ and $K$ zeropoints are varied within their uncertainties independently for each of the pointings that make up the UDS mosaic.
 Althoueh it is difficult to rule out significant envs iu the correlation functiou arising roni such variatkAm. it appears that they are not a dominant source o:OITOI.," Although it is difficult to rule out significant errors in the correlation function arising from such variations, it appears that they are not a dominant source of error."
 The redshift distribution of galaxies selected using the DRG color critcrioLis rather wide. extending from 2~1 to at least 2~3.52007b)..," The redshift distribution of galaxies selected using the DRG color criterion is rather wide, extending from $z \sim 1$ to at least $z
\sim 3.5$."
 We therefore pose the additional criterion 2Xtpher<5., We therefore impose the additional criterion $2 < z_{\rm phot} < 3$.
 Over this redshift ranec. the DRG criterion effectively selects galaxies with red rest-frame optical colors2007h).," Over this redshift range, the DRG criterion effectively selects galaxies with red rest-frame optical colors."
. We also limit the siuuple to objects with A.«21. ensuing that the catalog is complete and that objects have a sufficicut sisual-to-noise ratio to calculate accurate photometric redshifts.," We also limit the sample to objects with $K<21$, ensuring that the catalog is complete and that objects have a sufficient signal-to-noise ratio to calculate accurate photometric redshifts."
 We visually inspected each object in the sample. rejecting those that appeared to be artifacts in the images or were otherwise deemed to have unreliable photometry.," We visually inspected each object in the sample, rejecting those that appeared to be artifacts in the images or were otherwise deemed to have unreliable photometry."
 Objects were also rejected if the templates used for photometric redshifts gave poor fits., Objects were also rejected if the templates used for photometric redshifts gave poor fits.
 The final sample cousists of 1528 DRGs with A«21., The final sample consists of 1528 DRGs with $K<21$.
 Figure 1 shows the augulu correlation function (0) of DRGs., Figure \ref{fig:acf} shows the angular correlation function $w(\theta)$ of DRGs.
 Our measurements coufirin the strong aueular clustering previously reported byand(2007a).. but with reduced nucertaimtics and extending to much larger angular scales.," Our measurements confirm the strong angular clustering previously reported byand, but with reduced uncertainties and extending to much larger angular scales."
 Iuunediatelv apparent is the departure of w(t) πο a power law on small scales., Immediately apparent is the departure of $w(\theta)$ from a power law on small scales.
" To illustrate the sienificance of this excess clustering signal. aud to show how the results differ from what would be measured using a suuler field. we fit a power law μή7 over the range 2""cLO” aud over LO”<0«500""."," To illustrate the significance of this excess clustering signal, and to show how the results differ from what would be measured using a smaller field, we fit a power law $A_w\theta^{-\beta}$ over the range $2\arcsec < \theta
<40\arcsec$ and over $40\arcsec < \theta <500\arcsec$."
 The former fit vields ανν.ο)=(128.1.20.3) while the latter fit vields (Ay.3)=(1.12EO.8.0.170.11).," The former fit yields $(A_w,\beta) = (12 \pm 8, 1.2 \pm 0.3)$ while the latter fit yields $(A_w,\beta) = (1.1 \pm 0.8, 0.47 \pm 0.14)$."
" Note that there is sienificaut covariance between the two fitting parameters. so the uncertainties should be taken as represeutative only,"," Note that there is significant covariance between the two fitting parameters, so the uncertainties should be taken as representative only."
 As noted by previous authors2007a).. the larec- fit is preferred when comparing galaxy samples. as it better probes the clustering of dark matter halos that host the galaxies.," As noted by previous authors, the large-scale fit is preferred when comparing galaxy samples, as it better probes the clustering of dark matter halos that host the galaxies."
 The bottom panel of Figure 1 shows the residuals (0) from the large-scale power law., The bottom panel of Figure \ref{fig:acf} shows the residuals $w(\theta)$ from the large-scale power law.
" A simallscale excess has also been observed in large samples of optically selected ealaxies athieh redshift 20061, and for various galaxy samples at 0S2S2 2005)."," A small-scale excess has also been observed in large samples of optically selected galaxies athigh redshift , and for various galaxy samples at $0 \lesssim z \lesssim 2$ ."
 The spatial correlation leneth ry can be estimated from (0) using the Limber projection: see e.g., The spatial correlation length $r_0$ can be estimated from $w(\theta)$ using the Limber projection; see e.g.
The Hall effect is now acknowledged to be an important mechanism in the evolution of magnetic fields in the crusts of neutron stars.,The Hall effect is now acknowledged to be an important mechanism in the evolution of magnetic fields in the crusts of neutron stars.
 As derived by?.. the equation governing the magnetic field under the influence of both the Hall effect and Ohmic diffusion ts in which length is scaled by some characteristic length such as the depth 4 of the crust. time scaled by the Ohmic decay time πια/c7. where o is the conductivity. and c is the speed of light.," As derived by\cite{goldreich92}, the equation governing the magnetic field under the influence of both the Hall effect and Ohmic diffusion is in which length is scaled by some characteristic length such as the depth $d$ of the crust, time scaled by the Ohmic decay time $4 \pi \sigma
d^2 / c^2$, where $\sigma$ is the conductivity, and $c$ is the speed of light."
 The magnetic field is defined so that the Ohmic term. V-B. and the Hall term. -Vx|(VxB)B]. are formally of the same order: this is accomplished by scaling B by nec/o. where n is the electron number density. and e 1s the electron charge.," The magnetic field is defined so that the Ohmic term, $\nabla^2 {\bf B}$, and the Hall term, $-\nabla\times\left[(\nabla\times{\bf B})\times{\bf B}\right]$, are formally of the same order; this is accomplished by scaling $\bf B$ by $nec/\sigma$, where $n$ is the electron number density, and $e$ is the electron charge."
 We note though that the true magnetic field in neutron stars is typically several orders of magnitude greater than this. and correspondingly the Hall timescale is several orders of magnitude faster than the Ohmic timescale.," We note though that the true magnetic field in neutron stars is typically several orders of magnitude greater than this, and correspondingly the Hall timescale is several orders of magnitude faster than the Ohmic timescale."
 If Bo ts the non-dimensional amplitude of B in Eq. (1)).," If $B_0$ is the non-dimensional amplitude of $\bf B$ in Eq. \ref{eq:A}) ),"
 and ¢ is measured on the Ohmic timescale. then t=Bot is measured on the Hall timescale.," and $t$ is measured on the Ohmic timescale, then $t'=B_0 t$ is measured on the Hall timescale."
 We use both timescales in this work. as appropriate.," We use both timescales in this work, as appropriate."
 Arguing by analogy with ordinary. hydrodynamic turbulence. Goldreich and Reisenegger then conjectured that Eq. (1))," Arguing by analogy with ordinary, hydrodynamic turbulence, Goldreich and Reisenegger then conjectured that Eq. \ref{eq:A}) )"
" would generate a turbulent Hall cascade. transferring energy from large to small scales. with an energy spectrum FE,«&7 (where & is wavenumber). and a dissipative cutoff occurring at -Bo."," would generate a turbulent Hall cascade, transferring energy from large to small scales, with an energy spectrum $E_k \propto k^{-2}$ (where $k$ is wavenumber), and a dissipative cutoff occurring at $k \sim B_0$."
 They also suggested that the total energy could decay on the fast Hall timescale rather than the slow Ohmic timescale — despite the Hall term conserving energy and thus by itself being unable to cause decay on any timescale., They also suggested that the total energy could decay on the fast Hall timescale rather than the slow Ohmic timescale – despite the Hall term conserving energy and thus by itself being unable to cause decay on any timescale.
 Instead. by transferring energy from large to small scales. the Hall cascade enhances the efficiency of Ohmic decay. conjecturally by enough for the total energy to decay on the fundamentally different. and faster timescale.," Instead, by transferring energy from large to small scales, the Hall cascade enhances the efficiency of Ohmic decay, conjecturally by enough for the total energy to decay on the fundamentally different, and faster timescale."
 Following Goldreich and Reisenegger's seminal work. numerous authors have studied Eq. (1))," Following Goldreich and Reisenegger's seminal work, numerous authors have studied Eq. \ref{eq:A}) )"
 theoretically and numerically. in both the original spherical-shell geometry (????).. and," theoretically and numerically, in both the original spherical-shell geometry \citep{hollerbach02,hollerbach04,cumming04,pons07}, , and"
usage [9]...,usage \cite{pof2005}.
 Another contributing factor to the prevalence of particle schemes is their natural treatment of (he advection operator. which results in a numerical method that can easily handle and accurately capture traveling discontinuities in the distribution function |9]..," Another contributing factor to the prevalence of particle schemes is their natural treatment of the advection operator, which results in a numerical method that can easily handle and accurately capture traveling discontinuities in the distribution function \cite{pof2005}."
 An example of a particle method is the direct simulation Monte Carlo (DSMC) |10] which has become the prevalent simulation method for dilute gas flow., An example of a particle method is the direct simulation Monte Carlo (DSMC) \cite{bird} which has become the prevalent simulation method for dilute gas flow.
 One of the most important. disadvantages of particle methods lor solving the Boltzmann equation derives [rom their reliance on statistical averaging for extracting fiekd «quantities [rom particle data., One of the most important disadvantages of particle methods for solving the Boltzmann equation derives from their reliance on statistical averaging for extracting field quantities from particle data.
 When simulating processes close to equilibrium. thermal noise typically exceeds the available signal.," When simulating processes close to equilibrium, thermal noise typically exceeds the available signal."
 When coupled with (he slow convergence of statistical sampling (statistical error decreases wilh the square root of the number of samples). this often leads to computationally intractable problems |l1]..," When coupled with the slow convergence of statistical sampling (statistical error decreases with the square root of the number of samples), this often leads to computationally intractable problems \cite{jcp}."
 For example. to resolve a flow speed of the order of 1 m/s to statistical uncertainty in a dilute eas. on the order of 5x105 independent samples are needed 41]..," For example, to resolve a flow speed of the order of 1 m/s to statistical uncertainty in a dilute gas, on the order of $5\times 10^8$ independent samples are needed \cite{jcp}."
 In a recent paper. Baker and Hacdjiconstantinou have shown |9] that this rather severe limitation can be overcome wilh a form of variance reduction achieved by simulating only the deviation from equilibrium.," In a recent paper, Baker and Hadjiconstantinou have shown \cite{pof2005} that this rather severe limitation can be overcome with a form of variance reduction achieved by simulating only the deviation from equilibrium."
 By adopting Chis approach. it is possible to construct. Monte Carlo simulation methods that. can capture arbitrarily small deviations from equilibrium al a computational cost that is independent of the magnitude of this deviation.," By adopting this approach, it is possible to construct Monte Carlo simulation methods that can capture arbitrarily small deviations from equilibrium at a computational cost that is independent of the magnitude of this deviation."
 This is in sharp contrast to regular Monte Carlo methods. such as DSMC. whose computational cost [or the," This is in sharp contrast to regular Monte Carlo methods, such as DSMC, whose computational cost for the"
"a similar size of their svnthesised beams (1.01.0,84"" at 3.6 em. and 1.18""«1.117 at 2 em).","a similar size of their synthesised beams $1.01''\times
0.84''$ at 3.6 cm, and $1.18''\times 1.11''$ at 2 cm)."
 We have also retrieved publicly available submillimetre data at 450 and. S50 p/m. taken with SCUBA at the James Clerk Maxwell Telescope on 2003 November 6. corresponding to project MOSDUIS3.," We have also retrieved publicly available submillimetre data at 450 and 850 $\mu$ m, taken with SCUBA at the James Clerk Maxwell Telescope on 2003 November 6, corresponding to project M03BU13."
" The data at 500 sam have been presented in Curran&Chrevsostomou(2007) and Matthewset.al.(2009)... but in this paper we are using for our analvsis the original data (c14” angular resolution). without applying any spatial convolution (final angular resolution of 20"" in Matthews et al."," The data at 850 $\mu$ m have been presented in \cite{Curran&Chrysostomou07} and \cite{Matthews-etal09}, but in this paper we are using for our analysis the original data $\simeq 14''$ angular resolution), without applying any spatial convolution (final angular resolution of $20''$ in Matthews et al."
 2009)., 2009).
 Ehe data at 450 jin are presented. here for first time., The data at 450 $\mu$ m are presented here for first time.
" The angular resolution of the JCM'T at 450 pum is 8""", The angular resolution of the JCMT at 450 $\mu$ m is $\simeq 8''$ .
where and The normalization constant A is defined by the condition In eq. (16)),where and The normalization constant $K$ is defined by the condition In eq. \ref{flux}) )
 the only parameter that depends on (he grain size is E in the lower integration limit. wy.," the only parameter that depends on the grain size is $\Gamma$ in the lower integration limit, $u_1$."
" Ht follows that. if the values of 7, and Zi of models with different grain sizes are such that P is the same.SED."," It follows that, if the values of $\tau_0$ and $T_0$ of models with different grain sizes are such that $\Gamma$ is the same,."
 We show in Fieures 8 and 9 the SEDs and grain equilibrium temperatures for 20000Ix models with different erain sizes.," We show in Figures \ref{sedcs2} and \ref{tempc} the SEDs and grain equilibrium temperatures for $20000\,\rm K$ models with different grain sizes."
 Figure 9 shows that the temperature power-law exponent has a value s=0.4 at large r., Figure \ref{tempc} shows that the temperature power-law exponent has a value $s=0.4$ at large $r$.
 Other model parameters are listed in Table 3.., Other model parameters are listed in Table \ref{tab3}. .
" Note that the optical depths are very high (ranging [rom 7=3.4 to 5.0 for Ay=LOO jan). according to the assumption for case C. These optical depths were chosen so that (he parameter E is the same among the different models (i.e.. we measured 7, ancl adjusted 75 so that Tos""= constant)."," Note that the optical depths are very high (ranging from $\tau_0 = 3.4$ to 5.0 for $\lambda_0=100\,\mu\rm m$ ), according to the assumption for case C. These optical depths were chosen so that the parameter $\Gamma$ is the same among the different models (i.e., we measured $T_0$ and adjusted $\tau_0$ so that $T_0^s \tau_0^{1-n}=$ constant)."
 The resulting SEDs are very similar for À220jan. which indicates that the parameter T is an adequate scaling parameter for verv high optical depths.," The resulting SEDs are very similar for $\lambda\gtrsim 20 \mu\rm{m}$, which indicates that the parameter $\Gamma$ is an adequate scaling parameter for very high optical depths."
 For shorter wavelengths the SED invariance breaks down. because (he opacilies are no longer described by a power-law. and the SEDs are different.," For shorter wavelengths the SED invariance breaks down, because the opacities are no longer described by a power-law, and the SEDs are different."
 Also shown in Figure 8. is the analvtie expression for the emerging flux. eq. (16)).," Also shown in Figure \ref{sedcs2} is the analytic expression for the emerging flux, eq. \ref{flux}) )."
 This expression reproduces well the shape of the SED [ον long wavelengths: for shorter wavelengths (he agreement is not good for the reason stated above: our assumption lor the grain opacity fails in the well-known silicate spectral features al 10 and 20 jan. and as we go into the optical limit.," This expression reproduces well the shape of the SED for long wavelengths; for shorter wavelengths the agreement is not good for the reason stated above: our assumption for the grain opacity fails in the well-known silicate spectral features at 10 and 20 $\mu\rm m$, and as we go into the optical limit."
 We conclude that. egiven a dust egrain composition and spatial distribution. the shape of the IR. SED is controlled by a single parameter. EL.," We conclude that, given a dust grain composition and spatial distribution, the shape of the IR SED is controlled by a single parameter, $\Gamma$."
 From the observational point of view. we have a situation similar (to case D. where the IR emission bump does not convey. any inlormation about the dust grain size.," From the observational point of view, we have a situation similar to case B, where the IR emission bump does not convey any information about the dust grain size."
 However at shorter wavelengths. Al breaks down. so one can determineinformation about the dust grain size. if there is sufficient observable fIux al (hese wavelenethis.," However at shorter wavelengths, AI breaks down, so one can determineinformation about the dust grain size, if there is sufficient observable flux at these wavelengths."
1102480. using Hipparcos data as well as from. our observational data.,102480 using Hipparcos data as well as from our observational data.
 comparison of the observed colours of 999851 and 1102480 with those corresponding to the spectral types determined in the last section indicates that both stars sulfer little interstellar reddening., A comparison of the observed colours of 99851 and 102480 with those corresponding to the spectral types determined in the last section indicates that both stars suffer little interstellar reddening.
 Ehe Hipparcos parallaxes of 998851. ancl 1102480. are 5.84-EO.8T mas and 4.1741.04 mas. respectively.," The Hipparcos parallaxes of 98851 and 102480 are $\pm$ 0.87 mas and $\pm$ 1.04 mas, respectively."
 Phe corresponding distances of these stars are 1712525 ppe ancl 240460 ppe., The corresponding distances of these stars are $171\pm25$ pc and $240\pm60$ pc.
 At these distances the absolute magnitudes of 005551 and 1102480 are |1230.82 anc |1.50d:0.54. respectively. as given in ‘Table 4.," At these distances the absolute magnitudes of 98851 and 102480 are $+1.23 \pm 0.32$ and $+1.50 \pm 0.54$, respectively, as given in Table 4."
 Adopting a bolometric correction of -O.11 and -O.14 (Schmidt-Ixaler. 1982). for 998851 and 1102480 respectively. we derived the corresponding uminosities log(L/L.) as 1.12 and 1.36. respectively.," Adopting a bolometric correction of -0.11 and -0.14 (Schmidt-Kaler 1982) for 98851 and 102480 respectively, we derived the corresponding luminosities $\log (L/L_{\odot})$ as 1.12 and 1.36, respectively."
 The masses for both the stars were caleulated using the ollowing empirical relation (Schmidt-Ixaler 1982): ]t is worth noting that this relation is. determined rom well-observecl eclipsing ancl visual binaries. ancl so does not depend on evolutionary models.," The masses for both the stars were calculated using the following empirical relation (Schmidt-Kaler 1982): It is worth noting that this relation is determined from well-observed eclipsing and visual binaries, and so does not depend on evolutionary models."
 “Phe caleulated masses o£ 998851 ancl 1102480 are 2.247. and 2.1M. respectively. which are well within the range of masses for 9 Set stars.," The calculated masses of 98851 and 102480 are $2.2M_\odot$ and $2.1M_\odot$ respectively, which are well within the range of masses for $\delta$ Sct stars."
" Phe radii of the stars can be derived according to the racliation law (Schmidt-Ixaler 1982): eiving 361. and 35RR. [or I1D998851 and 1102450. respectively,"," The radii of the stars can be derived according to the radiation law (Schmidt-Kaler 1982): giving $_\odot$ and $_\odot$ for 98851 and 102480, respectively."
 From. the spectroscopic observations we determined. the effective temperatures of 998851. and. 1102480. to be 7000-250 WI ancl OT50#250KIN respectively. while the corresponding logg values are 3.5+0.5 and 3.00.5. respectively.," From the spectroscopic observations we determined the effective temperatures of 98851 and 102480 to be $\pm$ K and $\pm$ K respectively, while the corresponding $\log g$ values are $\pm$ 0.5 and $\pm$ 0.5, respectively."
 ALL the ohvsical parameters derived are listed in Table 4., All the physical parameters derived are listed in Table 4.
 The positions of 998851. and. 1102480. in a heoretical Εν diagram> are shown in Vie., The positions of 98851 and 102480 in a theoretical HR diagram are shown in Fig.
ὃν 6. alongo with he positions of other known pulsating Am stars taken from ‘Tureotte et al. (," 6, along with the positions of other known pulsating Am stars taken from Turcotte et al. ("
2000).,2000).
 Stellar evolutionary model paths for LOA. 2.0M. ancl 2.244. mass models are indicated. in he figure where it is clearly seen that both 998851 and 11024580 lie within the instability strip near the red edge.," Stellar evolutionary model paths for $1.9M_\odot$, $2.0M_\odot$ and $2.2M_\odot$ mass models are indicated in the figure where it is clearly seen that both 98851 and 102480 lie within the instability strip near the red edge."
 Fig., Fig.
 6 also indicates that both objects are probably near Ll core exhaustion. but. within the errors they could be past it.," 6 also indicates that both objects are probably near H core exhaustion, but, within the errors they could be past it."
 The age of the model corresponding the calculated Iuminosities and elfective temperatures is around. 1.0 Cyr., The age of the model corresponding the calculated luminosities and effective temperatures is around 1.0 Gyr.
 Given their unusual pulsation characteristics these would be interesting stars to model in more detail., Given their unusual pulsation characteristics these would be interesting stars to model in more detail.
 The unusual pulsations in the cool. marginal. evolved: Am stars 998851 and 1102480 are very interesting.," The unusual pulsations in the cool, marginal, evolved Am stars 98851 and 102480 are very interesting."
 The alternating high and low amplitude variations with nearly harmonic (or sub-harmonic) frequency ratio close to 2:1 are conspicuous., The alternating high and low amplitude variations with nearly harmonic (or sub-harmonic) frequency ratio close to 2:1 are conspicuous.
 In some 6 Sct stars. a 2:1 frequency. ratio is seen. and in others a pattern of alternating high and low amplitude peaks with longer periods is observed.," In some $\delta$ Sct stars, a 2:1 frequency ratio is seen, and in others a pattern of alternating high and low amplitude peaks with longer periods is observed."
 Lhe co-existence of both these phenomena has not previously been observed in chemically peculiar stars., The co-existence of both these phenomena has not previously been observed in chemically peculiar stars.
 Arentoft et al. (, Arentoft et al. (
2001) found that NN Pvx shows a near 2:1 frequency ratio. ancl suggested that NN Pyx is ad Set star in a binary system. in which the companion is low mass or compact object.,"2001) found that XX Pyx shows a near 2:1 frequency ratio, and suggested that XX Pyx is a $\delta$ Sct star in a binary system, in which the companion is low mass or compact object."
 A similar scenario cannot be excluded for 908851 and 1102480. since the spectroscopic and photometric evidence suggests that they are Jm stars. and it is known that Xm stars are almost all in short-period known binary systems.," A similar scenario cannot be excluded for 98851 and 102480, since the spectroscopic and photometric evidence suggests that they are Am stars, and it is known that Am stars are almost all in short-period known binary systems."
 In the case of these two stars. racial velocity observations are required to demonstrate their. binarity.," In the case of these two stars, radial velocity observations are required to demonstrate their binarity."
 This. in turn. would. constitute additional support for the Am classification for these two stars.," This, in turn, would constitute additional support for the Am classification for these two stars."
 The locations of 998851 and 1102480 in the Hi diagram are consistent with the theoretical work of Turcoe et al. (, The locations of 98851 and 102480 in the HR diagram are consistent with the theoretical work of Turcotte et al. (
2000). which shows that voung Am stars are stable against &-mechanism pulsation. and that as they evolve towards the red. edge of the instability strip they become unstable.,"2000), which shows that young Am stars are stable against $\kappa$ -mechanism pulsation, and that as they evolve towards the red edge of the instability strip they become unstable."
 Llowever the ratio of two prominent periods for 998851 and 1102480 has an intriguing value which is fav from the ratio 0.75 to 0.79 associated with the racial fundamental ancl first overtone pulsation for a large sample of double-mode à Set stars., However the ratio of two prominent periods for 98851 and 102480 has an intriguing value which is far from the ratio 0.75 to 0.79 associated with the radial fundamental and first overtone pulsation for a large sample of double-mode $\delta$ Sct stars.
 Among the low-amplituce 6 Set stars which exhibit near d frequeney ratios are AN Lyn (Itodriguez et al., Among the low-amplitude $\delta$ Sct stars which exhibit near 2:1 frequency ratios are AN Lyn (Rodriguez et al.
 1997. Zhou 102). V663 Cas (Mantegazza Porctti 1990) and 63 Ler (Breger οἱ al 1994).," 1997, Zhou 2002), V663 Cas (Mantegazza Poretti 1990) and 63 Her (Breger et al 1994)."
 In all these cases the authors concluded that a mixture of radial and non-radial modes is. needed to explain the non-standard. frequeney ratio., In all these cases the authors concluded that a mixture of radial and non-radial modes is needed to explain the non-standard frequency ratio.
 Therefore. in both of 998851 and 1102480 there may be mixture of racial ancl non-radial modes.," Therefore, in both of 98851 and 102480 there may be mixture of radial and non-radial modes."
the wake produced by the convergent shock on the rear side of the companion star.,the wake produced by the convergent shock on the rear side of the companion star.
 Similar plasma structures have been found in numerical simulations describing the outburst of RS Oph (e.g. ?)) and are generally expected as a result of the shielding of the blast by the companion star., Similar plasma structures have been found in numerical simulations describing the outburst of RS Oph (e.g. \citealt{2008A&A...484L...9W}) ) and are generally expected as a result of the shielding of the blast by the companion star.
" The X-ray lightcurves of the simulations in the [0.612.4] keV band are shown in Fig. 6,"," The X-ray lightcurves of the simulations in the $[0.6-12.4]$ keV band are shown in Fig. \ref{fig6},"
 together with the observed lightcurve., together with the observed lightcurve.
" The latter was based on the Swift/XRT count rates reported by ?,, considering a maximum X-ray luminosity of Lx=1.2x1034 erg s! at a distance of 2.7 kpc (?))."," The latter was based on the /XRT count rates reported by \citet{2011A&A...527A..98S}, considering a maximum X-ray luminosity of $L\rs{X} = 1.2\times 10^{34}$ erg $^{-1}$ at a distance of 2.7 kpc \citealt{2011A&A...527A..98S}) )."
" As found for U Sco by ?,, in NOCDE models (left panel in Fig. 6))"," As found for U Sco by \citet{2010ApJ...720L.195D}, in NOCDE models (left panel in Fig. \ref{fig6}) )"
" in general the X-ray luminosity, Lx, is larger for higher explosion energy (and ejecta mass), and wind gas density."," in general the X-ray luminosity, $L\rs{X}$, is larger for higher explosion energy (and ejecta mass), and wind gas density."
" For outburst energies and ejecta masses in the range of values typically observed for recurrent novae (i.e. 107""<Mey1079 Mo, 1043<Eo«107? erg; e.g. ?)), the synthetic Lx reaches its maximum very quickly, within few days after the explosion, so that none of these lightcurves fit the observed one which is characterized by a maximum around day 30."," For outburst energies and ejecta masses in the range of values typically observed for recurrent novae (i.e. $10^{-7}<
M\rs{ej} < 10^{-6}$ $M_\odot$, $10^{44} < E_0 < 10^{45}$ erg; e.g. \citealt{2010AN....331..160B}) ), the synthetic $L\rs{X}$ reaches its maximum very quickly, within few days after the explosion, so that none of these lightcurves fit the observed one which is characterized by a maximum around day 30."
" On the other hand, assuming outburst"," On the other hand, assuming outburst"
stellar evolution ages. reflecting a near common origin.,"stellar evolution ages, reflecting a near common origin."
 This model for the Local Group certainly has problems., This model for the Local Group certainly has problems.
 Perhaps most serious is the orbit of M33. which fits the measured distance. redshift ancl proper motion. within tolerable errors. along with the initial condition. but not the indications that M33. passed. close to AI21 at modest redshift.," Perhaps most serious is the orbit of M33, which fits the measured distance, redshift and proper motion, within tolerable errors, along with the initial condition, but not the indications that M33 passed close to M31 at modest redshift."
 This situation certainly requires further examination., This situation certainly requires further examination.
 The common anomalies in the redshifts and distances of the four misfits whose near common orbits are plotted as dashed lines in Figures 3. to 55. seem to be pointing to an error in the simplified phenomenological description of (he effect of the mass distribution outside the Local Group., The common anomalies in the redshifts and distances of the four misfits whose near common orbits are plotted as dashed lines in Figures \ref{Fig:3} to \ref{Fig:5} seem to be pointing to an error in the simplified phenomenological description of the effect of the mass distribution outside the Local Group.
 Investisating that will require a more detailed dynamical analvsis that takes explicit account of where (he larger external masses are observed to be now and where they are computed to have been., Investigating that will require a more detailed dynamical analysis that takes explicit account of where the larger external masses are observed to be now and where they are computed to have been.
 We hope to report on an exploration of tliis analvsis in due course., We hope to report on an exploration of this analysis in due course.
 The values of some ol the masses derived from the dvnanmics are curious. and may be in part the result of compensation for svstematic errors in the mocel.," The values of some of the masses derived from the dynamics are curious, and may be in part the result of compensation for systematic errors in the model."
 The verv real problems should not obscure the evidence that this dvnamical model has proved to be a useful approximation to reality. [rom (he success in fitting the considerable number of observational constraints on positions and velocities in Tables 1 and 2. and perhaps also the Milky Way. circular velocity ος in equation (3)). with orbits that. have the initial conditions required by a reasonable approximation to the cosmological growth of structure.," The very real problems should not obscure the evidence that this dynamical model has proved to be a useful approximation to reality, from the success in fitting the considerable number of observational constraints on positions and velocities in Tables 1 and 2, and perhaps also the Milky Way circular velocity $v_c$ in equation \ref{eq:vc}) ), with orbits that have the initial conditions required by a reasonable approximation to the cosmological growth of structure."
 It is notable also that the model has led to the considerable list of issues For further consideration presented in Section ??.., It is notable also that the model has led to the considerable list of issues for further consideration presented in Section \ref{sec:discussion}.
 In the past decade aclvances in detections of nearby galaxies. and in the measurements ol galaxy distances and velocities. have greatly enriched the study of dvnamiecs in the now densely sampled region within 1.5 Alpe clistance.," In the past decade advances in detections of nearby galaxies, and in the measurements of galaxy distances and velocities, have greatly enriched the study of dynamics in the now densely sampled region within 1.5 Mpc distance."
 We have the prospect of another advance of the same order [rom work in progress. or that seems (o be feasible. on measurements of ealaxy distances and proper motions.," We have the prospect of another advance of the same order from work in progress, or that seems to be feasible, on measurements of galaxy distances and proper motions."
 That will present us with a really serious but fascinating challenge lor the clevelopment of numerical methods of analysis of a tangled web of constraints on what happened in the local universe., That will present us with a really serious but fascinating challenge for the development of numerical methods of analysis of a tangled web of constraints on what happened in the local universe.
 We are grateful to Mark Reid and Jeremy Darling lor guidance to prospects for measurements of proper motions of Local Group galaxies., We are grateful to Mark Reid and Jeremy Darling for guidance to prospects for measurements of proper motions of Local Group galaxies.
orientation of clump spins is expected. while in the fragmentation scenario. clumps are expected to be coplanar.,"orientation of clump spins is expected, while in the fragmentation scenario, clumps are expected to be coplanar."
 In particular. one interesting predictions of these simulations is that the large-scale rotation in the underlying disk should be preserved during the fragmentation phase (Immelietal.2004:Bournaud2007).," In particular, one interesting predictions of these simulations is that the large-scale rotation in the underlying disk should be preserved during the fragmentation phase \citep{immeli04,bournaud07}."
 The achievement of integral field spectrograph working in the near infrared (e.g. SINFONI at the VLT. or OSIRIS at Keck). allowed several teams to gather spatially-resolved kinematic observations of 21 distant galaxies (see FórsterSchreiberal.2009 and references therein. as well as Wrightetal.2007.etal. 2009)).," The achievement of integral field spectrograph working in the near infrared (e.g., SINFONI at the VLT, or OSIRIS at Keck), allowed several teams to gather spatially-resolved kinematic observations of $>$ 1 distant galaxies (see \citealt{forsterschreiber09,law09} and references therein, as well as \citealt{wright07,wright08,bournaud08,vanstarkenburg08,epinat09}) )."
 Detection of an underlying rotation was claimed in several z-2 elumpy galaxies. which has been used as a support to the fragmentation scenario (Genzeletal.2008).," Detection of an underlying rotation was claimed in several $\sim$ 2 clumpy galaxies, which has been used as a support to the fragmentation scenario \citep{genzel08}."
. Surprisingly. analysis of integral field. spectroscopy observations suggest that two «different scenari might take place at z«l. and at 21.," Surprisingly, analysis of integral field spectroscopy observations suggest that two different scenarii might take place at $<$ 1, and at $>$ 1."
 It therefore becomes necessary to start investigating whether both scenarii are consistent. and whether or hot a transition between two different galaxy evolution drivers is occurring between these two epochs.," It therefore becomes necessary to start investigating whether both scenarii are consistent, and whether or not a transition between two different galaxy evolution drivers is occurring between these two epochs."
 To this aim. I take advantage of the IMAGES survey to study clumpy galaxies at z=0.6.," To this aim, I take advantage of the IMAGES survey to study clumpy galaxies at z=0.6."
 The goal is to investigate which of the merger or fragmentation scenario is the most consistent with z~0.6 clumpy galaxies. and investigate whether clumpy galaxies at z«l and at 221 are driven by a common physieal process.," The goal is to investigate which of the merger or fragmentation scenario is the most consistent with $\sim$ 0.6 clumpy galaxies, and investigate whether clumpy galaxies at $<$ 1 and at $>$ 1 are driven by a common physical process."
 This paper is organized as follows: In Sect., This paper is organized as follows: In Sect.
 2. I describe how the clumpy galaxies at z- 0.6 were selected: In Sect.," 2, I describe how the clumpy galaxies at $\sim$ 0.6 were selected; In Sect."
 3. I present their kinematic and dynamical properties: Sect.," 3, I present their kinematic and dynamical properties; Sect."
 + discusses the origin of the clumpiness in distant galaxies: Sect., 4 discusses the origin of the clumpiness in distant galaxies; Sect.
 5 discusses the results. while conclusions are drawn in Sect.," 5 discusses the results, while conclusions are drawn in Sect."
 6., 6.
 Throughout the paper. I adopt Hy=70 km/s/Mpe. ον=0.3. and O4=0.7. and the AB magnitude system.," Throughout the paper, I adopt $H_0=70$ km/s/Mpc, $\Omega _M=0.3$, and $\Omega _\Lambda=0.7$, and the $AB$ magnitude system."
 In this paper. I made use of the representative IMAGES sample of 63 emission line galaxies at z~0.6.," In this paper, I made use of the representative IMAGES sample of 63 emission line galaxies at $\sim$ 0.6."
 The sample is fully described in Yangetal.(2008)... and I refer the reader to this paper for details.," The sample is fully described in \cite{yang08}, and I refer the reader to this paper for details."
 This sample was observed by FLAMES/GIRAFFE at the VLT. which allowed us to derived spatially-resolved kinematies for all these galaxies.," This sample was observed by FLAMES/GIRAFFE at the VLT, which allowed us to derived spatially-resolved kinematics for all these galaxies."
 The dynamical state of these galaxies was studied. according to which they were classified into three classes: Rotating Disks (RDs) are galaxies showing regular rotation well-aligned along the morphological axis. Perturbed Rotators (PRs) show a large-scale rotation with a local perturbation in the velocity dispersion map that cannot be accounted for by rotation. while galaxies with Complex Kinematics (CKs) do not show any large-scale rotation. or a strong misalignment between the dynamical and morphological axes (Yangetal.2005).," The dynamical state of these galaxies was studied, according to which they were classified into three classes: Rotating Disks (RDs) are galaxies showing regular rotation well-aligned along the morphological axis, Perturbed Rotators (PRs) show a large-scale rotation with a local perturbation in the velocity dispersion map that cannot be accounted for by rotation, while galaxies with Complex Kinematics (CKs) do not show any large-scale rotation, or a strong misalignment between the dynamical and morphological axes \citep{yang08}."
.. This classification takes into account the residuals between the observed VF and o- and those predicted by a rotating-disk model (see Yangetal. 2008). which makes this classitication objective and reproducible.," This classification takes into account the residuals between the observed VF and $\sigma$ -map and those predicted by a rotating-disk model (see \citealt{yang08}) ), which makes this classification objective and reproducible."
 For reasons of homogeneity. and to benetit from the most exquisite images from the ACS camera on-board the HST. I limited myself to galaxies lying in the CDFS-GOODS field.," For reasons of homogeneity, and to benefit from the most exquisite images from the ACS camera on-board the HST, I limited myself to galaxies lying in the CDFS-GOODS field."
 This study is therefore based on a sample of 32 galaxies., This study is therefore based on a sample of 32 galaxies.
 Note that this sample is still representative of 70.6. intermediate-mass galaxies. as shown by Yangetal.(2008).," Note that this sample is still representative of $\sim$ 0.6, intermediate-mass galaxies, as shown by \cite{yang08}."
 I selected. clumpy galaxies following selection. criteria depicted by Elmegreen&(2005). and(2007)., I selected clumpy galaxies following selection criteria depicted by \cite{elmegreen05} and.
. Specifically. I required for a galaxy to be clumpy to fultil the following requirements: Obviously. such a selection. method is rather subjective.," Specifically, I required for a galaxy to be clumpy to fulfil the following requirements: Obviously, such a selection method is rather subjective."
 In order to stick to former studies of such galaxies at higher redshift. I nevertheless chose to keep these selection criteria and ignored further color and kinematic information from the IMAGES database.," In order to stick to former studies of such galaxies at higher redshift, I nevertheless chose to keep these selection criteria and ignored further color and kinematic information from the IMAGES database."
 Instead. this information was used to investigated whether the candidates selected this way were truly rotating and/or corresponded to Jeans fragmentation processes.," Instead, this information was used to investigated whether the candidates selected this way were truly rotating and/or corresponded to Jeans fragmentation processes."
 The selection was done independently by three member of the IMAGES consortium (including the author). and results compared until a consensus was reached for each object.," The selection was done independently by three member of the IMAGES consortium (including the author), and results compared until a consensus was reached for each object."
 I ended up with a sample of HI candidates. which are shown in Fig. |..," I ended up with a sample of 11 candidates, which are shown in Fig. \ref{morpho}."
 A visual comparison with Fig., A visual comparison with Fig.
 | of Elmegreen&(2005). contirms that all 7-0.6 candidates indeed look similar to higher-z clumpy galaxies., 1 of \cite{elmegreen05} confirms that all $\sim$ 0.6 candidates indeed look similar to higher-z clumpy galaxies.
 Since the IMAGES-CDFS sample is representative of z~0.6 emission line galaxies. I ean estimate the fraction of emission line clumpy galaxies at z-0.6. which is found to be 34218t.. with error-bars due to Poisson fluctuation in the parent sample.," Since the IMAGES-CDFS sample is representative of $\sim$ 0.6 emission line galaxies, I can estimate the fraction of emission line clumpy galaxies at $\sim$ 0.6, which is found to be $\pm$, with error-bars due to Poisson fluctuation in the parent sample."
 This is in very good agreement with results from the UDF at z«1. and a factor -2 less than at z~2-3 (Elmegreenetal.2007).," This is in very good agreement with results from the UDF at $<$ 1, and a factor $\sim$ 2 less than at $\sim$ 2-3 \citep{elmegreen07}."
 If one accounts for the fact that emission line galaxies represent of intermediate mass galaxies at z~0.6 (Hammeretal.1997:: see also Mignolietal. 2005). one finds that clumpy galaxies represent 20x of intermediate-mass galaxies at this redshift.," If one accounts for the fact that emission line galaxies represent of intermediate mass galaxies at $\sim$ 0.6 \citealt{hammer97}; see also \citealt{mignoli05}) ), one finds that clumpy galaxies represent $\pm$ of intermediate-mass galaxies at this redshift."
 Spatially-resolved kinematics was obtained. amongst others. for all of the Tl clumpy galaxies by Yangetal.(2008).," Spatially-resolved kinematics was obtained, amongst others, for all of the 11 clumpy galaxies by \cite{yang08}."
. Amongst them. five were found to be CKs. four corresponded to PR. and two were classified as RDs (see Appendix).," Amongst them, five were found to be CKs, four corresponded to PR, and two were classified as RDs (see Appendix)."
 Rotation requires a significant inclination angle to be detected. (.e.. a face-on disk would not show any rotation due to projection effects).," Rotation requires a significant inclination angle to be detected, (i.e., a face-on disk would not show any rotation due to projection effects)."
 I checked that such projection ettects do not atfect the relative fraction of CK galaxies: considering the full sample of 63 IMAGES galaxies from Yangetal.(2008)... the fraction of face-on galaxies (i.e. having an inclination angle smaller than 30 degrees). is found to be 10. 17. and 135cx8c amongst RD. PR. and CK galaxies. respectively.," I checked that such projection effects do not affect the relative fraction of CK galaxies: considering the full sample of 63 IMAGES galaxies from \cite{yang08}, the fraction of face-on galaxies (i.e., having an inclination angle smaller than 30 degrees), is found to be 10, 17, and $\pm$ amongst RD, PR, and CK galaxies, respectively."
 This rules out any significant bias due to such projection etfects. even if it cannot be excluded that a particular object might be affected.," This rules out any significant bias due to such projection effects, even if it cannot be excluded that a particular object might be affected."
 According to predictions from numerical models. Jeans- should preserve the rotation of the underlying disk (see. e.g.. Immelietal.2004:Bournaud 2007)).," According to predictions from numerical models, Jeans-fragmentation should preserve the rotation of the underlying disk (see, e.g., \citealt{immeli04, bournaud07}) )."
 Therefore. it is expected that distant clumpy galaxies should in principle," Therefore, it is expected that distant clumpy galaxies should in principle"
Tuneuska cosmic body. that max have decreased the inclination angle.,"Tunguska cosmic body, that may have decreased the inclination angle."
 Moreover. we have welected the litt effects. following Cliyba et al. (1993)).," Moreover, we have neglected the lift effects, following Chyba et al. \cite{CHYBA}) )."
 Among the authors quoted by Vasilvev. oulv Sckanina derived an angle lower than 57.," Among the authors quoted by Vasilyev, only Sekanina derived an angle lower than $5\degr$."
 Iuterestiugly. it was just Sckanina (1998)) who stronely favoured the conclusion of an asteroidal origin for the Tunsuska cosmic body.," Interestingly, it was just Sekanina \cite{SEKANINA}) ) who strongly favoured the conclusion of an asteroidal origin for the Tunguska cosmic body."
 The results obtained here provide additional support for Selkana’s conclusion., The results obtained here provide additional support for Sekanina's conclusion.
 Tn this paper. we have outlined a new analysis of the he Tuuguska event. starting from seine data obtained x BenMenahem (1975)) aucl iiproviue the relationship vetween body speed. mechanical streneth aud azirburst right.," In this paper, we have outlined a new analysis of the the Tunguska event, starting from seismic data obtained by Ben–Menahem \cite{MENAHEM}) ) and improving the relationship between body speed, mechanical strength and airburst height."
 The main conclusion is that the Tuuguska cosmic ουν was probably a stouv asteroid. with a diameter of about 60 ni. We have also sununarzed the properties of the ivpersonic flow around a small asteroid im the Earth's atimosplicre.," The main conclusion is that the Tunguska cosmic body was probably a stony asteroid, with a diameter of about 60 m. We have also summarized the properties of the hypersonic flow around a small asteroid in the Earth's atmosphere."
 We have shown that the stagnation eniperature is a direc measure of the body speed., We have shown that the stagnation temperature is a direct measure of the body speed.
 This introduces a multiplicative factor of VE(>1) iu the Eq. (1)).," This introduces a multiplicative factor of $\sqrt{\gamma/(\gamma-1)}$ in the Eq. \ref{e:velo}) ),"
 which is iustruinental to derive a reasonable solution for the Tuneuska event., which is instrumental to derive a reasonable solution for the Tunguska event.
 Eq. (6)), Eq. \ref{e:velo2}) )
 is consistent with the idea that the meteoroid’s fragmentation is due to the coupled action of thermodvuamical aud mechanical processes., is consistent with the idea that the meteoroid's fragmentation is due to the coupled action of thermodynamical and mechanical processes.
 This liud of aualvsis can he applicd whenever the body is lareeOo aud compact enouelC» to reach the lower atimosphere., This kind of analysis can be applied whenever the body is large and compact enough to reach the lower atmosphere.
 Further investigations are needed for application to other cases., Further investigations are needed for application to other cases.
 However. the crucial role of the stagnation temperature is probably a general feature of any realistic model of iieteoroid fight aud brealap.," However, the crucial role of the stagnation temperature is probably a general feature of any realistic model of meteoroid flight and breakup."
fit to the observed [requencies.,fit to the observed frequencies.
 Note that our model of the tachocline differs from that of Ixosovichev (1996) and. Charbonneau οἱ al. (, Note that our model of the tachocline differs from that of Kosovichev (1996) and Charbonneau et al. (
1999) and to compare the width we obtain to (hose obtained by them. (6 needs to be multiplied bv a [actor of 2.5.,"1999) and to compare the width we obtain to those obtained by them, $w$ needs to be multiplied by a factor of $2.5$."
 For each available data set in GONG and MDI data sets we fit rotation rate of the form eiven bx Eq. (1)), For each available data set in GONG and MDI data sets we fit rotation rate of the form given by Eq. \ref{eq:tach}) )
 to the observed splitting coellicients to calculate the 11 parameters of the model., to the observed splitting coefficients to calculate the 11 parameters of the model.
 hi this work we are only interested in the parameters (hat define the properties of ihe tachocline., In this work we are only interested in the parameters that define the properties of the tachocline.
 The time averaged parameters for the tachocline are listed in Table 1., The time averaged parameters for the tachocline are listed in Table 1.
 The average Lachocline properties were determined by averaging all results., The average tachocline properties were determined by averaging all results.
 GONG and MDI results were averaged. and are listed. separately.," GONG and MDI results were averaged, and are listed, separately."
 The average results are consistent with those of Basu Antia (2001) obtained using a subset of current. data ancl with a slightly different form for the rotation rate., The average results are consistent with those of Basu Antia (2001) obtained using a subset of current data and with a slightly different form for the rotation rate.
 As can be seen. that there are some systematic cillerences between the GONG and MDI results.," As can be seen, that there are some systematic differences between the GONG and MDI results."
 Similar differences have been seen in rotation inversion results (Schou et al., Similar differences have been seen in rotation inversion results (Schou et al.
 2002)., 2002).
 The discrepancy is mainly attributed to differences in the data processing pipelines of the two projects., The discrepancy is mainly attributed to differences in the data processing pipelines of the two projects.
 Studies of temporal variations in solar rotation rate (e.g..Antia et al.," Studies of temporal variations in solar rotation rate (e.g., Antia et al."
 2008a) suggest that the time-varving component of the rotation rate is not affected by the discrepancy between GONG ancl MDI data., 2008a) suggest that the time-varying component of the rotation rate is not affected by the discrepancy between GONG and MDI data.
 The fitted parameters can be used to calculate the properties of the tachocline at any eiven latitude., The fitted parameters can be used to calculate the properties of the tachocline at any given latitude.
 The time-averaged results al a few latitudes are shown in Table 2., The time-averaged results at a few latitudes are shown in Table 2.
 As can be seen clearly. (he tachocline is prolate as measured by the central radius defining the tachocline.," As can be seen clearly, the tachocline is prolate as measured by the central radius defining the tachocline."
 GONG data vields a difference of (0.01243:0.002)4... between the latitude of 60 and (he equator., GONG data yields a difference of $(0.012\pm0.002)R_\odot$ between the latitude of $60^\circ$ and the equator.
 The difference with MDI data is (0.040+0.003)..., The difference with MDI data is $(0.040\pm0.003)R_\odot$.
 Our results are in agreement wilh earlier results (Antia et al., Our results are in agreement with earlier results (Antia et al.
 L998: C'orbard et al., 1998; Corbard et al.
 1998. 1999: Charbonneau et al.," 1998, 1999; Charbonneau et al."
 1999: Basu Antia 2001. 2003).," 1999; Basu Antia 2001, 2003)."
 Table 2 also shows a clear variation of the thickness of the tachocline., Table 2 also shows a clear variation of the thickness of the tachocline.
 What we find is that the tachocline is (hie least thick at the equator. and (he (thickness increases steadily with latitude.," What we find is that the tachocline is the least thick at the equator, and the thickness increases steadily with latitude."
 While earlier investigations (Charbonneau et al., While earlier investigations (Charbonneau et al.
 1999: Basu Antia 2001. 2003) had indicated that this might be the case. the results were not statistically significant.," 1999; Basu Antia 2001, 2003) had indicated that this might be the case, the results were not statistically significant."
 The larger quantity ofdata now has resolved (is issue., The larger quantity of data available now has resolved this issue.
 We find that the thickness increases by ma(0.018dE“reae0.003)R. between the equator aud 607 latitude as per GONG and (0.023+t as per data., We find that the thickness increases by about $(0.018\pm0.003)R_\odot$ between the equator and $60^\circ$ latitude as per GONG and $(0.028\pm0.004)R_\odot$ as per MDI data.
are therefore dealing with the shortest apsidal motion period known among the EEBs today.,are therefore dealing with the shortest apsidal motion period known among the EEBs today.
 The whole light curve of V456 Oph was observed with the OMC camera onboard the INTEGRAL satellite. a description of which is given in Mas-Hesseetal.(2004).," The whole light curve of V456 Oph was observed with the OMC camera onboard the INTEGRAL satellite, a description of which is given in \cite{2004ESASP}."
. The standard V filter was used. but the optical telescope has an aperture of only 5 em in diameter.," The standard $V$ filter was used, but the optical telescope has an aperture of only 5 cm in diameter."
 We obtained several hundred observations. of which we used 449 for the analysis.," We obtained several hundred observations, of which we used 449 for the analysis."
 The programme (ver., The programme (ver.
 0.29. Prsaa Zwitter 2005). based on the Wilson-Devinney algorithm (Wilson Devinney 1971) was used for the analysis.," 0.29, Pršaa Zwitter 2005), based on the Wilson-Devinney algorithm (Wilson Devinney 1971) was used for the analysis."
 The “detached binary” mode (in Wilson Devinney mode 2) was used with several assumptions.," The ""detached binary"" mode (in Wilson Devinney mode 2) was used with several assumptions."
 First. the ephemerides (HJDy and P) and the apsidal motion parameters (e. w. and Ww) were adopted from the period analysis. because the minima times cover a longer time span. therefore these quantities are derived with higher precision.," First, the ephemerides $HJD_0$ and $P$ ) and the apsidal motion parameters $e$, $\omega$, and $\dot \omega$ ) were adopted from the period analysis, because the minima times cover a longer time span, therefore these quantities are derived with higher precision."
" Secondly. the mass ratio g and temperature of the primary component 7, were set and the other relevant parameters were adjusted for the best fit."," Secondly, the mass ratio $q$ and temperature of the primary component $T_1$ were set and the other relevant parameters were adjusted for the best fit."
" We changed the g and 7, values in the wide range of values to obtain the best fit according to the RMS value and also the physical plausibility of the fit.", We changed the $q$ and $T_1$ values in the wide range of values to obtain the best fit according to the RMS value and also the physical plausibility of the fit.
" This means during that the fitting process we scanned the parameter space ing ranging from 0.1 to 1.2 and in T, from 15400 K to 6500 K. We fitted the other light curve parameters. which are the luminosities { and Ls» in the V filer. the temperature of the secondary 7. the inclination /. the Kopal’s modified potentials OQ, and Q2. the synchronicity parameters Fy and F>. the third light /;."," This means during that the fitting process we scanned the parameter space in $q$ ranging from 0.1 to 1.2 and in $T_1$ from 15400 K to 6500 K. We fitted the other light curve parameters, which are the luminosities $L_1$ and $L_2$ in the $V$ filer, the temperature of the secondary $T_2$, the inclination $i$, the Kopal's modified potentials $\Omega_1$ and $\Omega_2$, the synchronicity parameters $F_1$ and $F_2$, the third light $l_3$."
" The limb-darkening coefficients were automatically interpolated by the programme from van Hamme's tables (see van Hamme 1993). using the linear cosine law for the values of 7,,; and logg of both components resulting from the analysis."," The limb-darkening coefficients were automatically interpolated by the programme from van Hamme's tables (see van Hamme 1993), using the linear cosine law for the values of $T_{eff}$ and $\log g$ of both components resulting from the analysis."
" The values of the gravity brightening and bolometric albedo coefficients were set at their suggested values for convective atmospheres (see Lucy 1968). Le. 0.32.A,;=A>0.5."," The values of the gravity brightening and bolometric albedo coefficients were set at their suggested values for convective atmospheres (see Lucy 1968), i.e. $G_1 = G_2 = 0.32$ , $A_1 = A_2 = 0.5$."
 The best fit was achieved with the light curve parameters given in Table 4.. and the figure with the final fit is plotted in Fig.3.," The best fit was achieved with the light curve parameters given in Table \ref{LCparam}, and the figure with the final fit is plotted in \ref{Fig1}."
. Nonzero eccentricity ts clearly visible from this plot. which ts quite surprising in a binary with such a short orbital period.," Nonzero eccentricity is clearly visible from this plot, which is quite surprising in a binary with such a short orbital period."
 No other EEB with a shorter period is known today., No other EEB with a shorter period is known today.
 The value of the third light is%.. which indicates that there is no other visible companion to the system in the V filter (under the assumption that this component ts also located on the main sequence).," The value of the third light is, which indicates that there is no other visible companion to the system in the $V$ filter (under the assumption that this component is also located on the main sequence)."
 We made several attempts with nonzero, We made several attempts with nonzero
Ποιά of view were removed.,field of view were removed.
 The final catalogue comprises a total of 43 X-ray sources to the limit fy(2SkeV)--L4010Mores.tem2 ," The final catalogue comprises a total of 43 X-ray sources to the limit $f_X(\rm 2 - 8 \,keV) \approx 7.7 \times10^{-15} \,erg \,s^{-1}
\,cm^{-2}$ ."
Count rates in the merged (PN|MOS) images as well as he individual PN and MOS images are estimated within an lSaarcsec aperture., Count rates in the merged (PN+MOS) images as well as the individual PN and MOS images are estimated within an arcsec aperture.
 For the background. estimation we use he backgrouncl maps generated. by the task of SAS., For the background estimation we use the background maps generated by the task of SAS.
 A small fraction of sources lie close to masked regions (CCD gaps or hot. pixels) on either the MOS or the PN detectors., A small fraction of sources lie close to masked regions (CCD gaps or hot pixels) on either the MOS or the PN detectors.
 This may introduce errors in the estimated source counts., This may introduce errors in the estimated source counts.
 To avoid this bias. the source count rates (and hence he hardness ratios and the flux) are estimated. using the detector (NOS or PN) with no masked pixels in the vicinity ofthe source.," To avoid this bias, the source count rates (and hence the hardness ratios and the flux) are estimated using the detector (MOS or PN) with no masked pixels in the vicinity of the source."
 To convert count. rates to [lux the Enerey Conversion Factors (ECE) of individual. detectors are calculated assuming a power law spectrum with P=1.7 and Galactic . ⋜↧∣⋡≱∖∪↓⋅↓≻↿↓∪⊔↓∖⊔∶−≽↓∪↾≼⇍⊔↓⋜↧↓≻↓≻↓⋅∪↓≻↓⋅↓⋜⋯⊾⇂∪↓⋅↿↓↕∢⊾↓↴↓↕∪⋖⋅⊔∟∖⇁.2σι . ⋅ ⋠ Ποιά (Dickey Lockman 1990).," To convert count rates to flux the Energy Conversion Factors (ECF) of individual detectors are calculated assuming a power law spectrum with $\Gamma=1.7$ and Galactic absorption $\rm N_H=2\times
10^{20} \rm {cm^{-2}}$ appropriate for the Phoenix field (Dickey Lockman 1990)."
 The mean ECE for the mosaic of all three detectors is estimated by weighting the ECEs of individual detectors by the respective exposure time., The mean ECF for the mosaic of all three detectors is estimated by weighting the ECFs of individual detectors by the respective exposure time.
 For the encireled energy. correction. accounting for the energy fraction outside the aperture within which source counts are accumulated. we adopt the calibration given by the Calibration Documentation," For the encircled energy correction, accounting for the energy fraction outside the aperture within which source counts are accumulated, we adopt the calibration given by the Calibration Documentation."
 The radio observations of the Phoenix Deep Survey (PDS*)) were carried out at the Australia Telescope Compact Array (ATCA) at νε during several campaigns between 1994 and 2001 in the GA. GB and GC array configurations.," The radio observations of the Phoenix Deep Survey ) were carried out at the Australia Telescope Compact Array (ATCA) at GHz during several campaigns between 1994 and 2001 in the 6A, 6B and 6C array configurations."
" The data cover a 4.56. square degree area centered at RAGJ2000)=01""11""13> 15745'00"". much larger than the aarcmin ciameter region covered by the observations."," The data cover a 4.56 square degree area centered at $01^{\rm h}11^{\rm m}13^{\rm s}$ $-45\degr45\arcmin00\arcsec$, much larger than the arcmin diameter region covered by the observations."
 A detailed. description of the radio observations. data reduction ancl source detection are discussed. by Hopkins et al. (," A detailed description of the radio observations, data reduction and source detection are discussed by Hopkins et al. ("
1998. 1999. 2003).,"1998, 1999, 2003)."
 The observational strategy adopted resulted in a radio map that is highly homogeneous within the central zος radius.," The observational strategy adopted resulted in a radio map that is highly homogeneous within the central $\rm \approx1\,deg$ radius."
 The le rms noise nevertheless increases from. 125v at the most sensitive region to about Opry close to the field. edge.," The $1\sigma$ rms noise nevertheless increases from $\rm
12\mu Jy$ at the most sensitive region to about $\rm 90\mu Jy$ close to the field edge."
 The final catalogue consists of a total of 2148 radio sources to a limit of 5/Jv (Llopkins et al., The final catalogue consists of a total of 2148 radio sources to a limit of $\rm \mu$ Jy (Hopkins et al.
 2003)., 2003).
 The pointing lies. within the most homogeneously covered region of the PDS but is ollset from the most sensitive area of the radio map by about dedeg., The pointing lies within the most homogeneously covered region of the PDS but is offset from the most sensitive area of the radio map by about deg.
 Therefore. the completeness limit of the radio observations in that aarcmin cdiameter region is =SOjp]y.," Therefore, the completeness limit of the radio observations in that arcmin diameter region is $\rm
\approx80\,\mu Jy$."
 The source detection was performed. using the False. Discovery. Rate algorithm described. by Hopkins et al. (, The source detection was performed using the False Discovery Rate algorithm described by Hopkins et al. (
2002) to robustly quantify the number of spurious sources in radio catalogues.,2002) to robustly quantify the number of spurious sources in radio catalogues.
 ‘This novel method is somewhat different from the traditional radio source detection algorithms that select. sources with peak Lux density above a user defined threshold expressed in multiples of the local background RAIS noise., This novel method is somewhat different from the traditional radio source detection algorithms that select sources with peak flux density above a user defined threshold expressed in multiples of the local background RMS noise.
 Nevertheless. in the region covered by theΔον the lux density limit of zδν corresponds to a detection. threshold. for the tracitional method of z5o.," Nevertheless, in the region covered by the the flux density limit of $\rm \approx80\,\mu Jy$ corresponds to a detection threshold for the traditional method of $\approx5\sigma$."
 X total of 204 radio sources overlap with the data., A total of 204 radio sources overlap with the data.
 Optical photometric observations of most of the PDS area in the VY and. Z?-bands were obtained at the Anelo-Australian Telescope. (AAT) during two observing runs in 1994 and 1995., Optical photometric observations of most of the PDS area in the $V$ and $R$ -bands were obtained at the Anglo-Australian Telescope (AAT) during two observing runs in 1994 and 1995.
 X detailed description of these observations including data reduction. photometric calibration. source detection ancl optical identification are presented by Cieorgakakis et al. (," A detailed description of these observations including data reduction, photometric calibration, source detection and optical identification are presented by Georgakakis et al. ("
1999).,1999).
 Phis dataset is complete to 2=22.5 mmasg., This dataset is complete to $R=22.5$ mag.
 The (-band data used in the present study are from the ESO mmm telescope using the WL on August 18 2001., The $U$ -band data used in the present study are from the ESO m telescope using the WFI on August 18 2001.
 These observations are described in Sullivan et al. (, These observations are described in Sullivan et al. (
2004) and reach a completeness limit of (zz22.5 mmag.,2004) and reach a completeness limit of $U\approx22.5$ mag.
 Near-infrared observations of the 30aarcmin diameter area covered. by the pointing were obtained using the OSIRIS multi-purpose instrument at the {19.0 focus of the C'PIO 1.5-m telescope., Near-infrared observations of the arcmin diameter area covered by the pointing were obtained using the OSIRIS multi-purpose instrument at the $f/13.5$ focus of the CTIO 1.5-m telescope.
 The observations were carried out in 2000 December 4-6 in the J ancl As bands., The observations were carried out in 2000 December 4-6 in the $J$ and $Ks$ bands.
 The OSIRIS is equipped with a 1024x1024 Hawaii LHeC'dTe array with a field of view of 11.I1aremin? and a pixel scale of aaresec for the f/2.8 camera focus position used in our program.," The OSIRIS is equipped with a 1024x1024 Hawaii HgCdTe array with a field of view of $\rm 11 \times
11 \, arcmin^{2}$ and a pixel scale of arcsec for the $f/2.8$ camera focus position used in our program."
 To cover the aarcmin diameter pointing we used 7 contiguous pointings., To cover the arcmin diameter pointing we used 7 contiguous pointings.
 In the As band for cach pointing a dithering pattern was acopted that consisted of a sequence of ss integrations ollowed. by an offset. of the telescope., In the $Ks$ band for each pointing a dithering pattern was adopted that consisted of a sequence of s integrations followed by an offset of the telescope.
 To. get accurate skv frames the olfset vector was not replicated. between successive exposures., To get accurate sky frames the offset vector was not replicated between successive exposures.
 The total integration time varied xtween 20 to mmin with a mean of mmin per pointing., The total integration time varied between 20 to min with a mean of min per pointing.
 In the J-band a similar dithering procedure was followed with an exposure time between olfsets of 30 to 60ss. The otal integration time was mmin per pointing., In the $J$ -band a similar dithering procedure was followed with an exposure time between offsets of 30 to s. The total integration time was min per pointing.
 The data reduction was carried out using tasks., The data reduction was carried out using tasks.
 The Hat. field. frame was constructed using both the target observations and the dome Lat following a method developed » Peter Witehalls and. Will Saunders ancl deseribed by Sullivan et al. (, The flat field frame was constructed using both the target observations and the dome flat following a method developed by Peter Witchalls and Will Saunders and described by Sullivan et al. (
2004).,2004).
 The sky frame to be subtracted from a given target [rame was generated. by median. combining he 10 images closest in time to the frame in question., The sky frame to be subtracted from a given target frame was generated by median combining the 10 images closest in time to the frame in question.
 MI images of the same pointing were registered and coadded to xoduce the final image., All images of the same pointing were registered and coadded to produce the final image.
 Photometric calibration was carried out using standard. stars from Persson et al. (, Photometric calibration was carried out using standard stars from Persson et al. (
1998).,1998).
 The xhotometrie solutions were stable from night-to-night with a zero-point variation of less than mamas., The photometric solutions were stable from night-to-night with a zero-point variation of less than mag.
 The estimated uncertainty in the zero-point is less than mumag., The estimated uncertainty in the zero-point is less than mag.
 Astrometric calibration for each field wasperformed. using the positions of at least 20 USNO-2 stars., Astrometric calibration for each field wasperformed using the positions of at least 20 USNO-2 stars.
 These solutions, These solutions
The launch of the and the missions has opened new opportunities in the study of the Χ-rav properties of normal galaxies.,The launch of the and the missions has opened new opportunities in the study of the X-ray properties of normal galaxies.
 Apart from targeted observations of local systems. surveys performed. by these telescopes have allowed. for the first time. study of normal ⋏∙≟⋜↧↓⋜∟∖⊲⊓⊾⊳∖⋜∐⇀∖−↓⋅⋜↧∙∖⇁∖∖⊽⋜↧∖⇁∢⊾↓∢⊾⊔⋏∙≟↿↓↕⊳∖∪⋯⊳∖⋠⊔⇂∢⊾↿↓↥⋖⊾⊔∢⋅⋜⊔⋅∣⋡∙∖⇁⇂⊽⊔⊲↓∖⇁∢⋅↓⋅≱∖∢⋅⋡ thus opening the. wav for evolutionary investigations.," Apart from targeted observations of local systems, surveys performed by these telescopes have allowed, for the first time, study of normal galaxies at X-ray wavelengths outside the nearby Universe, thus opening the way for evolutionary investigations."
 Stacking analysis studies using these missions for example. have provided. constraints on the mean X-ray properties (Lx. hardness ratios. Lx/ Lg) of star-forming svstems over the redshift range z=0.13 (Hornschemoeier et al.," Stacking analysis studies using these missions for example, have provided constraints on the mean X-ray properties $L_X$ , hardness ratios, $L_X/L_B$ ) of star-forming systems over the redshift range $z=0.1-3$ (Hornschemeier et al."
 2002: Nandra et al., 2002; Nandra et al.
 2002: Georgakakis ct al., 2002; Georgakakis et al.
 2003a. b: Laird et al.," 2003a, b; Laird et al."
 2004)., 2004).
 The stacking analysis results from the above independent studies when combinedtogether are consistent, The stacking analysis results from the above independent studies when combinedtogether are consistent
has to be satisfied.,has to be satisfied.
" Hore f,—(l is a geoinetrie factor. where 0,. is the maximo lab-frauie incidence angle between protons and photons."," Here f_g is a geometric factor, where $\theta_{p\gamma}$ is the maximum lab-frame incidence angle between protons and photons."
" The maxinuun potential drop of a magnetar with the rotation frequency. ο=2a/P (where P is the period). and surface magnetic field at the pole. By,=10:GD,as. B where R=1096R; is the stellar radius."," The maximum potential drop of a magnetar with the rotation frequency, $\Omega=2 \pi/P$ (where $P$ is the period), and surface magnetic field at the pole, $B_p=10^{15}{\rm G} B_{p,15}$, is where $R=10^6{\rm cm}R_6$ is the stellar radius."
" For a fraction f=O05f, (the nominal value is for the case assmuine a random distribution of the inchuation augles) of the magnetars Whose spin and magnetic moment poiut in opposite directions. ie. OQ-B,«0. the spin-inchiced parallel electric feld accelerates positive clarecs frou the surface."," For a fraction $f=0.5 f_{1/2}$ (the nominal value is for the case assuming a random distribution of the inclination angles) of the magnetars whose spin and magnetic moment point in opposite directions, i.e., ${\bf \Omega} \cdot {\bf B}_p < 0$, the spin-induced parallel electric field accelerates positive charges from the surface."
 The neutron star surface composition is poorly known., The neutron star surface composition is poorly known.
 Current modeling of the N-rav. data from sole nearby isolated neutron stars sugeests that it likely consists of Leht clements such as bydrogen or helimu (Zavlin. Pavlov Πήρα 1998: Sauwal et al.," Current modeling of the X-ray data from some nearby isolated neutron stars suggests that it likely consists of light elements such as hydrogen or helium (Zavlin, Pavlov Trümmper 1998; Sanwal et al."
 2002) rather than heavy elements such as iron., 2002) rather than heavy elements such as iron.
 Receutlv. Ioralina et al. (," Recently, Ibrahim et al. ("
2002. 2003) reported discoveries of the evclotron resonance features οι. SCR 1806-20 during outburst which are well interpreted as proton cyclotron features iu a magnetar euvironiment. leudiug credeuce that the surface conrposition of maguetars is hivdrogen.,"2002, 2003) reported discoveries of the cyclotron resonance features from SGR 1806-20 during outburst which are well interpreted as proton cyclotron features in a magnetar environment, lending credence that the surface composition of magnetars is hydrogen."
 Tere we assume a hydrogen composition of the magnuoetar surface., Here we assume a hydrogen composition of the magnetar surface.
 According to the standard pulsar theory. protons are accelerated from the polar cap region within charec-depleted gaps (ee. Ruderman Sutherland 1975: Arous Scharlemanu 1979).," According to the standard pulsar theory, protons are accelerated from the polar cap region within charge-depleted gaps (e.g. Ruderman Sutherland 1975; Arons Scharlemann 1979)."
 If surface teiiperatures of imagnetars are high enough to allow free cussion of protous. space-charge limited flow models of electron acceleration. (Arous Sharlemiaun 1979: Πάνοπιο Musliniov 1998. 2001. 2002) will apply to the acceleration of protous in this case.," If surface temperatures of magnetars are high enough to allow free emission of protons, space-charge limited flow models of electron acceleration (Arons Sharlemann 1979; Harding Muslimov 1998, 2001, 2002) will apply to the acceleration of protons in this case."
 Alaeuetars are also strong X-ray enütters., Magnetars are also strong X-ray emitters.
 Divine the carly epochs of their lives. maguetars cit thermal radiation. thought to be due to decay of the strong magnetic fields.," During the early epochs of their lives, magnetars emit thermal radiation, thought to be due to decay of the strong magnetic fields."
 This maintains a high XN-rav hDwunuinositv (typically 100107609st as observed in SCRs/ANPs) over a long period of time., This maintains a high X-ray luminosity (typically $10^{35}-10^{36} {\rm erg~s^{-1}}$ as observed in SGRs/AXPs) over a long period of time.
" The nominal Seld. decay law. dB,fd=ΜΕ 0. leads to a iaeuetic feld strength time dependence Bf)Βρυ1ltPME (Colpi et al."," The nominal field decay law, $d B_p/d t=-b B_p^{1+\delta}$ , leads to a magnetic field strength time dependence $B_p(t) = {B_{p,0}}/ 
{[1+t/\tau_{mag}]^{1/\delta}}$ (Colpi et al."
" 2000). where B,.9 is the initial surface field at the pole."," 2000), where $B_{p,0}$ is the initial surface field at the pole."
" Thus the maguetar surface field. and leuce its quiescent X-ray luminosity. remains almost constant for a typical decaying timescale 7,44(bàD)|o10! vr."," Thus the magnetar surface field, and hence its quiescent X-ray luminosity, remains almost constant for a typical decaying timescale $\tau_{mag} = (b \delta B_{p,0}^\delta)^{-1} \sim 10^4$ yr."
 The observed blackbody temperature for SCR/ANP quiescent enission is &Z4~(0.1.0.6) keV (e.g. IDuxlev 2000: Mereeheti 2001: Thompson 20013., The observed blackbody temperature for SGR/AXP quiescent emission is $k T_\infty \sim (0.4-0.6)$ keV (e.g. Hurley 2000; Meregheti 2001; Thompson 2001).
 The typical surface photon euergy is therefore 2SkT OL)zu eC.2.1) where (d|ty)Leb as the near-surface gravitational redshift., The typical near-surface photon energy is therefore = 2.8 k (1+z_g) (1.6-2.4) where $(1+z_g) \sim 1.4$ is the near-surface gravitational redshift.
 The threshold protou cnerev frou 0q.(2)) is therefore The potential across the polar cap [eq.(0))]. drops as Po? as the star spins down., The threshold proton energy from \ref{th}) ) is therefore The potential across the polar cap \ref{Phi}) )] drops as $P^{-2}$ as the star spins down.
" The typical eucrgv of the protons accelerated in the iuner gap is a fraction of e®, which also drops as the magnetar ages."," The typical energy of the protons accelerated in the inner gap is a fraction of $e\Phi$ , which also drops as the magnetar ages."
" The proton energv acquired in the iuuer gap can be written as €ye®. where yy, paraucterizes the uncertainties inpoc the utilization of the polu cap unipolar potential."," The proton energy acquired in the inner gap can be written as $\epsilon_{p} = \eta_p e\Phi$, where $\eta_p$ parameterizes the uncertainties in the utilization of the polar cap unipolar potential."
" The maxim efficicucy could be as high as 51,502(O15(0.55). (ο. and |7] of Zhang. Uarding ADushinov 2000).x if the electric feld parallel to the naenetic field line is not effectively screcued. and if the article acceleration is not radiatiou-reaction-Inuited."," The maximum efficiency could be as high as $\eta_{p,max} \sim
(0.15-0.85)$ (eqs.[2] and [7] of Zhang, Harding Muslimov 2000), if the electric field parallel to the magnetic field line is not effectively screened, and if the particle acceleration is not radiation-reaction-limited."
 For xoton accelerators. radiation reaction is negligible for he known magnetar candidates (SCRs/ANPs). mainly vecause of the small magnetic field curvature in slow rotators.," For proton accelerators, radiation reaction is negligible for the known magnetar candidates (SGRs/AXPs), mainly because of the small magnetic field curvature in slow rotators."
" Soft N-ravs could be Loreutz-boosted iu the xotons« rest frame and pair-produce in its C'oulonib field (Cheng Ruderinan 1977: Heitler 1951). with a οσους secflon σε~ΖασSi)|(2893InCE!ο}218/21]~Z2.7«102""en? for ενερ. Where Z Is atomic umber. o is the fine-structure coustaut. στ is the Thomson cross section. aud ἴ is the Lorentz-))osted. photon cuerev in the rest frame of the proton."," Soft X-rays could be Lorentz-boosted in the proton's rest frame and pair-produce in its Coulomb field (Cheng Ruderman 1977; Heitler 1954), with a cross section $\sigma_{\pm} \sim Z^2 \alpha \sigma_{T}
(3/8 \pi) [(28/9) \ln (E'_\gamma/m_e c^2)-218/27] \sim Z^2 \cdot
7\times 10^{-27}~{\rm cm}^2$ for $\epsilon_p \sim \epsilon_{p,th}$, where $Z$ is atomic number, $\alpha$ is the fine-structure constant, $\sigma_{T}$ is the Thomson cross section, and $E'_\gamma$ is the Lorentz-boosted photon energy in the rest frame of the proton."
 The typical pai-production mean free path is ἐς-(η.σε)$+1.35104 en comparable to that iu old uormal oulsaurs whose primary pairs are produced through non-resonant inverse Compton scattering.," The typical pair-production mean free path is $l_{\pm} \sim (n_\gamma \sigma_{\pm})^{-1} \sim
1.3\times 10^4$ cm, comparable to that in old normal pulsars whose primary pairs are produced through non-resonant inverse Compton scattering."
 Comparing with the muuerical results in the pulsar case (Tarding Muslimov 2002). the pairs produced iu such au cuviromment are likely ο be too few to fully screen the parallel electric field befor5 he protons reach the photomeson threshold.," Comparing with the numerical results in the pulsar case (Harding Muslimov 2002), the pairs produced in such an environment are likely to be too few to fully screen the parallel electric field before the protons reach the photomeson threshold."
 The cascade xir multiplicitv for old iagnuetars is small., The cascade pair multiplicity for old magnetars is small.
 As a result. Hp9Hpanas Could bo achieved in old inaguetars. so that Although iu voung maenetars the parallel electric fields may be screened before εν reaching ενας. this nonetheless happeus well above εφέ.," As a result, $\eta_p \sim \eta_{p,max}$ could be achieved in old magnetars, so that Although in young magnetars the parallel electric fields may be screened before $\epsilon_p$ reaching $\epsilon_{p,max}$, this nonetheless happens well above $\epsilon_{p,th}$."
" Thus. oue cau define a ""neutro death valley” for maguetars by requiring that €,,,4,. (7)) exceeds the threshold cuerey €, (6)). which gives The range of periods in the right haud side of eq.(8)) defines two diagonal Ties iu the PB, plane (Figure 1)"," Thus, one can define a “neutrino death valley” for magnetars by requiring that $\epsilon_{p,max}$ \ref{eps_pmax}) ) exceeds the threshold energy $\epsilon_{p,th}$ \ref{eps_th}) ), which gives The range of periods in the right hand side of \ref{valley}) ) defines two diagonal lines in the $P-B_p$ plane (Figure 1)."
 Photomeson interactions and neutrino enission cease when the magnetar crosses this valley from left to right caving its evolution., Photomeson interactions and neutrino emission cease when the magnetar crosses this valley from left to right during its evolution.
 Magnetars ling in the valley itself are mareinal neutrino enmütters. ie. they could be neutrino-loud for favorable paramcters.," Magnetars lying in the valley itself are marginal neutrino emitters, i.e., they could be neutrino-loud for favorable parameters."
 Tere we investigate the possibility of detecting neutrino enission frou the individual known maguctarcandidates (ie. SCBs and AXNPs). which aretypically slow rotators.," Here we investigate the possibility of detecting neutrino emission from the individual known magnetarcandidates (i.e. SGRs and AXPs), which aretypically slow rotators."
 As seen in eq.G)). the location of the ucutrino death valley depends ou the geometric factor fy (eq.[2]|).," As seen in \ref{valley}) ), the location of the neutrino death valley depends on the geometric factor $f_g$ \ref{fg}] ])."
 Below we will discuss the possible value of this factor., Below we will discuss the possible value of this factor.
 For the, For the
best-fit model with a more realistic PEOS will require a higher nunuber deusity. η. or a lower particle llbass.,"best-fit model with a more realistic PEOS will require a higher number density, $n$, or a lower particle mass."
 Another possible imiprovement is that without fixing the deusity ratio of total matter aud lieht ermions. tle deusities of light. ferinious aud colistouless cold dark matter particles can be treate(| separately in terms of the equation of hydrostatic equilibrium for the former aud the Jean equation for the atter. while they are coupled throug1 the total encircled mass aud the common velocivy clispersion at ag ]ven r.," Another possible improvement is that without fixing the density ratio of total matter and light fermions, the densities of light fermions and collisionless cold dark matter particles can be treated separately in terms of the equation of hydrostatic equilibrium for the former and the Jean equation for the latter, while they are coupled through the total encircled mass and the common velocity dispersion at a given $r$."
 Oue cosinological subject related particularity to lthe degeneracy of relie massive neutrinos. is its poential elTec ou “free streamiue”.," One cosmological subject related particularity to the degeneracy of relic massive neutrinos, is its potential effect on “free streaming”."
 Relic neutri105 a‘e hot dark matter. which were 'elativistic. when they clecou)ed.," Relic neutrinos are hot dark matter, which were relativistic, when they decoupled."
 Frol the picture of completely collisiouless particles. it is cousidered 1that s1uall-scale [luctuaious 1 ithe primordial neutrino density fluctuation spectrum should have )een wiped oul.," From the picture of completely collisionless particles, it is considered that small-scale fluctuations in the primordial neutrino density fluctuation spectrum should have been wiped out."
 However. the partial degeneracy of relic neutriunos implies that not all of lein coid have |j)ehaved as (ree. pa‘ticles.," However, the partial degeneracy of relic neutrinos implies that not all of them could have behaved as free particles."
" It is couceivable that the effect of Fermi-Dirac degeneracye is suppress the ""[ree streamine” or to preserve 1ie. initial [Iuctuation spectrum.", It is conceivable that the effect of Fermi-Dirac degeneracy is to suppress the “free streaming” or to preserve the initial fluctuation spectrum.
" Depeuding oi degree of this. ""[ree-streaulng suppression” of neutrinos or lermionic hot dark matter in geuer: our expectation on the iütial clensity [Iuctuation spectrum may need to be revised."," Depending on the degree of this “free-streaming suppression” of neutrinos or fermionic hot dark matter in general, our expectation on the initial density fluctuation spectrum may need to be revised."
 We plan to investigate this potential effect quantitatively 1 tlie near future., We plan to investigate this potential effect quantitatively in the near future.
 I1 this paper. we |ave shown that the observed mass profile at the center of the cluster ol gal:les. ALGSY. is reooduced by models assuming the preseuce of degenerate fermionic dark Inatte," In this paper, we have shown that the observed mass profile at the center of the cluster of galaxies, A1689, is reproduced by models assuming the presence of degenerate fermionic dark matter."
 In the case that oreviously unknown lermious with spit 1/2 dominate the dark matter. the acceptable particle mass range is between 2 aud [ eV. Iu tlie case that the dark matter cousists of the mixture of the degererate relic neutrinos aud classical collisionless cold dark matter particles. the mass range of ueutriuos is between 1 aud 2 eV. if the ratio of the two kiuds of dark matter partices is fixed to its €osmic value.," In the case that previously unknown fermions with spin 1/2 dominate the dark matter, the acceptable particle mass range is between 2 and 4 eV. In the case that the dark matter consists of the mixture of the degenerate relic neutrinos and classical collisionless cold dark matter particles, the mass range of neutrinos is between 1 and 2 eV, if the ratio of the two kinds of dark matter particles is fixed to its cosmic value."
 Both the pure fermionic dark matter inodels aud moclels reooduce the observations equally well., Both the pure fermionic dark matter models and neutrino-CDM-mixture models reproduce the observations equally well.
 We tlauk the anonymous referee for useful comments., We thank the anonymous referee for useful comments.
 This studs was motivated by a very interesting talk given by Tom Broadhurst in 2001 at ΝΑΟΙ. Mitaka. Japan.," This study was motivated by a very interesting talk given by Tom Broadhurst in 2004 at NAOJ, Mitaka, Japan."
here wis in racians ancl we have scaled {οι and (he supermassive black hole mass. Vp). bv tvpical values.,"here $\omega$ is in radians and we have scaled $L_{bol}$ and the supermassive black hole mass, $M_{BH}$, by typical values."
 Unless e<0.01. &>5° should be allowed for even verv powerful blazars.," Unless $\epsilon < 0.01$, $\omega > 5^{\circ}$ should be allowed for even very powerful blazars."
 It is interesting (o examine our picture in the context of the “spine + sheath” model ol the jet. e.g. as proposed by Chiaberge et ((2000) to bring the data into accord with 1e orientationr based unification model (e.g.. Gopal-Ixrishia 1995: Carry Padovani 1995).," It is interesting to examine our picture in the context of the “spine + sheath” model of the jet, e.g., as proposed by Chiaberge et (2000) to bring the data into accord with the orientation based unification model (e.g., Gopal-Krishna 1995; Urry Padovani 1995)."
 By considering the nuclear emission in radio. optical aud x-rav bands from BL Lacs and jejr presumed misaligned counterparts. (he FR I radio galaxies. (μον inferred (hat a mildly relativistic (D—1 2) sheath component is the prime contributor to the observed nuclear emission from ER. 1 RGs.," By considering the nuclear emission in radio, optical and x-ray bands from BL Lacs and their presumed misaligned counterparts, the FR I radio galaxies, they inferred that a mildly relativistic $\Gamma = 1 - 2$ ) sheath component is the prime contributor to the observed nuclear emission from FR I RGs."
 At the same time they showed that the spectral energy distribution ol the beamed counterparts. i.e.. BL Lacs. all the wav [rom radio to 5-rays can be modeled in terms of a relativistic spine component moving will a bulk-Lorenz [actor E15—20.," At the same time they showed that the spectral energy distribution of the beamed counterparts, i.e., BL Lacs, all the way from radio to $\gamma$ -rays can be modeled in terms of a relativistic spine component moving with a bulk-Lorenz factor $\Gamma \simeq 15-20$."
 Süll hieher DP values have been estimated in several other studies. particularly for TeV blazars. as mentioned above.," Still higher $\Gamma$ values have been estimated in several other studies, particularly for TeV blazars, as mentioned above."
 In this work we have only discussed the spine component ancl tried (o address the question: ean such large D values be reconciled with the statistics of superhuminal motion of the VLBI knots of blazars/BL Lacs. which suggest a (vpical D—3 ," In this work we have only discussed the spine component and tried to address the question: can such large $\Gamma$ values be reconciled with the statistics of superluminal motion of the VLBI knots of blazars/BL Lacs, which suggest a typical $\Gamma \sim 3$ "
inages D.. the flare corresponds to the λαο fat on Ib! phase of plume collapse.,"images , the flare corresponds to the `plume flat on limb' phase of plume collapse."
 The modeled helt curve declines very steeply after this time. reflecting our adopted sharp velocity cutoff.," The modeled light curve declines very steeply after this time, reflecting our adopted sharp velocity cutoff."
 The observed flare occurs somewhat earlier. iuplviug a slelthy smaller cutoff velocity for this impact.," The observed flare occurs somewhat earlier, implying a slightly smaller cutoff velocity for this impact."
 The more gradual decline in the observed flare also sugeests that the velocity cutoff in the real pluie was not quite so sharp as in the model. ic.. the outer edee of our mnodeled vanenard is too sharp.," The more gradual decline in the observed flare also suggests that the velocity cutoff in the real plume was not quite so sharp as in the model, i.e., the outer edge of our modeled vanguard is too sharp."
 Models which eliminate the vauguard cutirely do not produce an acceptable flare phase in their light curves., Models which eliminate the vanguard entirely do not produce an acceptable flare phase in their light curves.
 Figure 6. shows the light curve observed for the C impact by?.. plotted on a log scale iu comparison to our moceled Lelt curve.," Figure \ref{fig23lc} shows the light curve observed for the G impact by, plotted on a log scale in comparison to our modeled light curve."
 The morphologies of these light curves are very simular to each other. again indicating that the model is accounting to a large deeree for the splashback plysics.," The morphologies of these light curves are very similar to each other, again indicating that the model is accounting to a large degree for the splashback physics."
 Esscutially all features of the observed Lelt curve fiud counterparts im the model., Essentially all features of the observed light curve find counterparts in the model.
 These are marked PC3 (third precursor). ME (ain event). F (flare). D (bounce). and 2B (secoud bounce).," These are marked PC3 (third precursor), ME (main event), F (flare), B (bounce), and 2B (second bounce)."
 The PCS feature was identified iu the observations by?., The PC3 feature was identified in the observations by.
. The presence of third precursors on the ascending branches of Lelt curves is a natural consequence of the expansion of the lateral shock. as discussed iu refssecO9lc..," The presence of third precursors on the ascending branches of short-wavelength light curves is a natural consequence of the expansion of the lateral shock, as discussed in \\ref{ssec09lc}."
 In. rofüe23le.. the observed PC3 imunecdiately follows the second. precursor (from the fireball). which we do not model with ZEUS.," In \\ref{fig23lc}, the observed PC3 immediately follows the second precursor (from the fireball), which we do not model with ZEUS."
 Both the modeled. and observed Leht curves exhibit a broad main eveut phase between 500 aud 1000 sec., Both the modeled and observed light curves exhibit a broad main event phase between 500 and 1000 sec.
 Both ME features are similar in that they decline more sharply than they asceud., Both ME features are similar in that they decline more sharply than they ascend.
 This is due to the sudden collapse of the vanenard in the dare phase: iu the modeled light curve this creates a particularly sharp drop in intensity near 1000 sec., This is due to the sudden collapse of the vanguard in the flare phase; in the modeled light curve this creates a particularly sharp drop in intensity near 1000 sec.
 This bears an especially strong resemibleuce to observations of the I& impact(??).. but is uulike Keck observations of the simaller R impact(?).. which show a smoother and more sviuuetrie main event.," This bears an especially strong resemblence to observations of the K impact, but is unlike Keck observations of the smaller R impact, which show a smoother and more symmetric main event."
 A tendency toward smoother light curves for sanaller inüpacts was previously noted by ZM95. and may indicate the lack of a vanguard for aller plumes (see rofsseciruvle)).," A tendency toward smoother light curves for smaller impacts was previously noted by ZM95, and may indicate the lack of a vanguard for smaller plumes (see \\ref{ssecmrnvlc}) )."
 Both light curves exhibit a secoudary peak (the bouuce) about 500GOO sec after the main event peak., Both light curves exhibit a secondary peak (the bounce) about 500–600 sec after the main event peak.
 Uuderstaudiug this feature was the original motivation for our work., Understanding this feature was the original motivation for our work.
 It is discussed iu detail iu rofsseciurbounco.., It is discussed in detail in \\ref{ssecmrbounce}.
 There is one significant difference between the modeled and observed lisht curves of reffie23le.., There is one significant difference between the modeled and observed light curves of \\ref{fig23lc}.
 The modeled curve decays at a significantly ereater overall rate following the main event peak., The modeled curve decays at a significantly greater overall rate following the main event peak.
 To some extent this may be due to saturation iu the observed fluxes near peak(7).. but it is also likely that the model cooling rate. which occurs via radial expansion aud radiative damping. is too large.," To some extent this may be due to saturation in the observed fluxes near peak, but it is also likely that the model cooling rate, which occurs via radial expansion and radiative damping, is too large."
 Since the model does not iuchide viscous friction. the plumes expaud too rapidly in the model as compared to the observations (sce rofsseciurJ1r]).," Since the model does not include viscous friction, the plumes expand too rapidly in the model as compared to the observations (see \\ref{ssecmr34r}) )."
 Because radiation at this waveleneth is quite tempecrature-scusitive. a too-rapid rate of cooling will be quite noticeable.," Because radiation at this wavelength is quite temperature-sensitive, a too-rapid rate of cooling will be quite noticeable."
 Inclusion of viscosity aud waveleneth-dependent opacitics iu à future version of the model should allow us to refine the agreement with observations., Inclusion of viscosity and wavelength-dependent opacities in a future version of the model should allow us to refine the agreement with observations.
 Figure ? shows the light curve observed for the ID nupact. by2.. in comparison to our model.," Figure \ref{fig12lc} shows the light curve observed for the H impact by, in comparison to our model."
 We again see good agreement., We again see good agreement.
 The predominant feature of thermalIR Πο curves is their relatively slow rate of desceut following the main event peak., The predominant feature of thermal-IR light curves is their relatively slow rate of descent following the main event peak.
 This produces an asyiuuetry wherein the rise to the main event occurs more rapidly than the decline., This produces an asymmetry wherein the rise to the main event occurs more rapidly than the decline.
" The model reproduces this behavior: the iuteusitv near 2000 seconds. relative to the peak inteusity. is τον similar in the two Curves,"," The model reproduces this behavior; the intensity near 2000 seconds, relative to the peak intensity, is very similar in the two curves."
 Two differences between the modeled and observed light curves are noticeable ou reffiel2le.., Two differences between the modeled and observed light curves are noticeable on \\ref{fig12lc}.
 First. the main event is too broad iu the model relative to the observations.," First, the main event is too broad in the model relative to the observations."
 Second. a weal flare feature reais evident iu the modeled helt curve.," Second, a weak flare feature remains evident in the modeled light curve."
 While the dare is a dominant feature in both model aud observations at the shortest wavelengths reffigd9le)}. it does uot appear in the observations at 12pau.," While the flare is a dominant feature in both model and observations at the shortest wavelengths \\ref{fig09lc}) ), it does not appear in the observations at 12."
 The modeled flare exhibits the same qualitative behavior. being more prominent at shorter waveleneths. but does not fade sufficiently at this long thermal wavelength: this discrepancy could well derive from our crude οταν opacity.," The modeled flare exhibits the same qualitative behavior, being more prominent at shorter wavelengths, but does not fade sufficiently at this long thermal wavelength; this discrepancy could well derive from our crude gray opacity."
 Both the modeled aud observed light curves in Fies., Both the modeled and observed light curves in Figs.
 6 aud 7 exhibit a secondary maxiuun. or bounce. at ~600 seconds following the main event (—1200—L100 seconds post-npact).," \ref{fig23lc} and \ref{fig12lc} exhibit a secondary maximum, or bounce, at $\sim 600$ seconds following the main event $\sim 1200-1400$ seconds post-impact)."
 The modeled bounces are similar im amplitude to the observations. aud at least the first bounce occurs at the correct time in the model.," The modeled bounces are similar in amplitude to the observations, and at least the first bounce occurs at the correct time in the model."
 Civen the strong similarity between the model aud observations. we will look to the model to shed Lelt ou the plivsical cause of this phenomenon.," Given the strong similarity between the model and observations, we will look to the model to shed light on the physical cause of this phenomenon."
 We have investigated the nature of the modeled bouuces extensively., We have investigated the nature of the modeled bounces extensively.
 We extended the height of the top bouudary in the models (as high as 2000 kin). in order to be certain that material which rebounds upward is not lost via model boundary conditions.," We extended the height of the top boundary in the models (as high as $2000$ km), in order to be certain that material which rebounds upward is not lost via model boundary conditions."
 We used a variety of Jovian model atmospheres in order to evaluate the scusitivity to atmospheric parameters., We used a variety of jovian model atmospheres in order to evaluate the sensitivity to atmospheric parameters.
 We discuss several potential causes of this phenomenon., We discuss several potential causes of this phenomenon.
 A widely-accepted explanation for the bounce (c.g.?.. Z96) is matter deflected upward from the splasliback shocks. falling back ballistically at later times.," A widely-accepted explanation for the bounce (e.g., Z96) is matter deflected upward from the splashback shocks, falling back ballistically at later times."
 However. this matter is deflected upward at relatively shallow angles. reminiscent of a stoue skipped across the surface of a pond. aud its re-cutry is not the principal contributor to the bounce.," However, this matter is deflected upward at relatively shallow angles, reminiscent of a stone skipped across the surface of a pond, and its re-entry is not the principal contributor to the bounce."
 Iustead.," Instead,"
Bevoud the local Universe. however. the redshitted 21 jan baud probes wavelengths containing the aromatic enüssions (e.g. polycyclic aromatic lydrocarbons. hereafter PAID: the strongest PATT cinission complexcs at 6.2. 7.7. and 8.6 jnu (e.g.Sinithetal.2007). redshift into the21 gam baud by 2~2.,"Beyond the local Universe, however, the redshifted 24 $\micron$ band probes wavelengths containing the aromatic emissions (e.g., polycyclic aromatic hydrocarbons, hereafter PAH); the strongest PAH emission complexes at 6.2, 7.7, and 8.6 $\micron$ \citep[e.g.,][]{JDSmith07}
 redshift into the24 $\micron$ band by $z \sim 2$."
 Although they help boost the 21 gun flux and aid detection of galaxies at high :. ΡΑΠΕ pose two challenges to using sinele-band 21 jan observations to estimate aand SER.," Although they help boost the 24 $\micron$ flux and aid detection of galaxies at high $z$, PAHs pose two challenges to using single-band 24 $\micron$ observations to estimate and SFR."
 Fist. the PAID emission is iufluenced bv cuviromment (e... UV radiation. optical depth). introducing a 0.2. dex scatter to the aand SER estimates (Rousseletal.2001).," First, the PAH emission is influenced by environment (e.g., UV radiation, optical depth), introducing a $\sim 0.2-$ dex scatter to the and SFR estimates \citep{Roussel01}."
. Secoud. the PATI emission appears to strenethen intrinsically at high 2 colnpared to that fouud iu local galaxies with the same citepRigbvüs. Ewurahüs. TakagilO..," Second, the PAH emission appears to strengthen intrinsically at high $z$ compared to that found in local galaxies with the same \\citep{Rigby08, Farrah08, Takagi10}."
 Therefore. the bolometric corrections to determine nuneasured from local galaxies iu the PAID waveleneth reeious will not be applicable at high :.," Therefore, the bolometric corrections to determine measured from local galaxies in the PAH wavelength regions will not be applicable at high $z$."
 Receut studies indicate that applving the local bolometric corrections to lieh-: galaxies will overestimate their aand SER by up to an order of magnitude., Recent studies indicate that applying the local bolometric corrections to $z$ galaxies will overestimate their and SFR by up to an order of magnitude.
 This is reported as the “unic-IR excess” in recent far-IR studies usingSpitzer andBlerschel (ee.Papovichetal.2007: 2011).," This is reported as the “mid-IR excess” in recent far-IR studies using and \citep[e.g.,][]{Papovich07, Elbaz10, Nordon10, Barro11,
 Nordon11}."
. The SED evolution causes a that must be taken iuto account to use 21 jaa observations as aand SER indicators bevond the local Universe., The SED evolution causes a that must be taken into account to use 24 $\micron$ observations as and SFR indicators beyond the local Universe.
 Out to.~ 3. theSpitzer 21 pan observatious probe weaker SFR at anv given redshift than is possible with the far-IR bands (seeFigure{ofElbazetal.2011).," Out to $z \sim 3$ , the 24 $\micron$ observations probe weaker SFR at any given redshift than is possible with the far-IR bands \citep[see Figure 4 of][]{Elbaz11}."
. For JWST. 21 jnu is the longest wavelength band suitable for deep cosmological surveys.," For , 21 $\micron$ is the longest wavelength band suitable for deep cosmological surveys."
 Therefore. our nuderstanuding of star-forming galaxies at high + will depend critically on mid-IR SFR iudicators. where the curent state-of-the-art prescription to estimate ffromi sinele-band 21/421 observations still prescuts a O.L dex seatter (Nordouctal.2011).," Therefore, our understanding of star-forming galaxies at high $z$ will depend critically on mid-IR SFR indicators, where the current state-of-the-art prescription to estimate from single-band 24 $\micron$ observations still presents a $0.4-$ dex scatter \citep{Nordon11}."
 Additionally. theSpetcer 21 pau data are already available from deep legacy surveys (c.¢.. GOODS. FIDEL. COSMOS. SpUDS). aswell as from large-area surveys (c.g... SWIRE and the Doóttes field) that will uot be fully surveyed by current facilities to the same depth in terms of SFR.," Additionally, the 24 $\micron$ data are already available from deep legacy surveys (e.g., GOODS, FIDEL, COSMOS, SpUDS), as well as from large-area surveys (e.g., SWIRE and the Boöttes field) that will not be fully surveyed by current facilities to the same depth in terms of SFR."
 Exploration of mecaus to reduce the current 0.1 dex scatter in cestimates will allow utilization of the mid-IR observations to the fullest extent im the upcoming decade., Exploration of means to reduce the current $0.4-$ dex scatter in estimates will allow utilization of the mid-IR observations to the fullest extent in the upcoming decade.
 Iu this paper. we apply the results frou our previous study. Rujopakaruetal.(2011)... to take iuto account the SED evolution of star-forming galaxies aud refine the 21 ina aandSER incicators.," In this paper, we apply the results from our previous study, \citet{Rujopakarn11}, to take into account the SED evolution of star-forming galaxies and refine the 24 $\micron$ andSFR indicators."
 Rujopakarnetal.(2011) fiud that the IR huninositv surface deusitv.py). affords an accurate description ofthe SED aud subsequently allows accurate bolometric corrections out to Ligh z. specifically out to τ=2.8. the farthest redshift where theSpitzer 21 gan band traces predominantly PAIL euissious.," \citet{Rujopakarn11} find that the IR luminosity surface density, affords an accurate description of the SED and subsequently allows accurate bolometric corrections out to high $z$, specifically out to $z = 2.8$, the farthest redshift where the 24 $\micron$ band traces predominantly PAH emissions."
 The measurement of requires liel-vesolition nuagius of the star-forming regions. Which is only available for a small απο of ealaxies.," The measurement of requires high-resolution imaging of the star-forming regions, which is only available for a small number of galaxies."
 We use these galaxies as a tool to construct a simple formmla to estimate aand SER without measure ffor individual galaxies aud using only single-baud 2 1/4424 flux and redshift measurements. aud then compare the resulting cestimates with results using far-IR data.," We use these galaxies as a tool to construct a simple formula to estimate and SFR without measuring for individual galaxies and using only single-band 24 $\micron$ flux and redshift measurements, and then compare the resulting estimates with results using far-IR data."
 This paper is organized as follows., This paper is organized as follows.
 We ciscuss the evolution of the SEDs of star-forming galaxies aud the use of aas a guide to select the appropriate SED for high-: ealaxies in Section ??.., We discuss the evolution of the SEDs of star-forming galaxies and the use of as a guide to select the appropriate SED for $z$ galaxies in Section \ref{sec:sedevol}.
 The derivation of the siugle-baud indicator is described in Section ??.., The derivation of the single-band indicator is described in Section \ref{sec:indicator}.
 We then compare the new Hudicator with other independent measurements in Section ??and preseut the SSER indicator imu Section ??.., We then compare the new indicator with other independent measurements in Section \ref{sec:indicator_test} and present the SFR indicator in Section \ref{sec:final_SFR}.
 Lastly. we discuss the validity of rest-frame 821 ganas a tracer of aand SER. along with tuplicatious of our results in Section ??..," Lastly, we discuss the validity of rest-frame $8-24$ $\micron$ as a tracer of and SFR, along with implications of our results in Section \ref{sec:discuss}."
" Throughout this paper. we assumc a ACDAL cosinology with ©,,=0.3. O4.=0.7. aud Il,=TO lau Mpc +t."," Throughout this paper, we assume a $\Lambda$ CDM cosmology with $\Omega_m = 0.3$ , $\Omega_{\Lambda} = 0.7$, and $H_0 =
70$ km $^{-1}$ $^{-1}$."
 We follow the convention of Rickeetal.(2009) aud adopt the definition of oof Sanders&Mirabel(1996): these studies defined sslightly differently (ie. 5—1000 µια vs. 8S.1000. ju). but the resulting vvalues are closely consistent.," We follow the convention of \citet{Rieke09}
 and adopt the definition of of \citet{SandersMirabel96}; these studies defined slightly differently (i.e., $5-1000$ $\micron$ vs. $8 - 1000$ $\micron$ ), but the resulting values are closely consistent."
 We will refer to galaxies with lin the range of 1yt104? and those with Hot ans Lunuuous Infrared Calasxies (LIRGSs) ancl Ultraluminous Infrared Galaxies (ULIRGs}. respectively. or collectively as In this section. we will discuss the challenges in using Spitzer observations at 21 sau to estimate aand SER.," We will refer to galaxies with in the range of $10^{11}-10^{12}$ and those with $>10^{12}$ as Luminous Infrared Galaxies (LIRGs) and Ultraluminous Infrared Galaxies (ULIRGs), respectively, or collectively as In this section, we will discuss the challenges in using observations at 24 $\micron$ to estimate and SFR."
 Major issues are to determine an overall approach for using aas an indicator of a galaxy SED (Section ??3). aud the choice of SED library (Section ?7)). Iabelsecisedevol;otroU ," Major issues are to determine an overall approach for using as an indicator of a galaxy SED (Section \ref{sec:sedevol_intro}) ), and the choice of SED library (Section \ref{sec:sedevol_sedchoice}) )."
ntilinow. mnostestimationso ftheL(TIR aandS galavicsusiiigsiiigle FRof bandIBRobservationsrelgonanissumptionthatthe SE Do fastar forminggalarydocsnotecolvewithredshi ft. ic. thatbolometriccorr zgdlacios dl," Until now, most estimations of the and SFR of galaxies using single-band IR observations rely on an assumption that the SED of a star-forming galaxy does not evolve with redshift, i.e., that bolometric corrections measured from local SED templates can be applied to $z$ galaxies."
"owever, locall Ryalavics. f romwhichweconstructtheSI forminggalacics(quiescentdishs Πα ΜΗ... formationinduccdbygalacyinteractions"," However, local IR galaxies, from which we construct the SED libraries, comprise an inhomogeneous population of both normal star-forming galaxies (quiescent disks) and those with nuclear star-formation induced by galaxy interactions."
 DT hecomtributionof interac formationinercascs with }:: locally. theorcticalst induccdstarburststoachiceetheirLE(TIR ooutputstIHopkinsctal. 2010)... whichisconsistentuithobserca," The contribution of interaction-induced star formation increases with ; locally, theoretical studies suggest that ULIRGs must be dominated by interaction-induced starbursts to achieve their outputs \citep{Hopkins10}, , which is consistent with observations \citep[e.g.,][]{SandersMirabel96,
 Veilleux02}. ."
"tions( zygdlaciesthuscarricsanimplicittassamptionthath gstar forminggalaricsarclikcewisedominatedby ighinteraction induccdstarbursts,", The application of local bolometric corrections to $z$ galaxies thuscarries an implicit assumption that $z$ star-forming galaxies are likewise dominated by interaction-induced starbursts.
Observations have shown. however. that actively star- galaxies athigh : are different from their local counterparts in at least three major aspects.,"Observations have shown, however, that actively star-forming galaxies athigh $z$ are different from their local counterparts in at least three major aspects."
 First. the IR," First, the IR"
or from measurements of the small scale flux power spectrum 2001).,or from measurements of the small scale flux power spectrum .
. The picture arising from these measurements is unclear and when the temperaturc-densityv relationship is parametrized as a power-law {τον={ον the parameters are only loosely constrained in the range Z52(1.3)107Ix and Jz0.OS. depending on the technique and data set.," The picture arising from these measurements is unclear and when the temperature-density relationship is parametrized as a power-law $T_{\rm IGM}=T_0\Delta^\beta$ the parameters are only loosely constrained in the range $T_0\approx(1-3)\times10^4\K$ and $\beta\approx0-0.8$, depending on the technique and data set."
 Given the current level of the data and the importance of Hell reionization. we choose not to include temperature constraints in our analysis and only use these observations to set a reasonable prior on allowedvalues o£ Z5 and 2 (see [or an alternative approach).," Given the current level of the data and the importance of HeII reionization, we choose not to include temperature constraints in our analysis and only use these observations to set a reasonable prior on allowedvalues of $T_0$ and $\beta$ (see for an alternative approach)."
 Constraints on may be converted into constraints on the rate per unit volume at which sources. produce ionizing photonsNoone given assumptions about the source spectrum. the mean free path of ionizing photons. and the distribution of absorbing svstems.," Constraints on may be converted into constraints on the rate per unit volume at which sources produce ionizing photons, given assumptions about the source spectrum, the mean free path of ionizing photons, and the distribution of absorbing systems."
 These properties are poorly constrained by direct observation. às isdiscussed in more detail in re[sec:egamma2nedot..," These properties are poorly constrained by direct observation, as isdiscussed in more detail in \\ref{sec:gamma2ndot}."
 ‘Turning now to the future. more direct constraints on the ionization history may be possible with future 21 em observations.," Turning now to the future, more direct constraints on the ionization history may be possible with future 21 cm observations."
 Experiments such as MW. LOPAL. PAPER. and SIX will observe the 21 em power spectrum £1(4°) over a range of wavenumber A in cillerent redshift bins.," Experiments such as MWA, LOFAR, PAPER, and SKA will observe the 21 cm power spectrum $P_{21}(k)$ over a range of wavenumber $k$ in different redshift bins."
 Analytic niodels and simulations suggest that the shape of /S1(&) on scales A—OL1Alpe will provide good constraints on the bubble filling fraction ας} in cach redshift bin., Analytic models and simulations suggest that the shape of $P_{21}(k)$ on scales $k\sim0.1-1\iMpc$ will provide good constraints on the bubble filling fraction $x_i(z)$ in each redshift bin.
" Fisher matrix analvses sugeest that AIWA will be able to place constraints on rg in a fully neutral IGM at LcsNabthe c,—02 level. while SEA should be capable of placing constraints of m,,=0.07."," Fisher matrix analyses suggest that MWA will be able to place constraints on $x_H$ in a fully neutral IGM at $z=8$ at the $\sigma_{x_i}=0.2$ level, while SKA should be capable of placing constraints of $\sigma_{x_i}=0.07$."
 Adding in Planck €MD. data to break cegencracies with cosmological parameters improves these constraints by a factor of ~2., Adding in Planck CMB data to break degeneracies with cosmological parameters improves these constraints by a factor of $\sim2$.
 Similarly. by looking at the evolution of the amplitude ancl slope of {νι find. that AIWA could. constrain the midpoint of reionization with accuracy 0;=0.5+0.05.," Similarly, by looking at the evolution of the amplitude and slope of $P_{21}$, find that MWA could constrain the midpoint of reionization with accuracy $x_i=0.5\pm0.05$."
 21 cm experiments may also eive information about astrophysical parameters such as the minimum mass in which the first galaxies form3005)., 21 cm experiments may also give information about astrophysical parameters such as the minimum mass in which the first galaxies form.
.. The success of these instruments is dependent upon reionization occurring within their sensitivity window., The success of these instruments is dependent upon reionization occurring within their sensitivity window.
 Phese experiments are initially onlv likely to be sensitive to the »ower spectrum for frequencies above 100. MIlIg (although hey are designed to cover frequencies down to SO MlIZ). which corresponds to redshifts z13.," These experiments are initially only likely to be sensitive to the power spectrum for frequencies above 100 MHz (although they are designed to cover frequencies down to 80 MHz), which corresponds to redshifts $z\lesssim13$."
 At higher redshifts. svnchrotron emission from the galaxv becomes extrenie. errestrial racio interference becomes more problematic. aud he Earth's ionosphere is more of a challenge.," At higher redshifts, synchrotron emission from the galaxy becomes extreme, terrestrial radio interference becomes more problematic, and the Earth's ionosphere is more of a challenge."
 Lt is therefore important to gauge what the ionization history is doing in his range of recshilts., It is therefore important to gauge what the ionization history is doing in this range of redshifts.
 A second use of the 21 em line comes from looking at the evolution of the global brightness temperature1999).. rather its Uuetuations.," A second use of the 21 cm line comes from looking at the evolution of the global brightness temperature, rather its fluctuations."
 Experiments such as EDC 2007).. which look for a sharp transition in the 21 em signal brought about by reionization essentially constrain οας and so che; fds.," Experiments such as EDGES , which look for a sharp transition in the 21 cm signal brought about by reionization essentially constrain $\ud T_b/\ud z$ and so $\ud x_i/\ud z$ ."
 Beating experimental systematics down by a [actor of 10 or so should enable these experiments to begin. placing meaningful constraints on ο), Beating experimental systematics down by a factor of 10 or so should enable these experiments to begin placing meaningful constraints on $x_i(z)$.
 Again. it is important to ascertain the amplitude of the signal that these experiments may hope to sce.," Again, it is important to ascertain the amplitude of the signal that these experiments may hope to see."
 We next summarise our methodologv for inferring the ionization history., We next summarise our methodology for inferring the ionization history.
 For a summary of Bayesian inference see ee.(2003).. We wish to attempt predictions for 21 em observations eiven observational constraints from the and the CAIB., For a summary of Bayesian inference see e.g. We wish to attempt predictions for 21 cm observations given observational constraints from the and the CMB.
 Vo do this we make use of Dayes theorem where Al is à model for reionization with parameters ec. and 2 is the combination of observational constraints on and rcm.," To do this we make use of Bayes theorem where $M$ is a model for reionization with parameters $w$, and $D$ is the combination of observational constraints on and $\tau_{\rm CMB}$."
 The evidence pCD|M) provides an overall normalization for the posterior probability p(ie|D.AZ).," The evidence $p(D|M)$ provides an overall normalization for the posterior probability $p(w|D,M)$."
 We must specify our prior ΗΛ on the space of model xwameters., We must specify our prior $p(w|M)$ on the space of model parameters.
 Where measurements are. poor. we take flat alors over a specified range for each parameter. thus pGe|AM)=const.," Where measurements are poor, we take flat priors over a specified range for each parameter, thus $p(w|M)={\rm const}$."
 Our choice of prior should. only be important if the data only weakly. constrains the ionization ustory., Our choice of prior should only be important if the data only weakly constrains the ionization history.
 With these assumptions the posterior. probability is proportional to the likelihood p(Dlw.Al). which is straightforward to calculate.," With these assumptions the posterior probability is proportional to the likelihood $p(D|w,M)$, which is straightforward to calculate."
 Note that we will make he simplifving assumption of Caussian errors in the observational constraints on and When calculating he likelihood. although this is by no means guaranteed.," Note that we will make the simplifying assumption of Gaussian errors in the observational constraints on and when calculating the likelihood, although this is by no means guaranteed."
 The first stepof our inference procedure will be to use Eq., The first stepof our inference procedure will be to use Eq.
 and the data to constrain the model parameters for a set of dillerent. parametrizations., and the data to constrain the model parameters for a set of different parametrizations.
 Then. since cach model provides a definite prediction for .e;(2). we can calculate the probability cistribution for μη at a given redshift from ‘This gives us a bound on the ionization history given our choice of the form of the source model.," Then, since each model provides a definite prediction for $x_i(z)$, we can calculate the probability distribution for $x_i$ at a given redshift from This gives us a bound on the ionization history given our choice of the form of the source model."
 The same procedure can be applied to caleulate bounds on other quantities such as cle;/dz or the SYN fora 21 em experiment., The same procedure can be applied to calculate bounds on other quantities such as $\ud x_i/\ud z$ or the $S/N$ for a 21 cm experiment.
 The evolution of sources is likely to be complex. and to be resistant to description by a small set. o£ numbers., The evolution of sources is likely to be complex and to be resistant to description by a small set of numbers.
 Lowever. by looking at a handful of models and sccing if the predictions are relatively consistent. we can still hope to obtain meaningful predictions.," However, by looking at a handful of models and seeing if the predictions are relatively consistent, we can still hope to obtain meaningful predictions."
 We begin by mapping the constraints in into constraints on andN, We begin by mapping the constraints in into constraints on and.
owe This has been the subject of extensive work2008a).. with cach author making their own preferred. choice of assumptions about uncertain astrophysical parameters.," This has been the subject of extensive work, with each author making their own preferred choice of assumptions about uncertain astrophysical parameters."
 In order to control the assumptions used and to preserve the clfect of the non-lincar mapping of constraints. we perform our own analysis. but where possible we have checked that our. model. which is described in detail in Appendix A.. reproduces the results of previous work.," In order to control the assumptions used and to preserve the effect of the non-linear mapping of constraints, we perform our own analysis, but where possible we have checked that our model, which is described in detail in Appendix \ref{sec:modeling}, , reproduces the results of previous work."
 νο main parameters in our mocdel of reionization arethetwo porauneters describing the temperature-density relation Zi and 2. the spectral index of the sources as. and," The main parameters in our model of reionization arethetwo parameters describing the temperature-density relation $T_0$ and $\beta$ , the spectral index of the sources $\alpha_S$ , and"
sample (Lacy et 11998a) suggests some of the light may be inlluenced by the presence of the radio galaxy. however. so deep spectroscopy of these objects is really needed. to confirm this.,"sample (Lacy et 1998a) suggests some of the light may be influenced by the presence of the radio galaxy, however, so deep spectroscopy of these objects is really needed to confirm this."
 Six objects have no redshift information. ancl are 1plotted as open squares at the top left of 55 with arbitary redshifts.," Six objects have no redshift information, and are plotted as open squares at the top left of 5 with arbitary redshifts."
 Their most likely redshifts are in the range where no strong emission lines all in the optical., Their most likely redshifts are in the range where no strong emission lines fall in the optical.
 They mostly. have faint optical magnitudes ancl so do not allect the discussion below., They mostly have faint optical magnitudes and so do not affect the discussion below.
 There is one interesting [ος outlier on this plot., There is one interesting $z$ outlier on this plot.
 SCITA3|637 (2=21.3: 2= 0.324) is in the cluster AB2SO. close to SCLT42|637. but was originally thought to. be behind it from the point of view of both the radio structure," 8C1743+637 $R=21.3$; $z=0.324$ ) is in the cluster A2280, close to 8C1742+637, but was originally thought to be behind it from the point of view of both the radio structure"
? ? (27): (??).. 7.,"\citet{O50} \citet{HT86} \citep{DQT87,DWLD04}; \citep{CS06,CSS07}."
 ο «©+. (7?).. ?7..," \citet{CS08} $e$ $e^{-1}$ \citep{SZF95}. \citet{CS08},"
" 2. 3. { 5. ro(f). v,(t}, v(t}ri(f)=|re(f)|. απ)=πι)πω). "," \ref{secSingleEnc} \ref{secLevy} \ref{secGalacticTides} \ref{secConclusions} $\vec r_b(t)$ $\vec
v_b(t)$ $\vec r_b(t)$$r_b(t)=|\vec
r_b(t)|$ $\hat r_b(t)=\vec
r_b(t)/r_b(t)$ "
viewing periods. for which the detection significance exceeded the required tlieshold.,"viewing periods, for which the detection significance exceeded the required threshold."
 By minimizine ↕∐↩∖−−⊳∖⋯∐⋜↕⊳∖∑≟↥∖⇁↩∐⋯↕↢↲↽↓⋅∙⋝⋅∙⋝, By minimizing the $\chi^2$ -sum as given in Eq.
⋅−≻∖∖↽≺↵∩∣≻↕⋜⋃∐↥∩↕⋅↕∐⋜∏≺↵∐≺↵↕⋅∑≟⊽∖⊽↕⋅⋜↕∐∑∸≺↵E ⋅ ⋅  ∙∙⋅∣ ⋅ ⋅ The coordinates.H forH whichH the y-um> isH minimized.H aud theirH mareinalizedH 1o-errors are which is right inside the localization contour derived from teh combined data set for the same energy bail (Fig. 5)).," \ref{29} we obtain for that energy range The coordinates, for which the $\chi^2$ -sum is minimized, and their marginalized $\sigma$ -errors are which is right inside the localization contour derived from teh combined data set for the same energy band (Fig. \ref{f5}) )."
 The scatter iu the position estimates lor the 12 viewing periods aud the best-fit position thus determiued are therefore eutirely cousistent with the statistical uncertainties., The scatter in the position estimates for the 12 viewing periods and the best-fit position thus determined are therefore entirely consistent with the statistical uncertainties.
 There is uo evideuce for an acditioual systematic error in the localization., There is no evidence for an additional systematic error in the localization.
 Iu this paper I have re-analyzed EGRET data of the observing periods 1-1 to derive the variability. properties aud localization ofJ1716-2551., In this paper I have re-analyzed EGRET data of the observing periods 1–4 to derive the variability properties and localization of.
. For that purpose I have introduced correctious iu the way the diffuse foreground model is used treated tu the likelihood analysis. aud [ have also perlormed various cousistency checks using data for different energy bauds ancl frou different observing dates.," For that purpose I have introduced corrections in the way the diffuse foreground model is used treated in the likelihood analysis, and I have also performed various consistency checks using data for different energy bands and from different observing dates."
 Usiug only data for which I fouid uo evidence of systematic problems. e.g. excludiug the wide energy baud 100-10.000 MeV aud other low-euerey baucds. I obtained the following results:," Using only data for which I found no evidence of systematic problems, e.g. excluding the wide energy band 100–10.000 MeV and other low-energy bands, I obtained the following results:"
disturbed while the side remote from the interaction stavs relatively unaffected.,disturbed while the side remote from the interaction stays relatively unaffected.
 This behaviour can also he found in the rotation curve of the system (Ixrouberger ct al., This behaviour can also be found in the rotation curve of the system (Kronberger et al.
 2006)., 2006).
 By shifting the VF artificially to üeher redshifts we found that. although sunall scale structures iu the VF aro smieared out. distortions are still visible at intermediate redshifts for large galaxies.," By shifting the VF artificially to higher redshifts we found that, although small scale structures in the VF are smeared out, distortions are still visible at intermediate redshifts for large galaxies."
 Iu the case of sinall galaxies even strong distortions are not visible in the velocity field at 20.5 with currently available angular resolution., In the case of small galaxies even strong distortions are not visible in the velocity field at $z\approx0.5$ with currently available angular resolution.
 Severely distorted kinematics. be. velocity Belds with no or just a inall rotational componcut are identifiable also for nall galaxies at intermediate redshift.," Severely distorted kinematics, i.e. velocity fields with no or just a small rotational component are identifiable also for small galaxies at intermediate redshift."
 Thus they should not be mistaken as uudisturbed., Thus they should not be mistaken as undisturbed.
 If the galaxy is uudisturbed the quautities derived from the velocity field. as e.g. the maxima rotational velocity 44. do not show svstematic variations with resolution and are therefore useable for Tullv-Fisher studies.," If the galaxy is undisturbed the quantities derived from the velocity field, as e.g. the maximum rotational velocity $_{max}$, do not show systematic variations with resolution and are therefore useable for Tully-Fisher studies."
 We showed that with the help of adaptive optics ucar-iufrarec spectrographs (0.5. SINFONT) it is possible to study the internal kinematics of galaxies even at high-redshift (272). given that the observed fux is sufficient to construct a 2D velocity field.," We showed that with the help of adaptive optics near-infrared spectrographs (e.g. SINFONI) it is possible to study the internal kinematics of galaxies even at high-redshift $\sim$ 2), given that the observed flux is sufficient to construct a 2D velocity field."
 Ongoing merecr eveuts are then clearly visible in the VF., Ongoing merger events are then clearly visible in the VF.
 Disturbed velocity fields settle again to a relaxed state after about 1 Cr., Disturbed velocity fields settle again to a relaxed state after about 1 Gyr.
 However. already after several hundred Myr the distorious can get damped such that a duisclassification as 1xdisturbed rotation is possible at ligher redshifts.," However, already after several hundred Myr the distortions can get damped such that a misclassification as 'undisturbed rotation' is possible at higher redshifts."
 With the asstuuption of a constant iass-to-light ratio. we showed that 1o undisturbed model galaxies jio on a Tulh-Fisher relation with a slope of 2. which is consistent with observations of local galaxy salples.," With the assumption of a constant mass-to-light ratio, we showed that the undisturbed model galaxies lie on a Tully-Fisher relation with a slope of $~$ 3, which is consistent with observations of local galaxy samples."
 Using svstenis with uurecognized distortions or a Tulb-Fisher study will sienificautly increase the scatter iun the observed relation., Using systems with unrecognized distortions for a Tully-Fisher study will significantly increase the scatter in the observed relation.
 As it is more likely. hat an iuteraction ds uussed for μια. faint galaxies. he iuclusion of disturbed systems can iuinüc an evolution iu the fait cud of the Tully-Fisher relation.," As it is more likely, that an interaction is missed for small, faint galaxies, the inclusion of disturbed systems can mimic an evolution in the faint end of the Tully-Fisher relation."
 We thank the anouvimous referee for fruitful conuuenuts which improved the quality of the paper., We thank the anonymous referee for fruitful comments which improved the quality of the paper.
 The authors would like to thank Volker Springel for providing them with CADCET2 aud his initialconditious generator., The authors would like to thank Volker Springel for providing them with GADGET2 and his initial-conditions generator.
 We are also thankful to Davor hrajiovió for his Kiuecnmetryv software., We are also thankful to Davor Krajnović for his Kinemetry software.
 Thomas lrouberser is a recipient of a DOC fellowship of the Austrian Academy of Sciences;, Thomas Kronberger is a recipient of a DOC fellowship of the Austrian Academy of Sciences.
 The authors further acknowledge the Unuiblufrastrukturprogranun des BAIWE Forschungsprojekt Konsortiumn.— Hochleistuugsrechuen. the German Science Foundation (DEG) through Grant uunber Zi 663/6-1. and the Volkswagen Foundation το 520).," The authors further acknowledge the UniInfrastrukturprogramm des BMWF Forschungsprojekt Konsortium Hochleistungsrechnen, the German Science Foundation (DFG) through Grant number Zi 663/6-1, and the Volkswagen Foundation (I/76 520)."
 In addition. the authors acknowledge the ESO Mohbilitattsstipendicn des BAIWF (Austria). and the Tiroler Wisseuschaftsfounds: (Ceforrdert aus Mittelu des voii Laud Tirol eiugerichteten: Wisseuschaftsfonds).," In addition, the authors acknowledge the ESO Mobilitättsstipendien des BMWF (Austria), and the Tiroler Wissenschaftsfonds (Geförrdert aus Mitteln des vom Land Tirol eingerichteten Wissenschaftsfonds)."
 We thauk Elif Ikutdeimir for fruitful discussion., We thank Elif Kutdemir for fruitful discussion.
ux of 39.7£0.66 Jv tHellective beam size of 50.3«46.8 aremin) and a total mass of 4.6«1077 ML.,flux of $39.7 \pm 0.66$ Jy (effective beam size of $ 50.3 \times 46.8$ arcmin) and a total mass of $ 4.6 \times 10^{10}$ $_{\odot}$.
 Arecibo observations C2).. with a beam of 3.3x0.1 arcmin. gave an integrated Dux of Lid Jv which implies an mass of 141017 M. but these observations were based on a single pointing at the centre and therefore missed a substantial part of the belonging to 7765.," Arecibo observations \citep{shu94}, , with a beam of $3.3 \pm 0.1$ arcmin, gave an integrated flux of 11.1 Jy which implies an mass of $ 1.4 \times 10^{10}$ $_{\odot}$, but these observations were based on a single pointing at the centre and therefore missed a substantial part of the belonging to 765."
 The total mass derived for 1765 exceeds by one order of magnitude the median mass in the LSB survey hy ?.., The total mass derived for 765 exceeds by one order of magnitude the median mass in the LSB survey by \citet{one04}.
 lt also exeeeds by a factor of 3 the mean value obtained bv ?. for SabShe LSB spiral galaxies within the RCS-UGC sample., It also exceeds by a factor of 3 the mean value obtained by \cite{rob94} for Sab–Sbc HSB spiral galaxies within the RC3-UGC sample.
 The mass to luminosity ratio for 7765 is Mur/Ly1.6. Le. of the same order as the median Mgi/Lp=2.8 found for LSBs by 2.. and 510 times larger than the median for LISB galaxies (?).. NGC," The mass to luminosity ratio for 765 is $_{\rm HI}$ $_{\rm B} \sim 1.6$, i.e., of the same order as the median $_{\rm HI}$ $_{\rm B} = 2.8$ found for LSBs by \citet{one04}, and 5–10 times larger than the median for HSB galaxies \citep{rob94}."
7765 has an extended: distribution. with a diameter of 240 kpe., 765 has an extended distribution with a diameter of 240 kpc.
 In Fig., In Fig.
 Ὁ— (loft) we present contours overlaid on the optical image. highlighting the disproportionately larec disc compared to the optical.," \ref{765HI&OPT} (left) we present contours overlaid on the optical image, highlighting the disproportionately large disc compared to the optical."
 The structure shows an unresolved. central depression with a size ~ 50 30 aresec. better seen in the righthand panel of Fig.," The structure shows an unresolved central depression with a size $\sim$ 50 $\times$ 30 arcsec, better seen in the right–hand panel of Fig."
 5 where we present the distribution. overplotted with velocity contours.," \ref{765HI&OPT} where we present the distribution, overplotted with velocity contours."
 The inner eas distribution follows the spiral arms., The inner gas distribution follows the spiral arms.
 We can identify at least three of them: two to the East and. one to the Southwest., We can identify at least three of them: two to the East and one to the Southwest.
" Phe average column density out to 50 kpc. i.c. where the spiral arms dominate. is 54107"" atoms 7."," The average column density out to 50 kpc, i.e., where the spiral arms dominate, is $5 \times 10^{20}$ atoms $^{-2}$."
" To the East the column density drops bevonel that raclius to reach a plateau at 2.107"" atoms . populated with clumps (see below)."," To the East the column density drops beyond that radius to reach a plateau at $2 \times 10^{20}$ atoms $^{-2}$, populated with clumps (see below)."
 The outer disc of 7765 is asvmmoetrical or lopsided., The outer disc of 765 is asymmetrical or lopsided.
 This can be appreciated in Fig. 6..," This can be appreciated in Fig. \ref{profiles},"
 where we present average surface brightness profiles along the major axis: one along he receding side of the galaxy (hereafter West. profile) and he other along the approaching side of the galaxy (hereafter East profile)., where we present average surface brightness profiles along the major axis: one along the receding side of the galaxy (hereafter West profile) and the other along the approaching side of the galaxy (hereafter East profile).
 These were ereated using the ellipse integration ask available in (2)., These were created using the ellipse integration task available in \citep{vanhul92}.
. The profiles are derived rom wwide sectors centred around the major axis and inned every 40 aresec., The profiles are derived from wide sectors centred around the major axis and binned every 40 arcsec.
 This process has the advantage of compensating for the loss of signalto.noise when moving to arger galactic radii by averaging over larger areas. allowing us to recover information all the way out to the periphery of the diffuse disc.," This process has the advantage of compensating for the loss of signal–to–noise when moving to larger galactic radii by averaging over larger areas, allowing us to recover information all the way out to the periphery of the diffuse disc."
 In Fig. 6.," In Fig. \ref{profiles},"
 we plot column densities (in 7)ona logarithmic scale as a function of linear radius (in kpe)., we plot column densities (in $^{-2}$ ) on a logarithmic scale as a function of linear radius (in kpc).
 Each side ofthe galaxy. presents a dillerent signature., Each side of the galaxy presents a different signature.
" 3oth profiles start out at the same level of5107"" em.> and show a sudden drop near 50 kpe (2.5 aremin).", Both profiles start out at the same level of $5 \times 10^{20}$ $^{-2}$ and show a sudden drop near 50 kpc (2.5 arcmin).
" The eastern profile settles to a plateau of 2107"" em? followed by a rapid decline near LOO kpe out to the last measured point at 120 kpc.", The eastern profile settles to a plateau of $2 \times 10^{20}$ $^{-2}$ followed by a rapid decline near 100 kpc out to the last measured point at 120 kpc.
 The western profile declines smoothly from 70 kpe. extending also out to 120 kpe.," The western profile declines smoothly from 70 kpc, extending also out to 120 kpc."
 Both halves together result in a total extent of 240 kpc. measured at à column density ⋅ ⋅ ∪⇂⇉↓∪⊥⋅⋜⊔∪⊔↓⊳∖≼∼⊔↓ ⊳↘ ↾∐↕⋖⋅↓∪∖∖⊽≻⇂⇂↓," Both halves together result in a total extent of 240 kpc, measured at a column density of $2 \times 10^{19}$ atoms cm $^{-2}$."
⋅⇂⋅⋜↧≼⇍⋖⊾∠⋖⋅↓↕≻↕↿∙∖⇁⋯∐∢⋅↓⋅∠⇂⊲↓⊳∖≼∼∪⇂⋅↓∖⊽≺∶≺⊲⊤⊤∪⋅↱≻⊲↓⊳∖ ≻∪↓≻⊔↓⋜⋯⋅∠⇂∣⋡∙∖⇁≱∖⋖⋅∖⇁∢⊾↓⋅⋜↧↓⊔⊔↓⋅⋖⊾⊳∖∪↓∖⇁∢⊾∠⇂∪↓⋅⊔↓⋜⊔⋅⋏∙≟⊀↓⊔⋜↧∐∙∖⇁↓⋅⋖⊾⊳∖∪↓∖⇁∢⊾∠⇂ compact regions. notably in the NorthEast.," The low surface density outer disc of 765 is populated by several unresolved or marginally resolved compact regions, notably in the North–East."
 They. are prominent in the integrated map (Fig. 5)), They are prominent in the integrated map (Fig. \ref{765HI&OPT}) )
 as well as in he channel maps (Fig. 1))., as well as in the channel maps (Fig. \ref{765_CHAN_1}) ).
 Spectra taken at their location ave indistinguishable from neighbouring areas at the spatial and velocity resolution of the present data. except for their uigher surface brightness.," Spectra taken at their location are indistinguishable from neighbouring areas at the spatial and velocity resolution of the present data, except for their higher surface brightness."
 Integrated: properties of several of hese clumps are ὃνgiven in Table 3. masses. were derived using taskISPEC.. which retrieves the Dux of a region along individual velocity channels).," Integrated properties of several of these clumps are given in Table \ref{HICOMPACT_TABLE} masses, were derived using task, which retrieves the flux of a region along individual velocity channels)."
 Phe compact regions are labeled in Fig. 5.., The compact regions are labeled in Fig. \ref{765HI&OPT}.
 Given that they are at best mareinally resolved. their physical sizes must be less than LOkkpe: their masses are of order LO”M...," Given that they are at best marginally resolved, their physical sizes must be less than kpc; their masses are of order $10^8$."
 Together they represent ~ 10 per cent of the total mass of 7765., Together they represent $\sim$ 10 per cent of the total mass of 765.
 Their properties. are comparable to those of the companion galaxies detected in our field: 7765a and 11453 (see Sect. 4.3)), Their properties are comparable to those of the companion galaxies detected in our field: 765a and 1453 (see Sect. \ref{companions}) )
 and fall within the mass interval for rich companions found hy ?.., and fall within the mass interval for –rich companions found by \cite{pis03}.
 Phe masses of the compact regions are. however an order of magnitude higher than the range quoted hy ? for clumps in disces.," The masses of the compact regions are, however an order of magnitude higher than the range quoted by \cite{wil01} for clumps in discs."
 In addition. regions A and € present optical counterparts (Fig. 7)).," In addition, regions A and C present optical counterparts (Fig. \ref{blobs}) )."
 In both of these. the stellar counterparts and the gas are aligned but interestingly do not follow the spiral arms.," In both of these, the stellar counterparts and the gas are aligned but interestingly do not follow the spiral arms."
 Neither of the remaining compact regions coincides with an optical. radiocontinuum or counterpart.," Neither of the remaining compact regions coincides with an optical, radio–continuum or X--ray counterpart."
 The velocity field presented in Fig., The velocity field presented in Fig.
 5. (right) shows steeply rising rotation in the inner part (up to  40 aresee in raclius) hat is basically unresolved by our beam and a Hat rotation curve onwards., \ref{765HI&OPT} (right) shows steeply rising rotation in the inner part (up to $\sim$ 40 arcsec in radius) that is basically unresolved by our beam and a flat rotation curve onwards.
 Between 2.5 xr= 3.0 aremin there seems o be a kink in the isovelocity contours. best seen. on the ΟΡ side.," Between 2.5 $\leq r \leq$ 3.0 arcmin there seems to be a kink in the isovelocity contours, best seen on the North side."
 The positionvelocity (PV).diagrammes shown in lig., The position–velocity (PV)–diagrammes shown in Fig.
 S are taken along the major (PA= 256°)) and minor axis (PA= 1667)) of the galaxy (see below)., \ref{pv_diagrammes} are taken along the major $=256$ ) and minor axis $=166$ ) of the galaxy (see below).
 We see that x»uh Western and Eastern sides are fairly symmetric within he inner 3 aremin. with velocities rising quite steeply to a flat level of rotation after just 40 aresec. and a jump in velocity at  3 aremin.," We see that both Western and Eastern sides are fairly symmetric within the inner 3 arcmin, with velocities rising quite steeply to a flat level of rotation after just 40 arcsec, and a jump in velocity at $\sim$ 3 arcmin."
 The ellipse visible in the minor axis diagramme (Fig. S..," The ellipse visible in the minor axis diagramme (Fig. \ref{pv_diagrammes},"
 right) is a consequence of the presence of a bar. causing streaming motions in the central part. and our ~ 43 arcsec beam: emission from both the receding and approaching sides enters the beam simultaneously.," right) is a consequence of the presence of a bar, causing streaming motions in the central part, and our $\sim$ 43 arcsec beam: emission from both the receding and approaching sides enters the beam simultaneously."
 We used the task in (7) to modeΙ the velocity field of 7765., We used the task in \citep{beg87} to model the velocity field of 765.
 treats the observed velocity field as a series of tilted rings., treats the observed velocity field as a series of tilted rings.
 As already mentioned by ?. this task becomes increasingly unreliable for. low inclination(laceon) galaxies.," As already mentioned by \citet{beg87}, this task becomes increasingly unreliable for low inclination(face–on) galaxies."
 Ciüven the low inclination of 7765 our results should therefore be treated with some caution., Given the low inclination of 765 our results should therefore be treated with some caution.
 Llaving said that. 7765. provides a unique opportunity to probe its dynamical mass and hence the contribution from the Dark Matter. halo out to an unprecedented distance.," Having said that, 765 provides a unique opportunity to probe its dynamical mass and hence the contribution from the Dark Matter halo out to an unprecedented distance."
 Lt should also be noted that beam smearing will alfect, It should also be noted that beam smearing will affect
neutron stars as it will provide good timing ancl spectral resolution.,neutron stars as it will provide good timing and spectral resolution.
come from the dissipation of bulk and turbulent kinetic energy of gas driven by the galaxy interaction.,come from the dissipation of bulk and turbulent kinetic energy of gas driven by the galaxy interaction.
 Such a loss of Kinetic energy is a required step for this gas to become gravitationally bound., Such a loss of kinetic energy is a required step for this gas to become gravitationally bound.
 In the following. we quantify this interpretation which connects the energetics of the ggas to the formation of the CO complex by gas accretion.," In the following, we quantify this interpretation which connects the energetics of the gas to the formation of the CO complex by gas accretion."
 We consider that the CO complex evolves i à quasi-static way assuming that virial equilibrium applies at all times., We consider that the CO complex evolves in a quasi-static way assuming that virial equilibrium applies at all times.
" The virial equation relates the gravitational and turbulent kinetic energy. EQ, aNd Eu. and includes the external pressure P: where V is the volume."," The virial equation relates the gravitational and turbulent kinetic energy, $E_\mathrm{grav}$ and $E_\mathrm{turb}$, and includes the external pressure $P_\mathrm{ext} $: where $V$ is the volume."
 The turbulent energy is EusMeco σερ: where Mco is the mass of the CO complex and σοο Is the velocity dispersion along the line of sight derived from the integrated CO spectrum of SGMC 2.," The turbulent energy is $E_{\rm turb} = 3/2 \,
M_{\rm CO} \, \sigma_{\rm CO}^2$ , where $M_{\rm CO}$ is the mass of the CO complex and $\rm \sigma_{CO}$ is the velocity dispersion along the line of sight derived from the integrated CO spectrum of SGMC 2."
" For a fixed1 Po. the exchange of energy associated with gas accretion anc1 radiation ts associated with the derivative of the enthalpy Η of the complex (?):: where Fj, Buuis the energy input from gas accretion. Ly, the H» luminosity. and fy, the fraction of the gas bolometric luminosity that is radiated in the H» lines used to compute Ly,."," For a fixed $P_{\rm ext}$, the exchange of energy associated with gas accretion and radiation is associated with the derivative of the enthalpy $H$ of the complex \citep{huff07}: where $\dot{E}_\mathrm{in}$ is the energy input from gas accretion, $L_\mathrm{H_2}$ the $_2$ luminosity, and $f_\mathrm{H_2}$ the fraction of the gas bolometric luminosity that is radiated in the $_2$ lines used to compute $L_\mathrm{H_2}$."
 This fraction is smaller than | because some cooling occurs in lines which have not been measured., This fraction is smaller than 1 because some cooling occurs in lines which have not been measured.
" Combining H» and far-IR observations. ? find that fy, Is in the range ~0.25—0.5 for shocks associated with gas outflow from low mass stars."," Combining $_2$ and far-IR observations, \citet{maret09} find that $f_{\rm H_2}$ is in the range $\sim 0.25\, -
\, 0.5$ for shocks associated with gas outflow from low mass stars."
 We assume that the kinetic energy of the accreted gas is the main source of energy that balances radiative losses (?).., We assume that the kinetic energy of the accreted gas is the main source of energy that balances radiative losses \citep{klessen10}.
 In doing this we neglect the energy that comes from stellar feedback., In doing this we neglect the energy that comes from stellar feedback.
 The fact that the H» emission shows no enhancement around the SSC is an indication that stellar feedback does not have a significant contribution to the H» emission (Sect. ??))., The fact that the $_2$ emission shows no enhancement around the SSC is an indication that stellar feedback does not have a significant contribution to the $_2$ emission (Sect. \ref{sec:stef}) ).
 We follow ? in introducing the efficiency factor ε which represents the fraction of the accretion energy that drives turbulence., We follow \citet{klessen10} in introducing the efficiency factor $\epsilon $ which represents the fraction of the accretion energy that drives turbulence.
 The turbulent energy dissipates over a time scale tu(CO)=R/oco where R=~300pe is the radius of the SGMC 2 complex (?)..," The turbulent energy dissipates over a time scale $t_\mathrm{dis}({\rm CO}) \simeq R/\sigma_{\rm CO}$ where $R \simeq
300 \, {\rm pc}$ is the radius of the SGMC 2 complex \citep{mckee07}."
 The balance equation between energy input and dissipation ts: We compute the mass accretion rate necessary to drive the gas turbulence by equating the energy input and the radiative losses., The balance equation between energy input and dissipation is: We compute the mass accretion rate necessary to drive the gas turbulence by equating the energy input and the radiative losses.
 This corresponds to a solution of equation (2)) where the derivative of the gas kinetic energy ων and the term associated with V are both small with respect to the right-hand terms., This corresponds to a solution of equation \ref{eqH}) ) where the derivative of the gas kinetic energy $\dot{E}_{\rm turb}$ and the term associated with $\dot{V}$ are both small with respect to the right-hand terms.
" Thus. we write: where oy, is the velocity dispersion along the line of sight derived from the H» data."," Thus, we write: where $\sigma_{\rm H_2}$ is the velocity dispersion along the line of sight derived from the $_2$ data."
" To quantify Mj and e. we consider only the turbulent component of the velocity field (FWHM,~150 kms7!). assuming that aceretion occurs from turbulent gas with a mean velocity equal to that of the CO complex."," To quantify $\dot{M}_{\rm acc}$ and $\epsilon$, we consider only the turbulent component of the velocity field $\rm FWHM_{H_2} \sim 150 \, km \, s^{-1}$ ), assuming that accretion occurs from turbulent gas with a mean velocity equal to that of the CO complex."
" For oy,=65kms7! and σρο=30kms7!. we find Mae=5.3/fu, yr! and €2.1.8xfu."," For $\rm \sigma_{H_2} = 65 \, km \, s^{-1}$ and $\rm \sigma_{CO} = 30 \, km \, s^{-1}$, we find $\dot{M}_{\rm acc}
= 5.3 / f_{\rm H_2}\,$ $\, \mathrm{yr^{-1}} $ and $\epsilon = 1.8
\times f_{\rm H_2}$."
" For our interpretation to hold we must have e<|. and thus fy,«0.5."," For our interpretation to hold we must have $\epsilon < 1$, and thus $f_{\rm H_2} < 0.5$."
 The fact that the H» line width is larger than that of CO shows that this holds., The fact that the $_2$ line width is larger than that of CO shows that this holds.
" In the following. for numerical calculations. we consider that fy,~0.25."," In the following, for numerical calculations, we consider that $f_{\rm
H_2} \sim 0.25$."
" We scale Mj, by the turbulence dissipation time scale to estimate the mass of gas from which aceretion is occurring: Mace=XgsH3). where rg(H3)—R/oy, is the turbulent dissipation time scale."," We scale $\dot{M}_{\rm acc}$ by the turbulence dissipation time scale to estimate the mass of gas from which accretion is occurring: $M_{\rm
acc} = \dot{M}_{\rm acc} \times t_{\rm dis}({\rm H_2})$ , where $t_{\rm
dis}(\mathrm{H_2}) \simeq R/\sigma_{\rm H_2}$ is the turbulent dissipation time scale."
 The large scale flow can feed this mass reservoir. compensating for what is accreted by the CO complex. because the dissipation time scale is comparable to the dynamical time scale associated with the velocity gradient across SGMC 2 (3x10* yr).," The large scale flow can feed this mass reservoir, compensating for what is accreted by the CO complex, because the dissipation time scale is comparable to the dynamical time scale associated with the velocity gradient across SGMC 2 $3\times
10^6 \, $ yr)."
" We find Mj,=24x107/fu,M."," We find $M_{\rm acc} = 2.4 \times 10^7 / f_{\rm
H_2}$."
" For fu,-0.25. Mace is of Meo."," For $f_{\rm H_2} \sim 0.25$, $M_{\rm acc}$ is of $M_{\rm CO}$."
 Thus. the mean density in the CO complex may only be a few times larger than that of the accreted gas.," Thus, the mean density in the CO complex may only be a few times larger than that of the accreted gas."
 This agrees with our estimate of the efficiency factor e., This agrees with our estimate of the efficiency factor $\epsilon$.
 Based on their numerical simulations and theoretical arguments. ? argue that the efficiency factor εἰς roughly equal to the ratio between the density of the accreting gas to the mean density of the bound system.," Based on their numerical simulations and theoretical arguments, \citet{klessen10} argue that the efficiency factor $\epsilon$ is roughly equal to the ratio between the density of the accreting gas to the mean density of the bound system."
" Computing the ratio between the CO mass of SGMC 2 and the mass accretion rate. we obtain a time scale fy=7x107xfu, yr."," Computing the ratio between the CO mass of SGMC 2 and the mass accretion rate, we obtain a time scale $t_{\rm acc} = 7 \times 10^7
\times f_{\rm H_2}~\mathrm{yr}$ ."
" For fy,~0.25. this is slightly longer than the 10 Myr time scale over which the dynamics of the interaction may have been driving the convergent flow (?).."," For $f_{\rm H_2} \sim 0.25$, this is slightly longer than the 10 Myr time scale over which the dynamics of the interaction may have been driving the convergent flow \citep{renaud08}."
 However. within the uncertainties in this rough calculation. we consider that the present accretion rate 1s close to the mean rate needed to account for the formation of the SGMC 2 complex as a result of gravitational fragmentation and gas aceretion driven by the galaxy interaction.," However, within the uncertainties in this rough calculation, we consider that the present accretion rate is close to the mean rate needed to account for the formation of the SGMC 2 complex as a result of gravitational fragmentation and gas accretion driven by the galaxy interaction."
 Our interpretation introduces a dynamical view of the present state of the SGMC 2 complex., Our interpretation introduces a dynamical view of the present state of the SGMC 2 complex.
 The virtalized CO complex is physically associated with gas which is too turbulent to be bound., The virialized CO complex is physically associated with gas which is too turbulent to be bound.
 This unbound gas is dynamically fed by the convergent flow., This unbound gas is dynamically fed by the convergent flow.
 It contributes an accretion flow onto the CO complex. which has the necessary magnitude to drive the gas turbulence. re. to balance the dissipation of the gas turbulent kinetic energy.," It contributes an accretion flow onto the CO complex, which has the necessary magnitude to drive the gas turbulence, i.e. to balance the dissipation of the gas turbulent kinetic energy."
 ? have proposed a similar interpretation to account for the formation and subsequent growth of turbulent molecular clouds in the LMC., \citet{klessen10} have proposed a similar interpretation to account for the formation and subsequent growth of turbulent molecular clouds in the LMC.
 In Sect. ?2..," In Sect. \ref{sec:sfe},"
 we discuss the impact of this interpretation on the star formation efficiency within SGMC 2., we discuss the impact of this interpretation on the star formation efficiency within SGMC 2.
 In this section we extend our interpretation of the extended eemission to the compact ssource., In this section we extend our interpretation of the extended emission to the compact source.
 Like for the SGMC 2 complex we propose an interpretation where the luminosity traces the dissipation of the turbulent energy of the gas and gas accretion., Like for the SGMC 2 complex we propose an interpretation where the luminosity traces the dissipation of the turbulent energy of the gas and gas accretion.
 As a first test whether this 1s a plausible interpretation. we compare the dissipation rate of the turbulent energy of the gas with the lluminosityof the source.," As a first test whether this is a plausible interpretation, we compare the dissipation rate of the turbulent energy of the gas with the luminosityof the source."
 Combining equations (3)) and (4)). we obtain an equation that we can use to estimate the mass of," Combining equations \ref{eqD}) ) and \ref{eqA}) ), we obtain an equation that we can use to estimate the mass of"
obtained with the spots located on the less massive secondary.,obtained with the spots located on the less massive secondary.
 A model with spots on the primary fits the observations very well. but fits of similar quality can be obtained with different spot locations. and we could not decide definitely which of the possible solutions ts the most appropriate.," A model with spots on the primary fits the observations very well, but fits of similar quality can be obtained with different spot locations, and we could not decide definitely which of the possible solutions is the most appropriate."
 Having in mind the vicinity of the third visual component. which cannot be resolved spatially. we included a third light in the model.," Having in mind the vicinity of the third visual component, which cannot be resolved spatially, we included a third light in the model."
" The estimated third light contribution to the total light of the system (ναμετ)= 0.23) is significant. but somewhat lower than the contribution reported by Pribullaetal.(2009b) (Li/(L,+Ls)= 0.51. which equals L3/(Ly+LiLs)=0.33 for the luminosities of the eclipsing components in our model)."," The estimated third light contribution to the total light of the system ${\rm L_3/(L_h+L_c+L_3)}\approx 0.23$ ) is significant, but somewhat lower than the contribution reported by \citet{pribb09} $\rm L_3/(L_1+L_2)=0.51$ , which equals ${\rm L_3/(L_h+L_c+L_3)}\approx 0.33$ for the luminosities of the eclipsing components in our model)."
 There is also a considerable difference between our estimation of the distance of MR Del (63+ 9pc) and the one computed based on the new Hippareos parallax (49+ 6pe)., There is also a considerable difference between our estimation of the distance of MR Del $63 \pm 9 \rm pc$ ) and the one computed based on the new Hipparcos parallax $49 \pm 6 \rm pc$ ).
 The greater distance implies that the total luminosity of the system Is overestimated in our model. which may be a result of complications introduced by the third light.," The greater distance implies that the total luminosity of the system is overestimated in our model, which may be a result of complications introduced by the third light."
 We made trial runs with the Ls value fixed to the spectroscopic estimation of Pribullaetal.(2009b) and obtained plausible solutions. but the fits to the observations were significantly worse.," We made trial runs with the $\rm L_3$ value fixed to the spectroscopic estimation of \citet{pribb09} and obtained plausible solutions, but the fits to the observations were significantly worse."
" In comparison to the results that Cutispotoetal.(1997) obtained from the analysis of their UBV(RI) and soft X-ray light curves. our study gives somewhat lower masses of the components (Ate.cupo=0.64Me. Ath,cuis=0.72M; and Atethissway=0-69 Me. Λthissway=0.63 M. different radii (ReCuispow=9.70Ra. Rncup=0.77.R5 and anisμήν=0.83Re. Rothisstudy=0.05 As). and a slightly lower inclination Ceuispow*807 and Anisstudy= 797)."," In comparison to the results that \citet{cuti97} obtained from the analysis of their UBV(RI) and soft X-ray light curves, our study gives somewhat lower masses of the components ${\cal M}_{\rm c,\ Cutispoto}=0.64\ M_\odot$, ${\cal M}_{\rm h,\ Cutispoto}=0.72\ M_\odot$ and ${\cal M}_{\rm c,\ this\ study}=0.69\ M_\odot$ , ${\cal M}_{\rm h,\ this\ study}=0.63\ M_\odot$ ), different radii ${\cal R}_{\rm c,\ Cutispoto}=0.70\ R_\odot$, ${\cal R}_{\rm h,\ Cutispoto}=0.77\ R_\odot$ and ${\cal R}_{\rm c,\ this\ study}=0.83\ R_\odot$, ${\cal R}_{\rm h,\ this\ study}=0.65\ R_\odot$ ), and a slightly lower inclination $i_{\rm Cutispoto}\approx80^\circ$ and $i_{\rm this\ study}\approx79^\circ$ )."
 We note that their analysis was only based on photometric data., We note that their analysis was only based on photometric data.
 BD +07° 3142 (GSC 00380-00247. ASAS 162003407074. PPM 162491)) was found to be a variable star by Pojmanski(2002) during their analyses of the ASAS data.," BD $^{o}$ 3142 (GSC 00380-00247, ASAS 162003+0707.4, PPM 162491)) was found to be a variable star by \citet{poj02} during their analyses of the ASAS data."
 Since then. only one time of primary minimum has been published by Parimuchaetal.(2007).," Since then, only one time of primary minimum has been published by \citet{par07}."
. Rucinskietal.(2008) observed the star spectroscopically to determine its radial velocity variation., \citet{ruca08} observed the star spectroscopically to determine its radial velocity variation.
 They found a third component with a light contribution of ΕΕ + Ls) = 0.50., They found a third component with a light contribution of $_{3}$ $_{1}$ + $_{2}$ ) = 0.50.
 Because of the proximity of the radial velocity of the third component to that of the center of mass of the binary system. they argue that it appears to be physically attached to the eclipsing pair.," Because of the proximity of the radial velocity of the third component to that of the center of mass of the binary system, they argue that it appears to be physically attached to the eclipsing pair."
 From the analysis of broadening functions. they determined the mass ratio of 0.662 and spectral type as K2V. They classified the system as a W UMa type binary.," From the analysis of broadening functions, they determined the mass ratio of 0.662 and spectral type as K2V. They classified the system as a W UMa type binary."
 The results of our photometric analysis. based on spectroscopic elements of Rucinskietal.(2008). are given in Table 8..," The results of our photometric analysis, based on spectroscopic elements of \citet{ruca08}, are given in Table \ref{TabBD}."
 This ts the first ever published photometric solution of this system's light curves., This is the first ever published photometric solution of this system's light curves.
 Figure + shows the observed (LCO) and the synthetic (LCC) light curves in the B. V. and R filters (upper left). the B—V and V—B color indices (lower left). the ο—-C residuals (upper right) and the geometrical model of the system in representative phases 0.2 and 0.8 (lower right).," Figure \ref{fBD} shows the observed (LCO) and the synthetic (LCC) light curves in the B, V, and R filters (upper left), the $B-V$ and $V-B$ color indices (lower left), the $O-C$ residuals (upper right) and the geometrical model of the system in representative phases 0.2 and 0.8 (lower right)."
 As in the previous sections. the uncertainties given in Table 5. were estimated by combining the formal nonlinear least-squared fitting errors with the errors arising from the uncertainty of the spectroscopic mass ratio (¢=0.662+ 0.008). as described in Section 3..," As in the previous sections, the uncertainties given in Table \ref{TabQXAnd} were estimated by combining the formal nonlinear least-squared fitting errors with the errors arising from the uncertainty of the spectroscopic mass ratio $q=0.662 \pm 0.008$ ), as described in Section \ref{analysis}."
 According to our analysis. BD+773142 ts an overcontact binary (fas;= 10%) in which the hotter. less massive component is eclipsed in the deeper primary minimum. so the system belongs to the W subclass of W UMa type binaries.," According to our analysis, ${\rm BD+7^o3142}$ is an overcontact binary $f_{\rm over} \approx 10\% $ ) in which the hotter, less massive component is eclipsed in the deeper primary minimum, so the system belongs to the W subclass of W UMa type binaries."
 The light curve asymmetry can be explained by the presence of two cool spots. located near thepolar regions of each component.," The light curve asymmetry can be explained by the presence of two cool spots, located near thepolar regions of each component."
 Since the orbital period of the systemis short (P.= 028). its components are in fast rotation.," Since the orbital period of the systemis short $P\approx 0^d.28$ ), its components are in fast rotation."
 Owing to the tidal effects. we expect a synchronization of the star rotation with the," Owing to the tidal effects, we expect a synchronization of the star rotation with the"
manner includes as well some spatial correlations due to some lieht residual stripping aud thus does not correspond to white noise (at least iu the low f part).,manner includes as well some spatial correlations due to some light residual stripping and thus does not correspond to white noise (at least in the low $\ell$ part).
 For the aim of conrparison we compute the power spectrum of the iuput signal measured the same wav. Corr€ aud plotted both cep as well as Cm(ois| averaged in linear bauds of coustaut width Af=SO.," For the aim of comparison we compute the power spectrum of the input signal measured the same way, $C_{\ell}^{input}$, and plotted both $C_{\ell}^{input}$ as well as $C_{\ell}^{signal} -
C_{\ell}^{noise}$ averaged in linear bands of constant width $\Delta
\ell = 80$."
 The agreement is obviously very good as illustrate ou figure 9.., The agreement is obviously very good as illustrated on figure \ref{cl_mc}.
 The error bars take iuto account both the sampling iuduced variance as well as he beam spreading (I&nox1995)., The error bars take into account both the sampling induced variance as well as the beam spreading \cite{Kn95}.
. A few comments need o be made at this poiut., A few comments need to be made at this point.
 This is iu no way an appropriate Cy nmieasureiient since we are not dealing properly here with the sky coverage induced window function which rigecrs sole spurious correlations within the ys. We hus do not take iuto account the full-structure of the nap nolse covariance matrix., This is in no way an appropriate $C_{\ell}$ measurement since we are not dealing properly here with the sky coverage induced window function which triggers some spurious correlations within the $C_{\ell}$ s. We thus do not take into account the full-structure of the map noise covariance matrix.
 Nevertheless. the efficiency of such a rough method is eucouragiug for more detailed Huplementation and a full handling of the noise covariance natrix.," Nevertheless, the efficiency of such a rough method is encouraging for more detailed implementation and a full handling of the noise covariance matrix."
 Note finally that this is a naturally parallelised task which should therefore be feasible very quickly., Note finally that this is a naturally parallelised task which should therefore be feasible very quickly.
 The application to eenuine data could be problematic if our kev lapothesis were no to be fulfilled. thus we rave to discuss them., The application to genuine data could be problematic if our key hypothesis were not to be fulfilled thus we have to discuss them.
 Concerning the noise. we assumed hat it is Gaussian iu order to derive our map estimator and stationary in order to exploit the diagonuality of its iolse covariance matrix in Fourier space.," Concerning the noise, we assumed that it is Gaussian in order to derive our map estimator and stationary in order to exploit the diagonality of its noise covariance matrix in Fourier space."
 Both lypothesis are reasonable for a nominal mstrumneutal noise. at least ouwtiallv on (sufficiently lone) pieces of timeline.," Both hypothesis are reasonable for a nominal instrumental noise, at least partially on (sufficiently long) pieces of timeline."
 If not. we would have to cut the bad parts and replace them wea constrained realization of the noise in these parts.," If not, we would have to cut the bad parts and replace them by a constrained realization of the noise in these parts."
 Concerning the signal. no particular assumptions are receded in the method we are presenting.," Concerning the signal, no particular assumptions are needed in the method we are presenting."
 At this level. we reelected as wel the effect of the ton perfect sviunetiv of the instrumental beam.," At this level, we neglected as well the effect of the non perfect symmetry of the instrumental beam."
This effect should be quantified or a real experiment (Wietal.2000:Dorée2000).,"This effect should be quantified for a real experiment \cite{Wu00,Do00}."
. Another technical hypothesis is the neeligibility of the ακομαπο noise with respect to the instrunueutal noise mtosnee this is fully uuder coutrol of the user it should iot be a problem., Another technical hypothesis is the negligibility of the pixelisation noise with respect to the instrumental noise but since this is fully under control of the user it should not be a problem.
 Iu this paper. we have presented an origina fast uapanakiug iethod based on a multierid Jacobi algorithin.," In this paper, we have presented an original fast map-making method based on a multi-grid Jacobi algorithm."
 It naturally eutails the possibilitv of a joiut roisefsignal estimation., It naturally entails the possibility of a joint noise/signal estimation.
" We propose a wav to generate he noise covariance matrix aud illustrate its abilitv on a simple C, estimation.", We propose a way to generate the noise covariance matrix and illustrate its ability on a simple $C_{\ell}$ estimation.
 The efficiency. of this method las oen demoustrated in the case of two coming balloon experiments. namely ARCIIEOPS and Topllat but it raturally has a amore general rauge of application.," The efficiency of this method has been demonstrated in the case of two coming balloon experiments, namely ARCHEOPS and TopHat but it naturally has a more general range of application."
 This ool should be of natural interest for Planck and this will be demonstrated somewhere else., This tool should be of natural interest for Planck and this will be demonstrated somewhere else.
 Furthermore. due o the analogv of the formalisuas. this should have some applications as well iu the component separation problen.," Furthermore, due to the analogy of the formalisms, this should have some applications as well in the component separation problem."
 Wo lope to apply if very soon ou ARCIEOPS data., We hope to apply it very soon on ARCHEOPS data.
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though that depending on the shape of its SED this additional source will affect measurements of J1148—-5251] done with large (210) beams.,though that depending on the shape of its SED this additional source will affect measurements of J1148+5251 done with large $\gtrsim$ ) beams.
 Aperture photometry of J114845251 yields a flux density at 100 and 160m of 4.1+1.0 mmJy and 6.3€2.0 mmly. respectively (see reftablel for α summary of all flux measurements. including literature data).," Aperture photometry of J1148+5251 yields a flux density at 100 and $\mu$ m of $4.1 \pm 1.0$ mJy and $6.3 \pm 2.0$ mJy, respectively (see \\ref{table1} for a summary of all flux measurements, including literature data)."
 In the red band we removed the confusingσι source via PSF fitting prior to performing the photometry., In the red band we removed the confusing source via PSF fitting prior to performing the photometry.
 BR 1202-0725 is a well studied quasar at ς=4.69 with strong detections at sub-mm and mm wavelengths and in CO line emission2002)., BR $-$ 0725 is a well studied quasar at $z=4.69$ with strong detections at sub-mm and mm wavelengths and in CO line emission.
. At ~4” distance to the northwest of the quasar this source shows à secondary component in the dust continuum and in CO., At $\sim$ distance to the northwest of the quasar this source shows a secondary component in the dust continuum and in CO.
 A Lye extension is observed approximately 2733 northwest of the quasar2004)., A $\alpha$ extension is observed approximately 3 northwest of the quasar.
. We detected BR 1202-0725 at um and 160 m with fluxes of 15.0+2.0 mmJy and 39.8+3.7 mmJy. respectively (Fig. 2)).," We detected BR $-$ 0725 at $\mu$ m and $\mu$ m with fluxes of $15.0\pm2.0$ mJy and $39.8\pm3.7$ mJy, respectively (Fig. \ref{pic_br1202}) ),"
 which is consistent with the upper limits from ISO2001)., which is consistent with the upper limits from ISO.
. In theirSpitzer observations at jum. see slightly resolved emission and perform a two-component PSF fit to isolate the emission attributable to the QSO.," In their observations at $\mu$ m, see slightly resolved emission and perform a two-component PSF fit to isolate the emission attributable to the QSO."
 Although our PACS 70m data have slightly higher resolution than the MIPS jm observations. we donot detect à secondary component securely. probably due to the lower S/N in the jim maps compared to the jim image.," Although our PACS $\mu$ m data have slightly higher resolution than the MIPS $\mu$ m observations, we donot detect a secondary component securely, probably due to the lower S/N in the $\mu$ m maps compared to the $\mu$ m image."
" Although an apparent elongation may be visible for the red band in refpic,r 1202. atwo-component fitdoesnotgiveasignificantlybetter fitto."," Although an apparent elongation may be visible for the red band in \\ref{pic_br1202}, a two-component fit does not give a significantly better fit to the source profile."
 ," As we cannot reliably separate the two sources in the PACS observations, the quoted flux values refer to the sum of both components."
In Fig., In Fig.
 3 we show the spectral flux and energy distributions of J1148+5251 and 1202-0725., \ref{sed_j1148} we show the spectral flux and energy distributions of J1148+5251 and $-$ 0725.
 The SEDs of both high-redshift QSOs appear very similar in shape. and the strong emission in all infrared bands shows dust at a wide range of temperatures.," The SEDs of both high-redshift QSOs appear very similar in shape, and the strong emission in all infrared bands shows dust at a wide range of temperatures."
 This is particularly emphasised by the photometry. which fills the gap between the previously available MIR and sub-mm photometry.," This is particularly emphasised by the photometry, which fills the gap between the previously available MIR and sub-mm photometry."
 The shape of the SED also implies that large amounts of dust are present in the host galaxies already at high redshift and that the dust may be distributed on a wide range of scales: NIR emission from hot dust close to the nucleus2010).. warm MIR dust ~ ffew hundred K) on intermediate scales (or partly shielded). and colder dust in the FIR possibly distributed throughout the host galaxy.," The shape of the SED also implies that large amounts of dust are present in the host galaxies already at high redshift and that the dust may be distributed on a wide range of scales: NIR emission from hot dust close to the nucleus, warm MIR dust $\sim$ few hundred K) on intermediate scales (or partly shielded), and colder dust in the FIR possibly distributed throughout the host galaxy."
 Using the new Herschiel photometry to determine total infrarec luminosities for these, Using the new photometry to determine total infrared luminosities for these
The average mass of white dwarfs is approximately 0.6M... with the bulk of these stars having à mass near this average value.,"The average mass of white dwarfs is approximately $0.6\ M_\odot$, with the bulk of these stars having a mass near this average value."
 Additional smaller peaks at the ultramassive and low-mass ends of the mass distribution are observed., Additional smaller peaks at the ultramassive and low-mass ends of the mass distribution are observed.
 In the case of low-mass white dwarfs (M.«0.4 Μ.). the galaxy is not old enough to have produced such objects.," In the case of low-mass white dwarfs $M < 0.4\ M_\odot$ ), the galaxy is not old enough to have produced such objects."
 Several low-mass white dwarfs (M.<0.45 M.) have been checked for radial velocity variations (e.g..Maxtedetal..2000:Nelemansetal.. 2005): however. not all low-mass white dwarfs were found to be variable.," Several low-mass white dwarfs $M < 0.45\ M_\odot$ ) have been checked for radial velocity variations \citep[e.g.,][]{max2000,nel2005}; however, not all low-mass white dwarfs were found to be variable."
 Α statistical analysis of the incidence of binarity among low-mass white dwarfs was conducted by Naprwotzkietal.(2007) using the ESO SN Ia progenitor survey. which included 26 white dwarfs with masses lower than 0.42M...," A statistical analysis of the incidence of binarity among low-mass white dwarfs was conducted by \citet{nap2007} using the ESO SN Ia progenitor survey, which included 26 white dwarfs with masses lower than $0.42\ M_\odot$."
 From this sample they find only 11 with variable radial velocities showing that single low-mass white dwarfs do exist., From this sample they find only 11 with variable radial velocities showing that single low-mass white dwarfs do exist.
 It may be that some of these low-mass white dwarfs have indeed formed through single-star evolution., It may be that some of these low-mass white dwarfs have indeed formed through single-star evolution.
 Kilicetal.(2007c) argues that up to of low-mass helium core white dwarfs M—0.4M. could have formed from progenitors with supersolar metallicities., \citet{kil2007c} argues that up to of low-mass helium core white dwarfs $M \sim 0.4\ M_\odot$ could have formed from progenitors with supersolar metallicities.
 Although some low-mass white dwarfs may evolve through single-star evolution. no single-star evolutionary path is known for extremely low-mass (ELM) white dwarfs (M.«0.2 M..). which were initially identified as companions to pulsars (e.g..vanKerkwiyketal..1996:Bassa 2006).," Although some low-mass white dwarfs may evolve through single-star evolution, no single-star evolutionary path is known for extremely low-mass (ELM) white dwarfs $M < 0.2\ M_\odot$ ), which were initially identified as companions to pulsars \citep[e.g.,][]{van1996,bas2006}."
. Subsequently. several low-mass white dwarfs have been discovered in the Sloan Digitial Sky Survey (Kleinmanetal.2004;Etsen-steinetal.. 2006).. with SDSS J123410.37-022802.9 having the lowest mass at ~0.18M. (Liebertetal..2004).," Subsequently, several low-mass white dwarfs have been discovered in the Sloan Digitial Sky Survey \citep{kle2004,eis2006}, with SDSS $-$ 022802.9 having the lowest mass at $\sim 0.18\ M_\odot$ \citep{lie2004}."
.. Other spectroscopic identifications include the high-velocity white dwarf LP400-22. with a mass of 0.17+0.01M& (Kawkaetal. 2006). and SDSS JO917+46. which was discovered as part of a search for hypervelocity B-type stars (Kilicetal.. 2007a).," Other spectroscopic identifications include the high-velocity white dwarf LP400-22, with a mass of $0.17\pm0.01\ M_\odot$ \citep{kaw2006}, and SDSS $+$ 46, which was discovered as part of a search for hypervelocity B-type stars \citep{kil2007a}."
. Follow-up observations of the last two ofthese stars using the 6.5m MMT telescope found them to have variable radial velocities., Follow-up observations of the last two ofthese stars using the 6.5m MMT telescope found them to have variable radial velocities.
 SDSS J0917--46 was found to be in a binary with an orbital period of 7.5936+0.0024 hours and LP 400—22 was found to have an orbital period of approximately | day (Kilicetal..2009:Vennes2009).," SDSS $+$ 46 was found to be in a binary with an orbital period of $7.5936\pm0.0024$ hours and LP $-$ 22 was found to have an orbital period of approximately 1 day \citep{kil2009,ven2009}."
. In both cases. the companion remains unknown. but a white dwarf. neutron star. or a low-mass main-sequence star were considered as possible companions.," In both cases, the companion remains unknown, but a white dwarf, neutron star, or a low-mass main-sequence star were considered as possible companions."
 The recent discovery of ELM white dwarfs offers new challenges for binary evolution scenarios. as well as for calculating the internal structure of white dwarfs including the effect of residual burning and model atmospheres in a poorly explored range of stellar parameters.," The recent discovery of ELM white dwarfs offers new challenges for binary evolution scenarios, as well as for calculating the internal structure of white dwarfs including the effect of residual burning and model atmospheres in a poorly explored range of stellar parameters."
 We present our discovery and analysis of the ELM white dwarf NLTT 11748., We present our discovery and analysis of the ELM white dwarf NLTT 11748.
 Our spectroscopic and photometric observations are presented in. 32., Our spectroscopic and photometric observations are presented in 2.
 We present our model atmosphere analysis of the object in $3 and discuss our results in $4., We present our model atmosphere analysis of the object in 3 and discuss our results in 4.
 We summarise in $5., We summarise in 5.
 We obtained low-resolution spectroscopy of NLTT 11748 using the EFOSC2 (ESO Faint Object Spectrograph and Camera) attached to the New Technology Telescope (NTT) at La Silla Observatory., We obtained low-resolution spectroscopy of NLTT 11748 using the EFOSC2 (ESO Faint Object Spectrograph and Camera) attached to the New Technology Telescope (NTT) at La Silla Observatory.
 The 300 lines per mm grating with a blaze wavelength of 4000 ((Grism 11) was used to obtain two spectra with a resolution of 16A., The 300 lines per mm grating with a blaze wavelength of 4000 (Grism 11) was used to obtain two spectra with a resolution of 16.
. The exposure time for each spectrum was 20 minutes., The exposure time for each spectrum was 20 minutes.
 The slitwidth was set to | arcsecond., The slitwidth was set to 1 arcsecond.
 We acquired both spectra on UT 2008 October 21 in photometric conditions with the seeing varying between | and 1.3 areseconds throughout the night., We acquired both spectra on UT 2008 October 21 in photometric conditions with the seeing varying between 1 and 1.3 arcseconds throughout the night.
 A wavelength calibration spectrum of helium and argon immediately followed the NLTT11748 spectra., A wavelength calibration spectrum of helium and argon immediately followed the NLTT11748 spectra.
 The spectra were flux-calibrated using the flux standard Feige 110. which was obtained on the same night.," The spectra were flux-calibrated using the flux standard Feige 110, which was obtained on the same night."
 We then obtained ultraviolet photometry from the Galaxy Evolution Explorer (GALEX) all sky survey., We then obtained ultraviolet photometry from the Galaxy Evolution Explorer ) all sky survey.
GALEX provides photometry in two bands. FUV and NUV.," provides photometry in two bands, FUV and NUV."
 The FUV bandpass is 1344—1786 wwith an effective wavelength of 1528Á.. and the NUV bandpass is 1771—2831 wwith an effective wavelength of 2271 A.," The FUV bandpass is $1344 - 1786$ with an effective wavelength of 1528, and the NUV bandpass is $1771 - 2831$ with an effective wavelength of 2271 ."
. We also obtained infrared photometry from the Two Micron All Sky Survey (2MASS). which is available at VizieR at the Centre de," We also obtained infrared photometry from the Two Micron All Sky Survey ), which is available at VizieR at the Centre de"
"0.3-5 keV (?),, and also lies very close to the line f,cxv??.","0.3-5 keV \citep{rea09}, and also lies very close to the line $f_{\nu}\propto \nu^{-0.9}$."
" For comparison, the radio-to-optical slope of thepulsed radiation from the Crab has a much shallower slope of ~—0.2 (measured from Fig."," For comparison, the radio-to-optical slope of the radiation from the Crab has a much shallower slope of $\sim -0.2$ (measured from Fig."
 9.3 of ?))., 9.3 of \citealt{lyne05b}) ).
" For completeness, we have also calculated the 7 magnitude limit of J1819—1458 from the sum of all of the WHT observations."," For completeness, we have also calculated the $i'$ magnitude limit of $-$ 1458 from the sum of all of the WHT observations."
" We derive a magnitude limit of i’>21.9 at 5c confidence, significantly deeper than the limit on the persistent light of J=17.5 derived by ?.."," We derive a magnitude limit of $i'>21.9$ at $\sigma$ confidence, significantly deeper than the limit on the persistent light of $I=17.5$ derived by \citet{reynolds06}."
" Of course, our limit could have been substantially deeper had we decided not to look for bursts and instead taken just a few long exposures."," Of course, our limit could have been substantially deeper had we decided not to look for bursts and instead taken just a few long exposures."
" In paper I, we used aperture photometry centred on the X-ray position of J1819—1458 to search for the RRAT."," In paper I, we used aperture photometry centred on the X-ray position of $-$ 1458 to search for the RRAT."
" With no way of knowing in which frames the bursts occurred, this search strategy was relatively insensitive, as we were relying on detecting individual optical bursts above the noise."," With no way of knowing in which frames the bursts occurred, this search strategy was relatively insensitive, as we were relying on detecting individual optical bursts above the noise."
" In this paper, we present a far more sensitive technique: simultaneous radio observations of the RRAT bursts to tag the corresponding optical frames."," In this paper, we present a far more sensitive technique: simultaneous radio observations of the RRAT bursts to tag the corresponding optical frames."
 These tagged frames were then co-added to produce a summed ‘burst’ frame., These tagged frames were then co-added to produce a summed `burst' frame.
of the same order of magnitude as the expression in Eq. (6)).,of the same order of magnitude as the expression in Eq. \ref{eq:angle_var1}) ).
 This approximation is also computed for some telescopes and reported in Table 1.., This approximation is also computed for some telescopes and reported in Table \ref{tab:DC_toler}.
" For an on-axis source at a finite distance, the area of the circular corona between Ry and Ro—L4ó should be considered: the area is thereby increasedby the same term ~2zRoL46, for both Wolter and double cone."," For an on-axis source at a finite distance, the area of the circular corona between $R_M$ and $R_0- L_1 \delta$ should be considered: the area is thereby increasedby the same term $\sim 2\pi R_0L_1\delta$, for both Wolter and double cone."
" Hence, the approximation of the area is even better than the expression in Eq. (8))."," Hence, the approximation of the area is even better than the expression in Eq. \ref{eq:max_diam}) )."
 We now consider the third problem., We now consider the third problem.
" The case of a source on-axis, at a finite distance, represents the simplest case in which there is geometrical vignetting for double reflection."," The case of a source on-axis, at a finite distance, represents the simplest case in which there is geometrical vignetting for double reflection."
" If the source is on-axis, at infinite distance (with Lo> L4), all rays undergo two reflections and reach the focus of the mirror."," If the source is on-axis, at infinite distance (with $L_2 \geq L_1$ ), all rays undergo two reflections and reach the focus of the mirror."
" If the distance is finite, not all photons reflected by the primary segment also hit the secondary: for instance, photons reflected atz&+L, can miss being reflected by the secondary, so they are prevented from being focused, as depicted in Fig. 1.."," If the distance is finite, not all photons reflected by the primary segment also hit the secondary: for instance, photons reflected at $z \approx+L_1$ can miss being reflected by the secondary, so they are prevented from being focused, as depicted in Fig. \ref{fig:mirror_section}."
" We can quantify this geometric effect by means of the V, which is defined to be the ratio of the number of photons reflected twice to the number of photons reflected by the primary segment (assuming, ideally, that the mirror reflectivity is 1)."," We can quantify this geometric effect by means of the $V$, which is defined to be the ratio of the number of photons reflected twice to the number of photons reflected by the primary segment (assuming, ideally, that the mirror reflectivity is 1)."
" In this configuration, if the mirror profile is a double cone, we can compute V using the simple formula (see Appendix AppendixA: for a derivation) valid if Loo?«μαι. otherwise V=1."," In this configuration, if the mirror profile is a double cone, we can compute $V$ using the simple formula (see Appendix \ref{App_vign} for a derivation) valid if $L_2\alpha_2 < L_1\alpha_1$, otherwise $V =1$."
" If eithera;<0 ora»<O, then V=0."," If either $\alpha_1~<~0$ or $\alpha_2~<~0$, then $V =0$."
" In Sect. 3,,"," In Sect. \ref{EffArea},"
 we see that this result can even be generalized to the case of a source off-axis., we see that this result can even be generalized to the case of a source off-axis.
 If we consider that in this caseo; =a@ot+dandaz =αρ--δ. an alternative expression for Eq. (9))," If we consider that in this case $\alpha_1~=~\alpha_0+\delta$ and $\alpha_2~=~\alpha_0-\delta$, an alternative expression for Eq. \ref{eq:vign}) )"
 is where we have used the definitions of oy and 6., is where we have used the definitions of $\alpha_0$ and $\delta$.
 Using Eq. (3)), Using Eq. \ref{eq:f_num}) )
" and normalizing the L’s to the mirror diameter, we can also write Eq. (9))"," and normalizing the $L$ 's to the mirror diameter, we can also write Eq. \ref{eq:vign}) )"
 as It is interesting to note that a formula similar toEq. (10)), as It is interesting to note that a formula similar toEq. \ref{eq:vign2}) )
" for the effective area was empirically found by Van Speybroeck and Chase (1972)), but with a f term instead of f."," for the effective area was empirically found by Van Speybroeck and Chase \cite{VanSpeybroeck}) ), but with a $f$ term instead of $f$."
" Rigorously, Eqs. (9))"," Rigorously, Eqs. \ref{eq:vign}) )"
" to (11)) hold only for a double cone mirror, but they are expected to be applicable also to Wolter-I mirrors of a sufficiently large f#."," to \ref{eq:vign_fnum}) ) hold only for a double cone mirror, but they are expected to be applicable also to Wolter-I mirrors of a sufficiently large $f\#$."
" To determine the approximation that can be achieved, we computed by tracing theexact V factor of a Wolter-Imirror as a function of fft, in the case of Li=L5 0.5, for fixed 6 values."," To determine the approximation that can be achieved, we computed by ray-tracing the $V$ factor of a Wolter-Imirror as a function of $f\#$, in the case of $L'_1 = L'_2 = 0.5$ , for fixed $\delta$ values."
" In Fig. 3,,"," In Fig. \ref{fig:fin_dist}, ,"
 the results are compared with the findings of Eq. (11))., the results are compared with the findings of Eq. \ref{eq:vign_fnum}) ).
" The vignetting for the double cone is independent of L’, whereas forthe Wolter's it depends markedly on L’ only for small fs."," The vignetting for the double cone is independent of $L'$ , whereas forthe Wolter's it depends markedly on $L'$ only for small $f\#$ ."
 Inspection of Fig., Inspection of Fig.
 3 shows that, \ref{fig:fin_dist} shows that
factor. go=Da. where D is the standard linear growth factor for the matter density field in QCDALmodels. aud go.=dole1) denotes its value at the prescut day with scale factor a=1.,"factor, $g_Q=D/a$, where $D$ is the standard linear growth factor for the matter density field in QCDMmodels, and $g_{Q,0}=g_Q(a=1)$ denotes its value at the present day with scale factor $a=1$."
 We discuss cach piece of equation (1)) in turn., We discuss each piece of equation \ref{pl}) ) in turn.
 First we exanune the transfer function for the matter density feld., First we examine the transfer function for the matter density field.
 To isolate the effects of quintesscuce. we fiud it convenient and ilhuninating to compare a pair of QUDM ixl ACDM models that have the same set of cosmological parameters and differ only iu uy (recall a=l] for ACDMJ).," To isolate the effects of quintessence, we find it convenient and illuminating to compare a pair of QCDM and $\LCDM$ models that have the same set of cosmological parameters and differ only in $w_Q$ (recall $w=-1$ for $\LCDM$ )."
 We define the relative transfer function for a pair of such models to be Tox=To/Ta., We define the relative transfer function for a pair of such models to be $T_{Q\Lambda} = T_Q/T_\Lambda$.
 For T4. we follow the convention aud set the arbitrary amplitude of T4 to unity as &>O.," For $T_\Lambda$ , we follow the convention and set the arbitrary amplitude of $T_\Lambda$ to unity as $k \rightarrow 0$."
 The form of T4 is well πιο aud various fitting formulas have been published (c.e.. Bardeen ct al.," The form of $T_\Lambda$ is well known and various fitting formulas have been published (e.g., Bardeen et al."
 1986: Efstathiou. Boud. White 1992: Sugivama 1995).," 1986; Efstathiou, Bond, White 1992; Sugiyama 1995)."
" More complicated fits with higher accuracy have also beeu developed for higher barvou ratios (Q,,/0,,2 20%) and for the features due to barvouic oscillations aud damping (oe. Bunn White 1997: Eiseusteiu IIn 1998)."," More complicated fits with higher accuracy have also been developed for higher baryon ratios $\ob/\om\go 20\%$ ) and for the features due to baryonic oscillations and damping (e.g., Bunn White 1997; Eisenstein Hu 1998)."
 The trauster fiction Ty) for QCDAM models resenibles T4 for ACDAI but with one kev difference., The transfer function $T_Q$ for QCDM models resembles $T_\Lambda$ for $\LCDM$ but with one key difference.
" The lincar matter density field. 6,,. evolves according to the equation ὃν|Hà,=LGρωδνἄρω30po). where IE=afa and the dots denote differentiation with respect to proper time."," The linear matter density field, $\delta_m$, evolves according to the equation $\ddot\delta_m + 2
H\dot\delta_m = 4\pi G (\rho_m\delta_m +\delta\rho_Q + 3\,\delta
p_Q)$ , where $H=\dot a/a$ and the dots denote differentiation with respect to proper time."
" On small leneth scales. the Q-field is zuooth (6p.dpeOpn) and we recover the familiar equation for the evolution of à,, (Caldwell et al."," On small length scales, the Q-field is smooth $\delta\rho_Q,\,\delta p_Q \ll \delta\rho_m$ ) and we recover the familiar equation for the evolution of $\delta_m$ (Caldwell et al."
 1998)., 1998).
 On very huge length scales. however. the Q-field clusters and contributes to the oenergv deusifv and pressure perturbations.," On very large length scales, however, the Q-field clusters and contributes to the energy density and pressure perturbations."
" The result is a different growth rate for à,, on large and small scales once the quintessence starts to dominate the cosimological energv density.", The result is a different growth rate for $\delta_m$ on large and small scales once the quintessence starts to dominate the cosmological energy density.
 We cau determine the characteristic scale separating these two reehmes by examining the linear equation for the Q-field: à)|31160LO=Om[CLwoo!2 (where Vds the Q-field poteutialYoo)90 Vo.=dV/dQ?. aud po=Q3/2|Vj.," We can determine the characteristic scale separating these two regimes by examining the linear equation for the $Q$ -field: $\ddot{\delta Q} + 3 H\dot{\delta Q} + (k^2 + V_{,QQ})\delta Q
=\dot\delta_m [(1 + w_Q)\rho_Q]^{1/2}$ (where $V$ is the Q-field potential, $V_{,QQ}\equiv d^2V/dQ^2$, and $\rho_Q=\dot{Q}^2/2+V$ )."
" We see that 6Q itself beliaves as à scalar field with an effective. mass (Vo)!− and a Comptou waveunuber of ~(Vo)!=,"," We see that $\delta Q$ itself behaves as a scalar field with an effective mass $(V_{,QQ})^{1/2}$ and a Compton wavenumber of $k_Q\sim
(V_{,QQ})^{1/2}$."
 On small leneth scales (Lhe komho). hethe amplitude of 6Q aud hence dpe is damped aud docs not cuter the evolution equation for à.," On small length scales (i.e. $k\gg k_Q$ ), the amplitude of $\delta Q$ and hence $\delta\rho_Q$ is damped and does not enter the evolution equation for $\delta_m$ ."
" On large scales (&<<hey) 0 grows. so the Q-field clusters and in turn affects the evolution of 6,,."," On large scales $k\ll k_Q$ )$\delta Q$ grows, so the Q-field clusters and in turn affects the evolution of $\delta_m$."
" The change in the beliavior of 6, neu k~1? as a result of differing Q-clustering properties is(Wee illustrated in Figure 1 for Των vs. & for a vauge of we and Όμι"," The change in the behavior of $\delta_m$ near $k \sim (V_{,QQ})^{1/2}$ as a result of differing $Q$ -clustering properties is illustrated in Figure 1 for $T_{Q\Lambda}$ vs. $k$ for a range of $w_Q$ and $\om$."
 We have chosen to normalize To to unity at the high-end because both the Q-field and the cosmological constant are spatially smooth on these scales., We have chosen to normalize $T_{Q\Lambda}$ to unity at the $k$ end because both the Q-field and the cosmological constant are spatially smooth on these scales.
 The clustering property of Q is reminiscent of the case of massive neutrinos in C|IIDM models. which cannot cluster appreciably below he neutrino free-streamine scale but cau cluster with the same amplitude as the cold dark matter on large scales.," The clustering property of Q is reminiscent of the case of massive neutrinos in C+HDM models, which cannot cluster appreciably below the neutrino free-streaming scale but can cluster with the same amplitude as the cold dark matter on large scales."
" Analogous to the shape parameter D, that was introduced o model the neutrino streaming distances in C|IIDM uodels (Ala 1996). we introduce a new shape parameter hereto characterize the feature in Figure 1 in QCDM Γωnodels."," Analogous to the shape parameter $\Gamma_\nu$ that was introduced to model the neutrino streaming distances in C+HDM models (Ma 1996), we introduce a new shape parameter $\Gamma_Q$ hereto characterize the feature in Figure 1 in QCDM models."
 For a coustaut equation of state. we. oue can slow hat Vo29πα|pwop). where p aud p are the total energy density we)(2p aud pressure.," For a constant equation of state, $w_Q$, one can show that $V_{,QQ} = 6\pi G
(1 - w_Q) (2\rho + p + w_{Q}\rho)\,$, where $\rho$ and $p$ are the total energy density and pressure."
" We approximate and we use a sniple ratio of polvnomials to express the relative transfer function: where ko is d Mpe !2 and à is a scale-indepeudent but time-dependent coefücieutthat quantifies the relative amplitude of the matter densityfield à,, ou large aud small leugth scales."," We approximate and we use a simple ratio of polynomials to express the relative transfer function: where $k$ is in $^{-1}$ , and $\alpha$ is a scale-independent but time-dependent coefficientthat quantifies the relative amplitude of the matter densityfield $\delta_m$ on large and small length scales."
" We find a well approximated by where the matter density paramcter isGOq(u)= 772]. which reaches. the value Q,, at the present dav a= 1."," We find $\alpha$ well approximated by where the matter density parameter is$\om(a) =
\om/[\om + (1 - \om)\, a^{-3 w_Q}]\,$ , which reaches the value $\om$ at the present day $a=1$ ."
 Figure 1d illustrates the close agreecmeut (with errors < LO%) between the, Figure 1 illustrates the close agreement (with errors $\lesssim 10\%$ ) between the
 , 
When averaging over the relative velocity of interacting particles (eq. 8)).,"When averaging over the relative velocity of interacting particles (eq. \ref{integral}) ),"
" the slowly varying logarithmic factor can be considered constant. and the term rc(v,) in the logarithm’s argument can be replaced by = where ον)=347yo/u,e 18 the average of the square of the relative velocity (Spitzer 1956)."," the slowly varying logarithmic factor can be considered constant, and the term $r_{\rm C}(v_{ss^\prime})$ in the logarithm's argument can be replaced by )= where $\langle v_{ss^\prime}^2\rangle= 3kT_{ss^\prime}/\mu_{ss^\prime}$ is the average of the square of the relative velocity (Spitzer 1956)."
 The averaging of the remaining dependence of the cross section on the inverse fourth power of the relative velocity then gives. according to eq. (15)). IGE).," The averaging of the remaining dependence of the cross section on the inverse fourth power of the relative velocity then gives, according to eq. \ref{G_n}) ), ) ] ),"
 where ανν 18 given by eq. (11)), where $a_{ss^\prime}$ is given by eq. \ref{defa}) )
 and G4(é) by eq. (17))., and $G_4(\xi)$ by eq. \ref{g4}) ).
 If the particles are assumed to be pointlike (ri;= 0). and. in addition. tp>>7c (a condition generally fulfilled 1n the ISM). we have ," If the particles are assumed to be pointlike $r_{\rm
min}=0$ ), and, in addition, $\lambda_{\rm D}\gg \overline{r_{\rm C}}$ (a condition generally fulfilled in the ISM), we have ] )."
In Vthis case lefil(28)8]eq. (30)), In this case eq. \ref{sv_coul}) )
 for ΙΡ).€=0 gives the standard expression of the Coulomb momentum transfer rate (Spitzer 1956) where is thelogarithm., for $\xi=0$ gives the standard expression of the Coulomb momentum transfer rate (Spitzer 1956) where is the.
 In astrophysical applications. eq. (32))," In astrophysical applications, eq. \ref{spitz}) )"
 has often been adopted to compute the rate of momentum transfer between grain-ion and grain-eleetron collisions (Draine Salpeter 1979)., has often been adopted to compute the rate of momentum transfer between grain-ion and grain-electron collisions (Draine Salpeter 1979).
 However. a more accurate calculation requires a modification of the Coulomb logarithm to account for the finite size of the grains.," However, a more accurate calculation requires a modification of the Coulomb logarithm to account for the finite size of the grains."
" Taking Min=r, 1n Eq. (30))."," Taking $r_{\rm min}=r_g$ in Eq. \ref{sv_coul}) ),"
" and assuming as before Ap>rc. we obtain the modified Coulomb logarithm for ion Cor grain-electron) collisions which reduces to InClp/re) for ος0. and approaches InGlp/r.) for r,csre."," and assuming as before $\lambda_{\rm D} \gg
\overline{r_{\rm C}}$, we obtain the modified Coulomb logarithm for grain-ion (or grain-electron) collisions ], which reduces to $\ln(\lambda_{\rm D}/\overline{r_{\rm C}})$ for $r_g=0$, and approaches $\ln(\lambda_{\rm D}/r_g)$ for $r_g\gg \overline{r_{\rm
C}}$."
 The latter value of the Coulomb logarithm has been used. e.g.. by Benkadda et al. (," The latter value of the Coulomb logarithm has been used, e.g., by Benkadda et al. ("
1996) in their study of nonlinear instabilities in dusty plasmas.,1996) in their study of nonlinear instabilities in dusty plasmas.
" Following Draine Sutin (1987). 1t is easy to compute the ratio r,./rc for a variety of astrophysical conditions. since this quantity depends almost exclusively on the reduced temperature τιΕΚΤ er. Using the analytical approximations for the value of the average grain charge Z,e as a function of τι computed by Drame Sutin (1987). we obtain the modified Coulomb logarithm for grain-ion collisions given in Fig."," Following Draine Sutin (1987), it is easy to compute the ratio $r_g/\overline{r_{\rm C}}$ for a variety of astrophysical conditions, since this quantity depends almost exclusively on the reduced temperature $\tau_i\equiv r_g kT_i/(Z_ie)^2$ , Using the analytical approximations for the value of the average grain charge $Z_ge$ as a function of $\tau_i$ computed by Draine Sutin (1987), we obtain the modified Coulomb logarithm for grain-ion collisions given in Fig."
 21. for “light” and “heavy ions” (effective atomic weight | and 25. respectively).," \ref{coul} for “light” and “heavy ions” (effective atomic weight 1 and 25, respectively)."
 We see from this figure that the pointlike approximation for the Coulomb logarithm remains valid for 7; less than ~0.1., We see from this figure that the pointlike approximation for the Coulomb logarithm remains valid for $\tau_i$ less than $\sim 0.1$.
 For larger values of r; the value of the Coulomb logarithm 15 smaller than the pointlike value. but the effect is a reduction of the constant in Eq. (34)).," For larger values of $\tau_i$ the value of the Coulomb logarithm is smaller than the pointlike value, but the effect is a reduction of the constant in Eq. \ref{coul_log}) ),"
" from 9.42 to 8.97 or 9.18 at worst. for the ""light ion” and “heavy ton” cases. respectively."," from 9.42 to 8.97 or 9.18 at worst, for the “light ion” and “heavy ion” cases, respectively."
" For a more general caleulation of the Coulomb logarithm for grain-ion and grain-eleetron where rjj, depends on the finite size of the grain and the relative velocity of the collision. and. in addition. the approximation tpzFE Is removed. see Khrapak Morfill (2004)."," For a more general calculation of the Coulomb logarithm for grain-ion and grain-electron where $r_{\rm min}$ depends on the finite size of the grain and the relative velocity of the collision, and, in addition, the approximation $\lambda_{\rm D} \gg \overline{r_{\rm C}}$ is removed, see Khrapak Morfill (2004)."
 In Table | we list analytical fitting formulae for the momentum transfer rate coefficients computed numerically in Sects., In Table \ref{rates_vd0} we list analytical fitting formulae for the momentum transfer rate coefficients computed numerically in Sects.
" 3-5,", 3–5.
 For zero drift velocity. we have approximated the rates (or their logarithm) with third-order polynomials of the logarithm of the," For zero drift velocity, we have approximated the rates (or their logarithm) with third-order polynomials of the logarithm of the"
we select. several subsamples (candidates for a void »opulation). splitting according to a varicty of possibly relevant properties: Luminosity. colour. morphology.,"we select several subsamples (candidates for a void population), splitting according to a variety of possibly relevant properties: luminosity, colour, morphology."
 Also. we extract sets of haloes. binned by total mass.," Also, we extract sets of haloes, binned by total mass."
 We study the vpical densities (estimated on a 5 scale) in which hese various populations reside. with the goal of finding a irst signature of a void. population.," We study the typical densities (estimated on a 5 scale) in which these various populations reside, with the goal of finding a first signature of a void population."
 We recover the well-known pattern of halo bias., We recover the well-known pattern of halo bias.
 Blue. star-forming galaxies tend ο reside in underdense regions. a trend which is also present out is not so marked for bulge-Iess galaxies.," Blue, star-forming galaxies tend to reside in underdense regions, a trend which is also present but is not so marked for bulge-less galaxies."
 Next we estimate the environment density dependence ol the galaxy luminosity function. focussing on the variation of the shape and the normalization of the LE with the local mass density.," Next we estimate the environment density dependence of the galaxy luminosity function, focussing on the variation of the shape and the normalization of the LF with the local mass density."
 We compute Les in a series of density bins covering equal total volumes. and then in bins containing equal total mass.," We compute LFs in a series of density bins covering equal total volumes, and then in bins containing equal total mass."
 While the overall normalization. of the LE is quite different: between equal-volume bins. it is very similar between equal-mass bins.," While the overall normalization of the LF is quite different between equal-volume bins, it is very similar between equal-mass bins."
 We find some tendency for the faint-enel slope of the LE to vary with DM environment density., We find some tendency for the faint-end slope of the LF to vary with DM environment density.
" Finally. we consider ""multi-point probes like correlation functions anc nearest neighbour statistics."," Finally, we consider “multi-point” probes like correlation functions and nearest neighbour statistics."
 The correlation functions of our samples provide a check of the selection procedure., The correlation functions of our samples provide a check of the selection procedure.
 We recover the usual behaviour: in particular. red. galaxies are much more clustered. than blue ones. in agreement with WOOL," We recover the usual behaviour: in particular, red galaxies are much more clustered than blue ones, in agreement with K99."
 We stress. however. that it ds dillicult to deduce the relative spatial distribution of populations on the basis of their correlation functions.," We stress, however, that it is difficult to deduce the relative spatial distribution of populations on the basis of their correlation functions."
 Since Soneira&Peebles(C1977).. nearest neighbour gaatistics have been a common observational tool to discriminate between spatial distributions of objects.," Since \citet{So77}, nearest neighbour statistics have been a common observational tool to discriminate between spatial distributions of objects."
 We carry out such an analysis in real space lor our series of subsamples. taking L; spiral galaxies as the neighbour population (following POI).," We carry out such an analysis in real space for our series of subsamples, taking $L_{*}$ spiral galaxies as the neighbour population (following P01)."
" We find that blue. star-forming galaxies are more weakly associated with £, spirals than are other L. spirals but otherwise we find no significant ellect."," We find that blue, star-forming galaxies are more weakly associated with $L_{*}$ spirals than are other $L_{*}$ spirals but otherwise we find no significant effect."
 We conclude that there is no dwarf-tvpe void population in the simulation. cwn to our resolution limit.," We conclude that there is no dwarf-type void population in the simulation, down to our resolution limit."
 The disposition of our paper is as follows: in Section 2.. we begin with a qualitative. visual comparison between the galaxies that have formed. in a void and in a cluster environment.," 				 The disposition of our paper is as follows: in Section \ref{sec:Visual}, we begin with a qualitative, visual comparison between the galaxies that have formed in a void and in a cluster environment."
 The distribution of haloes in our void. does nol support the assertion of POI that the voids should contain small haloes., The distribution of haloes in our void does not support the assertion of P01 that the voids should contain small haloes.
 Then. in Section 3... we describe the characteristics of our galaxy ancl DM halo subsamples.," Then, in Section \ref{sec:Samples}, we describe the characteristics of our galaxy and DM halo subsamples."
 We show the distribution of our reference. spirals in a slice spanning the superealactic plane., We show the distribution of our reference spirals in a slice spanning the supergalactic plane.
 Section 4 evaluates the environment. densities for the various mock. samples. and computes the density. dependence of Les.," Section \ref{sec:Envi} evaluates the environment densities for the various mock samples, and computes the density dependence of LFs."
 We switch to (wo-point statistics in Section 5.., We switch to two-point statistics in Section \ref{sec:Clus}.
 We check the correlation functions of our subsamples ancl then carry oul a nearest neighbour analysis similar to that of POI., We check the correlation functions of our subsamples and then carry out a nearest neighbour analysis similar to that of P01.
 We summarize and conclude in Section 6.., We summarize and conclude in Section \ref{sec:CCL}.
 Figure 3 of MOL (see also the bottom right panel of Fig., Figure 3 of M01 (see also the bottom right panel of Fig.
" 1 below) plots the simulated. galaxy distribution within SOOO of the δνλίαν, in a slice of thickness 30 centred. on the supergalactic plane."," \ref{fig:4PanelsFig} below) plots the simulated galaxy distribution within 8000 of the Milky-Way, in a slice of thickness 30 centred on the supergalactic plane."
 “Pwo large voids are visible: the first at (-60.-30) in (SGN.SGY). with a diameter of ~201Alpe. the other nearly opposite to the MW. at (40.50). where a galaxy separates two smaller voids. of diameter 10h+Mpe each.," Two large voids are visible: the first at (-60,-30) in (SGX,SGY), with a diameter of $\sim 20 \: \hMpc$, the other nearly opposite to the MW, at (40,50), where a galaxy separates two smaller voids, of diameter $\sim 10\: \hMpc$ each."
 The sanie voids are seen in the real galaxy distribution for example the survey (consider Figure 2 of Plionis&Dasilalkos2001))., The same voids are seen in the real galaxy distribution for example the survey (consider Figure 2 of \citealt{Pl01}) ).
 As an example. we locus here on the void. located at (40.50) since it lies further inside high resolution region of the simulation. and may be less allectec by the transition from the high-resolution to the low-resolution zone.," As an example, we focus here on the void located at (40,50) since it lies further inside high resolution region of the simulation, and may be less affected by the transition from the high-resolution to the low-resolution zone."
 We excise a cubic region of side 24 centred on the micedle of the void., We excise a cubic region of side 24 centred on the middle of the void.
 The top left panel of Fig., The top left panel of Fig.
 1. shows the distribution of all galaxies within this cube brighter than our luminosity resolution limit., \ref{fig:4PanelsFig} shows the distribution of all galaxies within this cube brighter than our luminosity resolution limit.
 The colors of the svmbols scale with the DV index. and their sizes with the D-magnitude of the ealaxies.," The colors of the symbols scale with the $B-V$ index, and their sizes with the B-magnitude of the galaxies."
 Because of the galaxy. formation scheme adopted in MOI. all DM haloes in the simulation with 10 or more particles contain at least one galaxy.," Because of the galaxy formation scheme adopted in M01, all DM haloes in the simulation with 10 or more particles contain at least one galaxy."
 Furthermore. with our definition of the luminosity resolution limit (tho mean luminosity of the central galaxy. of a 10-particle halo). oneexpects that a fair fraction of the haloes more massive than 10Mss=357Lot)+AL. will have an associated ealaxy with Mgx16.27. and so will appear on the top Left picture of Fig. 1..," Furthermore, with our definition of the luminosity resolution limit (the mean luminosity of the central galaxy of a 10-particle halo), oneexpects that a fair fraction of the haloes more massive than $10\times M_{part}=3.57\times10^{10}
\hmsun$ will have an associated galaxy with $M_{B}\le-16.27$, and so will appear on the top left picture of Fig. \ref{fig:4PanelsFig}."
 In the region shown. of all haloes have a central galaxy at least this bright.," In the region shown, of all haloes have a central galaxy at least this bright."
 This is shown explicitly in the bottom left panel of Fig. 1..," This is shown explicitly in the bottom left panel of Fig. \ref{fig:4PanelsFig},"
 where we project the same low density region as in the top left panel. and. mark with circles the positions of all DM haloes with 10 or more particles.," where we project the same low density region as in the top left panel, and mark with circles the positions of all DM haloes with 10 or more particles."
 There are two obvious voids completely. depleted: of haloes. with diameters of order 10 and 8 and depths of 24h. This contradicts POL who claims that in a model. low mass haloes should. spread. fairly evenly through the voids defined by the larger ones.," There are two obvious voids completely depleted of haloes, with diameters of order 10 and 8 and depths of 24 This contradicts P01 who claims that in a model, low mass haloes should spread fairly evenly through the voids defined by the larger ones."
 In fact. eravity seems to have moved them out of these voids.," In fact, gravity seems to have moved them out of these voids."
 For comparison. we give in the top right panel of Fig.," For comparison, we give in the top right panel of Fig."
 1 the ealaxy distribution around the simulated. Virgo cluster Mound+10th5AL. ). also plotted in a cubic region of side 24h centeed on the cluster centre.," \ref{fig:4PanelsFig}
 the galaxy distribution around the simulated Virgo cluster $\Mvir \sim 4\times10^{14}\hmsun$ ), also plotted in a cubic region of side 24 centred on the cluster centre."
 The galaxy symbols follow the same rules and scales as in the left panel., The galaxy symbols follow the same rules and scales as in the left panel.
 Note that to avoid saturating the cluster region. galaxies are plotted down to Mg=17 only.," Note that to avoid saturating the cluster region, galaxies are plotted down to $M_{B}=-17$ only."
 The few isolatecl galaxies observed. in the Bobttes void are normal. late-type spirals (Szomoructal. 1996ab)..," The few isolated galaxies observed in the Boöttes void are normal, late-type spirals \citep{Szo96a,Szo96b}. ."
 Nevertheless. there is a widespread belief that voids should be filled by [ate-tvpe and. low surface brightness. chwarls.," Nevertheless, there is a widespread belief that voids should be filled by late-type and low surface brightness dwarfs."
where u*=B2/P* is measured at the disk surface.,where $\mu^+= B_z^2/P^+$ is measured at the disk surface.
 We plotted in Fig., We plotted in Fig.
 15 (down) the disk magnetization at the upper surface layers., \ref{sigmu} (down) the disk magnetization at the upper surface layers.
" It can be seen that o*(r) follows approximately the same trend as 4(r), namely a radial decrease."," It can be seen that $\sigma^+(r)$ follows approximately the same trend as $\mu^+(r)$, namely a radial decrease."
" In fact, analytical calculations done within the self-similar framework already pointed out the importance of the disk magnetization µ for launching super-FM jets."," In fact, analytical calculations done within the self-similar framework already pointed out the importance of the disk magnetization $\mu$ for launching super-FM jets."
" It was shown that isothermal (??) or adiabatic (?) magnetic surfaces require a field close to equipartition, namely u smaller but around unity."," It was shown that isothermal \citep{1995A&A...295..807F,1997A&A...319..340F} or adiabatic \citep{2000A&A...353.1115C} magnetic surfaces require a field close to equipartition, namely $\mu$ smaller but around unity."
 Our own results suggest that it is the disk magnetization that actually defines the ejecting zones., Our own results suggest that it is the disk magnetization that actually defines the ejecting zones.
" Beyond r=5, the magnetic field would be too small to allow a proper jet to be launched."," Beyond $r=5$, the magnetic field would be too small to allow a proper jet to be launched."
 Let us make a very crude approach by assuming a B; component almost constant in the vertical direction and an isothermal hydrostatic density profile., Let us make a very crude approach by assuming a $B_z$ component almost constant in the vertical direction and an isothermal hydrostatic density profile.
" In that case, one would have p(z)=μexp(z?/2h?) where µ isthe disk magnetization at the disk midplane."," In that case, one would have $\mu(z) = \mu
\exp(z^2/2h^2)$ where $\mu$ isthe disk magnetization at the disk midplane."
" It seems therefore dubious that u*, reaching a value 10 to 100 times µ at a few scale height, could ever reach a value of order unity if u is too small."," It seems therefore dubious that $\mu^+$, reaching a value 10 to 100 times $\mu$ at a few scale height, could ever reach a value of order unity if $\mu$ is too small."
" So, even in the presence of a numerical diffusion, no jets should be produced if the disk magnetization is too low."," So, even in the presence of a numerical diffusion, no jets should be produced if the disk magnetization is too low."
" In order to test this conjecture, we performed another numerical simulation with 4 decreased by one order of magnitude (namely starting at 2x107 at the disk inner radius)."," In order to test this conjecture, we performed another numerical simulation with $\mu$ decreased by one order of magnitude (namely starting at $2\times10^{-4}$ at the disk inner radius)."
" The other physical parameters, boundary and initial conditions, as well as the numerical resolution were otherwise identical to the reference simulation."," The other physical parameters, boundary and initial conditions, as well as the numerical resolution were otherwise identical to the reference simulation."
" We found that, in this case the super-Alfvénnic material was extremely sporadic and fragmented in the domain, and no super-fast-magnetosonic jet was observed."," We found that, in this case the super-Alfv\'ennic material was extremely sporadic and fragmented in the domain, and no super-fast-magnetosonic jet was observed."
 See Fig. 16.., See Fig. \ref{weakB}.
" In our view, this clearly confirms that the disk magnetization must be high enough in order to launch self- (super-FM) jets."," In our view, this clearly confirms that the disk magnetization must be high enough in order to launch self-confined (super-FM) jets."
 This result goes in the same direction as those obtained with self-similar solutions., This result goes in the same direction as those obtained with self-similar solutions.
" However, the latter claimed that only µ smaller but close to unity (namely a field close to equipartition) allows the launching of magnetized jets."," However, the latter claimed that only $\mu$ smaller but close to unity (namely a field close to equipartition) allows the launching of magnetized jets."
 The physical argument is the following., The physical argument is the following.
" For a jet to be launched, the lifted mass must cross the SM point around the disk surface."," For a jet to be launched, the lifted mass must cross the SM point around the disk surface."
 In a cold environment the only force able to do this is the magnetic one., In a cold environment the only force able to do this is the magnetic one.
" It turns out however, that it is much easier to do it if the accretion velocity is already not too far from the sound speed."," It turns out however, that it is much easier to do it if the accretion velocity is already not too far from the sound speed."
" This is the reason why isothermal or adiabatic jets require fields close to equipartition (Ferreira Casse, 2009, submitted)."," This is the reason why isothermal or adiabatic jets require fields close to equipartition (Ferreira Casse, 2009, submitted)."
" Apparently, this is in contradiction with our own result since we do obtain jets with µ of the order of a few 107°."," Apparently, this is in contradiction with our own result since we do obtain jets with $\mu$ of the order of a few $10^{-3}$."
 The reason for this discrepancy lies in the fact that the analytical models were obtained under the assumption of either isothermal or adiabatic magnetic surfaces., The reason for this discrepancy lies in the fact that the analytical models were obtained under the assumption of either isothermal or adiabatic magnetic surfaces.
" Here, as we showed, there is a numerical diffusion that allowed mass to leak from the disk to the open, rotating field lines."," Here, as we showed, there is a numerical diffusion that allowed mass to leak from the disk to the open, rotating field lines."
 This extra effect has been mimicked for instance in ? with the presence of a heating term at the disk surface., This extra effect has been mimicked for instance in \citet{2000A&A...361.1178C} with the presence of a heating term at the disk surface.
" New solutions, called “warm” in contrast to the previous “cold” ones, were found with an enhanced mass flux."," New solutions, called “warm” in contrast to the previous “cold” ones, were found with an enhanced mass flux."
 But these authors did not recognize that the required disk magnetization was indeed smaller than for cold jets., But these authors did not recognize that the required disk magnetization $\mu$ was indeed smaller than for cold jets.
" We report here that it is indeed the case, with some of the “warm” self-similar solutions found with yz= 0.08."," We report here that it is indeed the case, with some of the “warm” self-similar solutions found with $\mu=0.08$ ."
" In this paper, we performed four 2.5D numerical MHD simulations of a resistive viscous accretion disk threaded by"," In this paper, we performed four 2.5D numerical MHD simulations of a resistive viscous accretion disk threaded by"
been transtormed iuto internal kinetic energv due to conrpressionu work: the final volume of the fiuid is smaller than the initial volume. and the pressure has always opposed the ution induced by eravity.,"been transformed into internal kinetic energy due to compression work; the final volume of the fluid is smaller than the initial volume, and the pressure has always opposed the motion induced by gravity."
 This is similar to what happeus in the cosmological setting: volume elements would have equal tendency to compress or expand due to pressurelike forces. but the presence of eravity favors coutraction of volume clemeuts. especially when they are falling outo liehdensity regious.," This is similar to what happens in the cosmological setting: volume elements would have equal tendency to compress or expand due to pressure–like forces, but the presence of gravity favors contraction of volume elements, especially when they are falling onto high–density regions."
 The effect of transformation of caerey iuto internal velocity dispersion is large. because not oulv potential energy. but also bulk kinetic energy decreases. when the fluid elemoeuts crash outo highdensity walls.," The effect of transformation of energy into internal velocity dispersion is large, because not only potential energy, but also bulk kinetic energy decreases, when the fluid elements crash onto high–density walls."
 We have also shown that pressurelike forces may be relevant even on cosmological scales aud au ciective Jeans? leneth may be larger than the (comoving) elobular cluster scale; which is conunouly attributed to a lyvdrodvuamical (ideal gas) pressure (compare this to the effect of a gas pressure within pancakes investigated by Suuvaev Zeldovich 1972).," We have also shown that pressure–like forces may be relevant even on cosmological scales and an effective Jeans' length may be larger than the (comoving) globular cluster scale, which is commonly attributed to a hydrodynamical (ideal gas) pressure (compare this to the effect of a gas pressure within pancakes investigated by Sunyaev Zel'dovich 1972)."
 ITowever. we now understand clearly the limited status of this aud similar models;," However, we now understand clearly the limited status of this and similar models."
 We shall now elaborate on this critficisni dn more detail., We shall now elaborate on this criticism in more detail.
 The extrapolation of the relationship (16)) iuto the weakly nonlinear regunue is well justified for “dust”., The extrapolation of the relationship \ref{parallelity}) ) into the weakly non–linear regime is well justified for “dust”.
 If we define the weakly nonluear reegiue as sugeested im (Buchert 1989). namely. as a linearization iu Lagraugian space. Zeldovich’s extrapolation of the Eulerian linear perturbation solution iuto the nonlinear regime is not only consistent with Lagrangian luearization of the full EulerPoisson system (Duchert 1989. see also Doroshkevich ct al.," If we define the weakly non–linear regime as suggested in (Buchert 1989), namely, as a linearization in Lagrangian space, Zel'dovich's extrapolation of the Eulerian linear perturbation solution into the nonlinear regime is not only consistent with Lagrangian linearization of the full Euler--Poisson system (Buchert 1989, see also Doroshkevich et al."
 L973 for a ποconsistency test for part of the system). vut arises naturallv as a subcase of firstorder solutions in a Laerangian perturbation approach (Buchert 1992).," 1973 for a self–consistency test for part of the system), but arises naturally as a subcase of first–order solutions in a Lagrangian perturbation approach (Buchert 1992)."
" This selfcousistency of the ""dust model” is nirrored iu he fact that the trajectories of Lagrangian firstorder orturbations even provide a class of 3D exact solutious o the Euler.Poissou svstein (Buchert 1989).", This self–consistency of the “dust model” is mirrored in the fact that the trajectories of Lagrangian first–order perturbations even provide a class of 3D exact solutions to the Euler–Poisson system (Buchert 1989).
" If pressure is taken iuto account. extrapolation of the “cust” model trajectories (as we had to do in order ο recover the ""adhesion approximation) is no longer yernutted."," If pressure is taken into account, extrapolation of the “dust” model trajectories (as we had to do in order to recover the “adhesion approximation”) is no longer permitted."
 As we have shown it may be used as a approximation- for small velocity dispersion (and this is xeciselv the method followed iu Zeldovich Shaucdarin 1982 and Bharadwaj 1996)., As we have shown it may be used as a approximation for small velocity dispersion (and this is precisely the method followed in Zel'dovich Shandarin 1982 and Bharadwaj 1996).
 À similar remark applies o other assmuptious ofextrapolation like the FrozenPoteutialApproximation (Brainerd et al., A similar remark applies to other assumptions ofextrapolation like the Frozen--Potential--Approximation (Brainerd et al.
 1993. Bagla Padmanabhan 1991). which we may curplov instead of (16)).," 1993, Bagla Padmanabhan 1994), which we may employ instead of \ref{parallelity}) )."
" This approach simply follows the evolution of velocity dispersion along the trajectories dictated by the ""dust model its main drawback being that the ""backreaction” of velocity dispersion on these trajectories is uot taken iuto account."," This approach simply follows the evolution of velocity dispersion along the trajectories dictated by the “dust” model, its main drawback being that the ``back--reaction'' of velocity dispersion on these trajectories is not taken into account."
 Iu fact. the simple version of parallelisin (16)) is not even justified in the (Eulerian) linear regine (Buchert et al.," In fact, the simple version of parallelism \ref{parallelity}) ) is not even justified in the (Eulerian) linear regime (Buchert et al."
 1998)., 1998).
 The breakdown of the asstuuption (16)) after pancake formation is plausible aud it is wellknown for epochs after shellcrossing to which we want to apply the model (Doroshkevich 1973)., The break–down of the assumption \ref{parallelity}) ) after pancake formation is plausible and it is well–known for epochs after shell–crossing to which we want to apply the model (Doroshkevich 1973).
" Nevertheless. the ""adhesion approximation” works rather well. because the model is taken at face value aac rot restricted by the constraints which lec us to its derivation."," Nevertheless, the “adhesion approximation” works rather well, because the model is taken at face value and not restricted by the constraints which led us to its derivation."
 Accordingly. we expect the model we xoposed to improve on the “adhesion approximation. since it specifies the dependence of “adhesion” ou the loca density.," Accordingly, we expect the model we proposed to improve on the “adhesion approximation”, since it specifies the dependence of “adhesion” on the local density."
 The publication of this work was motivated bv lis gain of insight and iaprovenient., The publication of this work was motivated by this gain of insight and improvement.
 Let us παΊσο he constraümniug assumptions which restrict the gencra woblem to the ‘adhesion approximation”: roludente simall velocity dispersion: we keep oulv the first ©πιαοι in the hierarchy of velocity moments (AL)., Let us summarize the constraining assumptions which restrict the general problem to the “adhesion approximation”: $\bullet\;$ small velocity dispersion; we keep only the first equation in the hierarchy of velocity moments (A1).
 e isotropy of the dispersion tensor which implies isotropy of the ucan motion (A2)., $\bullet\;$ isotropy of the dispersion tensor which implies isotropy of the mean motion (A2).
 eparallelitv of peculiarvelocity aud acceleration. both defined relative to a elobal Hubbleflow CÀ3).," $\bullet\;$parallelity of peculiar–velocity and –acceleration, both defined relative to a global Hubble–flow (A3)."
 e a spatially constant relationship between initial density and initial pressure (AL)., $\bullet\;$ a spatially constant relationship between initial density and initial pressure (A4).
 e the GALcocfiicicut is coustaut in space and time as well ax positive (A5) @vhich could be achieved by a relationship-. p=D πρ].," $\bullet\;$ the GM–coefficient is constant in space and time as well as positive (A5) (which could be achieved by a relationship $p
= \kappa \varrho^2$ )."
" This list of assunptious shows that there is enough room for a eencralization of the ""adhlesion approxination.", This list of assumptions shows that there is enough room for a generalization of the “adhesion approximation”.
 Our main concern should be focussed on assumptions CAT) and (A2)., Our main concern should be focussed on assumptions (A1) and (A2).
 Although both of them may be cousidered a good working hvpothesis from a theoretical point of view. we kuow that the physical situation is not iu favor of these assiunptions. especially if we follow the average flow further into the nonlinear regine.," Although both of them may be considered a good working hypothesis from a theoretical point of view, we know that the physical situation is not in favor of these assumptions, especially if we follow the average flow further into the nonlinear regime."
 Velocity dispersion will become quickly huge in the inultistream regiae., Velocity dispersion will become quickly large in the multi–stream regime.
 A lhniut of this can be ποσα in observations of velocity dispersion im rich clusters of galaxies which is of the order of 1000 ans. Le. certainly not small compared with the bulk speed.," A hint of this can be seen in observations of velocity dispersion in rich clusters of galaxies which is of the order of 1000 km/s, i.e. certainly not small compared with the bulk speed."
 However. viewing assunuptious (AL) and (A2) together. we would obtain substantially similar uiodels by aremineg ou plenomenological grounds: assumption (A2) eives an idealized model for the dispersion tensor and it just necds a relationship p=(0) to close he hierarchy of velocity moments.," However, viewing assumptions (A1) and (A2) together, we would obtain substantially similar models by arguing on phenomenological grounds: assumption (A2) gives an idealized model for the dispersion tensor and it just needs a relationship $p = \beta
(\varrho)$ to close the hierarchy of velocity moments."
 In light ofa phenomenological approach asstuuption (AL) or. alternatively. assunptiou (À5) just specifv the function J. but we may as wel consider other relationships as discussed in (Buchert et al.," In light of a phenomenological approach assumption (A1) or, alternatively, assumption (A5) just specify the function $\beta$, but we may as well consider other relationships as discussed in (Buchert et al."
 1998)., 1998).
 This problei calls for a better model that relates the dispersio- tensor to the deusitv aud velocity fields. but it does uot invalidate the main ideas exempliied with assumption (Al).," This problem calls for a better model that relates the dispersion tensor to the density and velocity fields, but it does not invalidate the main ideas exemplified with assumption (A1)."
 Concerning assumption (A3) a first indication of how to proceed iu order to investigate a valid extrapolation into the weakly nonlinear regine is trnished by Eq. (15)), Concerning assumption (A3) a first indication of how to proceed in order to investigate a valid extrapolation into the weakly nonlinear regime is furnished by Eq. \ref{mastereq1}) )
 itself: the adhesive teri is proportional to Aw rather than Au., itself: the adhesive term is proportional to $\Delta {\bf w}$ rather than $\Delta {\bf u}$ .
 It is possible to derive an evolution equation for the secularC»eravitational field strenethOo w., It is possible to derive an evolution equation for the peculiar–gravitational field strength ${\bf w}$ .
 Its specification to, Its specification to
Creat progress has been mace in understanding jet acceleration. and. collimation.,Great progress has been made in understanding jet acceleration and collimation.
 Evicenee is found that the optical electric vector positional angle (IEEVPX) of à newly formecl feature. rotates while it moves through a region upstream of the radio cores in BL Lac and PIS 1510-089. which is attributed to helical magnetic fields accelerating and focusing the jet in that zone (Alarscheretal.2008. 2010).," Evidence is found that the optical electric vector positional angle (EVPA) of a newly formed feature rotates while it moves through a region upstream of the radio cores in BL Lac and PKS 1510-089, which is attributed to helical magnetic fields accelerating and focusing the jet in that zone \citep{mar08,mar10}."
. In these two cases. the accelerating and. focusing orocess lasts until the feature encounters a region of urbulence. where the magnetic fields. become chaotic.," In these two cases, the accelerating and focusing process lasts until the feature encounters a region of turbulence, where the magnetic fields become chaotic."
 Whereas further attempts to investigate the process. are still required. the jets in quite a few sources. are. found o be accelerated. downstream of the core (eg... Jorstadetal. 2005)). and helical magnetic fields. are argued to of present within the jets in many sources by identifving he Faraday rotation measure (IM). eraclicnts across the jet downstream from the core (ο-.. Zavala&Γανίοι2005:CGabuzelaotal. 2008)).," Whereas further attempts to investigate the process are still required, the jets in quite a few sources are found to be accelerated downstream of the core (e.g., \citealt{jor05}) ), and helical magnetic fields are argued to be present within the jets in many sources by identifying the Faraday rotation measure (RM) gradients across the jet downstream from the core (e.g., \citealt{zav05,gab08}) )."
 Lt is therefore likely that large-scale helical magnetic fields exist in jets. and. play an important role in jet acceleration and collimation downstream of the core às well. as suggested by Lyutikoyetal.(2005).," It is therefore likely that large-scale helical magnetic fields exist in jets, and play an important role in jet acceleration and collimation downstream of the core as well, as suggested by \cite{lyu05}."
. AIST ds one of the closest radio galaxies with a distance of 16.7 Mpce. which is active from radio to *5-rav waveband (e.g... Aeciarictal. 2009)).," M87 is one of the closest radio galaxies with a distance of 16.7 Mpc, which is active from radio to $\gamma$ -ray waveband (e.g., \citealt{acc09}) )."
" HIST-1 in ALS7 jet was defined. to be a feature. or component separated: by ~0.86"" from the central black hole. which. is resolved into several ‘subcomponents at mas resolution with the most east one stationary and some others moving down the jet superluminally (Cheungetal.2007)."," HST-1 in M87 jet was defined to be a feature or component separated by $\sim0.86''$ from the central black hole, which is resolved into several `subcomponents' at mas resolution with the most east one stationary and some others moving down the jet superluminally \citep{che07}."
.. I0 seems that new subcomponents are ejected. frequently. from. the stationary one., It seems that new subcomponents are ejected frequently from the stationary one.
 The Ilux density monitoring at multi-wavebands shows that a major outburst of LST-1 occurred from 2003 to 2007. of which the radiation was regarded to be svnchrotron fron radio to X-ray (Llarrisetal.2003.2009).," The flux density monitoring at multi-wavebands shows that a major outburst of HST-1 occurred from 2003 to 2007, of which the radiation was regarded to be synchrotron from radio to X-ray \citep{har03,har09}."
.. Alone with α 5-rav flare occurred in the period (Aharonianetal.2006).. LIS'T-1 was argued to show significant blazar-like properties (e.g.. Harrisetal. 2008)). of which the polarization nature has vet been rarely investigated.," Along with a $\gamma$ -ray flare occurred in the period \citep{aha06}, HST-1 was argued to show significant blazar-like properties (e.g., \citealt{har08}) ), of which the polarization nature has yet been rarely investigated."
 Llere we present the multi-epoch full polarization results oELIST-1 in AIST. which are derived from VL observations," Here we present the multi-epoch full polarization results of HST-1 in M87, which are derived from VLA observations"
" Photometrically distinet nuclear star clusters (NSCs) are ubiquitous in dynamically primitive galaxies. which are the galaxies that lack what can be interpreted as the structural signatures of major—""wet or ""dry surface brightness profiles with de Vaucouleurs-like or —mergers:.higher Sérsic(1963) indices. and. absent core scouring by binary massive black holes. power-law central luminosity cusps (e.g..Hop-kinsetal.2008.2009a.b.andreferences therein)..."," Photometrically distinct nuclear star clusters (NSCs) are ubiquitous in dynamically primitive galaxies, which are the galaxies that lack what can be interpreted as the structural signatures of major—“wet” or “dry”—mergers: surface brightness profiles with de Vaucouleurs-like or higher \citet{Sersic:63} indices, and, absent core scouring by binary massive black holes, power-law central luminosity cusps \citep[e.g.,][and references
 therein]{Hopkins:08,Hopkins:09a,Hopkins:09b}."
" Imaging with theTelescope CHST) has revealed that 75% of late-type (Sc-Sd) disk galaxies contain compact. luminous star clusters at their centers with masses in the range 8«10°-6107M ... effective radii in the range 1—9pe with a median of —3.5pc. and velocity dispersions in the range 13—34kms""! (e.g..Bokeretal.2002.2004:Walcheretal. 2005)."," Imaging with the (HST) has revealed that $75\%$ of late-type (Sc–Sd) disk galaxies contain compact, luminous star clusters at their centers with masses in the range $8\times10^5-6\times10^7\,M_\odot$ , effective radii in the range $1-9\,\textrm{pc}$ with a median of $\sim 3.5\,\textrm{pc}$, and velocity dispersions in the range $13-34\,\textrm{km}\,\textrm{s}^{-1}$ \citep[e.g.,][]{Boker:02,Boker:04,Walcher:05}."
". The clusters? ages range widely from 4«107 to 10""yr. the metallicities average at (Zi=0.015 with a significant scatter. the average star formation rate over the last LOOMyr is (“spp22«107M..yr7!. and older stars dominate the cluster mass."," The clusters' ages range widely from $4\times10^7$ to $10^{10}\,\textrm{yr}$, the metallicities average at $\langle Z\rangle=0.015$ with a significant scatter, the average star formation rate over the last $100\,\textrm{Myr}$ is $\langle\Sigma_{\rm
 SFR}\rangle=2\times10^{-3}\,M_\odot\,\textrm{yr}^{-1}$, and older stars dominate the cluster mass."
 Similarly. an HST survey of stellar nuclei in spheroidal galaxies in the Virgo and Fornax clusters (Ferrareseetal.2006a:Coté2006;Miller&Lotz2007) has revealed stellar nuclei that on average contain ~0.2% of the galaxy mass (Ferrareseetal.2006b;Wehner&ris2006) and have a median effective radius of 4.2pc also.Gehaetal. 2002).," Similarly, an HST survey of stellar nuclei in spheroidal galaxies in the Virgo and Fornax clusters \citep{Ferrarese:06a,Cote:06,Miller:07} has revealed stellar nuclei that on average contain $\sim 0.2\%$ of the galaxy mass \citep{Ferrarese:06b,Wehner:06} and have a median effective radius of $4.2\,\textrm{pc}$ \citep[see, also,][]{Geha:02}."
. Also. HST imaging of nearby dwarf galaxies with absolute magnitudes —16— (Georgievetal.2009) has revealed NSCs with masses ~(109—10) with a majority clustering at the lower end of the mass Mrange.... and half-light radii 1.5—4.7pe.," Also, HST imaging of nearby low-luminosity dwarf galaxies with absolute magnitudes $-16<M_V<-13$ \citep{Georgiev:09} has revealed NSCs with masses $\sim
(10^6-10^7)\,M_\odot$, with a majority clustering at the lower end of the mass range, and half-light radii $1.5-4.7\,\textrm{pc}$."
 It seems that photometrically distinct NSCs are common in galaxies that have not experienced major mergers m the epoch following the initial burst of star formation (see.e.g..Kodaetal. 2009).. but have a well-defined dynamical center.," It seems that photometrically distinct NSCs are common in galaxies that have not experienced major mergers in the epoch following the initial burst of star formation \citep[see, e.g.,][]{Koda:09}, but have a well-defined dynamical center."
 The faintest dwarf spheroidal satellites of the Milky Way. with the exception of the Sagittarius Dwarf (e.g..Bellazzinietal.2008.andreferences therein).. do not contain NSCs. even though they do seem to form a continuous structural family of stellar systems with the more massive dwarf spheroidals that do contain NSCs (e.g..Kormendyetal.2009:Wyse2010.andreferences thereimn).," The faintest dwarf spheroidal satellites of the Milky Way, with the exception of the Sagittarius Dwarf \citep[e.g.,][and references therein]{Bellazzini:08}, do not contain NSCs, even though they do seem to form a continuous structural family of stellar systems with the more massive dwarf spheroidals that do contain NSCs \citep[e.g.,][and references
 therein]{Kormendy:09,Wyse:10}."
. While the spheroidal galaxies in. Virgo are typically non-rotating (Gehaetal. 2003).. their nucleus-subtracted surface brightness profiles. with Sérrsic indices 2~|—2. resemble those of disks.," While the spheroidal galaxies in Virgo are typically non-rotating \citep{Geha:03}, their nucleus-subtracted surface brightness profiles, with Sérrsic indices $n\sim 1-2$, resemble those of disks."
 This motivates a unified approach to addressing NSC formation in these two morphological classes., This motivates a unified approach to addressing NSC formation in these two morphological classes.
" In disks. an excess surface brightness above the inward-extrapolated exponential law of the outer disk occurring at radit. ~100—500pe. variously characterized as ""central light excess” (Boker or a “pseudobulge” (Kormendy&Kennicutt2004:Fisher 2009). may provide additional clues. because the process that produced the NSC could also contribute to the formation of the pseudobulge."," In disks, an excess surface brightness above the inward-extrapolated exponential law of the outer disk occurring at radii $\sim 100-500\,\textrm{pc}$, variously characterized as “central light excess” \citep{Boker:03} or a “pseudobulge” \citep{Kormendy:04,Fisher:08,Fisher:10,Fisher:09,Weinzirl:09}, may provide additional clues, because the process that produced the NSC could also contribute to the formation of the pseudobulge."
 Pseudobulges can be distinguished from classical bulges through their cold. rapidly rotating kinematics and low Sérmsie index (seeKormendy&Fisher2008.andreferences therein)... Kormendy," Pseudobulges can be distinguished from classical bulges through their cold, rapidly rotating kinematics and low Sérrsic index \citep[see][and references therein]{Kormendy:08}."
etal.(2010). find that the pseudobulges in the sample of pure-disk (Sed) galaxies contain =3%of the galactie stellar mass. which is smaller than the stellar mass fraction in classical bulges.," \citet{Kormendy:10}
 find that the pseudobulges in the sample of pure-disk (Scd) galaxies contain $\lesssim 3\%$of the galactic stellar mass, which is smaller than the stellar mass fraction in classical bulges."
 Here we discuss the formation of NSCs in the context of normal star formation in which most stars form in clusters., Here we discuss the formation of NSCs in the context of normal star formation in which most stars form in clusters.
 A successful theory of NSC formation should explain the, A successful theory of NSC formation should explain the
GIU (e.g. 2). 2). 7). see also Sec. 4.4)).,"GR (e.g. \cite{2007MNRAS.378..852P}, , \cite{2007MNRAS.374.1527B}, \cite{2011MNRAS.412.1669T}, see also Sec. \ref{sec:implications}) )."
 Our approach corresponds to the model of ?).., Our approach corresponds to the model of \cite{2011arXiv1102.1014S}.
" For the cosmological model we will assume a .MCDM universe compatible with WALAP?T (2))) with O,,=0.272. Opp=0.728. οι=0.0456. ης=0.963. fh=0.704."," For the cosmological model we will assume a $\Lambda$ CDM universe compatible with WMAP7 \cite{2011ApJS..192...18K}) ) with $\Omega_m=0.272$, $\Omega_{DE}=0.728$, $\Omega_b=0.0456$, $n_s=0.963$, $h=0.704$."
 In our approach these parameters determine the shape of the real-space matter correlation function and the distance-redshift relation., In our approach these parameters determine the shape of the real-space matter correlation function and the distance-redshift relation.
 As shown in Sec., As shown in Sec.
 2.3 μαrlitting the data in multiple bins results in samples that are highly correlated., \ref{sec:photoz} splitting the data in multiple bins results in samples that are highly correlated.
 Εις. we decided to focus in a single redshift bin to present our main results and. clefer the study of narrow bins to Sec. 4.3..," Thus, we decided to focus in a single redshift bin to present our main results and defer the study of narrow bins to Sec. \ref{sec:systematics},"
 in the context of robustness and consistency. studies., in the context of robustness and consistency studies.
 Llenee we concentrate on the bin 0.50.0]., Hence we concentrate on the bin $[0.5-0.6]$.
 Lhis width is about. twice the tvpical photometric error while the center of the bin makes this data uncorrelated with the SDSS LRG spectroscopic sample (that we refer to later on)., This width is about twice the typical photometric error while the center of the bin makes this data uncorrelated with the SDSS LRG spectroscopic sample (that we refer to later on).
 Phe estimate for the true redshift distribution of galaxies in this bin is shown in Fig. 7.., The estimate for the true redshift distribution of galaxies in this bin is shown in Fig. \ref{fig:dndz0506}.
 Using a spline-fit to it we computed the observed: model. correlation. following Ies. (s..9--12))," Using a spline-fit to it we computed the observed model correlation following Eqs. \ref{eq:wstars}, \ref{eq:wtheta}- \ref{eq:wtheta2}) )"
 sampling the two-parameter space given by faxbos., sampling the two-parameter space given by $f\sigma_8-b\sigma_8$.
 We then performed a standard 72 minimization. where and Aw=Wobsmed(fox.Px)Uieasnrcd," We then performed a standard $\chi^2$ minimization, where and $\Delta w \equiv w_{obs,model}(f \sigma_8,b \sigma_8) -
w_{measured}$."
 Lo begin with we use the jack-knife covariance., To begin with we use the jack-knife covariance.
 The resulting hest-Lit values and 1(0 errors are listed in Table 1. (to convert to f and b one can use that ax(0.55)=0.611 for our fiducial cosmology σκίς=0) 0.8)., The resulting best-fit values and $1-\sigma$ errors are listed in Table \ref{table:best-fit} (to convert to $f$ and $b$ one can use that $\sigma_8(0.55)=0.611$ for our fiducial cosmology if $\sigma_8(z=0)=0.8$ ).
" The fit vields 2,5,=26.67 for 28 bins in 8 in the range 0.6%6°] and two fitting parameters.", The fit yields $\chi_{min}^2= 26.67$ for $28$ bins in $\theta$ in the range $\left[0.6^{\circ}-6^{\circ}\right]$ and two fitting parameters.
 Hence the quality of the fit is very good., Hence the quality of the fit is very good.
 The joint error distribution is displaved in Fig., The joint error distribution is displayed in Fig.
 14 and shows that these two parameters are not degenerate., \ref{fig:fbcontour} and shows that these two parameters are not degenerate.
 The significance of RSD in our data is ~1.266 (ie. the degree by which the recovered value for f£. is away from zero).," The significance of RSD in our data is $\sim 1.26 \, \sigma$ (i.e. the degree by which the recovered value for $f$ is away from zero)."
 Alternatively one can assume that the velocity growth rate is given. by General. Relativity. Le. f(z)ss(:)=ox(2eloD(z)/dlinea=0.45 (for our fiducial cosmology and σκίς=0) 0.8). and fit only for the overall amplitude of clustering.," Alternatively one can assume that the velocity growth rate is given by General Relativity, i.e. $f(z)\sigma_8(z) = \sigma_8(z){\rm d} \ln D(z) / {\rm d} \ln a = 0.45$ (for our fiducial cosmology and $\sigma_8(z=0)=0.8$ ), and fit only for the overall amplitude of clustering."
" Ln this case we recover a similar \5,;,(=26.7) and box=1.125Xx0.017."," In this case we recover a similar $\chi_{min}^2 (=26.7)$ and $b\,\sigma_8=1.125 \pm 0.017$."
 LHlence the [it does not. vary appreciable. showing that agreement with standard ACDAL is also very good and that the sensitivity of our data to redshift-cistortions is weak. as expected.," Hence the fit does not vary appreciable, showing that agreement with standard $\Lambda$ CDM is also very good and that the sensitivity of our data to redshift-distortions is weak, as expected."
 The results so far were obtained assuming JA star contamination (as discussed in Sec., The results so far were obtained assuming $4\%$ star contamination (as discussed in Sec.
 3.2. and Iq. (8))), \ref{sec:starcontamination} and Eq. \ref{eq:wstars}) )).
" M we now assume that our sample is perfectly. clean of stars Cfi,=O) we get bay=L08zE0.02 and fox=0.6630.39 and a XL;=26.87."," If we now assume that our sample is perfectly clean of stars $f_{stars}=0$ ) we get $b \,\sigma_8 = 1.08 \pm 0.02$ and $f \,\sigma_8 = 0.66 \pm 0.39$ and a $\chi^2_{min}=26.87$."
 This model is hardly. dillerentiable from the case with for.=4X (the X7 becomes only negligible worse). see Fig. 15..," This model is hardly differentiable from the case with $f_{stars}=4\%$ (the $\chi^2$ becomes only negligible worse), see Fig. \ref{fig:wtheta_z050-060}."
 Hence stellar contamination can mimick the effect. of RSD and become a worrisome source of svstematic ellect in future and better data., Hence stellar contamination can mimick the effect of RSD and become a worrisome source of systematic effect in future and better data.
 In the present case the agreement. between the best-fit value of f and the “GR” value degrades only slightly when [μα=0. with cdillerences well within the errors.," In the present case the agreement between the best-fit value of $f$ and the “GR” value degrades only slightly when $f_{star}=0$, with differences well within the errors."
 We conclude that the star contamination in our sample is sullicicnthy uncer control and does not drive our results., We conclude that the star contamination in our sample is sufficiently under control and does not drive our results.
 Figure 15. shows the resulting best-fit models discussed above against the measured angular correlation. function., Figure \ref{fig:wtheta_z050-060} shows the resulting best-fit models discussed above against the measured angular correlation function.
 In all cases the model matches the data quite well in all he angular range. with no signal of excess-clustering on arge scales.," In all cases the model matches the data quite well in all the angular range, with no signal of excess-clustering on large scales."
 This is to some extent at variance with the recent study by 2) (see also the follow up ?))) who finds he amplitude of the angular power spectra of photometric LRGs to have a 2-0 excess clustering away from the AC DAL wedietion at the lowest multipoles., This is to some extent at variance with the recent study by \cite{2011MNRAS.412.1669T} (see also the follow up \cite{2010arXiv1012.2272T}) ) who finds the amplitude of the angular power spectra of photometric LRGs to have a $\sigma$ excess clustering away from the $\Lambda$ CDM prediction at the lowest multipoles.
 Vhis result was obtained using the MegaZ catalog over three redshift bins in the range 0.4-0.65]., This result was obtained using the MegaZ catalog over three redshift bins in the range [0.4-0.65].
 Phe MegaZ catalog is also build upon the DIU SDSS photometric catalog., The MegaZ catalog is also build upon the DR7 SDSS photometric catalog.
 However the LRG selection criteria in MegaZ is cillerent from the one in this paper. most notably by the magnitude cuts (Atv<19.5 in MegaZ compared to pefror<21 in our case) but also in the color and meagso cuts intended to isolate stars from galaxies.," However the LRG selection criteria in MegaZ is different from the one in this paper, most notably by the magnitude cuts $i_{\rm deV} < 19.8$ in MegaZ compared to $petror < 21$ in our case) but also in the color and $mag_{50}$ cuts intended to isolate stars from galaxies."
 The photo-z code used for MegaZ is also dilferent. from. ours. although both are based in the ANNz code οἱ ?)..," The photo-z code used for MegaZ is also different from ours, although both are based in the ANNz code of \cite{2004PASP..116..345C}."
 7) also finds an excess clustering when searching for the xuwvon acoustic scale. but this result is dillicult to compare with ours as they use an over simplified. method to stack he signal from wider redshift bins calibrated with cillerent spectroscopic samples.," \cite{2009arXiv0912.0511S} also finds an excess clustering when searching for the baryon acoustic scale, but this result is difficult to compare with ours as they use an over simplified method to stack the signal from wider redshift bins calibrated with different spectroscopic samples."
 The constraints recovered so far. albeit. large. match he study of 2). that forecast how well redshift distortions can potentially be measured with galaxies selected from a οποιοιοεἰς survey.," The constraints recovered so far, albeit large, match the study of \cite{2011arXiv1102.0968R} that forecast how well redshift distortions can potentially be measured with galaxies selected from a photometric survey."
 This match can be clearly. seen. from. heir Fig., This match can be clearly seen from their Fig.
 δ alter scaling the results for a sample of unbiased ealaxies binned around 2=0.5 by our bias factor b~2 (since A(fox)x 6)., 8 after scaling the results for a sample of unbiased galaxies binned around $z=0.5$ by our bias factor $b\sim2$ (since $\Delta(f\sigma_8) \propto b$ ).
 For our redshift bin width the forecast vields ACfox)~O40.5. in very good. agreement with he results we obtain with the actual data (c.g. see Table 1).," For our redshift bin width the forecast yields $\Delta(f\sigma_8) \sim
0.4-0.5$, in very good agreement with the results we obtain with the actual data (e.g. see Table \ref{table:best-fit}) )."
 This implies that svstematics. photo-z errors. selection ancl modeling can be controlled sufficiently. well in. present xhotometric data to vield expected: constraints of recshilt distortions.," This implies that systematics, photo-z errors, selection and modeling can be controlled sufficiently well in present photometric data to yield expected constraints of redshift distortions."
 Near future data. provided with a better hanelle on systematics due to stars and. photo-z. should be able o complement shallower spectroscopicsurveys in. placing rounds to the growth of structure in the universe.," Near future data, provided with a better handle on systematics due to stars and photo-z, should be able to complement shallower spectroscopicsurveys in placing bounds to the growth of structure in the universe."
 More than thirty wears ago. Visvauathan Cuiersiuith (1977) found that earh-tvpe spiral ealaxies (SO0/a to Sab) exhibit) optical color-maenitude relation (CMB).. that is buniinous spirals show redder colors than less-bhuuinous SVSTCLUS.," More than thirty years ago, Visvanathan Griersmith (1977) found that early-type spiral galaxies (S0/a to Sab) exhibit optical color-magnitude relation (CMR), that is, luminous spirals show redder colors than less-luminous systems."
 Later. many studies confirmed this phenomena and extended the CMBs of spirals from optical to infrared bands (Tully ct al.," Later, many studies confirmed this phenomena and extended the CMRs of spirals from optical to infrared bands (Tully et al."
 1982: de Jong; 1996: Alobasher et al., 1982; de Jong 1996; Mobasher et al.
 1986: Tully et al., 1986; Tully et al.
 1998: Peleticr de Gays 1998: Bell de Joug 2000: Blanton tal., 1998; Peletier de Grijs 1998; Bell de Jong 2000; Blanton et al.
 2003: Baldry et al., 2003; Baldry et al.
 2001: Macarthur et al., 2004; Macarthur et al.
 2001)., 2004).
 Chang et al. (, Chang et al. (
2006a) collected a sample of late-type galaxies from Sloan Digital Sky Survey (SDSS) and Two-Micron Al-Sky Survey (2MASS) and found that. after correcting the effect of dust attenuation. the opticalinfrared (optical-IB) colors preseut tighter correlations with the absolute maguitucde than optical colors.,"2006a) collected a sample of late-type galaxies from Sloan Digital Sky Survey (SDSS) and Two-Micron All-Sky Survey (2MASS) and found that, after correcting the effect of dust attenuation, the optical-infrared (optical-IR) colors present tighter correlations with the absolute magnitude than optical colors."
 Based on a mnatcehed sample of Galaxy Evolution Explorer (GALEN) aud SDSS. Wryder et al. (," Based on a matched sample of Galaxy Evolution Explorer (GALEX) and SDSS, Wyder et al. ("
2007) and Schiminovich et al (2007) explored the ultraviolet-optical (UV-optical) CAIs for both disk aud bulee-doninated galaxies.,2007) and Schiminovich et al (2007) explored the ultraviolet-optical (UV-optical) CMRs for both disk and bulge-dominated galaxies.
 They also fouud that disk-domunated galaxies show a tight correlation between specific star formation rate (SF R/AL.) and ealactic stellar mass (AL.). which is denoted as the mass-SFR relation for convenience.," They also found that disk-dominated galaxies show a tight correlation between specific star formation rate $SFR/M_*$ ) and galactic stellar mass $M_*$ ), which is denoted as the mass-SFR relation for convenience."
 Ou the other laud. interpretation of these observed CMBs of disk galaxies is still au open question.," On the other hand, interpretation of these observed CMRs of disk galaxies is still an open question."
 vau deu Bosch (2002) and Bell ct al. (, van den Bosch (2002) and Bell et al. (
2003) found that. coutrary to the observed trends. their scmi-analvtic models (SAMS) predicted that faint spirals should be shehtly redder than bright spirals.,"2003) found that, contrary to the observed trends, their semi-analytic models (SAMs) predicted that faint spirals should be slightly redder than bright spirals."
 By introducing additional mechanics like active galactic nucleus (AGN) feedback. the more recent SAMS can predict positive slopes of the CARs of spirals. 16. lower mass systems being uer than massive svsteuis (6.8. I&aug et al.," By introducing additional mechanics like active galactic nucleus (AGN) feedback, the more recent SAMs can predict positive slopes of the CMRs of spirals, i.e., lower mass systems being bluer than massive systems (e.g. Kang et al."
 2005. Croton et al.," 2005, Croton et al."
 2006)., 2006).
 ILowever. the slopes of the oxedieted. CAIRs are still shallower than that of he observations (Ixaug ct al.," However, the slopes of the predicted CMRs are still shallower than that of the observations (Kang et al."
 2005)., 2005).
 This is mainly due to the fact that it is difficult for the SAMS to mild a scenario that permits low mass spiral to vc averagelv vounger than high mass svstem iu he hierarchical cosiogonies., This is mainly due to the fact that it is difficult for the SAMs to build a scenario that permits low mass spiral to be averagely younger than high mass system in the hierarchical cosmogonies.
 Paraietrized modelling. which adopts phenomenological descriptions for some complicated processes. is another fruitful tool to explore the ealactic formation aud evolution (Tinsley 1980).," Parametrized modelling, which adopts phenomenological descriptions for some complicated processes, is another fruitful tool to explore the galactic formation and evolution (Tinsley 1980)."
 In fact. parametric models have been successfully applied to investigate the formation and evolution of the Alilkv Wav disk (Chang et al.," In fact, parametric models have been successfully applied to investigate the formation and evolution of the Milky Way disk (Chang et al."
 1999: Chiappiui ct al., 1999; Chiappini et al.
 2001: Renda et al., 2001; Renda et al.
 2005: Naab Ostriker 2006: Fu et al., 2005; Naab Ostriker 2006; Fu et al.
 2009: Matteucci et al..," 2009; Matteucci et al.,"
 2009) aud other nearby disk ealaxies Boissicr 2000: Molla Diaz 2005: Dalcautou 2007: Yin et al., 2009) and other nearby disk galaxies Boissier 2000; Molla Diaz 2005; Dalcanton 2007; Yin et al.
 2009)., 2009).
 Besides these specific ealaxies. Doissier Prautzos (2000. hereafter BPOO) extended. them calibrated model on the Allkv Way disk fo study the chemical ane spectro-photometric evolution of disk galaxies and compared their model predictions with the observed CAIRs of de Jong (1996).," Besides these specific galaxies, Boissier Prantzos (2000, hereafter BP00) extended their calibrated model on the Milky Way disk to study the chemical and spectro-photometric evolution of disk galaxies and compared their model predictions with the observed CMRs of de Jong (1996)."
 They fou satisfactory agreement with local observations provided that the star formation time-scales of low-mass galaxies are longer than their higlhi-uass counterparts., They found satisfactory agreement with local observations provided that the star formation time-scales of low-mass galaxies are longer than their high-mass counterparts.
 However. thei model predictions have not been compared with observations a hieher redshift due to the limited observations ut that time.," However, their model predictions have not been compared with observations at higher redshift due to the limited observations at that time."
 Recently. several new scaling relations of disk oealaxies are statistically found aud huge amount of observed data are available. not ouly in the LLσα universe but also for intermediate andl —[umigh redshift ealaxies.," Recently, several new scaling relations of disk galaxies are statistically found and huge amount of observed data are available, not only in the local universe but also for intermediate and high redshift galaxies."
 Motivated bv the work : BPOO. we try to construct a new parametric," Motivated by the work of BP00, we try to construct a new parametric"
The kev result is that. for clusters with logCMa/M.)= and loe(Agesvr)21.80 (i.c.. for all but the two voungest subsamples). the CALF slopes are very close to unity (sce also 1103 for similar results for the high-mass wings of their age-limited CMESH)). which in this parameter space corresponcds to an initial CALF slope of à2— (see also the final column in Table 2. for a direct comparison).,"The key result is that, for clusters with $\log(M_{\rm cl}/{\rm
M}_\odot) \ge 3.0$ $\log({\rm Age/yr}) > 7.80$ (i.e., for all but the two youngest subsamples), the CMF slopes are very close to unity (see also H03 for similar results for the high-mass wings of their age-limited ), which in this parameter space corresponds to an initial CMF slope of $\alpha \simeq -2$ (see also the final column in Table \ref{massslopes.tab} for a direct comparison)."
 However. the voungest two subsamples exhibit significantly shallower slopes for the same mass range (see also Elmegreen Efremov 1997). while for log(A47M.)<3.0 the observed CMES turn down even more significantly below the linear slopes indicated.," However, the youngest two subsamples exhibit significantly shallower slopes for the same mass range (see also Elmegreen Efremov 1997), while for $\log(M_{\rm cl}/{\rm M}_\odot) < 3.0$ the observed CMFs turn down even more significantly below the linear slopes indicated."
 We note that the observed elfect goes well beyond that expected rom uncertainties in the data's completeness. limit: an uncertainty in the completeness limit of 0.5 mag translates o an uncertainty in mass of only z;Mlog(AM4GQ/M.)=0.2., We note that the observed effect goes well beyond that expected from uncertainties in the data's completeness limit: an uncertainty in the completeness limit of 0.5 mag translates to an uncertainty in mass of only $\Delta \log(M_{\rm cl}/{\rm M}_\odot) = 0.2$.
 As the Poissonian error bars indicate. this Dattening of he CALF slope is unlikely to be caused by. small-number statistics either. while cluster disruption or variable external idal elfects are expected to have had a more important ellect on the low-mass extreme ofolder clusters than on the vounger clusters for which we notice this ciscrepaney.," As the Poissonian error bars indicate, this flattening of the CMF slope is unlikely to be caused by small-number statistics either, while cluster disruption or variable external tidal effects are expected to have had a more important effect on the low-mass extreme of clusters than on the younger clusters for which we notice this discrepancy."
 Thus. his observation is Contrary to theoretical predictions if the initial. CME of the LMC cluster system were a power-law," Thus, this observation is contrary to theoretical predictions if the initial CMF of the LMC cluster system were a power-law"
 , 
enough to us to allow detailed model tests.,enough to us to allow detailed model tests.
 A significant portion of the more distant lenses nav produce [inite-source-size effects. providing constraints on the lens and/or the lensed source.," A significant portion of the more distant lenses may produce finite-source-size effects, providing constraints on the lens and/or the lensed source."
 On the other hand. the nearby lenses can be studied by a variety of complementary observations (hat can combine with measurements of the photometric event (o constrain the lens moclel ancl in some cases to obtain a measurement of the gravitational mass of the lens and/or of anv stellar companion it may orbit.," On the other hand, the nearby lenses can be studied by a variety of complementary observations that can combine with measurements of the photometric event to constrain the lens model and in some cases to obtain a measurement of the gravitational mass of the lens and/or of any stellar companion it may orbit."
 For each short-duration event. we would like to answer several kev questions.," For each short-duration event, we would like to answer several key questions."
 What is (he nature of the lens?, What is the nature of the lens?
 Is it à planet?, Is it a planet?
 If so. is it [ree-floating or bound to a star?," If so, is it free-floating or bound to a star?"
 Is the lens a brown dwarl?, Is the lens a brown dwarf?
 Is the lens a high-velocity object?, Is the lens a high-velocity object?
 If it is a high-velocity object. what is ib and [rom where did it come?," If it is a high-velocity object, what is it and from where did it come?"
 These questions can be answered by constraining the lens mass. testing for evidence that (he lens has a companion. and in some cases. through the direct detection of the lens system.," These questions can be answered by constraining the lens mass, testing for evidence that the lens has a companion, and in some cases, through the direct detection of the lens system."
 Observing teams are already sampling some events with the photometric sensitivity and cadence needed to allow meaningful model fits., Observing teams are already sampling some short-duration events with the photometric sensitivity and cadence needed to allow meaningful model fits.
 Especially il alerts are called (o encourage almost continuous monitoring. the light curves will quantilv the effects of blending. and identify those events with genuinely short Einstein diameter crossing times.," Especially if alerts are called to encourage almost continuous monitoring, the light curves will quantify the effects of blending, and identify those events with genuinely short Einstein diameter crossing times."
 Fits of the light curve will provide values of τι and b. the angle of closest approach. expressed in units of 65. Thev will also determine if there are signatures of [imite-source-size ellects or other deviations from the standard point-source/point-lens form.," Fits of the light curve will provide values of $\tau_E$ and $b,$ the angle of closest approach, expressed in units of $\theta_E.$ They will also determine if there are signatures of finite-source-size effects or other deviations from the standard point-source/point-lens form."
 When the angle of closest approach between (he source ancl lens. b@p. is comparable to the angular dimensions of the star. the light curve shape is altered (Witt Mao 1994: Lee et 2009 and references therein).," When the angle of closest approach between the source and lens, $b\, \theta_E$, is comparable to the angular dimensions of the star, the light curve shape is altered (Witt Mao 1994; Lee et 2009 and references therein)."
 Thus. when the Einstein angle is small. as is the case [or low-mass lenses al moderate to large distances. [inite-source-size effects are most likely to be detectable through their alteration of the light curve shape.," Thus, when the Einstein angle is small, as is the case for low-mass lenses at moderate to large distances, finite-source-size effects are most likely to be detectable through their alteration of the light curve shape."
 To set the scale. note that 85 [or an Earth-mass planet at 100 pe is 1.58xLO? aresec. only ~2.7 times larger (han (he angular size of a 102. star al 5 kpe.," To set the scale, note that $\theta_E$ for an Earth-mass planet at $100$ pc is $1.58 \times 10^{-5}$ arcsec, only $\sim 2.7$ times larger than the angular size of a $10\, R_\sun$ star at $8$ kpc."
 When. on the other hand. (he Einstein anele is large it may be possible to measure the astrometric effects of lensing.," When, on the other hand, the Einstein angle is large it may be possible to measure the astrometric effects of lensing."
" The Einstein angle of a neutron star. for example. is 0.01"" al D,=100 Under ideal circumstances. this is measurable. because the maximum shift in the source's centroid of light is nearly equal to the size of the Einstein rine."," The Einstein angle of a neutron star, for example, is $0.01^{\prime\prime}$ at $D_L = 100$ Under ideal circumstances, this is measurable, because the maximum shift in the source's centroid of light is nearly equal to the size of the Einstein ring."
 The maximum shift is attained when the lens-source separation is approximately equal to 85. The magnitude of the shift falls off as the inverse of the separation. (, The maximum shift is attained when the lens-source separation is approximately equal to $\theta_E.$ The magnitude of the shift falls off as the inverse of the separation. (
See Mivamoto Yoshii 1995: Hoe et 1995: Miralda-IEscude 1996; Dominik Sahu 2000: An Han 2002: Han 2002: Asada 2003: Han 2006.),See Miyamoto Yoshii 1995; Hog et 1995; Miralda-Escude 1996; Dominik Sahu 2000; An Han 2002; Han 2002; Asada 2003; Han 2006.)
 To detect the shift. we need good angular resolution. a," To detect the shift, we need good angular resolution, a"
although one has to be careful to make sure that louver<fay. the time taken to cover the entire sky available to PAST. which is given in terms of fon. by that is. the full 20000deg? will be covered in 5.5 davs for à system with 19 beams if one spent Ls per pointing.,"although one has to be careful to make sure that $t_{\rm survey}<t_{\rm sky}$, the time taken to cover the entire sky available to FAST, which is given in terms of $t_{\rm obs}$ by that is, the full $20000\,{\rm deg}^2$ will be covered in $5.5$ days for a system with 19 beams if one spent $1\,s$ per pointing."
 We remind the reader that we have assumed. that 1 day corresponds to LS hours of on-source integration timo., We remind the reader that we have assumed that 1 day corresponds to 18 hours of on-source integration time.
 Under the assumptions of this calculation. we deduce. [or example. that a survey of the whole sky available to FAST with a noise level of mis840yd. per beam is possible by performing 600« integrations in around 9 veers if ου=19 but that this would. take less than 2 vears with ng~100.," Under the assumptions of this calculation, we deduce, for example, that a survey of the whole sky available to FAST with a noise level of $\sigma_{\rm noise}\approx 40\,\mu{\rm Jy}$ per beam is possible by performing $600\,s$ integrations in around 9 years if $n_{\rm B}=19$, but that this would take less than 2 years with $n_{\rm B}\sim 100$ ."
" Such a survey would find around το10"" ealaxies with S/N=4 with an average redshift (2?20.15.", Such a survey would find around $7\times 10^{6}$ galaxies with $S/N=4$ with an average redshift $\langle z\rangle\approx 0.15$.
 In subsequent sections. we will improve the validitv. of the caleulation by modelling additional important effects.," In subsequent sections, we will improve the validity of the calculation by modelling additional important effects."
 Llowever. we shall see that this simple calculation holds true as a first order estimate.," However, we shall see that this simple calculation holds true as a first order estimate."
 Lt is already. clear that. if such a cosmologically important redshift survey is to be completed in an acceptable amount of time. a focal plane array with np100 is essential.," It is already clear that, if such a cosmologically important redshift survey is to be completed in an acceptable amount of time, a focal plane array with $n_{\rm B}\sim 100$ is essential."
 We can check«posteriori that the point source approximation that we have made ijs valid. by examining the limiting mass and expected number counts at very. low redshift., We can check that the point source approximation that we have made is valid by examining the limiting mass and expected number counts at very low redshift.
 We note first that the number counts tend to zero as 2 tendsto zero since the volume clement is xz., We note first that the number counts tend to zero as $z$ tendsto zero since the volume element is $\propto z^2$.
 Thus. even in the case of fan.=65. there are very few objects with z<0.02. the point at whieh Agi=10°AL. galaxies are resolved by FAST.," Thus, even in the case of $t_{\rm obs}=6\,s$, there are very few objects with $z<0.02$, the point at which $M_{\rm HI}=10^{10}\,M_{\odot}$ galaxies are resolved by FAST."
 Phe percentage of low redshift objects will decrease as £i increases., The percentage of low redshift objects will decrease as $t_{\rm obs}$ increases.
 We estimate that the percentage of galaxies that would. be significantly. resolved for fan.=65 is less than LO% and for £a>GOs the elfect will be negligible.," We estimate that the percentage of galaxies that would be significantly resolved for $t_{\rm obs}=6\,s$ is less than $10\%$ and for $t_{\rm obs}>60\,s$ the effect will be negligible."
 In the most extreme circumstances. one might also be sensitive to the effects of confusion even in a spectroscopic survey since the FAST beamsize is likely to be much Larger than the size of the galaxies we are trving to detect.," In the most extreme circumstances, one might also be sensitive to the effects of confusion even in a spectroscopic survey since the FAST beamsize is likely to be much larger than the size of the galaxies we are trying to detect."
" One can attempt to quantify the expected. amount of neutral hvdrogen in volume of the universe enclosed by a Gaussian beam with area Oy=πθτμι(1ος»}) and the fiducial linewidth AV,2200kms.+."," One can attempt to quantify the expected amount of neutral hydrogen in volume of the universe enclosed by a Gaussian beam with area $\Omega_{\rm B}=\pi\theta_{\rm FWHM}^2/(4\log2)$ and the fiducial linewidth $\Delta V_{\rm o}=200\,{\rm km}\,s^{-1}$."
 This is given by where conservatively we assume that ρηι(ς)Put pz)., This is given by where conservatively we assume that $\rho_{\rm HI}(z)=\rho_{\rm HI}(0)(1+z)^{3}$ .
 We take the local LEE density to be 107h! as found in the HIIPASS. survey (Zwaanal. 2005).. that is. pgi(0)z7.2LoopAL.Alpe7.," We take the local HI density to be $\Omega_{\rm HI}=2.6 \times 10^{-4}{h}^{-1}$ as found in the HIPASS survey \citep{hipass}, that is, $\rho_{\rm HI}(0)\approx 7.2\times 10^{7}\,h\,M_{\odot}\,{\rm Mpc}^{-3}$."
 The computed value of η ΕΕ compared to the mass limits for various values of £i in Fig. 2.., The computed value of $\langle M_{\rm HI}\rangle(z)$ is compared to the mass limits for various values of $t_{\rm obs}$ in Fig. \ref{fig:simple}.
 Ht is clear from this that the effects of confusion are only significant for large values of Fa: one might imagine that if (Alu)>Miu/3 then confusion could. lead. to an increase in the effective. noise making it more difficult to find ealaxies., It is clear from this that the effects of confusion are only significant for large values of $t_{\rm obs}$ ; one might imagine that if $\langle M_{\rm HI}\rangle > M_{\rm lim}/3$ then confusion could lead to an increase in the effective noise making it more difficult to find galaxies.
 From the above discussion it seems sensible toconsider values of 605«fon. 6000 in order to avoid any possible corrections to the statistics 5.due to resolving galaxiesat low >and confusion at high z.," From the above discussion it seems sensible toconsider values of $60\,s<t_{\rm obs}<6000\,s$ , in order to avoid any possible corrections to the statistics due to resolving galaxiesat low $z$ and confusion at high $z$ ."
The (wo shaded rectangles on Chis shell represent cxlinders in (hree-climensional space.,The two shaded rectangles on this shell represent cylinders in three-dimensional space.
 If the center of the absorber y lies in a cvlinder. it will be observed on the line of sight that passes through the center of that evlinder.," If the center of the absorber $\mathbf{y}$ lies in a cylinder, it will be observed on the line of sight that passes through the center of that cylinder."
" When (he absorber at x is αἱ the center. the contribution of the absorber αἱ y to the estimate of A needs to be corrected for boundary ellects. ffor the points in the shell 0D;,(x.|]y—x|) not probed by the lines of sight so that an absorber could be present there but not observable."," When the absorber at $\mathbf{x}$ is at the center, the contribution of the absorber at $\mathbf{y}$ to the estimate of $K$ needs to be corrected for boundary effects, for the points in the shell $\delta B_{du}(\mathbf{x},
|\mathbf{y}-\mathbf{x}|)$ not probed by the lines of sight so that an absorber could be present there but not observable."
 Using Ripleys method of edge correction. (he weight is the reciprocal of the probability of detecting an absorber. given (hat the absorber is present in the shell.," Using Ripley's method of edge correction, the weight is the reciprocal of the probability of detecting an absorber, given that the absorber is present in the shell."
" This weight is approximated by the ratio of the volume of the shell 9,;,(x.|y—x|) to the volume of the evlinders."," This weight is approximated by the ratio of the volume of the shell $\delta B_{du}(\mathbf{x},
|\mathbf{y}-\mathbf{x}|)$ to the volume of the cylinders."
 Specifically. the weight is given by where [yy is (he set of points in 0Box.|y—x|)1L. and & is the angle subtended by the line of sight that point p lies on. ancl the line joining point p aud absorber x.," Specifically, the weight is given by where $I_{\mathbf{x},\mathbf{y}}$ is the set of points in $\delta B_0(\mathbf{x},
|\mathbf{y}-\mathbf{x}|) \cap L$, and $\theta_\mathbf{p}$ is the angle subtended by the line of sight that point $\mathbf{p}$ lies on, and the line joining point $\mathbf{p}$ and absorber $\mathbf{x}$."
 In Figure μι [fyy just consists of the points y and b.," In Figure \ref{fig:pair}, $I_{\mathbf{x},\mathbf{y}}$ just consists of the points $\mathbf{y}$ and $\mathbf{b}$."
" The angles &, and 0, are also indicated in Figure l..", The angles $\theta_\mathbf{y}$ and $\theta_\mathbf{b}$ are also indicated in Figure \ref{fig:pair}.
 The variable d is the radius of an absorber in comoving units. and we assume it to be unknown. but fixed.," The variable $d$ is the radius of an absorber in comoving units, and we assume it to be unknown, but fixed."
 The value of d does not need to be specilied because it gets cancelled away aud does not appear in the estimator for A., The value of $d$ does not need to be specified because it gets cancelled away and does not appear in the estimator for $K$.
 For more details. see Section3. where (his cancellation also occurs for the estimator of the third-order statistic A.," For more details, see Section\ref{sect:thirdmoment} where this cancellation also occurs for the estimator of the third-order statistic $\mathcal{K}$."
 With weights specified by (1)). Lohetal.(2001) show that estimates for ΑΙ are unbiased.," With weights specified by \ref{eqn:2ndorderwt}) ), \citet{loh01} show that estimates for $\lambda^2 K$ are unbiased."
 It is well-known that second-order statistics do not completely describe (he clustering properties of point processes., It is well-known that second-order statistics do not completely describe the clustering properties of point processes.
 For example. Daddeley&Silverman(1934). provide an example of a non-Poisson process with the A funcüon identical to that of a homogeneous Poisson process.," For example, \citet{baddeley84} provide an example of a non-Poisson process with the $K$ function identical to that of a homogeneous Poisson process."
" Third- ancl higher-order statistics will allow a more detailed study of clustering (han just second-order statistics,", Third- and higher-order statistics will allow a more detailed study of clustering than just second-order statistics.
 Peebles&Groth(1975). defined a three-point correlation function ¢ ancl applied it to the Zwicky catalog., \citet{peebles75} defined a three-point correlation function $\zeta$ and applied it to the Zwicky catalog.
 With three volume elements dVi.νο and «γη that defime a triangle with lengths rp.re; ancl r43. Peebles&Groth(1975) wrote the probability of finding an object in each of these elements as," With three volume elements $dV_1,
dV_2$ and $dV_3$ that define a triangle with lengths $r_{12}, r_{23}$ and $r_{13}$ , \citet{peebles75} wrote the probability of finding an object in each of these elements as"
where we used the result of (he pressure balance across the laver in the last step.,where we used the result of the pressure balance across the layer in the last step.
" This means that the outflow velocity in the SweetParker reconnection is the Alfvénn velocity computed with the upstream reconnecting magnetic field 2, and the density p (which. we remind the reader. is considered to be uniform in the SweetParker mocel)."," This means that the outflow velocity in the Sweet–Parker reconnection is the Alfvénn velocity computed with the upstream reconnecting magnetic field $B_0$ and the density $\rho$ (which, we remind the reader, is considered to be uniform in the Sweet–Parker model)."
 In our situation. however. this is not the case and the tension force turns out to be much larger than (he pressure gradient force.," In our situation, however, this is not the case and the tension force turns out to be much larger than the pressure gradient force."
 Indeed. using half the value of reconnecting magnetic field By at the exhaust point [see eq. (2.2)))," Indeed, using half the value of reconnecting magnetic field $B_1$ at the exhaust point [see eq. \ref{eq-B_1}) )]"
" to estimate the characteristic reconnected magnetic field inside the laver. we see that the tension force —7.D,/cDg/8x Lis∙ by a factor. A1 greater than the pressure gradient- lorce. 0P/Ox2D~Dg/(8zL).>P"," to estimate the characteristic reconnected magnetic field inside the layer, we see that the tension force $-j_z B_y/c \sim (B_0/4\pi\delta) (B_0 A \delta/2L) = A\, B_0^2/{8\pi L}$ is by a factor $A\gg 1$ greater than the pressure gradient force $\partial P/\partial x \sim {B_0^2}/{(8\pi L)}$."
" By equating this tension force with the characteristic magnitude of the inertial term. pe,Ov,~pu/2L. 36 immediately follows that the outflow velocity at the end of the laver is comparable to the Alfvénn velocity (involving the density and the magnetic field just outside of the laver): novVio. Note that this is different [rom the usual assumption (which. we claim. is not justified) that the outflow velocity in the compressible case is that corresponding to the Alfvénn velocity delined with the compressed clensity inside the laver ancl the magnetic field outside the laver (?) a result that one would have obtained by taking into account the pressure eradient [orce only."," By equating this tension force with the characteristic magnitude of the inertial term, $\rho v_x \partial_x v_x \sim \rho u^2/2L$, it immediately follows that the outflow velocity at the end of the layer is comparable to the Alfvénn velocity (involving the density and the magnetic field just outside of the layer): u. Note that this is different from the usual assumption (which, we claim, is not justified) that the outflow velocity in the compressible case is that corresponding to the Alfvénn velocity defined with the compressed density inside the layer and the magnetic field outside the layer \citep{Parker-1963} — a result that one would have obtained by taking into account the pressure gradient force only."
 Once the compression ratio ο is found. solving the rest of the problem (i.e.. finding ihe reconnection rate) is rather stvaieht-lorwarcl. especially since the remaining equations that we need to complete this task [namely. eqs. (2.2))-(2.2))," Once the compression ratio $A$ is found, solving the rest of the problem (i.e., finding the reconnection rate) is rather straight-forward, especially since the remaining equations that we need to complete this task [namely, eqs. \ref{eq-mass-conserv}) \ref{eq-Ohm}) )"
 and (2.2))] are the same as in the classicalSweet-Parker analvsis., and \ref{eq-u}) )] are the same as in the classicalSweet-Parker analysis.
 First. combining Ohms law (2.2)) with the continuity equation (2.2)). we gel (eLcAud©uly/tye) A.," First, combining Ohm's law \ref{eq-Ohm}) ) with the continuity equation \ref{eq-mass-conserv}) ), we get $ v_{\rm rec} L \sim A u \delta \simeq u\, (\eta/v_{\rm rec})\, A $ ."
 Then. substituting &~Vay from equation (2.2)). we immediately. arrive al," Then, substituting $u\sim V_{A0}$ from equation \ref{eq-u}) ), we immediately arrive at"
,theoremTheorem[section] \begin{document}
The rate listed in the UMIST Data Base. BRate06. Dor the radiative association of C and 11 to give CIL is LO! em's 4.,"The rate listed in the UMIST Data Base, Rate06, for the radiative association of C and H to give CH is $10^{-17}$ $^{3}$ $^{-1}$."
 Such a low value can hardly start a chain of reactions leading ultimately to the formation of enough carbonaceous grains to fit observations., Such a low value can hardly start a chain of reactions leading ultimately to the formation of enough carbonaceous grains to fit observations.
 For want οἱ other experimental data. we resorted to chemical simulation.," For want of other experimental data, we resorted to chemical simulation."
 This consists in using molecular dvnamies to represent the collision of C and II atoms. for different relative velocities ancl impact parameters. (o vield the cross-section for association as a [Iunction of relative velocity.," This consists in using molecular dynamics to represent the collision of C and H atoms, for different relative velocities and impact parameters, to yield the cross-section for association as a function of relative velocity."
 At the end of a run. either the two atoms flv away in opposite directions. or the CII molecule is formed. with internal and external energies such that the initial energy and momentum are conserved.," At the end of a run, either the two atoms fly away in opposite directions, or the CH molecule is formed, with internal and external energies such that the initial energy and momentum are conserved."
 An order of magnitude cross-section is then deduced from these results: combined wilh a reference temperature (and hence a velocity of impact). this vields the rate. k. A deseription of the software package used [or simulation can be found in Papoular (2001).," An order of magnitude cross-section is then deduced from these results; combined with a reference temperature (and hence a velocity of impact), this yields the rate, k. A description of the software package used for simulation can be found in Papoular \cite{pap01}."
. The particular eode used here is AMI (Austin Model 1). as proposed. and optimized for carbou-vich molecules. by Dewar et al. (1985).," The particular code used here is AM1 (Austin Model 1), as proposed, and optimized for carbon-rich molecules, by Dewar et al. \cite{dew85}."
. This semi-empirical method combines a rigorous equantum-mechanical formalism wilh empirical parameters obtained [rom comparison wilh experimental results., This semi-empirical method combines a rigorous quantum-mechanical formalism with empirical parameters obtained from comparison with experimental results.
 It computes approximate solutions of Schroedinger's equation for each instantaneous spatial distribution of atomic nuclei. and uses experimental cata only when (he Q.M. calculations are too dillieult or too lengthy.," It computes approximate solutions of Schroedinger's equation for each instantaneous spatial distribution of atomic nuclei, and uses experimental data only when the Q.M. calculations are too difficult or too lengthy."
 This makes it more accurate (han poor ab initio methods. and faster than any of the latter.," This makes it more accurate than poor ab initio methods, and faster than any of the latter."
 The molecular dynamics relies on the Born-Oppenheimer approximation to determine the motions of atoms under nuclear and electronic forces due to their environment., The molecular dynamics relies on the Born-Oppenheimer approximation to determine the motions of atoms under nuclear and electronic forces due to their environment.
 At every step. all the svstem parameters are menmorized as snapshots so (hat. after completion of the run. a movie of the reactions can be viewed on (he screen. and (he trajectory of any. atom followed all along.," At every step, all the system parameters are memorized as snapshots so that, after completion of the run, a movie of the reactions can be viewed on the screen, and the trajectory of any atom followed all along."
 This makes for a better understanding of the details of mechanisms ancl outcomes., This makes for a better understanding of the details of mechanisms and outcomes.
 In particular. it becomes possible to measure the activation energy. if any.," In particular, it becomes possible to measure the activation energy, if any."
 Note that the dvnamies of bond dissociations ancl formation can only be simulated by using Unrestricted Hartree-Fock (UIIF) wave functions in the Q.M. part of the calculation (see Szabo and OstIund (1989)))., Note that the dynamics of bond dissociations and formation can only be simulated by using Unrestricted Hartree-Fock (UHF) wave functions in the Q.M. part of the calculation (see Szabo and Ostlund \cite{sza89}) ).
" The elementary caleulation step was set at. 10.1 s, senmi-enipirical simulation methods do not account for purely eqaantum-nmechanical phenomena like tunneling or zero-point energy.", The elementary calculation step was set at $10^{-16}$ s. Semi-empirical simulation methods do not account for purely quantum-mechanical phenomena like tunneling or zero-point energy.
 In (he present context of relatively high temperatures. (he former is irrelevant. and the latter hardly affects the probability ancl energetics of the physico-chemical reactions involved.," In the present context of relatively high temperatures, the former is irrelevant, and the latter hardly affects the probability and energetics of the physico-chemical reactions involved."
 The simulation does not treat absorption or emission, The simulation does not treat absorption or emission
ivdrogen at large distances above the photosphere (the hydrogen ionization state is frozen) and the delayed recombination in the optically thick layers. at and below the photosphere (the ionization ront is located further out).,"hydrogen at large distances above the photosphere (the hydrogen ionization state is frozen) and the delayed recombination in the optically thick layers, at and below the photosphere (the ionization front is located further out)."
 All these features apply to the helium ionization fraction as well. which we show in the right panel of Fig. 2..," All these features apply to the helium ionization fraction as well, which we show in the right panel of Fig. \ref{fig_h_he}."
 The population is [5 orders of magnitude lower han that of at intermediate heights. but it is comparable or stronger at large distances.," The population is 15 orders of magnitude lower than that of at intermediate heights, but it is comparable or stronger at large distances."
 Effects of time dependence can hus reach considerable levels. much beyond modulations that can be achieved through reasonable changes of (or uncertainties in) chemical abundance or ionizing fluxes.," Effects of time dependence can thus reach considerable levels, much beyond modulations that can be achieved through reasonable changes of (or uncertainties in) chemical abundance or ionizing fluxes."
 Because meal line-blanketing plays an increasingly important role at and aftertye recombination epoch. time-dependence effects on metals are also of special interest.," Because metal line-blanketing plays an increasingly important role at and after the recombination epoch, time-dependence effects on metals are also of special interest."
 We show in Fig., We show in Fig.
" 3. the ionization state o ""all species for model A on day 48.7.", \ref{fig_ion_state} the ionization state of all species for model A on day 48.7.
 Notice how the frozen-in ionization is readily visible at large distances above the photosjohere. with order of magnitude differences in ionization. level between the time-dependent and steady-state approaches.," Notice how the frozen-in ionization is readily visible at large distances above the photosphere, with order of magnitude differences in ionization level between the time-dependent and steady-state approaches."
 Ποιοe. the inclusion of time-dependence affects all species. irrespective. for example. of the atomic properties of the given ion. and wih a magnitude that is greatest at large distances above the photospjere Where the recombination timescale becomes increasingly large compared to the expansion timescale.," Hence, the inclusion of time-dependence affects all species, irrespective, for example, of the atomic properties of the given ion, and with a magnitude that is greatest at large distances above the photosphere where the recombination timescale becomes increasingly large compared to the expansion timescale."
 Finally. we show in Fig.," Finally, we show in Fig."
 4 the radial variation of the temperature and the electron density for model A at 48.7 days after explosion., \ref{fig_t_ed} the radial variation of the temperature and the electron density for model A at 48.7 days after explosion.
 The temperature in the time-dependent model is higher in the optically-thick regions. due to the reduction of density-sensitive line emission processes that dominate radiative cooling.," The temperature in the time-dependent model is higher in the optically-thick regions, due to the reduction of density-sensitive line emission processes that dominate radiative cooling."
 In the optically-thin regions. above the photosphere. the temperature is. however. quite similar between the two approaches.," In the optically-thin regions, above the photosphere, the temperature is, however, quite similar between the two approaches."
 This result stems from the poor coupling between gas and radiation in these scattering-dominated. tenuous. ejecta.," This result stems from the poor coupling between gas and radiation in these scattering-dominated, tenuous, ejecta."
 Departures from LTE are pronounced and the temperature is a secondary driver of the ejecta ionization. even in steady-state contigurations.," Departures from LTE are pronounced and the temperature is a secondary driver of the ejecta ionization, even in steady-state configurations."
 Models B-C-D offer valuable information on the sensitivity of the above effects on the ejecta velocity (models B and C) and the epoch (model Dy., Models B-C-D offer valuable information on the sensitivity of the above effects on the ejecta velocity (models B and C) and the epoch (model D).
 We find that varying the velocity of the ejecta within the range of values for slow/underluminous (model B: SN 1999br) and more standard luminosity Type II SN events (SN 1999em) has no sizeable effect on the main features of time dependence highlighted above., We find that varying the velocity of the ejecta within the range of values for slow/underluminous (model B; SN 1999br) and more standard luminosity Type II SN events (SN 1999em) has no sizeable effect on the main features of time dependence highlighted above.
 Recombination is inhibited. and a Significant over-ionization persists far above the photosphere.," Recombination is inhibited, and a significant over-ionization persists far above the photosphere."
 In Fig. 5..," In Fig. \ref{fig_ion_state_early},"
 we present the mean ionic charge for the six species (Π. He. ο. N. O. and Fe) included in model D tearly epoch). and at 6.7 days after explosion.," we present the mean ionic charge for the six species (H, He, C, N, O, and Fe) included in model D (early epoch), and at 6.7 days after explosion."
 While hydrogen remains fully ionized. in both time-dependent and steady-state approaches. all other species show a higher ionization at large distances above the photosphere: no difference is noticeable in the vicinity and below the photosphere. contrary to model A at 48.7 days.," While hydrogen remains fully ionized, in both time-dependent and steady-state approaches, all other species show a higher ionization at large distances above the photosphere; no difference is noticeable in the vicinity and below the photosphere, contrary to model A at 48.7 days."
 This supports the idea that time-dependence effects operate even at early times., This supports the idea that time-dependence effects operate even at early times.
 They do so at large distances because there. the recombination and expansion timescales are already sufficiently close. but they can also act at the photosphere because of optical depth effects and meta-stable levels (see refsectdfau)).," They do so at large distances because there, the recombination and expansion timescales are already sufficiently close, but they can also act at the photosphere because of optical depth effects and meta-stable levels (see \\ref{sect_ddt_tau}) )."
Halsoshowsthalover ionizalionallargedistancesdoesnolstembyessencef romlherecombinalionen densityfasterpansionofthes Nejectla., It also shows that over-ionization at large distances does not stem by essence from the recombination energy (since most of that energy is stored in hydrogen and hydrogen remains fully ionized in this whole sequence) but results exclusively from the low-density fast-expansion of the SN ejecta.
Acorollargisthallhisprogressive7or (Mitchell et al.," A corollary is that this progressive ``over-ionization'' does not require an extra source of ionization and excitation from, for example, $^{56}$ Ni (Mitchell et al."
 2001: Baron et al., 2001; Baron et al.
 2003). and operates instead. in a robust fashion. through the reduction of the effective recombination rates.," 2003), and operates instead, in a robust fashion, through the reduction of the effective recombination rates."
 Finally. this supports the idea that similar effects should apply in SN ejectageneral. since they generically possess similar low density and large velocities.," Finally, this supports the idea that similar effects should apply in SN ejecta, since they generically possess similar low density and large velocities."
 We concludethis section by presenting the evolution of the electron-scattering optical depth at the base radius Ay of each model in Fig., We concludethis section by presenting the evolution of the electron-scattering optical depth at the base radius $R_0$ of each model in Fig.
 6 (see Table 1. for model parameters).," \ref{fig_tau_es}
 (see Table \ref{table_model_param} for model parameters)."
 Rather than xiving a 1//7 dependence (dashed lines) for an expanding ejecta with a fixed mass absorption coefficient (as in model D where the dashed and solid lines overlap: fixed composition and ionization). he curves for models A. B. and C show a change of concavity. with a slow decrease at first and a turn-over to a faster decrease after 3— weeks.," Rather than having a $1/t^2$ dependence (dashed lines) for an expanding ejecta with a fixed mass absorption coefficient (as in model D where the dashed and solid lines overlap; fixed composition and ionization), the curves for models A, B, and C show a change of concavity, with a slow decrease at first and a turn-over to a faster decrease after 3--6 weeks."
 Recombination. which is delayed at first and permits the vrsistence of a high free-electron density in the ejecta. does in he longer term occur and triggers a fast drop of the base electron- optical depth.," Recombination, which is delayed at first and permits the persistence of a high free-electron density in the ejecta, does in the longer term occur and triggers a fast drop of the base electron-scattering optical depth."
 It remains. however. higher by up to a actor of about ten. than the value obtained in the corresponding steady-state model (for the same adopted lower bound radius in the optical-depth integral) at the ultimate time in the sequence. shown as filled colored circles.," It remains, however, higher by up to a factor of about ten, than the value obtained in the corresponding steady-state model (for the same adopted lower bound radius in the optical-depth integral) at the ultimate time in the sequence, shown as filled colored circles."
 Model D. whose ejecta contains species at least once ionized throughout the sequence. has a base optical depth of 51.9 in the time-dependent approach and 51.7 in the steady-state approach. the modest difference stemming from the slight over-ionization of sub-dominant species like helium and CNO elements. (," Model D, whose ejecta contains species at least once ionized throughout the sequence, has a base optical depth of 51.9 in the time-dependent approach and 51.7 in the steady-state approach, the modest difference stemming from the slight over-ionization of sub-dominant species like helium and CNO elements. ("
See the last columns of Table |. for a log of these values.),See the last columns of Table \ref{table_model_param} for a log of these values.)
 The various investigations and analyses discussed above indicate that time-dependence effects are robust. applyto all species. even prior to the recombination epoch of hydrogen.," The various investigations and analyses discussed above indicate that time-dependence effects are robust, applyto all species, even prior to the recombination epoch of hydrogen."
" Hence. for example. during the first week after explosion. such effects inhibit the recombination from Fe’ to Fe? and Fe?. or that of to He""."," Hence, for example, during the first week after explosion, such effects inhibit the recombination from $^{4+}$ to $^{3+}$ and $^{2+}$, or that of $^{+}$ to $^{0+}$."
 Different species block light at specific wavelengths. with often a tendency to bunch-up in specific spectral regions. but with no correlation. and therefore we anticipate that such modulations in the ionization state of the ejecta might lead to a Significant impact on the properties of the emergent light.," Different species block light at specific wavelengths, with often a tendency to bunch-up in specific spectral regions, but with no correlation, and therefore we anticipate that such modulations in the ionization state of the ejecta might lead to a significant impact on the properties of the emergent light."
 Both the optical. dominated by and lines. and the UV. dominated by metal line-blanketing due to. for example. iron. should be affected.," Both the optical, dominated by and lines, and the UV, dominated by metal line-blanketing due to, for example, iron, should be affected."
 In Fig. 7..," In Fig. \ref{fig_spec_modelA}, ,"
 we present a montage of synthetic spectra showing, we present a montage of synthetic spectra showing
as cGMpyBG]B5. leading to a semi-major axis only much smaller (han ecu.,"as $\sim GM_{BH}R_*/R_p^2$, leading to a semi-major axis only much smaller than $a_{\rm mean}$."
 Didal torques during the initial stage of the encounter could also lead to à somewhat larger binding energv lor the most bound material (Rees1983):: in terms of the discussion in this paper. the effects are indistinguishable from those due to a decrease in 53.," Tidal torques during the initial stage of the encounter could also lead to a somewhat larger binding energy for the most bound material \citep{rees88}; in terms of the discussion in this paper, the effects are indistinguishable from those due to a decrease in $\beta$."
 On (he assumption that the disrupted star has a unilorm distribution in orbital binding energy per wil mass. matter returns to the region ~Zi Hom the black hole at a rate dMf/dlx(1/09).7 (Phinney1989).," On the assumption that the disrupted star has a uniform distribution in orbital binding energy per unit mass, matter returns to the region $\sim R_p$ from the black hole at a rate $dM/dt \propto (t/t_0)^{-5/3}$ \citep{phinney89}."
. The characteristic timescale /;4 for initiation of this power-law accretion rate is the orbital period for the most bound matter the period for an orbit with semi-major axis (jg., The characteristic timescale $t_0$ for initiation of this power-law accretion rate is the orbital period for the most bound matter the period for an orbit with semi-major axis $a_{\rm min}$.
 More detailed calculations (e.g.. Lodatoοἱal. (2009))) indicate (hat. almost independent of the stars internal structure. the mass return rate does follow this behavior at late times. but at early times the rate rises rapidly after a delay ~{αμ} and then gradually rolls over to the /.5 proportionality. transitioning more slowly when the stars pre-disruption structure is more centrally concentrated.," More detailed calculations (e.g., \cite{lkp09}) ) indicate that, almost independent of the star's internal structure, the mass return rate does follow this behavior at late times, but at early times the rate rises rapidly after a delay $\sim P_{\rm orb}(a_{\rm min})$ and then gradually rolls over to the $t^{-5/3}$ proportionality, transitioning more slowly when the star's pre-disruption structure is more centrally concentrated."
" In this picture. once the matter returns to the vicinity of 2, itis captured into an accretion disk whose inflow time /j,«Du(eg)."," In this picture, once the matter returns to the vicinity of $R_p$ it is captured into an accretion disk whose inflow time $t_{\rm in} \ll P_{\rm orb}(a_{\rm min})$."
 As the matter moves inward through this disk. there is local dissipation of the conventional accretion disk variety. aud (he heat is radiated in the usual quasi-thermal fashion.," As the matter moves inward through this disk, there is local dissipation of the conventional accretion disk variety, and the heat is radiated in the usual quasi-thermal fashion."
 At the peak accretion rate estimated by the luminosity would be, At the peak accretion rate estimated by \cite{lkp09} the luminosity would be
domains. aud the cut-off energv measured usingFermi. can be explainecl as a cooling break and an acceleration limit. respectivelv.,"domains, and the cut-off energy measured using, can be explained as a cooling break and an acceleration limit, respectively."
 A prediction of this model is an extended emission region and hence non-variability of the signal., A prediction of this model is an extended emission region and hence non-variability of the signal.
 The non-variabilitv of the source in the 5-rav regime is supported by the observations., The non-variability of the source in the -ray regime is supported by the observations.
 While and do not provide sullicient angular resolution to resolve the blast wave (see Fig. 1)).," While and do not provide sufficient angular resolution to resolve the blast wave (see Fig. \ref{fig:etacar_sketch}) ),"
 high resolution radio observations will be able to resolve non-thermal radio emission [rom this acceleration site., high resolution radio observations will be able to resolve non-thermal radio emission from this acceleration site.
 Detection of a radio shell at. the location of the shock would support the proposed scenario., Detection of a radio shell at the location of the shock would support the proposed scenario.
 The time elapsed since the ejection of the very fast material by iis comparable to the age of the voungest Galactic SNR 11.920.3 and between the ages ol the remnants of 11987À and AA. It may provide important insights into particle acceleration and production of non-thermal radiation during (he early stages of remnant evolution where only very. lew objects can be studied., The time elapsed since the ejection of the very fast material by is comparable to the age of the youngest Galactic SNR 1.9+0.3 and between the ages of the remnants of 1987A and A. It may provide important insights into particle acceleration and production of non-thermal radiation during the early stages of remnant evolution where only very few objects can be studied.
 We would like to (hank (he anonymous referee for very helpful and informative comments which improved the quality of the paper., We would like to thank the anonymous referee for very helpful and informative comments which improved the quality of the paper.
 We would also like to thank Felix Aharonian. Werner Tohnann. Olaf Reimer and IxIaus Reitherger for (heir input ancl suggestions.," We would also like to thank Felix Aharonian, Werner Hofmann, Olaf Reimer and Klaus Reitberger for their input and suggestions."
 Furthermore. the support of the authors host institutions. anc additionalle support [rom the German Ministry of Eclucation and Research (BAIBF) is acknowledged.," Furthermore, the support of the authors host institutions, and additionally support from the German Ministry of Education and Research (BMBF) is acknowledged."
galactic dise fields in the constellation Carina. which are the only dise fields which have been observed for more than 400 epochs. all taken between January 1997 and July 2000.,"galactic disc fields in the constellation Carina, which are the only disc fields which have been observed for more than 400 epochs, all taken between January 1997 and July 2000."
 The relevant field parameters are given in Table1.., The relevant field parameters are given in Table\ref{fields}.
 For our analysis we used the photometric data obtained from the Difference Image Analysis (DIA: Alard Lupton 1998: Alard 2000)., For our analysis we used the photometric data obtained from the Difference Image Analysis (DIA; Alard Lupton 1998; Alard 2000).
 We concentrated on those stars that have good data for more thar of the epochs. and downloaded all stars with a mean magnitude of /<17.5 (52703 stars in the three fields).," We concentrated on those stars that have good data for more than of the epochs, and downloaded all stars with a mean magnitude of $I<17.5$ (52703 stars in the three fields)."
 In a later stage. our transit analysis only focuses on those stars with 13.0«I16.0 (15711 objects).," In a later stage, our transit analysis only focuses on those stars with $13.0<I<16.0$ (15711 objects)."
 Sys-Rem. a detrending algorithm designed to. remove systematic effects in a large set of light curves (Tamuz. Mazeh Zucker 2005). was used on the data from the three fields.," Sys-Rem, a detrending algorithm designed to remove systematic effects in a large set of light curves (Tamuz, Mazeh Zucker 2005), was used on the data from the three fields."
 The algorithm can detect any systematic. effect that appears linearly in many light curves obtained by the survey. without any prior knowledge of the origin of the effects.," The algorithm can detect any systematic effect that appears linearly in many light curves obtained by the survey, without any prior knowledge of the origin of the effects."
 It has become a standard tool in transit survey light curve processing (eg., It has become a standard tool in transit survey light curve processing (eg.
 Bakos et al., Bakos et al.
 2006: Collier Cameron et al., 2006; Collier Cameron et al.
 2006)., 2006).
 In Figure the scatter in the data is shown as a function of the mean stellar magnitude. for all stars in the CAR.SSCI field.," In Figure \ref{dispersion} the scatter in the data is shown as a function of the mean stellar magnitude, for all stars in the SC1 field."
 The light-grey symbols indicate the dispersion before the use of Sys-Rem. with the black erosses showing the average in each 0.1 magnitude bin.," The light-grey symbols indicate the dispersion before the use of Sys-Rem, with the black crosses showing the average in each 0.1 magnitude bin."
 The dark grey symbols indicate the same but after 10 cycles of Sys-Rem. with the black squares the resulting binned average dispersion.," The dark grey symbols indicate the same but after 10 cycles of Sys-Rem, with the black squares the resulting binned average dispersion."
 The algorithm is more efficient for brighter stars (for which the photometric errors are more dominated by systematic effects)., The algorithm is more efficient for brighter stars (for which the photometric errors are more dominated by systematic effects).
 For the brightest stars (4« 15) the mean dispersion has decreased from before Sys-Rem to after the use of Sys-Rem. showing the strength of this algorithm.," For the brightest stars $I<15$ ) the mean dispersion has decreased from before Sys-Rem to after the use of Sys-Rem, showing the strength of this algorithm."
 Several of the transits presented 1 section 3 would not have been found without the use of , Several of the transits presented in section 3 would not have been found without the use of Sys-Rem.
The final mean error as a function of magnitude is fittec with a three parameter function indicated by the solid black line. containing a residual noise term (0.22%)). and Poisso noise terms from the star and the background sky.," The final mean error as a function of magnitude is fitted with a three parameter function indicated by the solid black line, containing a residual noise term ), and Poisson noise terms from the star and the background sky."
 The few brightest stars seem to suffer from non-linearity or saturatio effects. resulting in an increase of the mean dispersion towards the brightest 7«13 objects.," The few brightest stars seem to suffer from non-linearity or saturation effects, resulting in an increase of the mean dispersion towards the brightest $I<13$ objects."
 After ten cycles of detrending. a box-fitting algorithm was used to search the light curves for periodic transits. based on the Box Least Squares (BLS) method (Kovaes et al.," After ten cycles of detrending, a box-fitting algorithm was used to search the light curves for periodic transits, based on the Box Least Squares (BLS) method (Kovacs et al."
 2002)., 2002).
 Transit periods between 0.1 and 10.0 days were searched for using a detection threshold of δσ., Transit periods between 0.1 and 10.0 days were searched for using a detection threshold of $\sigma$ .
 An estimate of the, An estimate of the
The detail study of this possibility will be discussed elsewhere.,The detail study of this possibility will be discussed elsewhere.
 For the high EBL flux a verv small photon index of D0.2 is required.," For the high EBL flux a very small photon index of $\Gamma_{\rm
 int}\simeq0.2$ is required."
 Remarkably. even such an unusual photon spectrum can be explained by internal absorption with a hieher target photon temperature and slightly larger optical depth (see Fig. 4..," Remarkably, even such an unusual photon spectrum can be explained by internal absorption with a higher target photon temperature and slightly larger optical depth (see Fig. \ref{fig:3c66a},"
 Fit 4)., Fit 4).
 With a certain combination of model parameters. the flux of the svnchrotron radiation [from secondary. electrons can match the levels detected in the optical band. as is demonstrated in Fig. 5..," With a certain combination of model parameters, the flux of the synchrotron radiation from secondary electrons can match the levels detected in the optical band, as is demonstrated in Fig. \ref{fig:3c66amwl},"
 Fits fa aud th., Fits 4a and 4b.
 For a small production region. the main fraction of (he secondary pairs are produced outside the blob.," For a small production region, the main fraction of the secondary pairs are produced outside the blob."
 Their radiation is not Doppler boosted and. therefore cannot be detected.," Their radiation is not Doppler boosted and, therefore cannot be detected."
 For the pairs produced inside (he blob. the secondary. svuchrotron radiation is Doppler boosted and thus it can contribute significantly to the observed [Iuxes.," For the pairs produced inside the blob, the secondary synchrotron radiation is Doppler boosted and thus it can contribute significantly to the observed fluxes."
 We note. however. that for this source we did not succeed to find a combination of parameters which could explain both the optical anc X-ray f[Iuxes together. bv svnchrotron raciation of secondary. electrons.," We note, however, that for this source we did not succeed to find a combination of parameters which could explain both the optical and X-ray fluxes together, by synchrotron radiation of secondary electrons."
 Since the internal absorption scenario requires a significant attenuation of the VIIE radiation over approximately (wo decades (see Figs., Since the internal absorption scenario requires a significant attenuation of the VHE radiation over approximately two decades (see Figs.
 3. and 5)). the secondary svnchrotroin component has to be at least 4 decades broad (wilh additional broadening related to the relativistic motion of the production region).," \ref{fig:0229}
 and \ref{fig:3c66amwl}) ), the secondary synchrotron component has to be at least 4 decades broad (with additional broadening related to the relativistic motion of the production region)."
 The strong magnetic [ield required in (he proton svnchrotron model provides [ast cooling of the pairs. thus the radiation spectrum will be featureless. without a cooling break.," The strong magnetic field required in the proton synchrotron model provides fast cooling of the pairs, thus the radiation spectrum will be featureless, without a cooling break."
 Ii the case of a small radius of the production region. the elective particle injection in the blob max be rather narrow.," In the case of a small radius of the production region, the effective particle injection in the blob may be rather narrow."
 Bit in this case the flux level would be significantly below the observed flux., But in this case the flux level would be significantly below the observed flux.
IX. were invisible to us.,"IX, were invisible to us."
 On the other hand. we did see many resolving objects near the Milky Way.," On the other hand, we did see many resolving objects near the Milky Way."
 The requirement is for some perceptible increase in surface brightness (possibly in the form of a stellar overdensitv) above the surrounding field., The requirement is for some perceptible increase in surface brightness (possibly in the form of a stellar overdensity) above the surrounding field.
 We note in particular. (hat compact star clusters of the tvpe recently reported by IIuxoretal.(2005).. if Chere were anv significant number associated with the \lilky Way. would have been seen and recorded.," We note in particular that compact star clusters of the type recently reported by \citet{H05}, if there were any significant number associated with the Milky Way, would have been seen and recorded."
 In that way and in others the satellite galaxy. svstems of Andromeda and the Milky Way must differ intrinsically. since the observed differences cannot now be attributed to selection effects.," In that way and in others the satellite galaxy systems of Andromeda and the Milky Way must differ intrinsically, since the observed differences cannot now be attributed to selection effects."
 For many vears the discrepancy between u-body simulations. which predict hundreds of ealaxies in structures like the Local Group. and the much smaller number of known galaxies has been noted.," For many years the discrepancy between n-body simulations, which predict hundreds of galaxies in structures like the Local Group, and the much smaller number of known galaxies has been noted."
" The simulations of Klvpinetal.(1999).. for instance. predict 300 satellite ealaxies within 1.5 Alpe of the center of a Local Group-like structure. an order-of magnitude disagreement ""unless a large fraction of the Local Group satellites has been nmüssed."," The simulations of \citet{KK99}, for instance, predict 300 satellite galaxies within 1.5 Mpc of the center of a Local Group-like structure, an order-of magnitude disagreement “unless a large fraction of the Local Group satellites has been missed.”"
 If that caveat is understood (o mean that observationally accessible galaxies (that is. those detectable over a large fraction of the sky in observations available at the (ime (he statement was made) had been somehow overlooked. it can now be removed.," If that caveat is understood to mean that observationally accessible galaxies (that is, those detectable over a large fraction of the sky in observations available at the time the statement was made) had been somehow overlooked, it can now be removed."
" The ""missing satellite problem persists in more recent simulations 2006).", The “missing satellite problem” persists in more recent simulations \citep{MM05}.
. A popular suggestion to deal with it involves reionization. which would make the conversion of small dark matter haloes into visible galaxies much more difficult.," A popular suggestion to deal with it involves reionization, which would make the conversion of small dark matter haloes into visible galaxies much more difficult."
 found that such a prescription allowed agreement between observed and simulated Local Group data for a large range of huminosities. producing a “large population of faint satellites.," \citet{BF02} found that such a prescription allowed agreement between observed and simulated Local Group data for a large range of luminosities, producing a “large population of faint satellites."
.. awaiting discovery.”, awaiting discovery.”
 In their simulation. the number of satellite galaxies brighter (han 22 magnitudes per square arc second (ini V) roughly matches observation: when the limit was lowered to 26. the number increased (ο 200.," In their simulation, the number of satellite galaxies brighter than 22 magnitudes per square arc second (in $V$ ) roughly matches observation; when the limit was lowered to 26, the number increased to 200."
" No numbers for intermediate levels are given. bul certainly the number per brightness bin must increase as one goes fainter, something we have shown does not happen."," No numbers for intermediate levels are given, but certainly the number per brightness bin must increase as one goes fainter, something we have shown does not happen."
 At anv rate. we have reached past 25 in fH: for the galaxies made of old stars such as Bensonetal...(2002). constructed this is most of the wav (to 26 in V.," At any rate, we have reached past 25 in $R$; for the galaxies made of old stars such as \citet{BF02} constructed this is most of the way to 26 in $V$."
" Their large population"" is simply not there. (", Their “large population” is simply not there. (
Note also that the detailed star formation histories of Local Group cdwarls appears to be inconsistent will supression by reionization (Grebel&Gallagher 2004))).,Note also that the detailed star formation histories of Local Group dwarfs appears to be inconsistent with supression by reionization \citep{GG04}) ).
 Another suggestion (both of these are found in IXlvpinetal.(1999).. though not original there) is that small dark matter haloes were never able to form a significant number of stars. retaining their barvonie matter in the form of atomic hydrogen. and that these have been detected as high-velocity clouds (IIVCs).," Another suggestion (both of these are found in \citet{KK99}, though not original there) is that small dark matter haloes were never able to form a significant number of stars, retaining their baryonic matter in the form of atomic hydrogen, and that these have been detected as high-velocity clouds (HVCs)."
 Aside [rom the fact that no stars or other optical signal have been identified with a HVC (for instance. Simon&Blitz (2002))). il appears that the IIVCs are too small and nearby to be identified with the missing dark halos. since," Aside from the fact that no stars or other optical signal have been identified with a HVC (for instance, \citet{SB02}) ), it appears that the HVCs are too small and nearby to be identified with the missing dark halos, since"
"canonical value of 10""(@preak/10~*), as a function of L with the pasta phases taken to be a Coulomb solid and a liquid.","canonical value of $10^{-7} (\bar{\sigma}_{\rm break}/10^{-2})$, as a function of $L$ with the pasta phases taken to be a Coulomb solid and a liquid."
" The maximum quadrupole ellipticity peaks at L=50 MeV; the peak occurs because of the competing trends of increasing stellar radius and decreasing crust-core transition density as L increases, a feature which only emerges when the crust and core are treated consistently."," The maximum quadrupole ellipticity peaks at $L \approx 50$ MeV; the peak occurs because of the competing trends of increasing stellar radius and decreasing crust-core transition density as $L$ increases, a feature which only emerges when the crust and core are treated consistently."
" For LS70 MeV, and taking the pasta phases to be solid, the maximum ellipticity is close to the canonical value."," For $L \lesssim 70$ MeV, and taking the pasta phases to be solid, the maximum ellipticity is close to the canonical value."
" However, if the pasta phases are assumed to be liquid, the maximum ellipticity is an order of magnitude lower ~1078(Gbreak/10~*)."," However, if the pasta phases are assumed to be liquid, the maximum ellipticity is an order of magnitude lower $\sim 10^{-8} (\bar{\sigma}_{\rm break}/10^{-2})$."
" At higher values of L, the pasta phases become insignificant, but the stiffening of the EoS and reduction of the crust-core transition density means that the crust thickness relative to the size of the star overall becomes smaller, and the maximum ellipticity sustainable falls by almost two orders of magnitude compared to the canonical value."," At higher values of $L$, the pasta phases become insignificant, but the stiffening of the EoS and reduction of the crust-core transition density means that the crust thickness relative to the size of the star overall becomes smaller, and the maximum ellipticity sustainable falls by almost two orders of magnitude compared to the canonical value."
 The equivalent gravitational wave strain is plotted in the bottom half of Fig., The equivalent gravitational wave strain is plotted in the bottom half of Fig.
 6 for a reference neutron star frequency of 300Hz and distance of 0.1kpc., 6 for a reference neutron star frequency of 300Hz and distance of 0.1kpc.
" The sensitivity of the most recent LIGO science run to such gravitational waves is indicated by the dashed line; this shows that, despite the reduced strain caused by a stiff EoS or the existence of liquid pasta phases, such GWs are still potentially observable with current or future detectors, offering the chance to distinguish between the various possibilities of crust properties."," The sensitivity of the most recent LIGO science run to such gravitational waves is indicated by the dashed line; this shows that, despite the reduced strain caused by a stiff EoS or the existence of liquid pasta phases, such GWs are still potentially observable with current or future detectors, offering the chance to distinguish between the various possibilities of crust properties."
 Fig., Fig.
" 7 shows the fundamental frequency |=2,n0 and first overtone |=2,n1 of the crustal torsional oscillation spectrum for a 1.4Mo star, where nm is the number of radial nodes and { the angular constant associated with the spherical harmonics Y;”."," 7 shows the fundamental frequency $l=2, n=0$ and first overtone $l=2, n=1$ of the crustal torsional oscillation spectrum for a $1.4 M_{\odot}$ star, where $n$ is the number of radial nodes and $l$ the angular constant associated with the spherical harmonics $Y_l^m$ ."
" Observed values of the frequency of QPOs at 16, 26, 28 and 30 Hz (Israeletal.2005;Watts&Strohmayer2006;Watts 2005a,b),, within the range of the calculated fundamental frequencies, are indicated by the dashed lines, as well as the observed 84, 92, 150, 155 and 625 Hz frequencies within the range of the first overtone."," Observed values of the frequency of QPOs at 16, 26, 28 and 30 Hz \citep{Israel2005, Watts2006, Strohmayer2005, Strohmayer2006}, within the range of the calculated fundamental frequencies, are indicated by the dashed lines, as well as the observed 84, 92, 150, 155 and 625 Hz frequencies within the range of the first overtone."
 The frequency generically decreases with increasing L as the crust becomes thinner relative tothe, The frequency generically decreases with increasing $L$ as the crust becomes thinner relative tothe
continuum contribution to each afterglow and associated uncertainty have been calculated in Sect. 3::,continuum contribution to each afterglow and associated uncertainty have been calculated in Sect. \ref{Sec:cont}:
 we evaluate the continuum at the original light-curve central time bins and combine the systematic uncertainty coming from the continuum evaluation procedure with a statistical uncertainty equal to the relative error on the original luminosity bins of that particular burst., we evaluate the continuum at the original light-curve central time bins and combine the systematic uncertainty coming from the continuum evaluation procedure with a statistical uncertainty equal to the relative error on the original luminosity bins of that particular burst.
 A grid of 545 flare peak time values tj; equally spaced in logarithmic units between 1 s and 10” s is generated together with a grid of 171 flare peak luminosity values in the interval 1035—10°’ergs”," A grid of 545 flare peak time values $t_{pk}$ equally spaced in logarithmic units between 1 s and $10^7$ s is generated together with a grid of 171 flare peak luminosity values in the interval $10^{40}-10^{57}\,\rm{erg\,s^{-1}}$."
" For each tx, we generate 171 flares with growing peak !.luminosity values: a Norrisetal.(2005) profile specified in terms of amplitude A, width w and peak time £j; is assumed."," For each $t_{pk}$, we generate 171 flares with growing peak luminosity values: a \cite{Norris05} profile specified in terms of amplitude $A$, width $w$ and peak time $t_{pk}$ is assumed."
" The Norris 2005 profile is defined as f=f(A,71,72,ts) where τι and Το are two flare shape parameters related to the rise and decay phases, respectively, while t, defines the pulse starting time."," The Norris 2005 profile is defined as $f\equiv f(A,\tau_1,\tau_2,t_s)$ where $\tau_1$ and $ \tau_2$ are two flare shape parameters related to the rise and decay phases, respectively, while $t_s$ defines the pulse starting time."
" We take advantage of previous studies on the flare phenomenology to reduce the flare parameter space and write τι~4w/3; T2~w/3; ts~toy—2w/3; w~10-994,0-125018) (see C10 and M10 for details).", We take advantage of previous studies on the flare phenomenology to reduce the flare parameter space and write $\tau_1\sim4w/3 $; $\tau_2\sim w/3$; $t_s\sim t_{pk}-2w/3$; $w\sim10^{-0.94}t_{pk}^{(1.12\pm0.18)}$ (see C10 and M10 for details).
 The signal of the fake flare is then integrated in the original light-curve bins and statistical error is assigned equal to the relative uncertaintya affecting the original luminosity curve at the same time., The signal of the fake flare is then integrated in the original light-curve bins and a statistical error is assigned equal to the relative uncertainty affecting the original luminosity curve at the same time.
" Finally, the flare contribution is added to the continuum and the uncertainties propagated into a final fake GRB afterglow."," Finally, the flare contribution is added to the continuum and the uncertainties propagated into a final fake GRB afterglow."
" To simulate the entire procedure, the fake afterglow is then processed by the automatic excess estimation routine described in the previous section: a new continuum is determined together with the associated flaring component, as before."," To simulate the entire procedure, the fake afterglow is then processed by the automatic excess estimation routine described in the previous section: a new continuum is determined together with the associated flaring component, as before."
" If the newly determined flaring component contains a fluctuation with a minimum 2c significance, the peak luminosity of the fake flare is recorded as 2c flare detection threshold at tp, and the following peak time value is considered."," If the newly determined flaring component contains a fluctuation with a minimum $2\sigma$ significance, the peak luminosity of the fake flare is recorded as $2\sigma$ flare detection threshold at $t_{pk}$ and the following peak time value is considered."
" The entire procedure is repeated for each tj, every time starting from the minimum peak luminosity value of the grid."," The entire procedure is repeated for each $t_{pk}$, every time starting from the minimum peak luminosity value of the grid."
 The result is shown in Fig., The result is shown in Fig.
" 1 for 0060607A taken as an example (blue solid line): due to the degrading quality of the signal, a growing flare-to-continuum flux contrast is required for the flare to be detected at later times."," \ref{Fig:example} for 060607A taken as an example (blue solid line): due to the degrading quality of the signal, a growing flare-to-continuum flux contrast is required for the flare to be detected at later times."
" The 2c detection luminosity curve of the different GRBs are then combined to produce the average detection threshold (L(t)) in the rest frame band 2.2-14.4 keV. In strict analogy with Eq. 1,,"," The $2\sigma$ detection luminosity curve of the different GRBs are then combined to produce the average detection threshold $\langle L_{th}(t) \rangle$ in the rest frame band 2.2-14.4 keV. In strict analogy with Eq. \ref{Eq:avelum}, ,"
 (Ltn(t)) is defined as: being N the number of GRBs displaying a significant positive flaring component in the time period t;—ty; Lit is the luminosity of the 2e detection threshold of the i' GRB at time t=(t;+t5)/2.," $\langle L_{th}(t) \rangle$ is defined as: being $N$ the number of GRBs displaying a significant positive flaring component in the time period $t_i-t_f$; $L_{i,th}$ is the luminosity of the $2\sigma$ detection threshold of the $i^{th}$ GRB at time $t=(t_i+t_f)/2$."
" The re-normalised average 2c threshold (Lj),(t*)) is computed as: where L;;,=Li,tn/(Li,y)."," The re-normalised average $2\sigma$ threshold $\langle L^{*}_{th}(t^{*}) \rangle$ is computed as: where $L_{i,th}^{*}\equiv L_{i,th}/\langle L_{i,\gamma} \rangle$."
" As before, N collects the GRBs with significant excesses in tj—t."," As before, $N$ collects the GRBs with significant excesses in $t_{i}^{*}-t_{f}^{*}$."
 The result is portrayed in Fig. 2::, The result is portrayed in Fig. \ref{Fig:avelum}:
 in both panels a thick green area marks the 2σ detection threshold region., in both panels a thick green area marks the $2\sigma$ detection threshold region.
 From panel (a) it is clear that the —1.2 slope is mainly due to the of bright fluctuations able to overshine the threshold., From panel $(a)$ it is clear that the $-1.2$ slope is mainly due to the of bright fluctuations able to overshine the threshold.
 The fact that (L(t)) tracks the (Lin(t)) temporal behaviour suggests that the sample of significant excesses detected in the time period 2-3000 ks are not representative of the real population of flares at those times., The fact that $\langle L(t) \rangle $ tracks the $\langle L_{th}(t)\rangle$ temporal behaviour suggests that the sample of significant excesses detected in the time period 2-3000 ks are not representative of the real population of flares at those times.
 The bright end of the flare luminosity distribution has been sampled and à biassed average flare luminosity has been re-constructed., The bright end of the flare luminosity distribution has been sampled and a biassed average flare luminosity has been re-constructed.
" Notably, the decay of the detection threshold at t104 s allows us to appreciate a steeper decay (a~ —2.5) of 30 significant positive fluctuations extending to t~6x104 s. No more than 2o significant flaring activity can be quoted for (1«t6)x10° s due to the flatter decay (and actually enhancement) of (L,4,(t))."," Notably, the decay of the detection threshold at $t\sim10^4$ s allows us to appreciate a steeper decay $\alpha\sim -2.5$ ) of $3\sigma$ significant positive fluctuations extending to $t\sim6\times 10^4$ s. No more than $2\sigma$ significant flaring activity can be quoted for $(1<t<6)\times 10^5$ s due to the flatter decay (and actually enhancement) of $\langle L_{th}(t) \rangle$."
 The unbiassed average flare luminosity curve is therefore likely to be steeper than —1.2 at late times., The unbiassed average flare luminosity curve is therefore likely to be steeper than $-1.2$ at late times.
" At early times the situation is different: for tX300 s the (L(t))οςt?* decay is not affected bythreshold effects; the extension of the same temporal law up to t~1000 s suggests that (Lin(t)) is likely to play a minor role in the time interval (0.3«t1) ks, as"," At early times the situation is different: for $t\lesssim 300$ s the $\langle L(t) \rangle \propto t^{-2.7}$ decay is not affected bythreshold effects; the extension of the same temporal law up to $t\sim1000$ s suggests that $\langle L_{th}(t) \rangle$ is likely to play a minor role in the time interval $(0.3<t<1)$ ks, as"
the two telescope's poiut of view with the angles Qu Which represent the star distauce frou the baseline. direction.,"the two telescope's point of view with the angles $\varphi_i$, which represent the star distance from the baseline direction."
 One παν reduce this to the ecoeraphic coordinates bv multiplying ro with s4 to ect cosa aud cosZ are close to unity in contemporary optical interferometry. so the actual implementation avoids the use of the cosine in favor of the haversine. to protect against cancellation of siguificaut digits. and returus frou there to the sine. if needed. The special coplanar case of refsec.zceui2D is included as Az=Z aud a=0 OY m. (," One may reduce this to the geographic coordinates by multiplying ${\bf r}_1+{\bf b}_{12}={\bf r}_2$ with ${\bf s}_1$ to get $\cos\sigma$ and $\cos Z$ are close to unity in contemporary optical interferometry, so the actual implementation avoids the use of the cosine in favor of the haversine, to protect against cancellation of significant digits, and returns from there to the sine, if needed, The special coplanar case of \\ref{sec.zeni2D} is included as $\Delta z=Z$ and $\sigma=0$ or $\pi$. ("
On some Linux svstenis. where the ibrary function docs not support the full accuracy of the 96 bit nuuber representation. it inakes sense to switch to alternative high- diuplemientatious of the cosine (?)..),"On some Linux systems, where the library function does not support the full accuracy of the 96 bit number representation, it makes sense to switch to alternative high-precision implementations of the cosine .)"
 Two calculations for the actual impact parameters I;=pnosinzy:EE starting: from- the given: :; would (0 done as in the preceding. “aligned” geometry of rofsec.zeni2D.. and these inserted iuto to calculate the projected baseline. 7 and to eeneralize (415).," Two calculations for the actual impact parameters $I_i=\rho n_0 \sin z_0^{(i)}$ starting from the given $z_i$ would be done as in the preceding, “aligned” geometry of \\ref{sec.zeni2D}, and these inserted into to calculate the projected baseline $P$ and to generalize \ref{eq.PofIs}) )."
 DP?—(hoDy| courprises the terms of ((56)) of the constrained ecometiy augmented by as read from ((71)) in combination with ((53)) (559)., $P^2=(I_1-I_2)^2+4I_1I_2\hav\sigma$ comprises the terms of \ref{eq.Pofdeltaz}) ) of the constrained geometry augmented by as read from \ref{eq.havsigma}) ) in combination with \ref{eq.I1ofDz}) \ref{eq.I2ofDz}) ).
We thus expect powerlaw time decays. a characteristic of strong explosions (see the Sedov analoev!).,"We thus expect power–law time decays, a characteristic of strong explosions (see the Sedov analogy!),"
 with time and spectralindices closely related., with time– and spectral–indices closely related.
 This is what is observed everywhere. from the X.ταν through the optical to the radio. (see Piro and Fruchter. this volume}. the few exceptions beiug discussed later on.," This is what is observed everywhere, from the X–ray through the optical to the radio, (see Piro and Fruchter, this volume), the few exceptions being discussed later on."
 Iu fact. the equality of the timedecay index of the XN.rav aud optical data in afterelows of individual sources has been taken as the key clement to show that emission in the different bands is due to the same source.," In fact, the equality of the time–decay index of the X–ray and optical data in afterglows of individual sources has been taken as the key element to show that emission in the different bands is due to the same source."
 Time mdiees in the X.rav are in the range 0.72.2 (Frouteraed... in preparation).," Time indices in the X–ray are in the range $0.7-2.2$ (Frontera, in preparation)."
 After having established that bursts are due to explosions. we happily learn that afterelows emit through svuchrotron processes.," After having established that bursts are due to explosions, we happily learn that afterglows emit through synchrotron processes."
 In fig., In fig.
 2 (Calamaal... 1998). we show the superposition of theoretical expectations for an optically thin svuchrotron spectrum (iucludine the cooling break at »zz1014 IT:) with obscrvatious for CRB 970508.," 2 (Galama, 1998), we show the superposition of theoretical expectations for an optically thin synchrotron spectrum (including the cooling break at $\nu \approx 10^{14}\; Hz$ ) with observations for GRB 970508."
 The remarkable agreement is even more exciting as we remark that observations are nof truly simultaneous. but are scaled back to the same time bv means of the theoretically expected laws for timedecay. thus simnultiueoulsv testing the correctness of our hwdro.," The remarkable agreement is even more exciting as we remark that observations are not truly simultaneous, but are scaled back to the same time by means of the theoretically expected laws for time–decay, thus simultaneoulsy testing the correctness of our hydro."
 Another piece of evidence comes from the discovery of polarization iu the optical afterglow of GRB 990510 (Fie., Another piece of evidence comes from the discovery of polarization in the optical afterglow of GRB 990510 (Fig.
 3. Covinoal... 1999. Wijersab... 1999).," 3, Covino, 1999, Wijers, 1999)."
 This polarization may appear small (zz:2 oC). but it is surely uot due to Galactic effects: stars in the same field show a comparable degree of polarization. but along an axis differcut by about 507.," This polarization may appear small $\approx 2\%$ ), but it is surely not due to Galactic effects: stars in the same field show a comparable degree of polarization, but along an axis different by about $50^\circ$."
 Also. polarization iu the source galaxy is unlikely. because of a very stringent upper limit ou the reddeniug due to this galaxy (Covino," Also, polarization in the source galaxy is unlikely, because of a very stringent upper limit on the reddening due to this galaxy (Covino"
"eft side of the S 6, phase diagram.",left side of the $S$ $\delta_g$ phase diagram.
" For instance. imposing he same entropy [loor at 2,=9 and at 2,=3 has quite different effects on the entropy pattern (see lower panels of Fie. 7)."," For instance, imposing the same entropy floor at $z_h=9$ and at $z_h=3$ has quite different effects on the entropy pattern (see lower panels of Fig. \ref{fi:srho}) )."
 Heating at ο=9 has the elfect of increasing the cooling time for most of the eas particles. so as to allow star formation to take place only quite recently (see Fig. 4).," Heating at $z_h=9$ has the effect of increasing the cooling time for most of the gas particles, so as to allow star formation to take place only quite recently (see Fig. \ref{fi:sfr}) )."
 The increased. time-scale for cooling causes this process to proceed in a more eradual wav., The increased time-scale for cooling causes this process to proceed in a more gradual way.
 For this reason. the entropy of gas particles undergoing cooling decreases slowly. thus making their removal from the hot clilluse phase less ellicient.," For this reason, the entropy of gas particles undergoing cooling decreases slowly, thus making their removal from the hot diffuse phase less efficient."
 The slope of the Ly Y relation also provides importan observational evidence for the lack of self.similar behaviour of the ICAL, The slope of the $L_X$ $T$ relation also provides important observational evidence for the lack of self–similar behaviour of the ICM.
 Since. the first. measurements of LCA temperatures for sizable sets of clusters. i0 has. been recognised that LyoxZ7. with ac3. although with a considerable scatter (e.g... White et. al.," Since the first measurements of ICM temperatures for sizable sets of clusters, it has been recognised that $L_X\propto T^\alpha$ with $\alpha\simeq 3$, although with a considerable scatter (e.g., White et al."
 1998. ane references therein)," 1998, and references therein)."
 Setter quality observations establishec that a significant contribution to this scatter is associatec with the cdillerent strength. of cooling Lows detected: in dilferent clusters., Better quality observations established that a significant contribution to this scatter is associated with the different strength of cooling flows detected in different clusters.
 Either excluding clusters with pronounced signatures of cooling Hows or correcting for their ellect (e.g. Alarkeviteh 1998: Allen Fabian 1998: Arnaud Evrard 1999: Ettori et al.," Either excluding clusters with pronounced signatures of cooling flows or correcting for their effect (e.g., Markevitch 1998; Allen Fabian 1998; Arnaud Evrard 1999; Ettori et al."
 2002) results in a much tighter Ly T relation. albeit still with a rather steep slope.," 2002) results in a much tighter $L_X$ $T$ relation, albeit still with a rather steep slope."
 At the same time. hints have also been found for a further steepening of this relation at Z2;1 keV (ο... Lelsclon Ponman 2000. and references therein). possibly indicating that the mechanism responsible for the Ly VY scaling should act in a dillerent way for clusters and. groups.," At the same time, hints have also been found for a further steepening of this relation at $T\mincir 1$ keV (e.g., Helsdon Ponman 2000, and references therein), possibly indicating that the mechanism responsible for the $L_X$ $T$ scaling should act in a different way for clusters and groups."
 Simulations that allow for nongravitational heating (e.g. Dialek et al.," Simulations that allow for non–gravitational heating (e.g., Bialek et al."
 2001: BOW) and radiative cooling (c.g. Pearce et al., 2001; BGW) and radiative cooling (e.g. Pearce et al.
 2001: Dave et al., 2001; Davè et al.
 2002: Muanwong ct al., 2002; Muanwong et al.
 2002) have been shown to be able to account for the observed Ly Yo relation., 2002) have been shown to be able to account for the observed $L_X$ $T$ relation.
 Llowever. a sometimes overlooked issue in determining the X rav luminosity of clusters in simulations concerns the contribution of metal lines to the emissivity.," However, a sometimes overlooked issue in determining the $X$ –ray luminosity of clusters in simulations concerns the contribution of metal lines to the emissivity."
 While this contribution is negligible above 2 keV. it becomes relevant at the scale of groups.," While this contribution is negligible above 2 keV, it becomes relevant at the scale of groups."
 For instance. neglecting the contribution from line enmissivity for an ICM enriched to a metallicity of Z=(.3Z. leads to an underestimate of the Nταν," For instance, neglecting the contribution from line emissivity for an ICM enriched to a metallicity of $Z=0.3Z_\odot$ leads to an underestimate of the X–ray"
(see Section 5.6)). the maximum percentage [Hux variation is from the secondary star alone.,"(see Section \ref{sec:distance}) ), the maximum percentage flux variation is from the secondary star alone."
 We plotted the relative flux. from cach spectroscopic observation against cach simultaneous photometric observation. to compare the fluxes.," We plotted the relative flux from each spectroscopic observation against each simultaneous photometric observation, to compare the fluxes."
 This was done bv summing the J0422]|32 spectra over the range of the shotometric pass-banel (7250-9500A)) and. dividing this by he sum of the comparison star over this same range and hen plotting this against the photometric observations (Fig. 4))., This was done by summing the J0422+32 spectra over the range of the photometric pass-band ) and dividing this by the sum of the comparison star over this same range and then plotting this against the photometric observations (Fig. \ref{fig:slitcor}) ).
 Phere is a large scatter about the best Lit line ο the spectroscopy verses photometry. showing that there are considerable cdillerences in the light losses at the slit oetween the two stars.," There is a large scatter about the best fit line to the spectroscopy verses photometry, showing that there are considerable differences in the light losses at the slit between the two stars."
 We therefore decided to correct. for hese slit losses. by multiplving cach J0422|32 spectrum w the ratio between the photometry and the spectroscopy.," We therefore decided to correct for these slit losses, by multiplying each J0422+32 spectrum by the ratio between the photometry and the spectroscopy."
 The second (local) star observed on the slit to be used as a Comparison star. was flux. calibrated using the Εαν standard. SP 22001178 (BD 117 4708). in the following wav.," The second (local) star observed on the slit to be used as a comparison star, was flux calibrated using the flux standard SP 2209+178 (BD +17 4708), in the following way."
 The instrument response function derived. from the ας standard was used to flux the weighted mean of the 53 spectra ofthe comparison star., The instrument response function derived from the flux standard was used to flux the weighted mean of the 53 spectra of the comparison star.
 The flux level was virtually --nsensitive to changes in the airmass., The flux level was virtually insensitive to changes in the airmass.
 A smooth curve was fitted to the [luxed mean comparison star spectrum., A smooth curve was fitted to the fluxed mean comparison star spectrum.
 The J0422|32 spectrum (cliviclecl by the comparison star) could then be Huxecl by multiplving by this curve and. remain corrected. for tellurie absorption., The J0422+32 spectrum (divided by the comparison star) could then be fluxed by multiplying by this curve and remain corrected for telluric absorption.
 Phe J0422{32 fluxed spectrum was cereddenec using an E(B-V)=0.3 (Llarlaltis et al. 1999 ancl Beekman et al. LOOT) and barvecntric raclial velocity corrected (Pig. 5)).," The J0422+32 fluxed spectrum was dereddened using an E(B-V)=0.3 (Harlaftis et al, 1999 and Beekman et al, 1997) and barycentric radial velocity corrected (Fig. \ref{fig:simpleadd}) )."
 We searched for the ellipsoidal modulation of the secondary star in two wavs., We searched for the ellipsoidal modulation of the secondary star in two ways.
 We measured Ες deficits for cach of the S phase bins. as described in Wade Horne (1988). but the errors were large (25-100%)).," We measured flux deficits for each of the 8 phase bins, as described in Wade Horne (1988), but the errors were large )."
 To σοι a higher signal-to-noise. varving proportions of the summed spectrum (Pig. 5))," To get a higher signal-to-noise, varying proportions of the summed spectrum (Fig. \ref{fig:simpleadd}) )"
 were subtracted from. each of the S phase binned spectra., were subtracted from each of the 8 phase binned spectra.
 The region wivas used. as it contains two strong TiO bands. but no Ca lines.," The region was used, as it contains two strong TiO bands, but no Ca lines."
 A line was fitted. to the residuals and. we used V to determine the best fitting proportion of the summed spectrum at each phase., A line was fitted to the residuals and we used ${\chi^2}$ to determine the best fitting proportion of the summed spectrum at each phase.
 This measure could also reveal ellipsoidal variation., This measure could also reveal ellipsoidal variation.
 Although the errors were smaller (10-20 »rcent). only a singlv-humped.e» lishteurveo was seen rather than a double-humped: curve.," Although the errors were smaller (10-20 percent), only a singly-humped lightcurve was seen rather than a double-humped curve."
 This is probably clue to irractiation of the front face of the M star. where the hotter temperatures diminish the quantity of TiO. Using the relationship between the amplitude. of the ellipsoidal variation anc the inclination. of the system (Shahbaz 1994). we etermined. the maximum. orbital inclination of J0422]|32 to be 45°. for the maximum amplitude: obtained. with the photometry (in which we can see evidence of ellipsoidal variation. see Sec. 4.2.2)).," This is probably due to irradiation of the front face of the M star, where the hotter temperatures diminish the quantity of TiO. Using the relationship between the amplitude of the ellipsoidal variation and the inclination of the system (Shahbaz 1994), we determined the maximum orbital inclination of J0422+32 to be $^{\circ}$, for the maximum amplitude obtained with the photometry (in which we can see evidence of ellipsoidal variation, see Sec. \ref{sec:photan2}) )."
 Using the semi-amplitude obtained from our radial velocity study. of ATSIXIÓ km s we obtained a mass function of 1.191-E0.021 M.," Using the semi-amplitude obtained from our radial velocity study of $\pm$ 16 km $^{-1}$, we obtained a mass function of $\pm$ $_\odot$."
.. We used this new mass function. the esinz of 90!22 km s+ ((larlaftisetal1999) and the inclination. to obtain a newq value of 9.0!22. where qis the mass of the primary (compact object) divided by the secondary.," We used this new mass function, the $v$ $i$ of $^{\scriptscriptstyle +22}_{\scriptscriptstyle
-27}$ km $^{-1}$ \cite{Harl99} and the inclination, to obtain a new value of $^{\scriptscriptstyle +2.2}_{\scriptscriptstyle
-2.7}$, where is the mass of the primary (compact object) divided by the secondary."
 From this we can constrain the mass of the primary to be greater than 2.2M.., From this we can constrain the mass of the primary to be greater than ${_\odot}$.
 Although this is below 2.9M.. (the upper bound on the mass of a neutron star. Ixalogera Εάν (1996)). it is well above the canonical value for a neutron star (1.44M..) ancl would. imply a super-massive neutron," Although this is below $_\odot$ (the upper bound on the mass of a neutron star, Kalogera Baym (1996)), it is well above the canonical value for a neutron star $_\odot$ ) and would imply a super-massive neutron"
red giant star would increase the surface metallicity by ~+0.3 dex in and decrease the Zn/Fe ratio on the stellar surface by a [Fe/H]similar amount.,red giant star would increase the surface metallicity by $\sim +0.3$ dex in [Fe/H] and decrease the Zn/Fe ratio on the stellar surface by a similar amount.
" There is no information on how much fraction of each element was depleted from the gas phase in the proto-planetary gas disks around metal-poor stars with the metallicity range of —2 Fe-group elements are depleted as (Weltymuch as in the 2001)::Galactic ISM, whereas a-elements such as Mg and Si are nearly undepleted, though Si has a higher condensation temperature than Fe."," There is no information on how much fraction of each element was depleted from the gas phase in the proto-planetary gas disks around metal-poor stars with the metallicity range of $-2$: Fe-group elements are depleted as much as in the Galactic ISM, whereas $\alpha$ -elements such as Mg and Si are nearly undepleted, though Si has a higher condensation temperature than Fe."
 This depletion pattern would support the scenario in which nearby low-a stars engulfed a fraction of dusts in the proto-planetary disks and low-a dSph stars engulfed giant planets., This depletion pattern would support the scenario in which nearby $\alpha$ stars engulfed a fraction of dusts in the proto-planetary disks and $\alpha$ dSph stars engulfed giant planets.
" Of course, the physical conditions could be quite different between proto-planetary gas disks and the observed ISM in the SMC."," Of course, the physical conditions could be quite different between proto-planetary gas disks and the observed ISM in the SMC."
 It is crucial to understand what determines the dust depletion pattern., It is crucial to understand what determines the dust depletion pattern.
 The most metal-poor star harboring planets to date is reported to have a metallicity of (Sadakaneetal., The most metal-poor star harboring planets to date is reported to have a metallicity of $-0.74$ \citep{Sadakane_02}.
 Thus the scenario proposed [Fe/H]=—0.74here has not been fully 2002)..supported by the ongoing search for extra-solar planets., Thus the scenario proposed here has not been fully supported by the ongoing search for extra-solar planets.
 Future observations will construct a much larger sample to answer if there are planets orbiting stars as metal-poor as [Fe/H]~—2., Future observations will construct a much larger sample to answer if there are planets orbiting stars as metal-poor as $\sim -2$.
 We have discussed the origin of low-a stars found in the Galaxy and dSph galaxies., We have discussed the origin of $\alpha$ stars found in the Galaxy and dSph galaxies.
 We have shown that the low-a stars found in the Galaxy are on the main-sequence and thus concludes that these stars have suffered from external pollution., We have shown that the $\alpha$ stars found in the Galaxy are on the main-sequence and thus concludes that these stars have suffered from external pollution.
 A comparison of the elemental abundance pattern of low-o stars with that of DLAs suggests that low-o stars have accreted dust grains., A comparison of the elemental abundance pattern of $\alpha$ stars with that of DLAs suggests that $\alpha$ stars have accreted dust grains.
" On the other hand, low-o dSph stars are located at the tip of the RGB."," On the other hand, $\alpha$ dSph stars are located at the tip of the RGB."
" To explain the origin of these low-a stars in different evolutionary stages, wehave proposed that these low-a stars harbor or have harbored planets and that stars engulfing planets and/or planetesimals show low a/Fe ratios."," To explain the origin of these $\alpha$ stars in different evolutionary stages, wehave proposed that these $\alpha$ stars harbor or have harbored planets and that stars engulfing planets and/or planetesimals show low $\alpha$ /Fe ratios."
" However, it should be noted that there is at least one subgiant with low a/Fe ratios, BD +80° 245 (Carneyetal."," However, it should be noted that there is at least one subgiant with low $\alpha$ /Fe ratios, BD $^\circ$ 245 \citep{Carney_97}."
 This star has a/Fe approximately equal to the solar 1997)..values and —2., This star has $\alpha$ /Fe approximately equal to the solar values and $\sim -2$.
" If our scenario is applied to this star, it must have[Fe/H]~ harbored a giant planet at the radius <0.014 AU."," If our scenario is applied to this star, it must have harbored a giant planet at the radius $\ltsim 0.014$ AU."
 This cannot necessarily discard our scenario., This cannot necessarily discard our scenario.
 The smallest semi-major axis of planets in the present catalogue of Schneider(2003) is ~0.02 AU comparable to the above value., The smallest semi-major axis of planets in the present catalogue of \citet{Schneider_03} is $\sim 0.02$ AU comparable to the above value.
" Recently, Ivansetal. reported some other metal-poor stars belonging (2003)to low-o stars that cannot be explained by the scenario presented here."," Recently, \citet{Ivans_03} reported some other metal-poor stars belonging to $\alpha$ stars that cannot be explained by the scenario presented here."
" These stars exhibit extremely large enhancements of Ti, Cr, Mn, Ni, Zn, and Ga relative to Fe and are deficient in Si and Mg."," These stars exhibit extremely large enhancements of Ti, Cr, Mn, Ni, Zn, and Ga relative to Fe and are deficient in Si and Mg."
 These elemental abundance features might witness the earliest SNe Ia as suggested by Ivansetal., These elemental abundance features might witness the earliest SNe Ia as suggested by \citet{Ivans_03}.
 We cannot expect a correlation of low-o stars(2003).. that have no sign of SNe Ia with any stellar kinematics from our scenario., We cannot expect a correlation of $\alpha$ stars that have no sign of SNe Ia with any stellar kinematics from our scenario.
 Therefore the relationship between low-a@ stars and stellar kinematics mentioned in 811 should be explained by some other mechanisms., Therefore the relationship between $\alpha$ stars and stellar kinematics mentioned in 1 should be explained by some other mechanisms.
 The correlation of ratios with the galactic radii was already found in the a/Fedisk stars (Edvardssonetal. 1993)., The correlation of $\alpha/$ Fe ratios with the galactic radii was already found in the disk stars \citep{Edvardsson_93}. .
". According to the standard chemical evolution model, this correlation has been explained as a result of the combination of SNe Ia products and longer evolution time scales of the chemical evolution at larger galactic distances (Tsujimotoetal. 1995b)."," According to the standard chemical evolution model, this correlation has been explained as a result of the combination of SNe Ia products and longer evolution time scales of the chemical evolution at larger galactic distances \citep{Tsujimoto_95b}."
. This argument could be applied to the correlation of ratios with the galactic radii found for the halo stars a/Fe(Nissen&Schuster1997;KingFulbright 2003b)..," This argument could be applied to the correlation of $\alpha/$ Fe ratios with the galactic radii found for the halo stars \citep{Nissen_97, King_97, Fulbright_00, Gratton_03b}. ."
 It is possible to test our conjecture by two kinds of observations., It is possible to test our conjecture by two kinds of observations.
 One is to search planets orbiting around a dwarfs in the metallicity range of—2 <[Fe/H]<—1 in, One is to search planets orbiting around $\alpha$ dwarfs in the metallicity range of$-2<$ $<-1$ in
ratio is proportional to the ratio of the seeimg fraction to he super-pixel size (κ).,ratio is proportional to the ratio of the seeing fraction to the super-pixel size $d_{sp}$ ).
 Cong from 1«1 to 3&2 super-κοιν iucreases the seeiug fraction by more than a factor 3., Going from $1\times 1$ to $3\times 3$ super-pixels increases the seeing fraction by more than a factor 3.
 Then iucreasing the super-pixel size further increases he secing fraction substantially less than the fluctuations of the sky backeround., Then increasing the super-pixel size further increases the seeing fraction substantially less than the fluctuations of the sky background.
 Figure 10. displays the variation with sccing of the signal to noise ratio for different super-κο] sizes., Figure \ref{fig:erf} displays the variation with seeing of the signal to noise ratio for different super-pixel sizes.
 It is clear that 3<3 super-pixels offer the best S/N ratio for our configuration., It is clear that $3\times 3$ super-pixels offer the best S/N ratio for our configuration.
 As discussed by Ausari et al (1997)... the alternative hat consists nu aking the average of the ucighbouring κος weighted wit[um je PSF is not appropriate here. as it amplifies the fiuctuatious due to ποσο variations.," As discussed by Ansari et al \nocite{Ansari:1997a}, the alternative that consists in taking the average of the neighbouring pixels weighted with the PSF is not appropriate here, as it amplifies the fluctuations due to seeing variations."
 Figure 11 show:μα 1e relative dispersion affecting the super-pixel fluxes for CCD 3: we measure in average in blue aud in red. which corresponds to about twice the estimated level of photon noise in blue aud in red).," Figure \ref{fig:stabi1} shows the relative dispersion affecting the super-pixel fluxes for CCD 3: we measure in average in blue and in red, which corresponds to about twice the estimated level of photon noise in blue and in red)."
 The comparison with Fie., The comparison with Fig.
 8 shows that the dispersion is reduced by a factor snaller than Vfn9(pixels)αχ.=3. because of. the correlationH between uciehbourmng pixels.," \ref{fig:stabpm} shows that the dispersion is reduced by a factor smaller than $\sqrt{ 9\, {\rm (pixels)}} = 3$ because of the correlation between neighbouring pixels."
 This stability can be iuproved even further by correcting for seeing variations., This stability can be improved even further by correcting for seeing variations.
 Despite of the stability discussed. above. fluctuations of super-pixel fluxes due to seeiug variations are still present.," Despite of the stability discussed above, fluctuations of super-pixel fluxes due to seeing variations are still present."
 For a star lviug in the ceutral pixel (of the DosD pixel). on average of the star flux euters on average the supor-pixel for a Gaussian PSF. but this sccing fraction is correlated with the changing secing.," For a star lying in the central pixel (of the $3 \times 3$ super-pixel), on average of the star flux enters on average the super-pixel for a Gaussian PSF, but this seeing fraction is correlated with the changing seeing."
 Iu this sub-section. we show that this correlation is linear and can be largely corrected for.," In this sub-section, we show that this correlation is linear and can be largely corrected for."
 Depending ou their position with respect to the nearest star. supor-pixel fiuxes cau significantly auti-correlate with the seeing if the super-pixel is in the seciug spot. or correlate if instead it lies in the tail of a star.," Depending on their position with respect to the nearest star, super-pixel fluxes can significantly anti-correlate with the seeing if the super-pixel is in the seeing spot, or correlate if instead it lies in the tail of a star."
" A correlation coefficient for cach super-pixel p can be defined using the usual formula: where o aud 5 are the mean values of the supor-pixel fux o aud scciug 5, ou uight ».", A correlation coefficient for each super-pixel $p$ can be defined using the usual formula: where $\phi$ and $S$ are the mean values of the super-pixel flux $\phi^p_n$ and seeing $S_n$ on night $n$.
 Iu Fig. 12..," In Fig. \ref{fig:corrsee},"
 we show the istributious ofcorrelation coefBcieuts pl in blue (left) aud Πα 1ος (right) or cach super-pixel p., we show the distributions of correlation coefficients $\rho^p$ in blue (left) and in red (right) for each super-pixel $p$.
 These histograms look quite different iu both colours but both distributions have a peak around p&0.8., These histograms look quite different in both colours but both distributions have a peak around $\rho \simeq -0.8$.
 This peak. which corresponds to the auti-correlation with seeime near the centre of resolved stars. is expected due to the large number of resolved stars.," This peak, which corresponds to the anti-correlation with seeing near the centre of resolved stars, is expected due to the large number of resolved stars."
 It is higher iu red than iu blue. which is consisteut with the EROS colouranagnitude diagram where most detected: stars have BHRo0 (Renault 1996).," It is higher in red than in blue, which is consistent with the EROS colour-magnitude diagram where most detected stars have $B-R > 0$ (Renault \nocite{Renault:1996}."
. The correlation with secing expected for star tails (p2 0) is less apparent., The correlation with seeing expected for star tails $\rho > 0$ ) is less apparent.
 However. a clear excess at high values of p (around px 1.6) appears in red. again consistent with the EROS color-magnitude diagram.," However, a clear excess at high values of $\rho$ (around $\rho \simeq 0.6$ ) appears in red, again consistent with the EROS colour-magnitude diagram."
 Figure 13 elves an example of such a correlation., Figure \ref{fig:clsee1} gives an example of such a correlation.
 The upper left panel, The upper left panel
irregular galaxies it is clear that an ellipse is a forced figure onto the isophotes.,irregular galaxies it is clear that an ellipse is a forced figure onto the isophotes.
 We will discuss the merit of this process in our structure paper., We will discuss the merit of this process in our structure paper.
" Final fits were visually inspected for robustness, converted to 1D surface brightness profiles and fit to an exponential disk."," Final fits were visually inspected for robustness, converted to 1D surface brightness profiles and fit to an exponential disk."
" A small number of objects had significant bulges, but r!/4 fits to the bulges did not significantly alter the disk fits due to the small bulge size."," A small number of objects had significant bulges, but $^{1/4}$ fits to the bulges did not significantly alter the disk fits due to the small bulge size."
" Both surface brightness profiles, aperture magnitudes and color profiles are based on these fits."," Both surface brightness profiles, aperture magnitudes and color profiles are based on these fits."
" For consistency, the V ellipse fits were applied to the B and Ha frames (after registration) to produce aperture luminosities and colors (i.e. the same pixels are integrated by aperture in all frames)."," For consistency, the $V$ ellipse fits were applied to the $B$ and $\alpha$ frames (after registration) to produce aperture luminosities and colors (i.e. the same pixels are integrated by aperture in all frames)."
" All the reduced values, single parameters (e.g., total luminosity) and array values (e.g., surface"," All the reduced values, single parameters (e.g., total luminosity) and array values (e.g., surface"
(85.2).,5.2).
 In 85.3 we discuss the variation of K with accretion rate (Fig., In 5.3 we discuss the variation of $K$ with accretion rate (Fig.
 2) in the context of the spectral models., 2) in the context of the spectral models.
 Suleimanov et al. (, Suleimanov et al. (
"2011b) calculated f. as a function of F/Fgaa for a range of surface gravities and atmospheric compositions, and discussed how these models could be applied to data.","2011b) calculated $f_c$ as a function of $F/F_{\rm Edd}$ for a range of surface gravities and atmospheric compositions, and discussed how these models could be applied to data."
" We follow their analysis, andfit the theoretical fc-F/Fraa curves to the observed relation between ΚΤΙ/4 and flux F."," We follow their analysis, andfit the theoretical $f_c$ $F/F_{\rm Edd}$ curves to the observed relation between $K^{-1/4}$ and flux $F$."
 The fitting parameters are Fgag and A=K/ fe: , The fitting parameters are $F_{\rm Edd}$ and $A=K^{-1/4}/f_c$ .
"Comparing with equation (4)), we see that If f. was a constant independent of flux, and K was constant in the cooling tail of the burst, fitting for A would be equivalent to using the measured normalization K and a value of f. to extract Rs;/d."," Comparing with equation \ref{eq:K}) ), we see that If $f_c$ was a constant independent of flux, and $K$ was constant in the cooling tail of the burst, fitting for $A$ would be equivalent to using the measured normalization $K$ and a value of $f_c$ to extract $R_\infty/d$."
" Instead, here we are using the entire cooling track to obtain A."," Instead, here we are using the entire cooling track to obtain $A$."
" In addition, by fitting the shape of the cooling track we can obtain the overall flux scale Fgaa even though the burst itself does not reach Eddington luminosity."," In addition, by fitting the shape of the cooling track we can obtain the overall flux scale $F_{\rm Edd}$ even though the burst itself does not reach Eddington luminosity."
 We start the data from wwith recurrenceby fittingtime 5.74 hr to the different models from Suleimanov et al. (, We start by fitting the data from with recurrence time $5.74$ hr to the different models from Suleimanov et al. (
2011b).,2011b).
" Below we use the fits to obtain an upper limit on R.,, which motivates us to start with the bursts with the largest value of K and therefore larger Rss values."," Below we use the fits to obtain an upper limit on $R_\infty$, which motivates us to start with the bursts with the largest value of $K$ and therefore larger $R_\infty$ values."
" At the end of this section, we discuss whether it is possible to include the 4.07hr recurrence time bursts which have smaller values of K (Fig."," At the end of this section, we discuss whether it is possible to include the $4.07\ {\rm hr}$ recurrence time bursts which have smaller values of $K$ (Fig."
 2) in a consistent picture., 2) in a consistent picture.
" When fitting, for we calculate X? based on ΚΙ and fc, and simplicitydo not include the errors in the flux comparingmeasurement."," When fitting, for simplicity we calculate $\chi^2$ based on comparing $K^{-1/4}$ and $f_c$, and do not include the errors in the flux measurement."
" This seems reasonable because the observational error in the overall flux scale is zz6F/V/N, where the individual flux error oFez10?ergcm?s!, smaller than the overall uncertainty in the parameter Fgaq that we obtain from our fits."," This seems reasonable because the observational error in the overall flux scale is $\approx \delta F/\sqrt{N}$, where the individual flux error $\delta F\approx 10^{-9}\ {\rm erg\ cm^{-2}\ s^{-1}}$, smaller than the overall uncertainty in the parameter $F_{\rm Edd}$ that we obtain from our fits."
 Suleimanov et al. (, Suleimanov et al. (
"2011b) calculate spectral models for pure H and pure He atmospheres, and solar H/He fractions with different metallicity.","2011b) calculate spectral models for pure H and pure He atmospheres, and solar H/He fractions with different metallicity."
" Based on the lightcurve comparison and energetics, Galloway et al. ("," Based on the lightcurve comparison and energetics, Galloway et al. ("
2004) and Heger et al. (,2004) and Heger et al. (
2007) concluded that the accreted layer has solar metallicity and a substantial amount of hydrogen (a solar H/He ratio in their models).,2007) concluded that the accreted layer has solar metallicity and a substantial amount of hydrogen (a solar H/He ratio in their models).
 Here we fit to the full range of models from Suleimanov et al. (, Here we fit to the full range of models from Suleimanov et al. (
2011b) to investigate how changing composition affects our derived limits on neutron star parameters.,2011b) to investigate how changing composition affects our derived limits on neutron star parameters.
" Also, the photospheric abundances could be different from the abundances near the base of the layer where the X-ray burst ignites."," Also, the photospheric abundances could be different from the abundances near the base of the layer where the X-ray burst ignites."
" For example, the metallicity in the burning layer could be enhanced by partially burned fuel left over from a previous burst."," For example, the metallicity in the burning layer could be enhanced by partially burned fuel left over from a previous burst."
" The hydrogen fraction in the accreted material could be lower than solar, and so it is useful to consider the pure He limit as a limiting case when the hydrogen fraction at the photosphere is reduced."," The hydrogen fraction in the accreted material could be lower than solar, and so it is useful to consider the pure He limit as a limiting case when the hydrogen fraction at the photosphere is reduced."
 Figure 3 shows example fits to the 5.74 hr recurrence time bursts., Figure \ref{fig:fit} shows example fits to the 5.74 hr recurrence time bursts.
" By varying A and Fgag, we are able to obtain good fits for fluxes down to approximately 1/3 of the peak flux."," By varying $A$ and $F_{\rm Edd}$ , we are able to obtain good fits for fluxes down to approximately 1/3 of the peak flux."
" At lower fluxes, the behavior of the model and observations is qualitatively similar, in that f. rises rapidly at low fluxes, but the detailed behavior does not match the models."," At lower fluxes, the behavior of the model and observations is qualitatively similar, in that $f_c$ rises rapidly at low fluxes, but the detailed behavior does not match the models."
" At late times or low fluxes, K~!/* rises more rapidly than "," At late times or low fluxes, $K^{-1/4}$ rises more rapidly than predicted."
We therefore confine the fit to the initial part of thepredicted. cooling tail and use it to derive A and Fgaa., We therefore confine the fit to the initial part of the cooling tail and use it to derive $A$ and $F_{\rm Edd}$.
" To do this, we fit Ata seconds of data starting at the time of peak flux."," To do this, we fit $\Delta t_{\rm fit}$ seconds of data starting at the time of peak flux."
" The time Δίῃι is chosen so that as much data is included in the fit as possible while still giving an adequate fit (with the late time data excluded, we find reduced X? values in the range 0.23—0.47 for the models listed in Table 1)."," The time $\Delta t_{\rm fit}$ is chosen so that as much data is included in the fit as possible while still giving an adequate fit (with the late time data excluded, we find reduced $\chi^2$ values in the range 0.23–0.47 for the models listed in Table 1)."
" For all except the solar metallicity models we take Ag,=35s, corresponding to fluxes down to z1/3 of the peak flux."," For all except the solar metallicity models we take $\Delta t_{\rm fit}=35\ {\rm s}$, corresponding to fluxes down to $\approx 1/3$ of the peak flux."
 The solar metallicity models begin to deviate from the data after ΔίῃιΞ10s (about 1/2 of the peak flux) because of the in f. at lowfluxes., The solar metallicity models begin to deviate from the data after $\Delta t_{\rm fit}=10\ {\rm s}$ (about 1/2 of the peak flux) because of the dip in $f_c$ at lowfluxes.
 The results of the fits for differentdip spectral models are listed in Table 1.., The results of the fits for different spectral models are listed in Table \ref{tab:fcfits}.
" We used Markov Chain Monte Carlo implemented with the Metropolis-Hastings algorithm (Gregory 2005, "," We used Markov Chain Monte Carlo implemented with the Metropolis-Hastings algorithm (Gregory 2005, Chap."
"12) to sample the parameter space and find the distributionsChap. for the fitting parameters, A and Fgaa."," 12) to sample the parameter space and find the distributions for the fitting parameters, $A$ and $F_{\rm Edd}$."
 Those distributions were then each fitted by a Gaussian profile in order to derive their respective central values and lo uncertainties., Those distributions were then each fitted by a Gaussian profile in order to derive their respective central values and $1\sigma$ uncertainties.
" In certain cases, as noted in Table 1,, the distributions were not well describedby a single Gaussian profile, due to the presence of more than one peak."," In certain cases, as noted in Table \ref{tab:fcfits}, , the distributions were not well describedby a single Gaussian profile, due to the presence of more than one peak."
" In those cases, we fit a Gaussian profileto the peak at the lowest values of Fgaa and A."," In those cases, we fit a Gaussian profileto the peak at the lowest values of $F_{\rm Edd}$ and $A$ ."
 Note that these parameters are correlated since, Note that these parameters are correlated since
 The Be/X-ray systems consist of a Be star and a neutron star.,"}\tikzmark{mainBodyEnd1} 

% \offprints{}

 \institute{Physics Department, Middle East Technical University, Ankara 06531, Turkey
 }


 \date{Received -- / Accepted --}



\abstract
 {}

{ The main goal of this study is to investigate the relationship 
between the optical and X-ray behaviours of the 
Be/X-ray binary system \object{\ee}~. 
 }\tikzmark{mainBodyStart2}}\tikzmark{mainBodyEnd2} 
{ We present the relations between the H$\alpha$ equivalent width, 
optical brightness and X-ray flux of the system, 
by analysing the optical photometric
and spectroscopic observations together with the X-ray observations~.
 }\tikzmark{mainBodyStart3}}\tikzmark{mainBodyEnd3}
{The X-ray outburst of the system occurred just after the optical outburst.
The nearly symmetric H$\alpha$ emission line profiles observed during the
beginning of the optical outburst turn into asymmetric profiles with increased EW 
values during the dissipation of the Be disc.
 A correlation between the H$\alpha$ emission line strength
and the X-ray flux was found. The H$\alpha$ line indicates the existence of
an equatorial disc around the Be star. The 
H$\alpha$ line changed from emission to absorption during the observation 
period.
The observed double peaked HeI emission lines might come from the accretion
disc of neutron star which is temporarily formed at the time of X-ray outburst.
 }\tikzmark{mainBodyStart4}}\tikzmark{mainBodyEnd4}
 {}\tikzmark{mainBodyStart5}} The Be/X-ray systems consist of a Be star and a neutron star."
 Be stars are close to the main-sequence., Be stars are close to the main-sequence.
 They exhibit Balmer lines in emission and show strong infrared excess (compared to normal stars of the same stellar spectral type) in their spectra., They exhibit Balmer lines in emission and show strong infrared excess (compared to normal stars of the same stellar spectral type) in their spectra.
 The disappearance and appearance of emission lines suggest a disc structure around the Be star (Slettebak 1988; Okazaki & Negueruela 2001)., The disappearance and appearance of emission lines suggest a disc structure around the Be star (Slettebak 1988; Okazaki $\&$ Negueruela 2001).
 The disc formation is related to the material lost from the rapidly rotating Be star (Porter & Rivinius 2003)., The disc formation is related to the material lost from the rapidly rotating Be star (Porter $\&$ Rivinius 2003).
 The motions of dises are rotationally dominated and quasi Keplerian (Hanuschik 1996)., The motions of discs are rotationally dominated and quasi Keplerian (Hanuschik 1996).
 The interaction of neutron star with the Be disc (Stella et al., The interaction of neutron star with the Be disc (Stella et al.
 1956: Negueruela et al., 1986; Negueruela et al.
 1998) causes X-ray outbursts either at the periastroi passage (type I outburst with X-ray luminosity -107? - 107 erg/s) or at any orbital phase (type II outburst with X-ray luminosity greater than 1077 erg/s)., 1998) causes X-ray outbursts either at the periastron passage (type I outburst with X-ray luminosity $\sim 10^{36}$ - $10^{37}$ erg/s) or at any orbital phase (type II outburst with X-ray luminosity greater than $10^{37}$ erg/s).
 There is also X -ray quiescence phase where the accretion onto the NS ts partially or completely stopped., There is also X -ray quiescence phase where the accretion onto the NS is partially or completely stopped.
 In some Be/X-ray systems the Be disc forms. grows. causes an X-ray outburst and disappears (Negueruela et al.," In some Be/X-ray systems the Be disc forms, grows, causes an X-ray outburst and disappears (Negueruela et al."
 2001: Reig et al., 2001; Reig et al.
 2001: Reig et al., 2001; Reig et al.
 2005b: Baykal et al., 2005b; Baykal et al.
 2005: Reig et al., 2005; Reig et al.
 2007)., 2007).
 Some systems display an X-ray outburst and do not go to a dise loss phase (Wilson et al., Some systems display an X-ray outburst and do not go to a disc loss phase (Wilson et al.
 2002: Wilson et al., 2002; Wilson et al.
 2005; Kizilogllu et al., 2005; loğllu et al.
 2007; Baykal et al., 2007; Baykal et al.
 2008)., 2008).
 [tis also possible to observe long pertods of X-ray quiscence even when a Be dise is present (Kizilogllu et al., It is also possible to observe long periods of X-ray quiscence even when a Be disc is present loğllu et al.
 2007)., 2007).
 The favoured dise model is the viscous dise model (Okazaki 2001: Okazaki & Negueruela 2001)., The favoured disc model is the viscous disc model (Okazaki 2001; Okazaki $\&$ Negueruela 2001).
 Okazaki & Negueruela (2001) applied their resonantly truncated viscous dise model to some Be/X-ray systems., Okazaki $\&$ Negueruela (2001) applied their resonantly truncated viscous disc model to some Be/X-ray systems.
 They found that systems with high orbital eccentricity show regular type I outbursts while systems with low orbital eccentricity display occasional giant type [I outbursts., They found that systems with high orbital eccentricity show regular type I outbursts while systems with low orbital eccentricity display occasional giant type II outbursts.
 Be/X-ray systems with moderate eccentricity and relatively close orbits have efficient dise truncation which allows the material to be stored in the Be disc., Be/X-ray systems with moderate eccentricity and relatively close orbits have efficient disc truncation which allows the material to be stored in the Be disc.
 When the dise radius ts larger than the Roche lobe radius. type I outbursts result.," When the disc radius is larger than the Roche lobe radius, type I outbursts result."
 If the Be disc 1s smaller than the eritical lobe radius. type II outbursts are displayed and temporary type I| outbursts are seen when the Be disc ts strongly disturbed.," If the Be disc is smaller than the critical lobe radius, type II outbursts are displayed and temporary type I outbursts are seen when the Be disc is strongly disturbed."
 If the Be dise is elongated toward the periastron then material is transferred around the periastron forming an accretion disc around the NS., If the Be disc is elongated toward the periastron then material is transferred around the periastron forming an accretion disc around the NS.
 The transient X-ray source was discovered on February 1997 by theBeppoSAX Wide Field Camera (Hulleman et al., The transient X-ray source was discovered on February 1997 by the Wide Field Camera (Hulleman et al.
 1998)., 1998).
 The source was active between February and September 1997. with a pulse period of 358.61 s. The next activity was two years later on November 1999 and was detected by the all-sky monitor (ASM) on board ofRXTE.," The source was active between February and September 1997, with a pulse period of 358.61 s. The next activity was two years later on November 1999 and was detected by the all-sky monitor (ASM) on board of."
 Using X-ray data fromRXTE. Baykal et al. (," Using X-ray data from, Baykal et al. ("
2000) determined the orbital parameters of the system by Doppler shift analysis of the pulsations.,2000) determined the orbital parameters of the system by Doppler shift analysis of the pulsations.
 These observations indicated a high mass companion m the system., These observations indicated a high mass companion in the system.
 had another X-ray activity in. July 2002., had another X-ray activity in July 2002.
 Orbital solutions for showed that this system had a moderately eccentric orbit (e20.4) with an orbital period of 12.66 d (Baykal et al. (, Orbital solutions for showed that this system had a moderately eccentric orbit (e=0.4) with an orbital period of 12.66 d (Baykal et al. (
2007)).,2007)).
 Baykal et al. (, Baykal et al. (
2007) obtained a distance of 3520.5 kpe to this source.,2007) obtained a distance of $\pm$ 0.5 kpc to this source.
 They reported a correlation between spin-up/down trend and an X-ray flux which could be explained by accretion from an accretion dise onto the NS., They reported a correlation between spin-up/down trend and an X-ray flux which could be explained by accretion from an accretion disc onto the NS.
 Reig et al. (, Reig et al. (
2004) identified a highly reddened optical counterpart (V214.2 mag) toJ2103.,2004) identified a highly reddened optical counterpart (V=14.2 mag) to.
5+4545.. They suggested a spectral type of BOVe for the optical companion whose distance was suggested to be 6.540.9 kpe., They suggested a spectral type of B0Ve for the optical companion whose distance was suggested to be $\pm$ 0.9 kpc.
" During their spectroscopic observations. a double peak H,, emission line was observed which was an indication of a dise surrounding Be star."," During their spectroscopic observations, a double peak $_\alpha$ emission line was observed which was an indication of a disc surrounding Be star."
" Less than a month later the H, line was seen in absorption (after MJD 52870).", Less than a month later the $_\alpha$ line was seen in absorption (after MJD 52870).
 Blay et al. (, Blay et al. (
"2006) presented the evolution of H, line profiles between August 2003 and September 2004.",2006) presented the evolution of $_\alpha$ line profiles between August 2003 and September 2004.
" They noticed that the transition of the H, line from emission to absorption occurs in coincidence with the transition of the Be/X-ray system from a higher to lower X-ray activity state.", They noticed that the transition of the $_\alpha$ line from emission to absorption occurs in coincidence with the transition of the Be/X-ray system from a higher to lower X-ray activity state.
 Although Be star lost its disc after the August 2003. the X-ray emission was detected due to the accretion of material on to the NS.," Although Be star lost its disc after the August 2003, the X-ray emission was detected due to the accretion of material on to the NS."
 Reig et al. (, Reig et al. (
20053) concluded that the X-ray emission with a luminosity of 3-5 x10? erg/s during the quiescent phase must originate in the stellar wind.,2005a) concluded that the X-ray emission with a luminosity of 3-5 $\times 10^{35}$ erg/s during the quiescent phase must originate in the stellar wind.
 The system was in a low activity state since the end of 2003., The system was in a low activity state since the end of 2003.
" The first detection of the H, emission since August 2003 was made by Manousakis et al. (", The first detection of the $_\alpha$ emission since August 2003 was made by Manousakis et al. (
2007) on MJD 54234 and,2007) on MJD 54234 and
AGTOS Li I (5.39 eV) and A7698 IN I (4.34 eV) lines. two alkali elements with a similar atomic structure.,"$\lambda 6708$ Li I (5.39 eV) and $\lambda 7698$ K I (4.34 eV) lines, two alkali elements with a similar atomic structure."
 In fact. svnthetic fits to the Ix I line give low potassium abundances (undersolar).," In fact, synthetic fits to the K I line give low potassium abundances (undersolar)."
 No nuclear mechanism able (o destroy. potassium in stars is known., No nuclear mechanism able to destroy potassium in stars is known.
 This ellect with IX however. was not found bv Plez.Smith&Lambert(1993) in M giants of the Magellanic Clouds.," This effect with K however, was not found by \citet{ple93} in M giants of the Magellanic Clouds."
 Non-LTE effects (overionization) could be the reason for this as in the case of the strong resonance Li 1 line in some C-stars (Abia. 1999)., Non-LTE effects (overionization) could be the reason for this as in the case of the strong resonance Li I line in some C-stars \citep{abi99}.
. N-LTE corrections on this line might extend to until. +0.6 dex in the sense of N-LTE minus LTE abundances., N-LTE corrections on this line might extend to until +0.6 dex in the sense of N-LTE minus LTE abundances.
 Similar unexpected low |Wb/M] ratios were also derived bv Plez.Smith&Lambert(1993) and Lambertetal.(1995)., Similar unexpected low [Rb/M] ratios were also derived by \citet{ple93} and \citet{lam95}.
.. However. these authors conclude Chat NLTE effects on Rb are probably weak. from analvsis of the Rb line in Deltegeuse and. Aldebaran. two galactic M supergiants whose atmospheres are presumed to be similar to those of the O-rich ACB stars they studied.," However, these authors conclude that NLTE effects on Rb are probably weak, from analysis of the Rb line in Beltegeuse and Aldebaran, two galactic M supergiants whose atmospheres are presumed to be similar to those of the O-rich AGB stars they studied."
 A quantitative N-LTE study of the formation of the Rb line in cool C-rich atmospheres is needed before this question can be answered., A quantitative N-LTE study of the formation of the Rb line in cool C-rich atmospheres is needed before this question can be answered.
 Typical maximum uncertainties in the atmosphere parameters (ATay=£200 I: Af=+1kms|: ACNO/LE=£0.3 dex: AC/O=£0.05 dex) and ~5% in (he continu added quacratically. represent a maxiniun uncertainty of 0.4 dex in the absolute abundance of Rb derived.," Typical maximum uncertainties in the atmosphere parameters $\Delta \rm{T_{eff}=\pm 200}$ K; $\Delta\xi=\pm 1~\rm{kms}^{-1}$; $\Delta \rm{CNO/H}=\pm 0.3$ dex; $\Delta \rm{C/O=\pm 0.05}$ dex) and $\sim 5\%$ in the continuum added quadratically, represent a maximum uncertainty of $\sim 0.4$ dex in the absolute abundance of Rb derived."
 The formal uncertainty concerning the [Rb/M)] value is probably less (han this. because some of these sources of uncertainty affect the [M/II] values in a similar wav.," The formal uncertainty concerning the [Rb/M] value is probably less than this, because some of these sources of uncertainty affect the [M/H] values in a similar way."
 Thus. when deriving the |[Rb/M] ratio. many errors would be canceled out.," Thus, when deriving the [Rb/M] ratio, many errors would be canceled out."
 However. we avoid to estimate neutron densities in (he s-process site from the derived lb/Zr or Rb/Sr ratios due to the uncertain Rb abundauces as mentioned above.," However, we avoid to estimate neutron densities in the s-process site from the derived Rb/Zr or Rb/Sr ratios due to the uncertain Rb abundances as mentioned above."
the GLAST sources.,the GLAST sources.
 At least two FR I galaxies have been detected by EGRET (e.g..Padovani2007).. and more FR | galaxies were predicted to be detected by GLAST (e.g..Ghisellinietal. 2005).," At least two FR I galaxies have been detected by EGRET \citep*[e.g.,][]{p07}, and more FR I galaxies were predicted to be detected by GLAST \citep*[e.g.,][]{gtc05}."
. In the unification scheme. FR I galaxies are BL Lac objects with misaligned jets.," In the unification scheme, FR I galaxies are BL Lac objects with misaligned jets."
 The 5-rày PR I galaxies/BL Lac objects are beyond the scope of this work., The $\gamma$ -ray FR I galaxies/BL Lac objects are beyond the scope of this work.
 The EGRET quasars identified at high confidence (64 sources) have radio flux densities =1 Jy and are about of the total FSRQs (~ 300) above the same flux density limit (seeFig.6inPadovanietal.2007)., The EGRET quasars identified at high confidence (64 sources) have radio flux densities $\ga 1$ Jy and are about of the total FSRQs $\sim 300$ ) above the same flux density limit \citep*[see Fig. 6 in][]{pg07}.
. Assuming the radio flux density limit of the GLAST quasars to be ~30 times lower than the EGRET limit. i.e... ~30 mJy. the total all-sky number of the FSRQs above this limit is about 20000.," Assuming the radio flux density limit of the GLAST quasars to be $\sim 30$ times lower than the EGRET limit, i.e., $\sim 30$ mJy, the total all-sky number of the FSRQs above this limit is about $20000$."
 So. about 4000 FSRQs will be detected by GLAST. if the same percentage (~ 20%)) is adopted as the EGRET blazars.," So, about 4000 FSRQs will be detected by GLAST, if the same percentage $\sim 20$ ) is adopted as the EGRET blazars."
 This rough estimate is about a factor of two higher than our model calculations., This rough estimate is about a factor of two higher than our model calculations.
 Considering that some unidentified ~-ray sources are likely to be blazars. our estimates on >-ray quasars to be detected by GLAST are only lower limits.," Considering that some unidentified $\gamma$ -ray sources are likely to be blazars, our estimates on $\gamma$ -ray quasars to be detected by GLAST are only lower limits."
 Both the duty cycle in the 5-ray band and identification rate may affect the detection rate. and we implicitly assume them to be similar to those of EGRET blazars.," Both the duty cycle in the $\gamma$ -ray band and identification rate may affect the detection rate, and we implicitly assume them to be similar to those of EGRET blazars."
 The EGRB integrated above 100 MeV was determined to be 1.45(+0.05)«107 photons em? s! si! from the EGRET data (Sreekumaretal.1998).," The EGRB integrated above 100 MeV was determined to be $1.45(\pm
0.05)\times 10^{-5}$ photons $^{-2}$ $^{-1}$ $^{-1}$ from the EGRET data \citep{s98}."
". Strongetal.(2004) used a new model of the Galactic background. and obtained a slightly smaller value of the EGRB. 1.14(£0.12)«107? photons em s! σσ,"," \citet{s04} used a new model of the Galactic background, and obtained a slightly smaller value of the EGRB, $1.14(\pm 0.12)\times 10^{-5}$ photons $^{-2}$ $^{-1}$ $^{-1}$."
 We sum up the >-ray emission from all radio quasars/FR IIs with our derived 7-ray LF and find that they contribute ~30% of the EGRB (see Table 1). which are only lower limits. because our calculations are limited to radio quasars/FR II galaxies.," We sum up the $\gamma$ -ray emission from all radio quasars/FR IIs with our derived $\gamma$ -ray LF and find that they contribute $\sim$ of the EGRB (see Table 1), which are only lower limits, because our calculations are limited to radio quasars/FR II galaxies."
 The BL Lac objects/FR Is must contribute some part to the EGRB., The BL Lac objects/FR Is must contribute some part to the EGRB.
 The detailed calculation of their contribution to the EGRB ts beyond the scope of thisLetter., The detailed calculation of their contribution to the EGRB is beyond the scope of this.
 We thank the anonymous referee for his/her helpful comments/suggestions. and Paolo Padovani for providing us the data of the LF of FSRQs.," We thank the anonymous referee for his/her helpful comments/suggestions, and Paolo Padovani for providing us the data of the LF of FSRQs."
 This work 1s supported by the NSFC (grants 10325314. 10333020. 10573030 and 10773020). and the CAS (grant KJCX2-YW-TO3).," This work is supported by the NSFC (grants 10325314, 10333020, 10573030 and 10773020), and the CAS (grant KJCX2-YW-T03)."
detected by modern studies.,detected by modern studies.
ft a theoretical ZAMS to the Uwades. aud derive a helium coutent 3=0.26+0.02. which agrees closely with the value obtained by Lebretonetal.(2001) of Y=0.255+40.013 as well as with the solar value of Ξ0.2659.,"fit a theoretical ZAMS to the Hyades, and derive a helium content $Y = 0.26 \pm 0.02$, which agrees closely with the value obtained by \citet{LeFeLe01} of $Y = 0.255 \pm 0.013$ as well as with the solar value of $Y = 0.2659$."
 Tsochroue fitting to the ITvades main sequence turn off then eives a cluster ago of 6254:50 MM. in excellent agreement with previous deteruiuatious of MM (οαστοίbel1990). aud NX (Torresctal.1997).," Isochrone fitting to the Hyades main sequence turn off then gives a cluster ago of $625 \pm 50$ My, in excellent agreement with previous determinations of My \citep{CS90} and My \citep{ToStLa97}."
. deDruijnueetal.(2001) used the Iipparcos proper motions to derive secular parallaxes for IEvades members which they claimed were sienificautly better than the Iipparcos trigonometric parallaxes., \citet{dBHodZ01} used the Hipparcos proper motions to derive secular parallaxes for Hyades members which they claimed were significantly better than the Hipparcos trigonometric parallaxes.
 While this analysis does indeed narrow the width of the Ibvades main sequence. it does uot really alter the theoretical fit to the main sequence or the derived values of Y and Z.," While this analysis does indeed narrow the width of the Hyades main sequence, it does not really alter the theoretical fit to the main sequence or the derived values of $Y$ and $Z$."
 However. the very narrow width of the lain sequence does support the lack of significant star mnctallicity variations amoueg Ivades cwarfs.," However, the very narrow width of the main sequence does support the lack of significant star-to-star metallicity variations among Hyades dwarfs."
 While the Ivades are a nearby. homogeneous sample of stars. are they really suitable targets for high-precision radial velocity work?," While the Hyades are a nearby, homogeneous sample of stars, are they really suitable targets for high-precision radial velocity work?"
 Achieving hieh velocity precision requires SNR of about 300 in a relatively short exposure on stars that are slowly rotating and are relatively inactive., Achieving high velocity precision requires SNR of about 300 in a relatively short exposure on stars that are slowly rotating and are relatively inactive.
 Ivades dwarts range iu maeuitucde from about V.—7.5 for alate F star through V—9.11 for IX stars., Hyades dwarfs range in magnitude from about $V = 7.5$ for a late F star through $V = 9 - 11$ for K stars.
 The M chwarts range down through siguiicautlv fainter magnitudes., The M dwarfs range down through significantly fainter magnitudes.
 The necessary SNR can be achieved on these stars iun exposures of 15 miuutes or less with the Ixeck WIRES spectrograph (Vostctal.1991)., The necessary SNR can be achieved on these stars in exposures of 15 minutes or less with the Keck HIRES spectrograph \citep{VoAlBi94}.
. The IIvades dwarts. at their age of NIMS.. have slowed their rotation rates significautlv from the rotation rates found in vouuger clusters such as the Pleiades CTerudrupetal. 2000)...," The Hyades dwarfs, at their age of My, have slowed their rotation rates significantly from the rotation rates found in younger clusters such as the Pleiades \citep{TeStPi00}. ."
 ITvacdles. cvarfs show a good correlation between rotation rates and activity (Staufferctal.1997:Terudrupet2000).," Hyades dwarfs show a good correlation between rotation rates and activity \citep{StBaKr97,TeStPi00}."
.. We have selected stars of low cesin/ for our Ibvades sample., We have selected stars of low $v \sin i$ for our Hyades sample.
" The uajor question that we face in the use of IEvades dwarts for oxecise radial velocity studies is whether the level of stellar activity iu theses stars will be low enoush that we will still be able to detect RV variations due to orbital reflex notion ou top of the level of intrinsic RV jitter in these stars. Saar&Donahue(100Τὸ, Saaretal. (1998)."," The major question that we face in the use of Hyades dwarfs for precise radial velocity studies is whether the level of stellar activity in theses stars will be low enough that we will still be able to detect RV variations due to orbital reflex motion on top of the level of intrinsic RV jitter in these stars. \citet{SaDo97}, \citet{SaBuMa98},"
.. aud Santosetal.(2000) have derived enirical relationships )etyeen the observed radial velocity “jitter” of single stars westuned uot to have planetary mass compamious and observable stellar properties such as esin/ and Ri)..., and \citet{SaMaNa00} have derived empirical relationships between the observed radial velocity “jitter” of single stars presumed not to have planetary mass companions and observable stellar properties such as $v \sin i$ and $^{\prime}_{HK}$.
 The applicability of such relationships to our IHyades data set is he major subject of Paper IT in this series (Paulson 2002)., The applicability of such relationships to our Hyades data set is the major subject of Paper II in this series \citep{PaSaCo02}.
. Since 1996 we have been using the heck 1 telecope with its TIRES (Vostetal.1991). spectrograph to conduct high precision radial velocity observations of a sample of 98 TIvacdes dwart stars., Since 1996 we have been using the Keck 1 telecope with its HIRES \citep{VoAlBi94} spectrograph to conduct high precision radial velocity observations of a sample of 98 Hyades dwarf stars.
 An Is eas absorption cell provides the velocity metric (Libbrecht1988:Butleretal.1996).," An $_2$ gas absorption cell provides the velocity metric \citep{Li88,BuMaWi96}."
. The HIRES spectrograph is set to include the stellar Ca IT ASS lines (3933.66. 3968.17À).," The HIRES spectrograph is set to include the stellar Ca II K lines (3933.66, )."
. These lines fall in a wavelength region free of Is eas cell absorption. and thus provide a direct and simultancous measurement of stellar chromospheric activity for every radial velocity observation.," These lines fall in a wavelength region free of $_2$ gas cell absorption, and thus provide a direct and simultaneous measurement of stellar chromospheric activity for every radial velocity observation."
 This automatic ability to monitor stellar activity has turned out to be extremely important for the relatively voune. chromospherically active Hades cliwarts.," This automatic ability to monitor stellar activity has turned out to be extremely important for the relatively young, chromospherically active Hyades dwarfs."
 The red limit of the spectrograph setting is6188À. near the effective red lait of useful I» gas absorption.," The red limit of the spectrograph setting is, near the effective red limit of useful $_2$ gas absorption."
" The HIRES -D2"" cutrauce aperture (0.571""«7.0"" projected on the sky) gives a nominal resolving power of 60.000 and an actual measured resolving power rangiug up to 67.000 depending on position on the CCD."," The HIRES “B2” entrance aperture $0.574\arcsec \times 7.0\arcsec$ projected on the sky) gives a nominal resolving power of 60,000 and an actual measured resolving power ranging up to 67,000 depending on position on the CCD."
 The image derotator is used to keep the slit vertical ou the sky., The image derotator is used to keep the slit vertical on the sky.
 This keeps atmospheric dispersion perpendicular to the spectrographara echelle dispersion. aud presents a constant orientation of the telescope pupil to the spectrograph.," This keeps atmospheric dispersion perpendicular to the spectrograph echelle dispersion, and presents a constant orientation of the telescope pupil to the spectrograph."
 The Ίνους TIRES exposure meter. installed im 2000. is used to teruiunate exposures when the desired SNR is reached. and to give an accurate determination of the photon-weighted mid-exposure time.," The Keck HIRES exposure meter, installed in 2000, is used to terminate exposures when the desired SNR is reached, and to give an accurate determination of the photon-weighted mid-exposure time."
 We did not apply the empirical correction for problems iu the WIRES CCD readout clectronics derived by Vostetal.(2000)., We did not apply the empirical correction for problems in the HIRES CCD readout electronics derived by \citet{VoMaBu00}.
. The target stars were selected in 1996. before the Tipparcos data for the Πνασος stars were released.," The target stars were selected in 1996, before the Hipparcos data for the Hyades stars were released."
 After he analysis bv Perrvinanetal.(1998) om target List was sienificautly modified., After the analysis by \citet{PeBrLe98} our target list was significantly modified.
 Several nonanenibers were dropped. aud a few new stars were added.," Several non-members were dropped, and a few new stars were added."
 The observing ist now comprises 98 Tvades dwarts: 10 F stars. 21 CG stars. Ll IN stars and 20 M stars.," The observing list now comprises 98 Hyades dwarfs: 10 F stars, 24 G stars, 44 K stars and 20 M stars."
 Since the object of this program is to explore the dependence of plauct ormation on stellar mass. we have taken ereat care to include a representative sample of lower mass stars.," Since the object of this program is to explore the dependence of planet formation on stellar mass, we have taken great care to include a representative sample of lower mass stars."
 Due o the steepuess of the stellar miass-luninosity relation. he overwhehuing majority of the telescope time is spent observing these late Ἡ aud AL dwarfs.," Due to the steepness of the stellar mass-luminosity relation, the overwhelming majority of the telescope time is spent observing these late K and M dwarfs."
 The stars have )ocn selected to mect the following criteria: 1) classified as a Πνααος mneniber by Perrvinanetal. (1998). 2) commonly accepted by other studies as à IEvades member owed ou radial velocity. parallax. and proper motion. 3) iof in a spectroscopic binary svsteni with orbital period ess than ~100 vears. aud £4) stellar rotational velocity esin/€d5kkmss +.," The stars have been selected to meet the following criteria: 1) classified as a Hyades member by \citet{PeBrLe98}, , 2) commonly accepted by other studies as a Hyades member based on radial velocity, parallax, and proper motion, 3) not in a spectroscopic binary system with orbital period less than $\sim 100$ years, and 4) stellar rotational velocity $v \sin i \leq 15$ $^{-1}$."
 The motivation for criteria 1 and 2 is obvious: this is what euarautees the homogeneity of our suuple., The motivation for criteria 1 and 2 is obvious; this is what guarantees the homogeneity of our sample.
 The prohibition ou spectroscopic binary stars is a conmauon feature of virtually all RV planet detection xogranis., The prohibition on spectroscopic binary stars is a common feature of virtually all RV planet detection programs.
 It simply guarantees that there will not be a stellar companion in an orbit simall enough to preveut he presence of planetary conipauionus., It simply guarantees that there will not be a stellar companion in an orbit small enough to prevent the presence of planetary companions.
 The limit to siuall csnnj is also a conunon feature of ligh precision RV xoerans. and is predicated on practical consideratious.," The limit to small $v \sin i$ is also a common feature of high precision RV programs, and is predicated on practical considerations."
 Narrow stellay lines are necessary to achievethe very veh radial velocity precision required to detect planetary companions., Narrow stellar lines are necessary to achievethe very high radial velocity precision required to detect planetary companions.
 This is the one selection criterion which might introduce some small biases iu our stellar saaple., This is the one selection criterion which might introduce some small biases in our stellar sample.
 Low rotational velocities will probably select for stars with ower levels of stellar activity (Staufferetal.1997)., Low rotational velocities will probably select for stars with lower levels of stellar activity \citep{StBaKr97}.
. Iu addition. if the stellar rotation rate at the MM age of he ILvades is related to the degree of rotational braking of he star by maenetic coupling to a remnant circiunstellar disk at earlier times. we may be selectiug for stars that dad more massive disks. or stars for which the disks survived longer.," In addition, if the stellar rotation rate at the My age of the Hyades is related to the degree of rotational braking of the star by magnetic coupling to a remnant circumstellar disk at earlier times, we may be selecting for stars that had more massive disks, or stars for which the disks survived longer."
 However. in practice. our esin;criterion was actually more of a spectral class criterion.," However, in practice, our $v \sin i$criterion was actually more of a spectral class criterion."
We rejected cight F5 stars. three FG stars. aud oue G5 star based solely ou esin/.,"We rejected eight F5 stars, three F6 stars, and one G5 star based solely on $v \sin i$ ."
 Of the ten F stars remaining in the survey. we," Of the ten F stars remaining in the survey, we"
also has a number of commercial MMIC: amplifier chips. as well as a custom double balanced mixer and a LO tripler fabricated by United Monolithie Semiconductor using their GaAs Schottky diode technology.,"also has a number of commercial MMIC amplifier chips, as well as a custom double balanced mixer and a LO tripler fabricated by United Monolithic Semiconductor using their GaAs Schottky diode technology."
 In order to avoid the adverse effects of non-linearities which may arise in anv of the active components (such as amplifiers. mixers or solid state attenuators) in (he sienal path when driven into compression. the EVLA front-ends and downstream IF circuitry were designed to ensure that the standard operating point is well below the saturation level of any device.," In order to avoid the adverse effects of non-linearities which may arise in any of the active components (such as amplifiers, mixers or solid state attenuators) in the signal path when driven into compression, the EVLA front-ends and downstream IF circuitry were designed to ensure that the standard operating point is well below the saturation level of any device."
 This is a delicate balance since the signal must also be well above the noise floor of the various amplifiers in the RE/IF chain to ensure (he overall svstem temperature is not degraded., This is a delicate balance since the signal must also be well above the noise floor of the various amplifiers in the RF/IF chain to ensure the overall system temperature is not degraded.
 The headroom specifieation for the EVLA is to have all active components operating at least 20 dB below their compression points (which corresponds to being about 32 dB below the 1 dB compression point) when the antenna is looking at cold sky., The headroom specification for the EVLA is to have all active components operating at least 20 dB below their compression points (which corresponds to being about 32 dB below the 1 dB compression point) when the antenna is looking at cold sky.
 The impact of insufficient headroom will only become a major limitation in the presence of strong RFI. where unwanted harmonics and intermodulation products generated in the amplifiers or mixers may cause spurious signals to arise.," The impact of insufficient headroom will only become a major limitation in the presence of strong RFI, where unwanted harmonics and intermodulation products generated in the amplifiers or mixers may cause spurious signals to arise."
 It is expected that the 20 dB headroom specilication will provide adequate dvnamie range so that the effects of both existing and future REI will not cause saturation., It is expected that the 20 dB headroom specification will provide adequate dynamic range so that the effects of both existing and future RFI will not cause saturation.
 The LO and IF svstems in the EVLA are highly flexible systems designed to support observations in (wo modes of operation., The LO and IF systems in the EVLA are highly flexible systems designed to support observations in two modes of operation.
 During the early construction phase of the project. ihe EVLA is required to operate in a transition mode where the new electronies must operate in conjunction with unmocdified VLA antennas and provide signals io (he existing VLA correlator.," During the early construction phase of the project, the EVLA is required to operate in a transition mode where the new electronics must operate in conjunction with unmodified VLA antennas and provide signals to the existing VLA correlator."
 Later in the project. as more VLA antennas are converted to the EVLA design and the WIDAR correlator is brought online. the systems will be able to easily switch between {his transition mode and the final EVLA mode of operation.," Later in the project, as more VLA antennas are converted to the EVLA design and the WIDAR correlator is brought online, the systems will be able to easily switch between this transition mode and the final EVLA mode of operation."
 The transition mode will no longer be required once the original VLA correlator is decommissioned., The transition mode will no longer be required once the original VLA correlator is decommissioned.
 The EVLA IF system (Figure 3)) consists of three primary signal paths: one for the hieher frequency bands (Ixu. IX. Ixa and QJ). one for the lower Irequency bands (4. P. L. $5 anc," The EVLA IF system (Figure \ref{fig:if}) ) consists of three primary signal paths: one for the higher frequency bands (Ku, K, Ka and Q), one for the lower frequency bands (4, P, L, S and"
 The EVLA IF system (Figure 3)) consists of three primary signal paths: one for the hieher frequency bands (Ixu. IX. Ixa and QJ). one for the lower Irequency bands (4. P. L. $5 ancl," The EVLA IF system (Figure \ref{fig:if}) ) consists of three primary signal paths: one for the higher frequency bands (Ku, K, Ka and Q), one for the lower frequency bands (4, P, L, S and"
" Matter effects due to the difference between the weak interaction effective potential for v, with respect to sterile neutrinos. vield different total umuber aud different zenith distributions of upthroughgoiug nmous."," Matter effects due to the difference between the weak interaction effective potential for $\nu_{\mu}$ with respect to sterile neutrinos, yield different total number and different zenith distributions of upthroughgoing muons."
" 5,4; hxtween the eveuts withL<cosO«0.7 and with OL«xcosO<0 was used [9]..", $R_{meas}$ between the events with$-1 < \cos \Theta < -0.7$ and with $-0.4 < \cos \Theta < 0$ was used \cite{macro98}.
 The measured ratio. PE1.98. is comparedwith AR;1.61 and Reterite=2.03.," The measured ratio, $R_{meas}=1.38$, is comparedwith $R_\tau=1.61$ and $R_{sterile}=2.03$."
" One concludes that vy,<fSeele oscillations are excluded at the cl", One concludes that $\nmns$ oscillations are excluded at the c.l.
 An estimate of the muon energy was made through their Multiple Coulomb Scattering (AICS) in the absorbers |19[.., An estimate of the muon energy was made through their Multiple Coulomb Scattering (MCS) in the absorbers \cite{mulsca}.
" The ratios 11yoose¢ as a function of (L/L) ave in aereciueuto with 1,H<>Ho oscillations |20]..", The ratios $_{\mbox{no~osc}}$ as a function of $(L/E_\nu)$ are in agreement with $\nmnt$ oscillations \cite{mac52}.
 Tn order to reduce the effects of svstematic uncertaiuties in the MCs. MACRO used the following three indc“peident ratios (it was checked that FLUINA. IHISMOL and Bartol96 MC vield the same predictions): (1) Iligh Enerey Data: zenith distribution ratio: 1—NuvtNor Gi) Tigh Euergy Data: ij energy micasurement ratlo: HR»=NowNhigh Gi) Low Euerev Data: Rs=(Data/ALC)y1-(DateUMC)IDPUGS.," In order to reduce the effects of systematic uncertainties in the MCs, MACRO used the following three independent ratios (it was checked that FLUKA, HKKM01 and Bartol96 MC yield the same predictions): (i) High Energy Data: zenith distribution ratio: $R_1 = N_{vert}/N_{hor}$ (ii) High Energy Data: $\nu_{\mu}$ energy measurement ratio: $R_2 = N_{low}/N_{high}$ (iii) Low Energy Data: $R_3 = (Data/MC)_{IU}/(Data/MC)_{ID+UGS}$."
 The uo oscillation hypothesis had a probability P~3 De10* aud is ruled out by ~Sa., The no oscillation hypothesis had a probability $\sim$ 3 $\cdot 10^{-7} $ and is ruled out by $ \sim 5 \sigma$.
 Fitting the 3 ratios to the wy<>ve oscillation .ornmlae. one obtained+ siuPI20=1.Au?2.3-107 eV?.," Fitting the 3 ratios to the $\nmnt$ oscillation formulae, one obtained $\stheta = 1,~\Dm = 2.3
\cdot 10^{-3}$ $^2$."
 Using Bartol96. one adds the information on absolute fluxes: (iv) Iligh euerey data (systematic error z17 }: HRNineNsNacE (," Using Bartol96, one adds the information on absolute fluxes: (iv) High energy data (systematic error $\simeq$ $\%$ ): $R_4 = N_{meas}/N_{MC}$. ("
"v) Low energv senicoΜαιος, nmons (scale error 21 ks=Nincas/Nate",v) Low energy semicontained muons (scale error $21 \%$ ): $R_5 = N_{meas}/N_{MC}$.
 These leave the best ft vaues uuchauged aud inprove the significance to 69. (SIX) is a large evliudrica water Cherenkov detector of 39 ur diameter and 11 i height containiug 50 kt of water (fiducial mass 22.5 st): itis seen by 50-cni-dianeter iuner-faciug phototubes (PAITS). Fig.," These leave the best fit values unchanged and improve the significance to $6\sigma$ (SK) is a large cylindrical water Cherenkov detector of 39 m diameter and 41 m height containing 50 kt of water (fiducial mass 22.5 kt); it is seen by 50-cm-diameter inner-facing phototubes (PMTs), Fig."
 Tb., 1b.
 The 2 mi thick outer laver of water acts as an anticoincidence: it is seen by smaller PMTs., The 2 m thick outer layer of water acts as an anticoincidence; it is seen by smaller outward-facing PMTs.
 The detector is located iu the Izuuxka mine. Japan.," The detector is located in the Kamioka mine, Japan."
 Atmospheric neutrinos are detected iu Sk bv measuring the Chereukov light generated w the charged particles produced in the neutrino CC interactions with water nuclei., Atmospheric neutrinos are detected in SK by measuring the Cherenkov light generated by the charged particles produced in the neutrino CC interactions with water nuclei.
 Thanks to the high PMT coverage. the experiment detects eveits of euergies as low as 5 MeV. The large detector mass allows to collect a hig1 statistics sample ofcontained oveuts (FC) up to 5 GeV. TheFC events vield rings of Cherenkov light on the PAITs.," Thanks to the high PMT coverage, the experiment detects events of energies as low as $\sim$ 5 MeV. The large detector mass allows to collect a high statistics sample of events ) up to $\sim$ 5 GeV. The events yield rings of Cherenkov light on the PMTs."
FC events can be stbdivided into Vand events. with energies below aud above 1:PH»»GcV.," events can be subdivided into and events, with energies below and above 1.33."
FC eveuts include ouly sinele-ving eveuts. while ones (ALRING) are tyoatec as a separate category.," events include only single-ring events, while ones ) are treated as a separate category."
 Thecontained events (PC) ave CC interactions with vertex within the fiducial volume and at least a charged paricle. typically the ji. exits the detector (the light pattern is a fille circle).nmons (UPAMCU).," The events ) are CC interactions with vertex within the fiducial volume and at least a charged particle, typically the $\mu$, exits the detector (the light pattern is a filled circle). ),"
" produced by 1, from below interacting in the rock. are subdivided into stopping (CEL)οη~Y GeV) aud nuons UE,~702NO GeV) 1ο."," produced by $\nu_{\mu}$ from below interacting in the rock, are subdivided into $\langle E_\nu \rangle \sim 7$ GeV) and $\langle E_\nu \rangle \sim 70 \div 80$ GeV) \cite{skam}."
 Particle identification is performed using likeihood functions to parametrize the sharpuess of the Cherenkov rings. which are more diffused for clectrous than for πιο».," Particle identification is performed using likelihood functions to parametrize the sharpness of the Cherenkov rings, which are more diffused for electrons than for muons."
 The zenith angle distributions, The zenith angle distributions
"There several possible photon emission mechanisms operating in the vacuum breakdown. each leading to a dillerent. ""death fine”. or more precisely. “death boundary.","There several possible photon emission mechanisms operating in the vacuum breakdown, each leading to a different “death line”, or more precisely, “death boundary”."
 We consider them separately., We consider them separately.
" ns ... ⋅ . EN . ⋅∙3∕−∣∙⇁↾∐↕⋖⋅≼⇍↓↕⋜⊔⋅⋯↛↿∢⊾↓⋅↓⊳∖⇂⊔∼⋖⋅⊔⋖⊾↓⋅⋏∙≟∙∖⇁∪⇂⋜↧↓≻↓↥∪⋯⊔⋖⊾⊔∐∏∢⊾∠⇂↿↓⊔⋅∪⊔⋏∙≟↓↕≼⇍⊔↓⋅∖⇁⋜⋯⊔⋅⋖⊾↓⋅⋯⊔⋜∐↓∪⊔↓⊳∖↙∿↿∖⇀∫≻−≽⊐⋎⋟∣⊓⋅⊼∖⊽∕∡∶≟⋅⋀≟↓∪⊰≪⋟≻∣∣⊺≟∫↗∩⋟⊼∖⊽∕∩toy, where Re=R,/(10""em). and we have used 5ση, ‘Eq. (53))]."," The characteristic energy of a photon emitted through curvature radiation is $\epsilon \sim (3/2)\gamma^3 \hbar c/{\cal R}_c =
4.74\times 10^9 \beta_Q^3 h_3^6 P_0^{-3} {\cal R}_6^{-1}$ eV,where ${\cal R}_6={\cal R}_c/(10^6~{\rm cm})$, and we have used $\gamma\sim\gamma_m$ [Eq. \ref{gammaeq}) )]."
" The emission angle is 8,—+ which is typically much less ↿↓↕⋜⋯∣∣⊤∖⊽∕∡↿∖↿↓↥⋠↓≱∖≼∙⋜⋯∣⋡∢⊾⋖⊾⋜↧⊳∖∐∙∖⇁≼∼↓↕⋖⋅≼∙↳⋖⋅"," The emission angle is $\theta_e\sim \gamma^{-1}$, which is typically much less than $h/{\cal
R}_c$ (this can be easily checked )."
∠⇂∣∣∣↗∩∣∖∣↙∣⊳∣∩∣⊳∣↴∣∸⊲⊏↥⇂⇂⋜↧⇂↕⋖≱↓↥↿∖≼≨↓∃∃⇂↓↥∢⊾⊔↓⋅⋖⋅∠⇂⋯∙∢⋅≱∖∣∪ h Pda: ΟΛ)οORE The rate of energy loss of an electron or positron emitting curvature radiation is Loy=2Y.ghus the number of photons emitted through curvature radiation bv a single electron or positron across the gap is Toaen (2 hepoak," Equation \ref{chieq}) ) then reduces to h > = 546 The rate of energy loss of an electron or positron emitting curvature radiation is $P_{\rm CR} = 2e^2\gamma^4/(3c^3)(c^2/{\cal R}_c)^2$, thus the number of photons emitted through curvature radiation by a single electron or positron across the gap is = 17.6 _Q h_3^3."
 The condition Nin2A Eq. (57))] then gives ⋅ü∣∣↓⊔⋯⊽∣⇁∶∆⊰↖∖≟∕∖∶, The condition $N_{\rm ph}>\lambda$ [Eq. \ref{secondceq}) )]
↓⇀⊰⊰↓⊽⋧∫↗⋟⊼∖⊽∕⋟≼∼⊔↓⊳ Thus the minimum gap height required for vacuum breakdown is {στmax(huiuipn.μίμος).," then gives = 384 Thus the minimum gap height required for vacuum breakdown is $h \simeq \max(h_{\rm min,ph},h_{\rm min,e})$."
 Combining Eqs. (55)). (62)).," Combining Eqs. \ref{hmaxeq}) ), \ref{hphCReq}) ),"
 and (64)). we have + ," and \ref{heCReq}) ), we have })< ."
This gives a necessary concition for pulsar emission ancl defines the pulsar vcleath line”., This gives a necessary condition for pulsar emission and defines the pulsar “death line”.
 bor all relevant parameter regimes. Minin2Maui. and Eq. (65))," For all relevant parameter regimes, $h_{\rm min,ph}>h_{\rm min,e}$, and Eq. \ref{CRdeatheq}) )"
 simply. becomes μμ< fines.," simply becomes $h_{\rm min,ph}< h_{\rm max}$ ."
" Phe critical pulsar spin period is thenPon where the dipole polar field is Dj»=20€D0,D,5:)! ?owith P,=Pp(10Issy"," The critical pulsar spin period is then = where the dipole polar field is $B_{12} = 2.0 (P_0 \dot{P}_{15})^{1/2}$, with $\dot{P}_{15} = \dot{P}/(10^{-15}~\rm s~s^{-1})$."
 For oX1/15 this is the same as the result of Ruclerman&Sutherland(1975)., For $\beta_Q\la 1/15$ this is the same as the result of \citet{ruderman75}.
. In Fig. S..," In Fig. \ref{deathfig},"
 we show the death lines determined. from Eq. (65)), we show the death lines determined from Eq. \ref{CRdeatheq}) )
" for the cases of AR;=1 and Ri=10017? (pure dipole field at the polar cap). with 5,=1."," for the cases of ${\cal R}_6=1$ and ${\cal R}_6=100P_0^{1/2}$ (pure dipole field at the polar cap), with $b_d=1$."
 Lere the high-enerey photons in the cascade are produced by Compton upscatterings of thermal photons from the neutron star surface., Here the high-energy photons in the cascade are produced by Compton upscatterings of thermal photons from the neutron star surface.
 Resonant scattering in strong magnetic fields (e.g. Llerold 1979)) can be thought of as resonant absorption (where the electron makes a transition from the ground. Landau level to the first excited level) followed by radiative decay.," Resonant scattering in strong magnetic fields (e.g., \citealt{herold79}) ) can be thought of as resonant absorption (where the electron makes a transition from the ground Landau level to the first excited level) followed by radiative decay."
 Resonance occurs when the photon energy in the electron rest [rame satislies €στ€.—h(eBmc)=Jomoc., Resonance occurs when the photon energy in the electron rest frame satisfies $\epsilon'\simeq \epsilon_c= \hbar (eB/m_ec)= \beta_Q m_ec^2$.
 The resonant photon energy (in the “lab” frame) before scattering is ei—€σα.cos8;)]. where 8; is the incident angle (the angle between the incident photon momentum and the electron. velocity).," The resonant photon energy (in the “lab” frame) before scattering is $\epsilon_{\rm i}=\epsilon_c/[\gamma(1-\cos\theta_i)]$, where $\theta_i$ is the incident angle (the angle between the incident photon momentum and the electron velocity)."
 After absorbing5 a photon. the electron Lorentz factor dropsἹ to 2omσα|2359) 77. and then radiatively decays isotropically in its rest frame.," After absorbing a photon, the electron Lorentz factor drops to $\gamma_e=\gamma/(1+2\beta_Q)^{1/2}$ and then radiatively decays isotropically in its rest frame."
" Thecharacteristic photon energy after resonant scattering is therefore (c.g.. Beloborodov&""Thompson 2007)) with typical emission angle 8,— 1/5."," Thecharacteristic photon energy after resonant scattering is therefore (e.g., \citealt{beloborodov07}) ) = ) m_ec^2, with typical emission angle $\theta_e\sim 1/\gamma_e$ ."
 The condition fy<h see Eq. (661))], The condition $l_{\rm ph}<h$ [see Eq. \ref{chieq}) )]
 becomes, becomes
similar calcium abundance. no evidence of a cise is found in JUN and Spitzer observations (αρetal.2010b).,"similar calcium abundance, no evidence of a disc is found in JHK and Spitzer observations \citep{far2010b}."
. In conclusion. our analysis of the abundance pattern in GALIZNJI931|O11 and the absence of a close companion. as well as à comparison7 with other. similar cases such as the DAZ PGIO0I5]|161 support the Likely presence of a dusty disc around CLXALIZNJ1931|O11.," In conclusion, our analysis of the abundance pattern in GALEXJ1931+0117 and the absence of a close companion, as well as a comparison with other, similar cases such as the DAZ PG1015+161 support the likely presence of a dusty disc around GALEXJ1931+0117."
 S.V. and ur are by GA AV erant. numbers LAA300030908 -ane LX:NOM‘in30901. respectively. and by CA ο grant P20!Contre0967.," S.V. and A.K. are supported by GA AV grant numbers IAA300030908 and IAA301630901, respectively, and by GA ČRR grant number P209/10/0967."
" AN. also acknowledges support from the for ""Theoretical Astrophysics (LOCOGOLL).", A.K. also acknowledges support from the Centre for Theoretical Astrophysics (LC06014).
 We thank M. Wilie for useful discussions and the referee D. Homeier for helpful comments., We thank M. Kilic for useful discussions and the referee D. Homeier for helpful comments.
structure core.,structure core.
A The structure persistence mechanism allows laree-amplitucle fluctuations to persist for many eddy-turnover times.,  The structure persistence mechanism allows large-amplitude fluctuations to persist for many eddy-turnover times.
A As the turbulent clecavs these structures eventually dominate the statistics of the system.,  As the turbulent decays these structures eventually dominate the statistics of the system.
A The spatial structure of the density field associated with localized cirenlarly svinmetrie current filaments was shown [rom analvtical theory (Terry&Smith2007). to vield a Lévvsy-distributed density gradient field.,  The spatial structure of the density field associated with localized circularly symmetric current filaments was shown from analytical theory \citep{terry-smith07} to yield a Lévvy-distributed density gradient field.
A The kurtosis for the enrrent field was significantly larger than the Gaussian-valued kurtosis of 3. indicating enhanced tails.,"  The kurtosis for the current field was significantly larger than the Gaussian-valued kurtosis of 3, indicating enhanced tails."
A The electron density and magnetic field kurtosis values were not significantly ereater than 3A. However. just as the current is non-Gaussian when the magnetic field is not. it is expected that numerical solutions should show non-Gaussian behavior for the clensily eracient.,"  The electron density and magnetic field kurtosis values were not significantly greater than 3.  However, just as the current is non-Gaussian when the magnetic field is not, it is expected that numerical solutions should show non-Gaussian behavior for the density gradient."
A In the present paper. densitv gradient statistics are measured ancl found to be non-Gaussian.,"  In the present paper, density gradient statistics are measured and found to be non-Gaussian."
A Rather than relving on kurtosis values alone. the probability density fanctions (PDF's) are computed from ensembles of numerical solutions. showing Gaussian PDFs for the density gradient field.,"  Rather than relying on kurtosis values alone, the probability density functions (PDFs) are computed from ensembles of numerical solutions, showing non-Gaussian PDFs for the density gradient field."
 The previous stuclies of filament generation in KAW turbulence leave significant unanswered «questions relating to structure morphology ancl its effect on scintillation., The previous studies of filament generation in KAW turbulence leave significant unanswered questions relating to structure morphology and its effect on scintillation.
À It is well established that MILD turbulence admits structures that are both filament-like and sheet-like.,  It is well established that MHD turbulence admits structures that are both filament-like and sheet-like.
"AÀ Can sheet-like structures arise in KAW turbulence?A Iso. what are the conditions or parameters favoring one tvpe of απότο versus the other?A. If sheet-like structures dominate in some circumstances, what are the statistics of the density eracient?A Can (μον be sulliciently non-Gaussian (o be compatible with pulsar scintillation scaling?","  Can sheet-like structures arise in KAW turbulence?  If so, what are the conditions or parameters favoring one type of structure versus the other?  If sheet-like structures dominate in some circumstances, what are the statistics of the density gradient?  Can they be sufficiently non-Gaussian to be compatible with pulsar scintillation scaling?"
 It is desirable to consider such questions prior to calculation of rf wave scattering properties in the clensily gradient fields obtained [rom numerical solutions., It is desirable to consider such questions prior to calculation of rf wave scattering properties in the density gradient fields obtained from numerical solutions.
 In (his paper. we show that both current filaments aud current sheet structures naturally arise in numerical solutions of a decaving KAW turbulence model.," In this paper, we show that both current filaments and current sheet structures naturally arise in numerical solutions of a decaying KAW turbulence model."
A Each has a structure of (he same tvpe and at the same location in (he electron clensity gradient,  Each has a structure of the same type and at the same location in the electron density gradient.
A. These structures become prevalent as the numerical solutions progress in time. and each is associated. with highly non-Gaussian PDFs.,"  These structures become prevalent as the numerical solutions progress in time, and each is associated with highly non-Gaussian PDFs."
 Moreover. we show that small-seale current filaments and current sheets. along with their associated density. structures. are highly sensitive to tlie magnitude of resistive damping and diffusive damping of density fluctuations.," Moreover, we show that small-scale current filaments and current sheets, along with their associated density structures, are highly sensitive to the magnitude of resistive damping and diffusive damping of density fluctuations."
 Current filaments persist provided that resistivity 7 is small: similarly. electron density fluctuations and gradients are diminished by large diffusive camping in (he electron continuity equation.," Current filaments persist provided that resistivity $\eta$ is small; similarly, electron density fluctuations and gradients are diminished by large diffusive damping in the electron continuity equation."
 The latter results from collisions assuming clensity fIuctuations are subject to a Fick’s law for diffusion., The latter results from collisions assuming density fluctuations are subject to a Fick's law for diffusion.
 The magnitude of the resistivity affects (1) whether current filaments can become large in amplitude. (2) (heir spatial scale. and (3) the preponderance of these filaments as compared to sheets.," The magnitude of the resistivity affects (1) whether current filaments can become large in amplitude, (2) their spatial scale, and (3) the preponderance of these filaments as compared to sheets."
 The maenitude of the diffusive damping parameter. j. similarly influences the amplitude of density gracients ancl. (o a lesser degree. influences (he extent to which electron density structures are non-localized.," The magnitude of the diffusive damping parameter, $\mu$, similarly influences the amplitude of density gradients and, to a lesser degree, influences the extent to which electron density structures are non-localized."
the PINTof[ALE package. and the CIILANTTI project for making publicly available the results of their substantial effort in assembling atomic data useful for coronal plasma analysis.,"the PINTofALE package, and the CHIANTI project for making publicly available the results of their substantial effort in assembling atomic data useful for coronal plasma analysis."
 Bracl Wareelin is thanked for useful comments., Brad Wargelin is thanked for useful comments.
 JJD was supported by NASA contract to theCenter curing the course of this research., JJD was supported by NASA contract NAS8-39073 to the during the course of this research.
 X-ray data analysis in Tiibbingen is supported by the DLR. under grant 00201., X-ray data analysis in Tübbingen is supported by the DLR under grant 0201.
The radial metallicity gradient — re. the behaviour of [Fe/H] with Galactocentric radius ως) — provides a powerful observational constraint on models of Galactic disk formation and chemical evolution (e.g.. Tosi 1996:; Twarog. Ashman. Anthony-Twarog 1997:; Freeman Bland-Hawthorn 2002)),"The radial metallicity gradient – i.e. the behaviour of [Fe/H] with Galactocentric radius $R_{\rm gc}$ ) – provides a powerful observational constraint on models of Galactic disk formation and chemical evolution (e.g., Tosi \cite{tosi96}; Twarog, Ashman, Anthony-Twarog \cite{twarog}; ; Freeman Bland-Hawthorn \cite{freeman}) )."
 At the same time. the ratios of the abundances of different elements to Fe. are fundamental for understanding the role of stars with different masses in the Galactic chemical enrichment. and thus to get insights on the initial mass function (IMP) at the epoch of disk formation.," At the same time, the ratios of the abundances of different elements to Fe, are fundamental for understanding the role of stars with different masses in the Galactic chemical enrichment, and thus to get insights on the initial mass function (IMF) at the epoch of disk formation."
 Open clusters are widely recognized as ideal targets for the investigation of chemical abundances and their evolution. due to their internal homogeneity in age and composition. to the fact that their distance is in principle easy to measure. and to their broad coverage of ages and Galactocentric distances.," Open clusters are widely recognized as ideal targets for the investigation of chemical abundances and their evolution, due to their internal homogeneity in age and composition, to the fact that their distance is in principle easy to measure, and to their broad coverage of ages and Galactocentric distances."
 For this reason. many spectroscopic determinations of abundances in open clusters have been carried out during the last years.," For this reason, many spectroscopic determinations of abundances in open clusters have been carried out during the last years."
 Among them we cite those based on low resolution spectroscopy by Friel (1995)) and Friel et al. (2002)), Among them we cite those based on low resolution spectroscopy by Friel \cite{friel95}) ) and Friel et al. \cite{friel02}) )
 which focused on metallicity ([Fe/H]). and those based on high resolution. by Friel et al. (2003:; 2005))," which focused on metallicity ([Fe/H]), and those based on high resolution, by Friel et al. \cite{friel03}; \cite{friel05}) )"
 Carretta et al. (2004... 2005)).," Carretta et al. \cite{carretta04}, \cite{carretta05}) ),"
 Carraro et al. (2004.. 2006.. 2007)).," Carraro et al. \cite{carraro04}, \cite{carraro6791}, \cite{carraro07}) ),"
 Yong. Carney. de Almeida (2005)). most of which present the analysis of several other elements besides iron.," Yong, Carney, de Almeida \cite{yong05}) ), most of which present the analysis of several other elements besides iron."
 In. spite of these observational efforts. various problems/discrepancies persist. especially regarding the abundance gradients.," In spite of these observational efforts, various problems/discrepancies persist, especially regarding the abundance gradients."
 For example. a consensus on the shape of the gradient has not been reached yet.," For example, a consensus on the shape of the gradient has not been reached yet."
 Early works suggested an unique negative slope (e.g.. Friel 1995:: Carraro. Ng. Portinari 1998.. Friel et al. 2002)):," Early works suggested an unique negative slope (e.g., Friel \cite{friel95}; Carraro, Ng, Portinari \cite{carraro98}, Friel et al. \cite{friel02}) );"
 but Twarog et al. (1997)), but Twarog et al. \cite{twarog}) )
 and Corder Twarog (2001)) favor a step-like distribution of the Fe content with Galactocentric distance., and Corder Twarog \cite{corder}) ) favor a step-like distribution of the Fe content with Galactocentric distance.
 Very recent results (e.g.. Carraro et al. 2004.. 2007::," Very recent results (e.g., Carraro et al. \cite{carraro04}, \cite{carraro07};"
 Yong et al. 2005)), Yong et al. \cite{yong05}) )
 suggest that the gradient might become flat for Νις larger than about 14 kpe., suggest that the gradient might become flat for $R_{\rm gc}$ larger than about 14 kpc.
 In this context we have carried out two VLT/FLAMES programs. aimed at deriving the gradient. based on the homogeneous analysis of a large sample of open clusters.," In this context we have carried out two VLT/FLAMES programs, aimed at deriving the gradient, based on the homogeneous analysis of a large sample of open clusters."
 With the fiber link to UVES we observed giant stars i orderto obtain accurate element abundances based on high resolution data., With the fiber link to UVES we observed giant stars in orderto obtain accurate element abundances based on high resolution data.
 The first program was based on Guaranteed Time Observations (GTO; Pallavicini et al. 2006)), The first program was based on Guaranteed Time Observations (GTO; Pallavicini et al. \cite{pallavic06}) )
 and included three open clusters. while within the other one (Randich et al. 2005))," and included three open clusters, while within the other one (Randich et al. \cite{messenger}) )"
 nine clusters were observed., nine clusters were observed.
 The targets span a wide range of ages (from ~ | Gyrup to ~7 Gyr). metallicities (|Fe/H from ~—O0.5 up to ~+0.40) and Galactocentrie radi (~6-22 kpe): In Sestito et al. (," The targets span a wide range of ages (from $\sim$ 1 Gyr up to $\sim$ 7 Gyr), metallicities ([Fe/H] from $\sim-0.5$ up to $\sim$ +0.40) and Galactocentric radii $\sim$ 6–22 kpc); In Sestito et al. ("
2006 — hereafter Paper 1) and Bragaglia et al. (,2006 – hereafter Paper ) and Bragaglia et al. (
2008 — hereafter Paper m). the analysis of Fe and other elements in five of the observed clusters (NGC 2660. NGC 2324. NGC 2477. NGC 3960. and Be 32) was presented: in Sestito et al.,"2008 – hereafter Paper ), the analysis of Fe and other elements in five of the observed clusters (NGC 2660, NGC 2324, NGC 2477, NGC 3960, and Be 32) was presented; in Sestito et al."
 (2007. — hereafter S07) we reported on abundances in the very metal rich cluster NGC 6253., \cite{ngc6253} – hereafter S07) we reported on abundances in the very metal rich cluster NGC 6253.
 In this paper we present the analysis of the two far. old. anticenter clusters in the sample by Randich et al. (2005::," In this paper we present the analysis of the two far, old, anticenter clusters in the sample by Randich et al. \cite{messenger};"
 Berkeley 20. and Berkeley 29). and of two old clusters observed within the GTO program (Cr 261 and Mel 66; Pallavicini et al. 2006)).," Berkeley 20, and Berkeley 29), and of two old clusters observed within the GTO program (Cr 261 and Mel 66; Pallavicini et al. \cite{pallavic06}) )."
 Be 20 and Be 29are the most distant clusters in the sample. while Cr 261 is the one with the lowest A...," Be 20 and Be 29are the most distant clusters in the sample, while Cr 261 is the one with the lowest $R_{\rm gc}$ ."
 For all these clusters, For all these clusters
C. whose value must be fitted cupirically to the observed ταις cllipticity of galaxies.,"$C$, whose value must be fitted empirically to the observed rms ellipticity of galaxies."
. ThisTEM means that fe]~£?.D , This means that $| \epsilon | \propto L^2$.
Calaxies acquire angular momentum during formation by tidal torques., Galaxies acquire angular momentum during formation by tidal torques.
 Tt can be shown !that the augular momentum acquired to first order iu the gravitational potential P(x). is L;«XειLpéh;cb(x). where Dj is the protogalaxy’s inertia tensor.," It can be shown \cite{white84}
 that the angular momentum acquired to first order in the gravitational potential $\Phi({\bf x})$, is $ L_{i} \propto \epsilon_{ijk} I_{kl}\partial_{l} \partial_{j}\Phi({\bf x}) \;,$ where $I_{ij}$ is the protogalaxy's inertia tensor."
 The ellipticitv componueuts >; cau thus be calculated in terms of the eravitational potential., The ellipticity components $\gamma_i$ can thus be calculated in terms of the gravitational potential.
 I then decompose the spin-2 cllipticity field iuto scalar (electric-type E-anuode) aud pseudo-scalar GQuaguetic-tvpe B-mode) fields iu Fourier space., I then decompose the spin-2 ellipticity field into scalar (electric-type $E$ -mode) and pseudo-scalar (magnetic-type $B$ -mode) fields in Fourier space.
 This enables the construction of 3D power spectra for the ellipticities. which are convolutious over the density power spectrun.," This enables the construction of 3D power spectra for the ellipticities, which are convolutions over the density power spectrum."
 Finally. the Liniber approximation is used to obtain the predicted angular power spectrum for different source galaxy. distributions.," Finally, the Limber approximation is used to obtain the predicted angular power spectrum for different source galaxy distributions."
" To demonstrate the results I use a low aud a high redshift source distribution (with mean source redshifts (24,5=OL aud 1.0).", To demonstrate the results I use a low and a high redshift source distribution (with mean source redshifts $\langle z_{\rm src}\rangle=0.1$ and 1.0).
 The cllipticity-cllipticity aneular power spectra obtained are shown in Fie. 1..," The ellipticity-ellipticity angular power spectra obtained are shown in Fig. \ref{fig:l2cl},"
 along with the corresponding weal. lensing prediction., along with the corresponding weak lensing prediction.
 The most obvious feature in Fie., The most obvious feature in Fig.
 1 is that for a low redshift survey. intrinsic correlations are expected to dominate over weak lensing signal. while with Cae?=140. they are a very small contaiuiuant to weak leusiug measureimoeuts.," \ref{fig:l2cl} is that for a low redshift survey, intrinsic correlations are expected to dominate over weak lensing signal, while with $\langle z_{\rm src}\rangle=1.0$, they are a very small contaminant to weak lensing measurements."
 The reasons for this result are twofold: the lensing sigual is proportional to the projected density which increases with survey depth. while the intrinsic signal is increasingly washed out by projection effects for deeper survevs.," The reasons for this result are twofold: the lensing signal is proportional to the projected density which increases with survey depth, while the intrinsic signal is increasingly washed out by projection effects for deeper surveys."
 The predicted amplitudes are in good agrecment with the caleulatious of Crittenden alt. and comparable to the observational results of Brownal over augular scales from LOO100%," The predicted amplitudes are in good agreement with the calculations of Crittenden \cite{cnpt00a}, and comparable to the observational results of Brown \cite{brownetal00} over angular scales from $10'-100'$."
" They are also comparable to fiudiues frou. uiuucerical sinniations 9,"," They are also comparable to findings from numerical simulations \cite{CroMet00,HeaRefHey00}."
 The shape of the angular power spectrum reflects the shape of the uuderling deusity⋅ power spectrum AZ)(E). going. roughly flat ou sinall scales.," The shape of the angular power spectrum reflects the shape of the underlying density power spectrum $\Delta^{2}_{m}(k)$, going roughly flat on small scales."
 On: large scales the loe-slope is 2. cousistent with shot noise.," On large scales the log-slope is $2$, consistent with shot noise."
 The shapes of the intrinsic and lensing power spectra are quite snülur. makine it difficult to use this to separate the two componcuts.," The shapes of the intrinsic and lensing power spectra are quite similar, making it difficult to use this to separate the two components."
 A potentially better discriminant is given by the relative levels of E- iud B- power (this was also investigated by Crittenden af? im their model)., A potentially better discriminant is given by the relative levels of $E$ - and $B$ -mode power (this was also investigated by Crittenden \cite{cnpt00b} in their model).
 Weak leusiug shear has no handeduess aud can therefore produce only £-, Weak lensing shear has no handedness and can therefore produce only $E$ -modes.
 Fie., Fig.
 P. showsthat the iutriusic power spectra have equal E- aud B- power on large scales. but ο23.5 on small scales.," \ref{fig:l2cl} showsthat the intrinsic power spectra have equal $E$ - and $B$ -mode power on large scales, but $C_{\ell}^{\epsilon\epsilon} / C_{\ell}^{\beta\beta} \simeq 3.5$ on small scales."
 This has, This has
in5.,in.
. is α welbstudied galaxy at a distance of(Capacciolietal.1992:Lee1993)... which places it in the outskirts of the Local Group of galaxies (vandeuBere1991).," is a well-studied galaxy at a distance of\citep{capaccioli92, lee93}, which places it in the outskirts of the Local Group of galaxies \citep{vandenbergh94}."
. The radial velocity of the galaxy is (Jobin&Carignan1990)., The radial velocity of the galaxy is \citep{jobin90}.
. has been variously classified as hr (Saudage1961:vaucenBere 1991).. Sin (Sandage&Tanunann1987) and SD(shu (deVaucouleursetal.1991).," has been variously classified as Irr \citep{sandage61,
 vandenbergh94}, Sm \citep{sandage87} and SB(s)m \citep{devaucouleurs91}."
. In this paper. we consider to be a LSB. Magellanic. late-tvpe galaxy secu edee-on at au inclination of 75 to SO (Carignan1985:Jobin&Carignan 1990).," In this paper, we consider to be a LSB, Magellanic, late-type galaxy seen edge-on at an inclination of $75^\circ$ to $80^\circ$ \citep{carignan85, jobin90}."
. appears to be. darkanatter dominated. having a dark-to-hunuinous mass ratio of 510 (Jobin&Carignan 1990).," appears to be dark-matter dominated, having a dark-to-luminous mass ratio of 5–10 \citep{jobin90}."
. Stars im the disk of have agey. as voune asvr. aud a halo of stars wit[um ages excceding las been observed (Minnitietal. 1999).," Stars in the disk of have ages as young as, and a halo of stars with ages exceeding has been observed \citep{minniti99}."
. This is consistent with models of LSB ealaxies which exhibit star formation rates that are on average low but erratic (e.e.Gerritsen&deBlok1999)., This is consistent with models of LSB galaxies which exhibit star formation rates that are on average low but erratic \cite[\eg][]{gerritsen99}.
. Since these ealaxies evolve slowly. the surrounds of may be relatively vich in agelomerations of wwhose densities have not reached those required for star formation. and which have not vet settled iuto the disk of the galaxy.," Since these galaxies evolve slowly, the surrounds of may be relatively rich in agglomerations of whose densities have not reached those required for star formation, and which have not yet settled into the disk of the galaxy."
" Already. one possible companion to is known: the (hereafter, Autlia). at a distance of from the Milkv. Way (Sarajedinietal.1997:Aparicioet 1997)."," Already, one possible companion to is known: the (hereafter, Antlia), at a distance of from the Milky Way \citep{sarajedini97,
 aparicio97}."
. Like 3109... Autlia contains stellar populations of all ages (Sarajedinictal. 1997).. with most recent star formation taking place slowly in the central regions (Aparicioctal.1997).," Like , Antlia contains stellar populations of all ages \citep{sarajedini97}, with most recent star formation taking place slowly in the central regions \citep{aparicio97}."
. In this respect. it is typical of the dwarf spleroidals in the Local Group TTucana. Carina}.," In this respect, it is typical of the dwarf spheroidals in the Local Group Tucana, Carina)."
 The radial velocity of Autla is (Fouquéetal., The radial velocity of Antlia is \citep{fouque90}.
1990).. Aparicioetal.(1997) calculate the physical separation of and Autlia to be between and180kpe.. with a maxim separation of for the pair to be gravitationally bound.," \cite{aparicio97} calculate the physical separation of and Antlia to be between and, with a maximum separation of for the pair to be gravitationally bound."
" We have acquired a deep Mage= of a field square centered on aud iucludiug Αμα,", We have acquired a deep image of a field square centered on and including Antlia.
 Our observations utilise the new A[ultibeau svstem at the prime focus of the Parkes With narrowhaucd filters installed. the instantancous observing bandwidth isMITz.. which is correlated into WwS chaunels for two products NN and YY.," Our observations utilise the new Multibeam system at the prime focus of the Parkes With narrowband filters installed, the instantaneous observing bandwidth is, which is correlated into 2048 channels for two products XX and YY."
 We centered the observing baud atMITz.. corresponding to a radial velocity ofLl. with chaunels extending cither side iu increments of|.," We centered the observing band at, corresponding to a radial velocity of, with channels extending either side in increments of."
 To image and its surroundings. scans at a rate of were made separately in the Declination and Right Ascension directions. across a square field of side leneth 5.67.," To image and its surroundings, scans at a rate of were made separately in the Declination and Right Ascension directions, across a square field of side length $5.67^\circ$."
 Adjacent scans were separated by a iaxinmni ofLo. so that each beam of the Multibemu system sampled the sskv at better than the Nyquist rate of the beam.," Adjacent scans were separated by a maximum of, so that each beam of the Multibeam system sampled the sky at better than the Nyquist rate of the beam."
 A total of 208 scans were mace. comprising ~200000 sinele-polarisation spectra.," A total of 208 scans were made, comprising $\sim 200000$ single-polarisation spectra."
" The individual spectra. recorded every5s, were baudpass-corrected. calibrated and imaged using the software developed for the PParkes All Sky Survev (IIIPASS,Barnesetal.2001).. with the HIPASS paramcters slightly modified to acconmmodate the shorter scan leneth of this survey."," The individual spectra, recorded every, were bandpass-corrected, calibrated and imaged using the software developed for the Parkes All Sky Survey \citep[HIPASS,][]{barnes01}, with the HIPASS parameters slightly modified to accommodate the shorter scan length of this survey."
 The final image. having pixels of side leneth spatially and spectrally. had aw RAIS noise level of1113s.," The final image, having pixels of side length spatially and spectrally, had an RMS noise level of."
.. The image beam (resolution) is FWITIM spatially. aud FWIIM spectrally.," The image beam (resolution) is FWHM spatially, and FWHM spectrally."
 The spectrum of comission integrated spatially over is shown ind., The spectrum of emission integrated spatially over is shown in.
 The overall profile shape is markedly asvuuuetric. having a surplus of dux in the lower velocity component of the disk.," The overall profile shape is markedly asymmetric, having a surplus of flux in the lower velocity component of the disk."
 This property of has already been noted by Πιοetal.(1980).. aud will be diseussed below.," This property of has already been noted by \cite{huchtmeier80}, and will be discussed below."
 The line centre is located at 1... beiug thecoincident midpoint of the and peak flux points on the edees of the profile.," The line centre is located at , being thecoincident midpoint of the and peak flux points on the edges of the profile."
 This is in excellent aereeimneut with all previously published values. measured bv Jobin&Carignan(1990) and measured by Huchtineier(1973).," This is in excellent agreement with all previously published values, measured by \cite{jobin90} and measured by \cite{huchtmeier73}."
" The fluxaveighted mean velocity isή, and the peak fux is detected atL|."," The flux-weighted mean velocity is, and the peak flux is detected at."
 The peak projected rotation velocity of the ealaxyv is1., The peak projected rotation velocity of the galaxy is.
 The integrated Πας of3109.. measured by the traditional zeroth-order inoinent of the spectruni and verified with an equivalent but robust statistic. ists the nucertaimty iu this value is dominated by the mucertaity in the image beam (seeBarnesctal.2001).," The integrated flux of, measured by the traditional zeroth-order moment of the spectrum, and verified with an equivalent but robust statistic, is; the uncertainty in this value is dominated by the uncertainty in the image beam \citep[see][]{barnes01}."
. This value compares well to measured frou standard TIPASS data. aud ds also recovered when the narrowband spectral data is re-imaged in the traditional wav with a Gaussian smoothing kerucl and using the mean estimator iustead of the median).," This value compares well to measured from standard HIPASS data, and is also recovered when the narrowband spectral data is re-imaged in the traditional way with a Gaussian smoothing kernel and using the mean estimator instead of the median)."
 Compared to the spectrum published by Jobin&Carignan(1990).. frou which an integrated fux of is measured. we cdotect substantially muore flux in atf all velocities.," Compared to the spectrum published by \cite{jobin90}, from which an integrated flux of is measured, we detect substantially more flux in at all velocities."
" This is expected. since Jobin&Carignan(1990) used ιο αν Large Array (VLA) in a configuration bliud © Cluission on scales much greater than on the sky,"," This is expected, since \cite{jobin90} used the Very Large Array (VLA) in a configuration blind to emission on scales much greater than on the sky."
 Of the many sinele dish studies made of 3109.. ιο flux measured here is consistent only with the value 1. measured by Epstein(1961). with je ITarvard 60-ff autenna. whose beam is sufficiently large o κος all o£ 3109.," Of the many single dish studies made of , the flux measured here is consistent only with the value , measured by \cite{epstein64} with the Harvard 60-ft antenna, whose beam is sufficiently large to “see” all of ."
. Other single dishmeasurements.," Other single dishmeasurements,"
diffuse cuuission.,diffuse emission.
 With such a plivsical configuration. it is difücult to uniquely isolate the individual chips by some cluup finding algorithin (e.g...Williams.ctal.1991).," With such a physical configuration, it is difficult to uniquely isolate the individual clumps by some clump finding algorithm \citep[e.g.,][]{will94}."
. Moreover. the detection of clumps is depenudoeut ou the aneular resolution which is available.," Moreover, the detection of clumps is dependent on the angular resolution which is available."
 Iu this paper. we select the peaks of the main structures in the -CO(J = 21) inteerated intensity map. in order to locate the individual Chuups within the ring and the molecular ridges.," In this paper, we select the peaks of the main structures in the $^{12}$ CO(J = 2–1) integrated intensity map, in order to locate the individual clumps within the ring and the molecular ridges."
 We define the clumps located outside of the 10” radius as the dust lane chips. and those within the 107 radius as the rive clips.," We define the clumps located outside of the $\arcsec$ radius as the dust lane clumps, and those within the $\arcsec$ radius as the ring clumps."
 The typical size of Ciant Molecular Clouds (GAICSs} is on the order of a few to few tens pc (Scoville 1987)., The typical size of Giant Molecular Clouds (GMCs) is on the order of a few to few tens pc \citep{scov87}.
. The size of the chumps iu the ring of NGC 1097. as detected with our svuthesized beam. is at least ~200 pc.," The size of the clumps in the ring of NGC 1097, as detected with our synthesized beam, is at least $\sim$ 200 pc."
 We are therefore detecting molecular clumps arecr than GAICs. most likely a eroup of GAICS. nuncly Caant Molecular Cloud Associations (GATAs:Vogeletal. 1998).," We are therefore detecting molecular clumps larger than GMCs, most likely a group of GMCs, namely Giant Molecular Cloud Associations \citep[GMAs;][]{vog98}."
. Tere we still use the term “chumps” to describe he individual peaks at the scale of CALA.," Here we still use the term “clumps"" to describe the individual peaks at the scale of GMA."
 We expect hat the chumps would be resolved into iudividual CMCs when higher augular resolution is available., We expect that the clumps would be resolved into individual GMCs when higher angular resolution is available.
 In Table 2.. we show the quautitics measured within oue svuthesized juna to study the kinematic properties with a hieh resolution (i.e.. one svuthesize beam) in the following sections.," In Table \ref{t.clumps}, we show the quantities measured within one synthesized beam to study the kinematic properties with a high resolution (i.e., one synthesize beam) in the following sections."
 The observed peak brightucss temperatures of oeidividual Clamps are in the range from —2 al. To show the geueral properties of the CALAs. we also measured the plysical properties integrated over their size iu Table 3.," The observed peak brightness temperatures of individual clumps are in the range from $\sim$ 2 – 8 K. To show the general properties of the GMAs, we also measured the physical properties integrated over their size in Table \ref{t.vir}."
 We define the area of the GALA by measuring the nunuber of pixels above the tweshold intensity of bo in the 2CO(J=21) integrated iintensity map., We define the area of the GMA by measuring the number of pixels above the threshold intensity of $\sigma$ in the $^{12}$ CO(J = 2–1) integrated intensity map.
" We calculated the ""equivaleut radius if the measured areas are modeled as spherical clumps.", We calculated the “equivalent radius” if the measured areas are modeled as spherical clumps.
" The results arereported in Table 3. together with the resulting Aj, integrated over the area."," The results arereported in Table \ref{t.vir}
 together with the resulting $M_{\rm H_{2}}$ integrated over the area."
" The derived values of are therefore larger than that in Table 2.. which oulv Mj,measured the mass within one svuthesized bean at the intensity peak."," The derived values of $M_{\rm H_2}$ are therefore larger than that in Table \ref{t.clumps}, , which only measured the mass within one synthesized beam at the intensity peak."
" The method to derive the Afi, will be described iu Sect. 3.2.0..", The method to derive the $M_{\rm H_{2}}$ will be described in Sect. \ref{sect-mass}.
 Several factors are essential to consider for fair conrpaurisons of our {Ας with other galaxies. such as beam size. filline factor. ete.," Several factors are essential to consider for fair comparisons of our GMAs with other galaxies, such as beam size, filling factor, etc."
 A rough comparisons show that our CALAs in the starburst rine have physical extent of ~200300 pe. which have simile order with the CALAs in other ealaxies (e.g.Rand&Ixullzuni1990:TosakiAfuraokaetal.," A rough comparisons show that our GMAs in the starburst ring have physical extent of $\sim$ 200–300 pc, which have similar order with the GMAs in other galaxies \citep[e.g.][]{rand90,tosaki07,m83}."
 2009).. We show in Figure 5. the spectra of the individual chumps selected above at their peak positions (1.0... within a svuthesized beam).," We show in Figure \ref{fig-spectra} the spectra of the individual clumps selected above at their peak positions (i.e., within a synthesized beam)."
 Most of the clumps show single eaussian profiles., Most of the clumps show single gaussian profiles.
 The gaussian EWIIM dine widths of the chumps are determined by eaussian fitting. aud are listed in Table 2..," The gaussian FWHM line widths of the clumps are determined by gaussian fitting, and are listed in Table \ref{t.clumps}."
" The line widtls we observed are quadratic suus of the intrinsic line widths aud the ealactic motion across the svuthesized beam. and the relation can be expressed as where op. is the observed EWIIM line width. 65, is the line width due to the galactic rotation aud the radial motions within the beam. and the eg; is the iutriusic lue width."," The line widths we observed are quadratic sums of the intrinsic line widths and the galactic motion across the synthesized beam, and the relation can be expressed as where $\sigma_{\rm obs}$ is the observed FWHM line width, $\sigma_{\rm rot}$ is the line width due to the galactic rotation and the radial motions within the beam, and the $\sigma_{\rm int}$ is the intrinsic line width."
 Thesystematic ealactic motion would be icelieible if the spatial resolution is small cnough., Thesystematic galactic motion would be negligible if the spatial resolution is small enough.
 We subtract σι as derived from a circular rotation nodel by fitting the rotation curve in Sect. 3.5.., We subtract $\sigma_{\rm rot}$ as derived from a circular rotation model by fitting the rotation curve in Sect. \ref{sect-kinematics}.
 The derived intrinsic line width of the chumps are reported im Table 2.., The derived intrinsic line width of the clumps are reported in Table \ref{t.clumps}.
 Iu general. there is ~10 differeuce between he observed aud intrinsic line widths. which suggests hat the ealactic cieular motions do not donünuate he line broadening at this high resolution.," In general, there is $\sim$ difference between the observed and intrinsic line widths, which suggests that the galactic circular motions do not dominate the line broadening at this high resolution."
 However. he large “intrinsic” velocity dispersion thus derived by subtracting circular motions could still be dominated by ion-circular motious. especially in the twin-peak region and in the dust lanes.," However, the large “intrinsic” velocity dispersion thus derived by subtracting circular motions could still be dominated by non-circular motions, especially in the twin-peak region and in the dust lanes."
 We will meution this effect in the cliscussion., We will mention this effect in the discussion.
 Iun Table 2.. we define the chumps based on their location aud their velocity dispersions.," In Table \ref{t.clumps}, we define the clumps based on their location and their velocity dispersions."
 The clumps in the ving are further designated by their velocity dispersion being broader or narrower than 30 km +. and muued respectively as Bl... D3. and Ni... NIL. respectively.," The clumps in the ring are further designated by their velocity dispersion being broader or narrower than 30 km $^{-1}$, and named respectively as B1,..., B3, and N1,..., N11, respectively."
 The chuups located at the dust lanes are named DI..... Dh.," The clumps located at the dust lanes are named D1,..., D5."
 adio interferometers have a discrete sampling of the ae coverage limited by both the shortest aud longest basclues., Radio interferometers have a discrete sampling of the $uv$ coverage limited by both the shortest and longest baselines.
 ere. we estimate the effects due to the nüssius information.," Here, we estimate the effects due to the missing information."
 Iu Fieure 6.. we convolved our newly combined SALA data to match the beam size of the JCNIT data (217). and overlaid the spectra to compare the fux.," In Figure \ref{fig-single}, we convolved our newly combined SMA data to match the beam size of the JCMT data $\arcsec$ ), and overlaid the spectra to compare the flux."
 The integrated COW = 21) fluxes of the JOAMT aud SALA data ave ~120 IKlans. band ~71 Kan +. respectively.," The integrated $^{12}$ CO(J = 2–1) fluxes of the JCMT and SMA data are $\sim$ 120 Kkm $^{-1}$ and $\sim$ 71 K km $^{-1}$, respectively."
" The integrated. PCO(J = 21) fluxes of the JOMT aud SMA data ave ~15 I lan + and ~7 I hans + respectively,", The integrated $^{13}$ CO(J = 2–1) fluxes of the JCMT and SMA data are $\sim$ 15 K km $^{-1}$ and $\sim$ 7 K km $^{-1}$ respectively.
 Therefore our SALA data recover ~60% and ~17% of the 12C0(J = 21) and ο = 21) fluxes 1icasured by the JCAIT. respectively.," Therefore our SMA data recover $\sim60\%$ and $\sim47\%$ of the $^{12}$ CO(J = 2–1) and $^{13}$ CO(J = 2–1) fluxes measured by the JCMT, respectively."
 If the mussing flux is attributed to the exteuded cussion. then the derived fiuxes for the compact clumps will remain reliable.," If the missing flux is attributed to the extended emission, then the derived fluxes for the compact clumps will remain reliable."
 However. the spectra of the SALA data seeii to have similar hue profiles as that of the JCAIT.," However, the spectra of the SMA data seem to have similar line profiles as that of the JCMT."
 Part of the inconsisteney could then possibly be due to the uucertaintv of the flux calibration. which are 15% and 20% for the JCAT aud SALA. respectively.," Part of the inconsistency could then possibly be due to the uncertainty of the flux calibration, which are $15\%$ and $20\%$ for the JCMT and SMA, respectively."
 We note herefore that the following quantities nieasured from the Hux will have uncertainties of at least 20% from the flux calibration., We note therefore that the following quantities measured from the flux will have uncertainties of at least $20\%$ from the flux calibration.
 There are several wavs to calculate the molecular eas nass., There are several ways to calculate the molecular gas mass.
 The conventional Πο (Lee (Neu) conversion factor is one of the methods to derivethe eas mass assuiiug hat the molecular clouds are virialized., The conventional $_{2}$ $I_{\rm CO}$ $X_{\rm CO}$ ) conversion factor is one of the methods to derivethe gas mass assuming that the molecular clouds are virialized.
" First. we use he CO emission to calculate the molecular IT? cohunn density, Nyy. aswhere Ny, for cach ehuup is calculated by adopting the Veg conversion factor of 102 cin2 (IK kins 1) 1 (Solomonetal. 1987).. Ti, is the brightuess temperature. aud deis the line width."," First, we use the $^{12}$ CO emission to calculate the molecular $_{2}$ column density, $N_{\rm H_{2}}$ , aswhere $N_{\rm H_{2}}$ for each clump is calculated by adopting the $X_{\rm CO}$ conversion factor of $\times$ $^{20}$ $^{-2}$ (K km $^{-1}$ $^{-1}$ \citep{sol87}, , $T_{\rm b}$ is the brightness temperature, and $dv$ is the line width."
 The values are listed iu Table 2.. , The values are listed in Table \ref{t.clumps}. .
Note that the Yoo las wide range of ἐς 1020 cm? (S lans !) | (Young&Scoville1991:StroneALat- 2008)...," Note that the $X_{\rm CO}$ has wide range of $\times$ $^{20}$ $^{-2}$ (K km $^{-1}$ $^{-1}$ \citep{young91,strong96,dame01,draine07,meier01,meier08}. ."
 We adopted 341079 em2 US dnos 5|) 4 to be consistent with Paper I., We adopted $\times$ $^{20}$ $^{-2}$ (K km $^{-1}$ $^{-1}$ to be consistent with Paper I.
"atmospheres (Paquetteetal.1986;Zuckerman2003) and low luminosities, allowing for the detection of cool low mass counterparts, such as debris or circumstellar disks, as well as planets or brown dwarfs (Probst1983;Zuckerman&Becklin1987a;Farihietal.2005;Burleighet 2006).","atmospheres \citep{paquette86,zuckerman03} and low luminosities, allowing for the detection of cool low mass counterparts, such as debris or circumstellar disks, as well as planets or brown dwarfs \citep{probst83,zuckerman87a,farihi05,burleigh06}."
". One form of evidence for the existence of these features is the presence of IR excess (Zuckerman&Becklin1987b;Grahametal.etal.2008,among"," One form of evidence for the existence of these features is the presence of IR excess \citep[][among others]{zuckerman87b,graham90,kilic06,vonhippel07,farihi08}."
" The likelihood of a WD having an IR excess others)..is correlated with its Teg and/or presence of photospheric metals (Farihietal.2009),, the most well studied of which is calcium."," The likelihood of a WD having an IR excess is correlated with its $_{\rm{eff}}$ and/or presence of photospheric metals \citep{farihi09}, the most well studied of which is calcium."
 The pollution of hydrogen-rich photospheres of WDs with metals could be caused by circumstellar accretion., The pollution of hydrogen-rich photospheres of WDs with metals could be caused by circumstellar accretion.
" If a planetesimal or asteroid is tidally disrupted, a circumstellar dust disk could arise (Jura2003, and likely pollute the WD’s atmosphere with detectable2008) metals."," If a planetesimal or asteroid is tidally disrupted, a circumstellar dust disk could arise \citep{jura03,jura08} and likely pollute the WD's atmosphere with detectable metals."
" With no prior information about the calcium abundance of the WDs in our study, we examine their Τομ."," With no prior information about the calcium abundance of the WDs in our study, we examine their $_{\rm{eff}}$ ."
" Of our 59 WDs, all but 3 have fitted 10000 K, with the remaining three having Teg=9684+1041 K, Το=9804+921 K, and 9924+841 K. We can estimate how many cases of IR excess we expect to see in our sample using upper and lower bounds provided by various studies in the literature."," Of our 59 WDs, all but 3 have fitted $_{\rm{eff}}>10 000$ K, with the remaining three having $_{\rm{eff}}=9684\pm1041$ K, $_{\rm{eff}}=9804\pm921$ K, and $9924\pm841$ K. We can estimate how many cases of IR excess we expect to see in our sample using upper and lower bounds provided by various studies in the literature."
 Recent work by Zuckermanetal.(2010) show that ~30% of the photospheres of white dwarfs with Teg between 13500—19500 K show evidence for being polluted with metals., Recent work by \cite{zuckerman10} show that $\sim30\%$ of the photospheres of white dwarfs with $_{\rm{eff}}$ between $13 500-19 500$ K show evidence for being polluted with metals.
" In the same temperature range, approximately 20% of DA WDs with metal bearing photospheres also show evidence for warm circumstellar material (Kilicetal."," In the same temperature range, approximately $20\%$ of DA WDs with metal bearing photospheres also show evidence for warm circumstellar material \citep{kilic08}."
" This indicates a fraction of dusty WDs to be 6%, 2008)..which we take as an upper limit."," This indicates a fraction of dusty WDs to be $6\%$, which we take as an upper limit."
" A lower limit to the fraction of warm DA WDs with dusty disks is ~1% etal. obtained from recent large WD surveys (Farihi(Liebert2009),,etal.2005;KoesterFar-ihietal.2005;Hoard2007;McCook&Sion 1999)."," A lower limit to the fraction of warm DA WDs with dusty disks is $\sim1\%$ \citep{farihi09}, obtained from recent large WD surveys \citep{liebert05,koester05,farihi05,hoard07,mccook99}."
". From our 171 (59) WDs, we therefore expect between 2—10 WDs to have IR excess."," From our 171 (59) WDs, we therefore expect between $\sim2-10$ $\sim0-4$ ) WDs to have IR excess."
" However, our reddest photometric(~0—4) measurement is F160W, a wide-H band filter with a central wavelength of 1536.9 nm."," However, our reddest photometric measurement is F160W, a wide-H band filter with a central wavelength of $1536.9$ nm."
" This is just where the spectral energy distribution of a typical warm WD would begin to show evidence for an IR excess given a companion or disk with Tg~800 K. In addition, we must also consider our photometric uncertainties."," This is just where the spectral energy distribution of a typical warm WD would begin to show evidence for an IR excess given a companion or disk with $_{\rm{eff}}\sim800$ K. In addition, we must also consider our photometric uncertainties."
 It thus seems unlikely that we would find clear evidence for an IR excess in our best sample of WDs., It thus seems unlikely that we would find clear evidence for an IR excess in our best sample of WDs.
" However, if we do find indication for IR-excess in the H band, it could be confirmed with additional observations in redder filters."," However, if we do find indication for IR-excess in the H band, it could be confirmed with additional observations in redder filters."
" We have searched the SEDs in Figure 6 of all our WD candidates for evidence of an IR. excess, focusing on the examination of the F160W band."," We have searched the SEDs in Figure \ref{fig:sed_prob} of all our WD candidates for evidence of an IR excess, focusing on the examination of the F160W band."
" Any evidence that we found for an enhanced F160W measurement appeared to be due to contamination by neighboring stars, diffraction spikes, or cosmic rays."," Any evidence that we found for an enhanced F160W measurement appeared to be due to contamination by neighboring stars, diffraction spikes, or cosmic rays."
" In addition, for each filter, we plot the observed magnitude minus the expected model magnitude as a function of effective temperature for each of our 59 WDs, shown in Figure 12.."," In addition, for each filter, we plot the observed magnitude minus the expected model magnitude as a function of effective temperature for each of our 59 WDs, shown in Figure \ref{fig:obs_exp}."
" The few outlying points in the F160W filter at lower temperature are objects that are too faint compared to the models, driven by large photometric uncertainties."," The few outlying points in the F160W filter at lower temperature are objects that are too faint compared to the models, driven by large photometric uncertainties."
" If IR. excess is present, we would expect to see points with negative values on the y-axis."," If IR excess is present, we would expect to see points with negative values on the y-axis."
" Figure 12 indicates that in the F160W band, our WDs are at most in excess of 0.2 magnitudes compared to the model."," Figure \ref{fig:obs_exp} indicates that in the F160W band, our WDs are at most in excess of 0.2 magnitudes compared to the model."
 We question whether this could indicate a cool brown dwarf companion or debris disk., We question whether this could indicate a cool brown dwarf companion or debris disk.
" Using the evolutionary models for cool brown dwarfs and exosolar giant planets of Baraffeetal.(2003),, we combine a typical WD in our sample with the contribution from a 10 Gyr dust free cool brown dwarf."," Using the evolutionary models for cool brown dwarfs and exosolar giant planets of \cite{baraffe03}, we combine a typical WD in our sample with the contribution from a 10 Gyr dust free cool brown dwarf."
" Here, we convert the model IR data from the CIT system to that of 2MASS using the transformations of Carpenter(2001), as 2MASS is a nearly identical photometric system to that in WFC3 and provides the closest comparison possible."," Here, we convert the model IR data from the CIT system to that of 2MASS using the transformations of \cite{carpenter01}, as 2MASS is a nearly identical photometric system to that in WFC3 and provides the closest comparison possible."
" In Figure 13,, we show that for masses of a companion below 0.06 Mo (60 Mj), we would be unable to detect it."," In Figure \ref{fig:sed_companion}, we show that for masses of a companion below $0.06$ $_{\odot}$ $\sim60$ $_{J}$ ), we would be unable to detect it."
" However, a cool brown dwarf with 0.06 Mo, may be detectable, as this leads to a difference in magnitude in the F160W band of >0.23 magnitudes."," However, a cool brown dwarf with $0.06$ $_{\odot}$, may be detectable, as this leads to a difference in magnitude in the F160W band of $>0.23$ magnitudes."
" Examining Figure 12,, it is clear that none of our 59 WDs have F160W in excess of 0.2 from the models."," Examining Figure \ref{fig:obs_exp}, it is clear that none of our 59 WDs have F160W in excess of 0.2 from the models."
 We conclude that we do not find any case in our sample that shows strong evidence for a genuine IR. excess., We conclude that we do not find any case in our sample that shows strong evidence for a genuine IR excess.
" Our search for planets in 47 'Tuc continues, however, with this cycle 17 dataset."," Our search for planets in 47 Tuc continues, however, with this cycle 17 dataset."
" We have time series data of two fields outside the core of 47 Tuc, one is the stare field described in Section 2 consisting of F110W and F160W observationsover ~20 days with 10000 stars, and the other field has ACS images"," We have time series data of two fields outside the core of 47 Tuc, one is the stare field described in Section \ref{sec:dataset} consisting of F110W and F160W observationsover $\sim20$ days with $10000$ stars, and the other field has ACS images"
In this paper. we make more detailed: study on the inverse-Compton processes in GIA. afterglows. especially the X-ray afterglow emission from jets in SSC-dominatec regimes.,"In this paper, we make more detailed study on the inverse-Compton processes in GRB afterglows, especially the X-ray afterglow emission from jets in SSC-dominated regimes."
 We first introduce in section 2 the whole dvnamica evolution of a beaming afterglow., We first introduce in section 2 the whole dynamical evolution of a beaming afterglow.
 In section 3 we calculate the lieht curve of svnchrotron emission and the mode parameter constraint on the SSC-dominatecl case., In section 3 we calculate the light curve of synchrotron emission and the model parameter constraint on the SSC-dominated case.
 We then discuss in section 4 the case when SSC emission. emerges clirectly in. X-rays., We then discuss in section 4 the case when SSC emission emerges directly in X-rays.
 Section 5 is a brief summary ane cdiscussion., Section 5 is a brief summary and discussion.
 Consider a beaming outllow [rom the GRB source. so called a jet. which decelerates as sweeping up the ambient medium and sideways expanding in the local sound speed.," Consider a beaming outflow from the GRB source, so called a jet, which decelerates as sweeping up the ambient medium and sideways expanding in the local sound speed."
 If the radiation energv is negligible compared to the jet kinetic energy. the jet can be regarded. as adiabatie when considering its dynamic evolution.," If the radiation energy is negligible compared to the jet kinetic energy, the jet can be regarded as adiabatic when considering its dynamic evolution."
 This is awavs the case provided the energy fraction that goes into shocked elections is ¢--- which is the common value from mocel fit to observational daa.," This is always the case provided the energy fraction that goes into shocked elections is $\eps_e\la0.1$, which is the common value from model fit to observational data."
" For a higher ος. the jet will undergo first an carly raciative stage. in whieh the afterglow light curve index is reevant to c, (Botttcher Dernier 2000: Li. Dai Lu 2002)."," For a higher $\eps_e$, the jet will undergo first an early radiative stage, in which the afterglow light curve index is relevant to $\eps_e$ (Bötttcher Dermer 2000; Li, Dai Lu 2002)."
 We consider only acliabatic cvnanmies here., We consider only adiabatic dynamics here.
" A jet. with equivalent isotropic energy Le. coasts first with initial Lorentz factor 5o until it sweeps up enough material at a deceleration time fy=(21132adrm,yt5 with » the ambient medium censitv."," A jet, with equivalent isotropic energy $E$, coasts first with initial Lorentz factor $\gamma_0$ until it sweeps up enough material at a deceleration time $t_0=(3E/32\pi\gamma_0^8nm_pc^5)^{1/3}$, with $n$ the ambient medium density."
 After fy the jet begins to decelerate., After $t_0$ the jet begins to decelerate.
 Lhe deceleration. of the jet. includes. three stages: First. when the sideways expansion is not significant compared to the initial jet open angle dy. the jet undergoes a spherical-like phase where the je Lorentz [actor decreases as οκ77 (Blandford Melxee 1976). with / the observer's time. and the jet open angle is 6=65.," The deceleration of the jet includes three stages: First, when the sideways expansion is not significant compared to the initial jet open angle $\theta_0$, the jet undergoes a spherical-like phase where the jet Lorentz factor decreases as $\gamma\propto t^{-3/8}$ (Blandford McKee 1976), with $t$ the observer's time, and the jet open angle is $\theta\simeq\theta_0$."
 Secondly. when the sideways expansion begins to dominate the dynamical evolution a ἐν=du(sg84) 77. we have 6=αι and the jet turns into a spreading phase where 5x/1/2 (Rhoads 1999).," Secondly, when the sideways expansion begins to dominate the dynamical evolution at $t_j=t_0(\gamma_0\theta_0)^{8/3}$ , we have $\theta\simeq
1/\gamma$, and the jet turns into a spreading phase where $\gamma\propto t^{-1/2}$ (Rhoads 1999)."
 Here we have assumed trat the sound. speed in the relativistic stage is comparable to light speed. ὃς€.," Here we have assumed that the sound speed in the relativistic stage is comparable to light speed, $c_s\sim c$."
 Finally. the jet. becomes non-relativistic. 7=1l. at ἐν=hyο”...i)Cy56 and 8~ 1.," Finally, the jet becomes non-relativistic, $\gamma\approx1$, at $t_n=t_0^{3/4}t_j^{1/4}\gamma_0^2$ and $\theta\sim 1$ ."
 In the non-relativisticoai phase. the sicleways expansion is not important to allect the dynamical evolution. and the shocks velocity οx/ its raclius rexdo," In the non-relativistic phase the sideways expansion is not important to affect the dynamical evolution, and the shock's velocity $v\propto t^{-3/5}$, its radius $r\propto
t^{2/5}$."
" The three special times are calculated in the following: where we have used the convention €=107C, and c.g.s units."," The three special times are calculated in the following: , where we have used the convention $U=10^xU_x$ and c.g.s units."
" Hereafter. by ""sphere we mean the fy<[om phase. by jet the £j<fἐν phase and by “NRO the ££, phase."," Hereafter, by “sphere"" we mean the $t_0<t<t_j$ phase, by “jet"" the $t_j<t<t_n$ phase and by “NR"" the $t>t_n$ phase."
" The shock accelerates the ambient electrons to high energies. with electron Lorentz factors described. by a power-law distribution: dV,ds,x5."""" [or 5.m5nte"," The shock accelerates the ambient electrons to high energies, with electron Lorentz factors described by a power-law distribution: $dN_e/d\gamma_e\propto\gamma_e^{-p}$ for $\gamma_e>\gamma_m$."
" The typical Lorentz factor of electrons is proportional to the internal energy density ο ""the shock as >eX51.", The typical Lorentz factor of electrons is proportional to the internal energy density of the shock as $\gamma_m\propto\gamma-1$.
" At the beginning it is approximated as 5,,8O610«e55 at fy. and evolves as 4,OX5 dn the relativistic regime since +]ms while in the NR phase it. becomes aXU2κ πιησο. ]xc for di."," At the beginning it is approximated as $\gamma_m\approx610\eps_e\gamma_0$ at $t_0$, and evolves as $\gamma_m \propto\gamma$ in the relativistic regime since $\gamma-1\approx\gamma$, while in the NR phase it becomes $\gamma_m\propto v^2\propto t^{-6/5}$ since $\gamma-1\propto v^2$ for NR."
 The magnetic field is also created. by he shock. commonly assumed. to carey a fraction cg of he total internal energy. behind the shock front.," The magnetic field is also created by the shock, commonly assumed to carry a fraction $\eps_B$ of the total internal energy behind the shock front."
 Vhus the energy density of magnetic field. D7ftir. is also proportional o51.," Thus the energy density of magnetic field, $B^2/4\pi$, is also proportional to $\gamma-1$."
" At the deceleration time fy. the magnetic field is B=(3igmcnnmy,c)75u.⋅−!⊐⊥⊐ later on it: evolves as Dx5 in: he relativistic regime x(77 (sphere) and xftee (jet)]. while Dxext36h in NR phase."," At the deceleration time $t_0$, the magnetic field is $B=(32\pi\eps_Bnm_pc^2)^{1/2}\gamma_0$, later on it evolves as $B\propto\gamma$ in the relativistic regime $\propto t^{-3/8}$ (sphere) and $\propto t^{-1/2}$ (jet)], while $B\propto v\propto
t^{-3/5}$ in NR phase."
 Under these conditions the svnchrotron. radiation is xoduced. with the instantaneous spectrum. described: as vower-law segments (Sari. Piran Naravan 1998).," Under these conditions the synchrotron radiation is produced, with the instantaneous spectrum described as power-law segments (Sari, Piran Narayan 1998)."
 The vpical frequeney of svnchrotron photons is relevant to the vpical electron energy. where αν is defined by. Wijers Galama (1999) and of order of unity.," The typical frequency of synchrotron photons is relevant to the typical electron energy, where $x_p$ is defined by Wijers Galama (1999) and of order of unity."
 The electrons lose. energy through both svnchrotron and SSC. and the Compton parameter Y. Ee. the ratio between the inverse-Compton to svnchrotron Luminosity. is calculatec as (Sari Din 2001) where j is the fraction of electron energy that is raciated away (bv both svnchrotron and SSC).," The electrons lose energy through both synchrotron and SSC, and the Compton parameter $Y$, i.e., the ratio between the inverse-Compton to synchrotron luminosity, is calculated as (Sari Esin 2001) where $\eta$ is the fraction of electron energy that is radiated away (by both synchrotron and SSC)."
 The synachrotron cooling frequency. Le. the frequency of the synchrotron photons raciated by those electrons. which cool on the cwnarmiucal time of the shock. is given by sincethe. electrons responsible τοsvnchrotron frequencies abovei. lose energy quickly. the racdiated fraction of electron energy. is therefore The equations (5))-(7)) show that Y and ο are correlated. and these three equations should be combined," The synchrotron cooling frequency, i.e. the frequency of the synchrotron photons radiated by those electrons which cool on the dynamical time of the shock, is given by Sincethe electrons responsible tosynchrotron frequencies above$\nu_c$ lose energy quickly, the radiated fraction of electron energy is therefore The equations \ref{eq:Y}) \ref{eq:eta}) ) show that $Y$ and $\nu_c$ are correlated, and these three equations should be combined"
To resolve this issue we require data that is precise enough to discern tiny variations in the clark energy.,To resolve this issue we require data that is precise enough to discern tiny variations in the dark energy.
 Lt is also required that data be available at a large number of redshifts o constrain the detailed temporal behaviour of dark energy., It is also required that data be available at a large number of redshifts to constrain the detailed temporal behaviour of dark energy.
 At present the only data that comes reasonably close to hese requirements is provided. by the observations of the vigh redshift tvpe la supernovae (SNe). which are believed ο be standard: candles a claim that may be doubtful considering the relatively poorly understood. physics. and he possibility. of physical mechanisms. such as dust in he inter-galactic medium that svstematically dims them.," At present the only data that comes reasonably close to these requirements is provided by the observations of the high redshift type Ia supernovae (SNe), which are believed to be standard candles — a claim that may be doubtful considering the relatively poorly understood physics, and the possibility of physical mechanisms such as dust in the inter-galactic medium that systematically dims them."
 Another concern is that the SN data sets are usually collated rom several different experiments that might have slightly dilferent. svstematics. due either to instrumental elfects or he fact that they observe different directions in the sky: for example svstematic errors in correcting Galactic or source galaxy dust extinction might leave residual anisotropies in data.," Another concern is that the SN data sets are usually collated from several different experiments that might have slightly different systematics, due either to instrumental effects or the fact that they observe different directions in the sky; for example systematic errors in correcting Galactic or source galaxy dust extinction might leave residual anisotropies in data."
 These considerations implv that to. obtain precise information about the behaviour of dark energy. we should first have a good. knowledge of the statistical properties of SNe. both random as well as systematic.," These considerations imply that to obtain precise information about the behaviour of dark energy we should first have a good knowledge of the statistical properties of SNe, both random as well as systematic."
 The standard. model of cosmology is. based on. the Cosmological Principle (Peebles1993)... which states that the Universe is homogeneous and isotropic on Large scales.," The standard model of cosmology is based on the Cosmological Principle \citep{pee93}, which states that the Universe is homogeneous and isotropic on large scales."
 Recent work provides some evidence for what is known as the Hubble Bubble (Zehavietal.1998:Jhaοἱ2007) indicating that we might be living inside a laree void that has a different value of the Llubble constant inside [roni what is outside the bubble.," Recent work provides some evidence for what is known as the Hubble Bubble \citep{zeh98,jha07} indicating that we might be living inside a large void that has a different value of the Hubble constant inside from what is outside the bubble."
 There is evidence for such large scale voids in the €MD. maps as well (Cruzctal.2005:deOliveira-Costa&Tegmark.2006:Cruzetal.2006. 2007).. suggesting that such large voids are not implausible.," There is evidence for such large scale voids in the CMB maps as well \citep{cruz05,olv06,cruz06,cruz07}, suggesting that such large voids are not implausible."
 There is no reason for us to believe that we are living at the precise centre of such a void., There is no reason for us to believe that we are living at the precise centre of such a void.
 Fherefore. if we happen to be living in such a void then various distance measures. such as the luminosity distance and the angular diameter distance. may not be isotropic.," Therefore, if we happen to be living in such a void then various distance measures, such as the luminosity distance and the angular diameter distance, may not be isotropic."
 Gupta.Saini&Laskar(2008) (hereafter GSLOS) used the extreme value statistics to show hat the two SN data sets. Riessetal.(2004). (0014) and Riessetal.(2007) (CDOT). do show some evidence or direction. dependence.," \cite{gup08} (hereafter GSL08) used the extreme value statistics to show that the two SN data sets, \cite{rie04} (GD04) and \cite{rie07} (GD07), do show some evidence for direction dependence."
 Several other works have also indicated either svstematie problems with the high-redshift SN data or directional dependence in the SN data and other xobes (Ixolatt&Lahay2001:NesserisPerivolaropoulosDell 2010)..," Several other works have also indicated either systematic problems with the high-redshift SN data or directional dependence in the SN data and other probes \citep{kol01,ness04,ness05,ness07, jain106,jain206,coo09}."
 We recently found that our analysis in GSLOS contained mistakes due to a coding error., We recently found that our analysis in GSL08 contained mistakes due to a coding error.
 In this paper our main task is 1) to provide revised. results and to extend the previous work by introducing a statistic that marginalises over the Llubble parameter instead of using its best fit. value 2) to introduce a new statistic that has a greater sensitivity to the signatures of living oll-centre inside a large void., In this paper our main task is 1) to provide revised results and to extend the previous work by introducing a statistic that marginalises over the Hubble parameter instead of using its best fit value 2) to introduce a new statistic that has a greater sensitivity to the signatures of living off-centre inside a large void.
 We also give asymptotic form for this statistic. which makes it easier to use it for large data sets.," We also give asymptotic form for this statistic, which makes it easier to use it for large data sets."
 The plan of the paper is as follows., The plan of the paper is as follows.
 In 2 we introduce the likelihood. function. marginalised. over. the Llubble constant ancl provide a comparison with what we obtain using the Bavesian statistic., In 2 we introduce the likelihood function marginalised over the Hubble constant and provide a comparison with what we obtain using the Bayesian statistic.
 In 3 we introduce the two statistics., In 3 we introduce the two statistics.
 In 4 we present our results and 5 we end with conclusions., In 4 we present our results and 5 we end with conclusions.
 lor a 5given SN the measured quantity.1 the distance modulus [ids the dillerence between the apparent and the absolute magnituce where the apparent magnitude mz) is à function of the intrinsic luminosity of a SN. the recshilt z and the cosmological parameters: and A is the absolute magnitude ofa type la SN.," For a given SN the measured quantity, the distance modulus $\mu$, is the difference between the apparent and the absolute magnitude where the apparent magnitude $m(z)$ is a function of the intrinsic luminosity of a SN, the redshift $z$ and the cosmological parameters; and $M$ is the absolute magnitude of a type Ia SN."
 The distance modulus can be expressed in terms of the luminosity distance D; as where the luminosity distance is given by where h(z:QiOx)=ffs:QuOx)/Ho. ancl thus depends only on the cosmological parameters: the matter density. ©; and. the dark energy. density Ox.," The distance modulus can be expressed in terms of the luminosity distance $D_L$ as where the luminosity distance is given by where $h(z; \Om,\Ox)=H(z; \Om,\Ox)/H_0$, and thus depends only on the cosmological parameters; the matter density $\Om$ and the dark energy density $\Ox$."
 The prescription for the variation of dark energy. with redshift has to be specifiecl separately., The prescription for the variation of dark energy with redshift has to be specified separately.
 For example. in the AC DAL model the energy density in the dark energy Qy remains a constant.," For example, in the $\Lambda$ CDM model the energy density in the dark energy $\Ol$ remains a constant."
 In Eq 1. the dependence of the measured quantity ος AL is linear., In Eq \ref{eq:mu1} the dependence of the measured quantity $\mu$ on $M$ is linear.
 Since ji depends on the logarithm of the luminosity distance it is clear that yr depends Lnearly on the logarithm of the Llubble parameter. fy., Since $\mu$ depends on the logarithm of the luminosity distance it is clear that $\mu$ depends linearly on the logarithm of the Hubble parameter $H_0$.
 Usually the data are given in terms of Eq 2.. where the constant Al has already been mareinaliseedl over.," Usually the data are given in terms of Eq \ref{eq:mu}, where the constant $M$ has already been marginalised over."
 Thus. instead. of two nuisance parameters we are left with only one the Hubble constant {10," Thus, instead of two nuisance parameters we are left with only one — the Hubble constant $H_0$."
 The Hubble parameter could in principle be measured using the observations of the low-redshift. distance-recshift relationship. however. the quantities of real interest. are the cosmological parameter. O3; and. Ox.," The Hubble parameter could in principle be measured using the observations of the low-redshift distance-redshift relationship, however, the quantities of real interest are the cosmological parameter $\Om$ and $\Ox$."
 Thus it. would be useful to marginalise over the nuisance parameter. Ly., Thus it would be useful to marginalise over the nuisance parameter $H_0$.
 Although this can be done numerically while estimating the cosmological parameters. Ht turns out that our statistic in - does not allow this.," Although this can be done numerically while estimating the cosmological parameters, it turns out that our statistic in 3.2 does not allow this."
 Ht is clear that the Llubble constant term can be eliminated by considering the dillerence of two magnituces It the error on ya; I8 0; then the standard error for pe;prj is given by a=\a;|mi. where we have assumed that the errors on the two magnitudes are statistically independent.," It is clear that the Hubble constant term can be eliminated by considering the difference of two magnitudes If the error on $\mu_i$ is $\sigma_i$ then the standard error for $\mu_i - \mu_j$ is given by $\sigma = \sqrt{\sigma^2_i +
\sigma^2_j}$, where we have assumed that the errors on the two magnitudes are statistically independent."
 If we have a laree number. IN. of SNe in the sample then dividing it into two equal halves we can form IN/2 such dilferences (assuming an even IN). and estimate parameters that would be independent. of {ο," If we have a large number, $N$, of SNe in the sample then dividing it into two equal halves we can form $N/2$ such differences (assuming an even $N$ ), and estimate parameters that would be independent of $H_0$."
 However. the degrees of freeclom reduce by half in this process. thereby reducing the amount of information that can be extracted. from data.," However, the degrees of freedom reduce by half in this process, thereby reducing the amount of information that can be extracted from data."
 Another reason for our not choosing this method is the [act that we are interested in quantifving direction dependence in data and therefore our statistics depend on the direction each SN is observed., Another reason for our not choosing this method is the fact that we are interested in quantifying direction dependence in data and therefore our statistics depend on the direction each SN is observed.
 Specifically. we consider statistics that depend on the dillerence of quantities computed. on. the opposite hemispheres and then maximizing this dillerence over all the directions.," Specifically, we consider statistics that depend on the difference of quantities computed on the opposite hemispheres and then maximizing this difference over all the directions."
 IH is clear that such a procedure does not allow the use of above method., It is clear that such a procedure does not allow the use of above method.
of cosmic structure.,of cosmic structure.
 As in 7.. we parameterize the growth rate. f(a). usingwhere Ona)=Oye EtaYy.," As in , we parameterize the growth rate, $f(a)$, usingwhere $\Omega_{\rm m}(a)=\Omega_{\rm m}
a^{-3}/E(a)^{2}$ ."
" Here. O4, is the present. mean matter densitv in units of the critical density. E(a)=Ha)fldy the evolution parameter. ££(@) the Llubble parameter (with ££, its present-day value). and 7? the growth index."," Here, $\Omega_{\rm m}$ is the present, mean matter density in units of the critical density, $E(a)=H(a)/H_0$ the evolution parameter, $H(a)$ the Hubble parameter (with $H_0$ its present-day value), and $\gamma$ the growth index."
 We calculate the evolution of the linear matter density contrast. ὃOpmféfra. Le. the ratio of matter. density IIuctuations to the cosmic mean. by solving Equation 2..," We calculate the evolution of the linear matter density contrast, $\delta\equiv \delta\rho_{\rm m}/\bar{\rho}_{\rm m}$, i.e. the ratio of matter density fluctuations to the cosmic mean, by solving Equation \ref{eq:fa}."
" As an initial condition for this differential equation. we match the value of 0 at carly times (2,= 30) tothat of GR."," As an initial condition for this differential equation, we match the value of $\delta$ at early times $z_{\rm
 t}=30$ ) tothat of GR."
 Given of). we obtain the growth factor between οι and 2. Note that ο30 is well within the dark matter dominated era. when theexpansion is expected to be decelerating.," Given $\delta(z)$, we obtain the growth factor between $z_{\rm t}$ and $z$, Note that $z_{\rm t}=30$ is well within the dark matter dominated era, when theexpansion is expected to be decelerating."
 At this time f(r) is close to 1. independent of 5.," At this time $f(a)$ is close to 1, independent of $\gamma$."
" Using (+). we calculate the linear matter power spectrum for a given wavenumboer. A. and recshilt. where Z(À&.z) is the matter transfer. function. of Cli at redshift, σι, anc n. the primordial scalar spectral index."," Using $D(z)$, we calculate the linear matter power spectrum for a given wavenumber, $k$, and redshift, where $T(k,z_{\rm t})$ is the matter transfer function of GR at redshift $z_{\rm t}$, and $n_{\rm s}$ the primordial scalar spectral index."
 Given Pk.z). the variance of the linearly evolved. density field. smoothed by a spherical top-hat window function of comoving radius {0 enclosing mass AL=πριμ9. is then where Waph) is the Fourier transform of the window function.," Given $P(k,z)$, the variance of the linearly evolved density field, smoothed by a spherical top-hat window function of comoving radius $R$ enclosing mass $M=4\pi\bar{\rho}_{\rm m}R^3/3$, is then where $W_{\rm M}(k)$ is the Fourier transform of the window function."
 Using this value for στ). in Equation 1.. we obtain nCM.z).," Using this value for $\sigma(M,z)$ in Equation \ref{eq:nmz}, we obtain $n(M,z)$."
 Equation 2. depends on the evolution of the background enerey density. through the parameter £(ea).," Equation \ref{eq:fa} depends on the evolution of the background energy density, through the parameter $E(a)$."
 Assuming an expansion history appropriate for CDM. it has been shown that. for 5~0.55. 1e growth rate calculated using Equation 2 corresponds remarkably accurately tothat. predicted. by GR.," Assuming an expansion history appropriate for $\Lambda$ CDM, it has been shown that, for $\gamma\sim0.55$, the growth rate calculated using Equation \ref{eq:fa} corresponds remarkably accurately tothat predicted by GR."
 also showed that > and ew are only weakly correlated (see also 2)) ancl. therefore. that one can measure the cosmic growth and expansion rates almost independently.," also showed that $\gamma$ and $w$ are only weakly correlated (see also ) and, therefore, that one can measure the cosmic growth and expansion rates almost independently."
 These convenient. properties of the 5-mocel provide us with a simple but. powerful consistency test. for CI., These convenient properties of the $\gamma$ -model provide us with a simple but powerful consistency test for GR.
 Llere. we perform this test using the new cluster SLE data (Papers Land LL) in combination with fi. SNla. BAO and CMD constraints.," Here, we perform this test using the new cluster XLF data (Papers I and II) in combination with $f_{\rm gas}$, SNIa, BAO and CMB constraints."
 The latter four data sets primarily constrain the background evolution of the Universe. ic. the expansion history.," The latter four data sets primarily constrain the background evolution of the Universe, i.e. the expansion history."
 As a background. expansion model. we where a is the dark energy. equation of state parameter and iw=1 corresponds to flat ACDAL," As a background expansion model, we where $w$ is the dark energy equation of state parameter and $w=-1$ corresponds to flat $\Lambda$ CDM."
 Note that we use this modelon£y as a phenomenological ceseription of the expansion history., Note that we use this model as a phenomenological description of the expansion history.
 We do assume the existence of dark energy., We do assume the existence of dark energy.
 Our model allows us to simultaneously test for departures from CDM in the expansion history («=1) ancl departures from. GR in the growth history (5=0.55)., Our model allows us to simultaneously test for departures from $\Lambda$ CDM in the expansion history $w\neq-1$ ) and departures from GR in the growth history $\gamma\neq0.55$ ).
 Deviations from these benchmarks can be caused either by clustering dark energv or mocified gravity., Deviations from these benchmarks can be caused either by clustering dark energy or modified gravity.
 The fitting formula f(a) from has been tested for a wide range of background dark energy. cosmologies and has been shown to provide a “universal” description. within 20 per cent uncertaintymodel.," The fitting formula $f(\sigma)$ from has been tested for a wide range of background dark energy cosmologies and has been shown to provide a “universal” description, within $\sim 20$ per cent uncertainty."
.. Hlowever. in these tests. typically ον he background evolution was modified. while the linear density perturbations of the new component (dark energy) were neglected.," However, in these tests, typically only the background evolution was modified, while the linear density perturbations of the new component (dark energy) were neglected."
 Recently. have presented: results tha account for the clfectsof the density Lluctuations provided by quintessence models on the linear matter power spectrum. finding only a small οσο on f(a) for the masses anc redshifts relevant here.," Recently, have presented results that account for the effectsof the density fluctuations provided by quintessence models on the linear matter power spectrum, finding only a small effect on $f(\sigma)$ for the masses and redshifts relevant here."
 The present study aims to test the consistency of current observations with CGlt (and ACDAL)., The present study aims to test the consistency of current observations with GR (and $\Lambda$ CDM).
 Therefore. we use an f(a) corresponding to GR. seeking to determine whether values of 570.55 (and we= 1) are preferrec bv the data.," Therefore, we use an $f(\sigma)$ corresponding to GR, seeking to determine whether values of $\gamma\sim0.55$ (and $w=-1$ ) are preferred by the data."
 Bevone the mass function of?.. several authors have investigated the inclusion of an explicit dependence on redshift in the fitting formula f(a.z)(2?7)..," Beyond the mass function of, several authors have investigated the inclusion of an explicit dependence on redshift in the fitting formula $f(\sigma,z)$."
 FollowingL. we use the f(e.z) Function of7.. determined from a large suite of simulations. llere. parameters have the redshift) dependence w(2)=voll| z)785. with ov[LA ecb.el.," Following, we use the $f(\sigma,z)$ function of, determined from a large suite of simulations, Here, parameters have the redshift dependence $x(z)=x_0(1+z)^{\varepsilon \alpha_{\rm x}}$ , with $x \in
\{A,a,b,c\}$ ."
 Thevalues for. ο and os are given by 2., Thevalues for $x_0$ and $\alpha_{\rm x}$ are given by .
. As described. in. L. we use the parameter ¢ to account for residual svstematic uncertainties in the evolution of fta. z).," As described in , we use the parameter $\varepsilon$ to account for residual systematic uncertainties in the evolution of $f(\sigma,z)$ ."
 We mareinalize over the uncertainties in the parameters of Equation 7.. accounting for their covariance.," We marginalize over the uncertainties in the parameters of Equation \ref{eq:mf}, , accounting for their covariance."
 We also account lor acditional svstemiatie uncertainties due to e.g. the presence, We also account for additional systematic uncertainties due to e.g. the presence
"axially symmetric, presumably relaxed mass distribution.","axially symmetric, presumably relaxed mass distribution."
" Although there is still room for potentially important further improvements, our algorithm should now be ready for reliable recoveries of lensing potentials, density distributions and density profiles of real galaxy clusters."," Although there is still room for potentially important further improvements, our algorithm should now be ready for reliable recoveries of lensing potentials, density distributions and density profiles of real galaxy clusters."
 The matrix representation of the finite (Eqs. ," The matrix representation of the finite differences (Eqs. \ref{Kij}, ,"
1]. and allows a simplification of the y? differencesminimisation., \ref{G1ij} and \ref{G2ij}) ) allows a simplification of the $\chi^2$ minimisation.
" For weak |2323))lensing, it reads containing one term each for the two ellipticity components."," For weak lensing, it reads containing one term each for the two ellipticity components."
" Of course, we have to perform the minimisation for both components, but we show only the calculation for one component."," Of course, we have to perform the minimisation for both components, but we show only the calculation for one component."
" Correspondingly, the quantities e,y,C,G represent εἰ,ΥΙ,C!,g! and «’,y”,C”,G’, respectively."," Correspondingly, the quantities $\varepsilon,\gamma,\mathcal{C},\mathcal{G}$ represent $\varepsilon^{1},\gamma^{1},\mathcal{C}^{1},\mathcal{G}^{1}$ and $\varepsilon^{2},\gamma^{2},\mathcal{C}^{2},\mathcal{G}^{2}$, respectively."
 One y? contribution is where we combined the non-linear factors (1—Zx) in the matrix prefactor ¥;;., One $\chi^2$ contribution is where we combined the non-linear factors $(1-Z\kappa)$ in the matrix prefactor $\mathcal{F}_{ij}$.
 We deal with this term by the iterative approach described in Sect. 4.2]., We deal with this term by the iterative approach described in Sect. \ref{innerleveliteration}.
" In each iteration, this term is simply kept constant."," In each iteration, this term is simply kept constant."
" Now, we minimise this equation with respect to the potential values y, The derivative of y? with respect to y is Using y;=Gira, Ki=""Ki, and AKi=Kiron, we can replace the remaining derivatives, The last equation shows that we can now write Eq."," Now, we minimise this equation with respect to the potential values $\psi_{l}$, The derivative of $\chi^2$ with respect to $\psi_l$ is Using $\gamma_{i}=\mathcal{G}_{ik}\psi_{k}$, $\kappa_{i}=\mathcal{K}_{ik}\psi_{k}$ and $\frac{\partial}{\partial\psi_{l}}\mathcal{K}_{ik}\psi_{k}=\mathcal{K}_{ik}\delta_{kl}$, we can replace the remaining derivatives, The last equation shows that we can now write Eq."
" [A3] as a linear system ofequations,"," \ref{minimum} as a linear system ofequations,"
"second passage, but still without any direct interaction yet.","second passage, but still without any direct interaction yet."
 We also investigate a major merger between two large galaxies A where the velocity field spreads over more pixels., We also investigate a major merger between two large galaxies A where the velocity field spreads over more pixels.
" In this case the situation is, as expected, slightly better than for the small galaxy B. Especially for the kinematical analysis more ellipses can be fitted even at redshift z—0.5."," In this case the situation is, as expected, slightly better than for the small galaxy B. Especially for the kinematical analysis more ellipses can be fitted even at redshift z=0.5."
" Nevertheless a similar trend is observable as for galaxy B. While the distortions are clearly visible at low redshift, they are subsequently smeared out when the galaxy is placed at higher redshift."," Nevertheless a similar trend is observable as for galaxy B. While the distortions are clearly visible at low redshift, they are subsequently smeared out when the galaxy is placed at higher redshift."
 In Fig., In Fig.
" 13 we show the 2D velocity field and the radial profiles of the kinemetric properties as seen at redshift z—0.5 for model galaxy A disturbed by an equal mass merger, 100 Myr after the first passage of the two discs."," \ref{equalAA} we show the 2D velocity field and the radial profiles of the kinemetric properties as seen at redshift z=0.5 for model galaxy A disturbed by an equal mass merger, 100 Myr after the first passage of the two discs."
 The appearance of the velocity field and the radial behaviour of mainly the flattening q point towards an interaction (compare to Figs., The appearance of the velocity field and the radial behaviour of mainly the flattening $q$ point towards an interaction (compare to Figs.
 1 and 2 for the undisturbed case)., \ref{regular0105} and \ref{plotsreg} for the undisturbed case).
" Small scale distortions are, however, smeared out."," Small scale distortions are, however, smeared out."
" Complementary to these investigations of VFs at intermediate redshift, we study the appearance of the 2D velocity field at z=2 as observed with SINFONI at VLT."," Complementary to these investigations of VFs at intermediate redshift, we study the appearance of the 2D velocity field at z=2 as observed with SINFONI at VLT."
" Due to the use of adaptive optics, these observations are not seeing limited and have a high angular resolution (we adopt 0.15"" for our study)."," Due to the use of adaptive optics, these observations are not seeing limited and have a high angular resolution (we adopt 0.15"" for our study)."
 As a test case we take galaxy A during an ongoing minor merger in simulation 2., As a test case we take galaxy A during an ongoing minor merger in simulation 2.
 The velocity field of this snapshot for seeing limited observations and various redshifts was shown in Fig. 7.., The velocity field of this snapshot for seeing limited observations and various redshifts was shown in Fig. \ref{distevo}.
 The galaxy has a radial exponential disc scale length of 4.5 kpc and is therefore comparable in size to the one observed by Genzel et al. (, The galaxy has a radial exponential disc scale length of 4.5 kpc and is therefore comparable in size to the one observed by Genzel et al. (
2006).,2006).
 The velocity field of this snapshot as seen at redshift z=2 with an instrument such as SINFONI is presented in Fig. 14.., The velocity field of this snapshot as seen at redshift z=2 with an instrument such as SINFONI is presented in Fig. \ref{sinfoni}.
 Peculiarities in the VF caused by the merger are clearly visible., Peculiarities in the VF caused by the merger are clearly visible.
"and (QQ.On.)=(0.15.0.55,0.017.0.65). for the we=1/3 model.","and $(\om,\oq,\ob,h)=(0.45, 0.55, 0.047, 0.65)$ for the $w_Q=-1/3$ model."
 The N-body code used is a parallel version of the particle-particle particle-uesh aleoritlin (Dertschinger Gelb 1991: Ferrell Dertschiuger 1991)., The $N$ -body code used is a parallel version of the particle-particle particle-mesh algorithm (Bertschinger Gelb 1991; Ferrell Bertschinger 1994).
 Each simulation uses 128° particles in a box of comoving volue 1007 Mpc?., Each simulation uses $128^3$ particles in a box of comoving volume $100^3$ $^3$.
 The Pliner softening leugth is 50 kpc comoving. which allows us to compute the nonlinear power spectruu in highly clustered regions with kx10h Mpe+ aud Ay1000.," The Plummer softening length is 50 kpc comoving, which allows us to compute the nonlinear power spectrum in highly clustered regions with $k\lo 10\,h$ $^{-1}$ and $\Dnl \lo 1000$."
 Since the ()-field clusters ο] ou scales umch above the box size. the presence of the quintessence only affects the initial conditions aud the evolution of the scale factor a.," Since the Q-field clusters only on scales much above the box size, the presence of the quintessence only affects the initial conditions and the evolution of the scale factor $a$."
 Figure 3 compares the linear power spectrum aud the fully evolved. spectrum from both. the N-body ruus and the approximations of PD96 aud Ma9s8., Figure 3 compares the linear power spectrum and the fully evolved spectrum from both the $N$ -body runs and the approximations of PD96 and Ma98.
 Five redshifts are shown for each of three QCDAL models., Five redshifts are shown for each of three QCDM models.
 Overall. we find that the PD96 formula works well at :=0 when the factor g in their formmla is set to g=ge. which is the appropriate growth factor for the density feld for QCDM. modelsleq. (5))].," Overall, we find that the PD96 formula works well at $z=0$ when the factor $g$ in their formula is set to $g=g_Q$, which is the appropriate growth factor for the density field for QCDM models[eq. \ref{gql}) )]."
 At earlier times. however. the PD96 forunilaunderestimates qj at kh=Lh * in tho wo=2/3 and 1/2 models by up to30%.," At earlier times, however, the PD96 formula $\Dnl$ at $k\go 1\,h$ $^{-1}$ in the $w_Q=-2/3$ and $-1/2$ models by up to."
. We have attempted less plysically-notived combinations of growth factors (o... g=0go aud gx) but did not find a way to make PD96 fit.," We have attempted less physically-motived combinations of growth factors (e.g., $g=\alpha\,g_Q$ and $g_\Lambda$ ) but did not find a way to make PD96 fit."
 We find that the Ma98 formula. can be easily extended to QCDAM models.," We find that the Ma98 formula, can be easily extended to QCDM models."
 Specifically. we propose to keep eq=1.08&10.1 and ey=2.10.10. used for ACDM in Maes. but adopt gy=go. which is the appropriate QCDAL erowth factor |eq. (5))].," Specifically, we propose to keep $c_1= 1.08\times 10^{-4}$ and $c_2 = 2.10 \times 10^{-5}$ used for $\LCDM$ in Ma98, but adopt $g=g_Q\,$, which is the appropriate QCDM growth factor [eq. \ref{gql}) )],"
 an gyοιο ," and $g_0=|w_Q|^{1.3\,|w_Q|-0.76}\,g_{Q,0}\,$ , where $g_{Q,0}\equiv
g_Q(a=1)$."
described in Ma98. the parameter in equation (6)) is introduced to approximate the power-law dependence of the effective spectral iudex πμ|3 in previous work ou σε: dlu(noag|3)/dluosX 25.," As described in Ma98, the parameter $\beta$ in equation \ref{dnl}) ) is introduced to approximate the power-law dependence of the effective spectral index $n_{\rm eff}+3$ in previous work on $\sigma_8$ : $d\ln(n_{\rm eff}+3)/d\ln\sigma_8\propto \beta\,$ ."
 We find ο)=0.83 an excelleu approximation for all three QCDAM models that we tested (see the lower-vight pauel of Figure 3)., We find $\beta=0.83$ an excellent approximation for all three QCDM models that we tested (see the lower-right panel of Figure 3).
 Other panels of Figure 3 illustrate the close agrecuent (rus errors ~LOY )) between N-body results aud equation (6))., Other panels of Figure 3 illustrate the close agreement (rms errors $\sim 10$ ) between $N$ -body results and equation \ref{dnl}) ).
 We have presented simple foxiuulas to approximate both the linear and nonlinear power spectra for matter density perturbations in viable quintessence cosmological models with au equation of state LosοX19.," We have presented simple formulas to approximate both the linear and nonlinear power spectra for matter density perturbations in viable quintessence cosmological models with an equation of state $-1\le w_Q
\lo -1/3$."
 Equatious (2)).(3)). and CL) together specify the ratio of the linear transfer fuuctions Ty aud T4 for the matter density field for a given pair of QCDALT and ACDAL models with the same cosmological parameters.," Equations \ref{kq}) \ref{tql}) ), and \ref{alpha}) ) together specify the ratio of the linear transfer functions $T_Q$ and $T_\Lambda$ for the matter density field for a given pair of QCDM and $\LCDM$ models with the same cosmological parameters."
 Equation (5)) specifies the ratio ofthe linear erowth suppression factors ge aud ga in QCDAL and ACDATL models., Equation \ref{gql}) ) specifies the ratio of the linear growth suppression factors $g_Q$ and $g_{\Lambda}$ in QCDM and $\LCDM$ models.
 Equation (6)) approximates the nonlinear niass power spectruni., Equation \ref{dnl}) ) approximates the nonlinear mass power spectrum.
 A kev difference between gravitational clustering iu QCDM and ACDM models is that A is spatially siiootli on all leneth scales. whereas the Q-field can cluster above a certain leugth scale.," A key difference between gravitational clustering in QCDM and $\LCDM$ models is that $\Lambda$ is spatially smooth on all length scales, whereas the Q-field can cluster above a certain length scale."
 We characterize this leneth scale by the shape parameter Po of equation (2)). which is derived from the Compton waveleneth of the O-field.," We characterize this length scale by the shape parameter $\Gamma_Q$ of equation \ref{kq}) ), which is derived from the Compton wavelength of the Q-field."
" The QCDM matter power spectruni therefore changes shape at two characteristic scales: Do. aud the familiar Pox04,7 that corresponds to the cross-over frou radiation- to matterdominated cra."," The QCDM matter power spectrum therefore changes shape at two characteristic scales: $\Gamma_Q$ , and the familiar $\Gamma
\propto \om\,h$ that corresponds to the cross-over from radiation- to matter-dominated era."
" For the QCDMI models studied in thisLetter νο, coustaut wo). the Compton wavelength of the Q-field ix very large: keyc0.001 to 0:010 +."," For the QCDM models studied in this (i.e. constant $w_Q$ ), the Compton wavelength of the Q-field is very large: $k_Q \sim 0.001$ to $0.01\,h$ $^{-1}$."
 Ou scales of galaxy clusters aud below. therefore. the linear QCDAL power spectrum has.identical shape as in the correspouding ACDM 1iiodel and differs only iu the overall iuplitude by a factor of Clo/i4)(os Gor).," On scales of galaxy clusters and below, therefore, the linear QCDM power spectrum has shape as in the corresponding $\LCDM$ model and differs only in the overall amplitude by a factor of $(A_Q/A_\Lambda) (g_{Q\Lambda}/g_{Q\Lambda,0})^2$ ."
 This realization should simplify comparisons between QCDM and ACDAM moclels., This realization should simplify comparisons between QCDM and $\LCDM$ models.
 Forthe fully evolvednonlinearpower spectrum. we find that PD96 works well at 2=0 but underestimates its amplitude bv up to~30% at carlicr times.," Forthe fully evolvednonlinearpower spectrum, we find that PD96 works well at $z=0$ but underestimates its amplitude by up to$\sim 30$ at earlier times."
 The formula of Ma98. on the other hand. cau be easily extended," The formula of Ma98, on the other hand, can be easily extended"
"further spectral fit was tried, tying together the ionisation parameters of the absorbing and emitting matter in XSTAR.","further spectral fit was tried, tying together the ionisation parameters of the absorbing and emitting matter in XSTAR."
" The fit was again acceptable (x?—72/74), for a tied ionisation parameter £=10+2 erg cm s!."," The fit was again acceptable $\chi^2$ =72/74), for a tied ionisation parameter $\xi$ $\pm$ 2 erg cm $^{-1}$."
" This alternative “best fit model had an unchanged ‘normal’ power law with photon index [~1.9, now absorbed by a column Ng 1x10? cm? of warm (moderately ionised) gas."," This alternative `best fit' model had an unchanged `normal' power law with photon index $\Gamma$$\sim$ 1.9, now absorbed by a column $_{H}$$\sim$ $\times 10^{22}$ $^{-2}$ of warm (moderately ionised) gas."
 Figure 6 illustrates this model (hereafter Model 2)., Figure 6 illustrates this model (hereafter Model 2).
" Although not quite such a good fit as Model 1, the common warm absorber/re-emitter model for the EEPIC spectrum of 2M23 has the attraction of involving one less physical component."," Although not quite such a good fit as Model 1, the common warm absorber/re-emitter model for the EPIC spectrum of 2M23 has the attraction of involving one less physical component."
 Visual comparison of figures 5 and 6 shows the soft X-ray emission to be somewhat weaker in Model 2., Visual comparison of figures 5 and 6 shows the soft X-ray emission to be somewhat weaker in Model 2.
" Specifically, the 0.2-2 keV luminosity in the ‘soft’ emission component of Model 2 fell to 2.5x10* erg s~’, while the observed 0.2-10 keV luminosity of the power law continuum of 2M23 was 7x10? erg s1 or 1x10“ ergs“? correcting for intrinsic absorption."," Specifically, the 0.2-2 keV luminosity in the `soft' emission component of Model 2 fell to $2.5\times 10^{42}$ erg $^{-1}$, while the ${\it observed}$ 0.2-10 keV luminosity of the power law continuum of 2M23 was $7\times 10^{43}$ erg $^{-1}$, or $1\times 10^{44}$ erg $^{-1}$ correcting for intrinsic absorption."
" In Model 2 the luminosity of the soft emission component is then ~8% of that removed from the power law continuum by absorption, a factor ~2 less than in Model 1."," In Model 2 the luminosity of the soft emission component is then $\sim$ of that removed from the power law continuum by absorption, a factor $\sim$ 2 less than in Model 1."
 The sspectrum of 2M23 has a typical Type 1 AGN power law continuum slope ((~1.9) when allowance is made for absorption., The spectrum of 2M23 has a typical Type 1 AGN power law continuum slope $\Gamma$$\sim$ 1.9) when allowance is made for absorption.
" The nature of the line-of-sight column is unclear from our fitting, the low energy opacity (larger than usual for a Type 1 QSO) being modelled by a column Ny~3x10?! cm? of ‘cold’ gas, or Ngc1x107 οι of moderately ionised or ‘warm’ gas."," The nature of the line-of-sight column is unclear from our fitting, the low energy opacity (larger than usual for a Type 1 QSO) being modelled by a column $_{H}$$\sim$ $\times 10^{21}$ $^{-2}$ of `cold' gas, or $_{H}$$\sim$ $\times 10^{22}$ $^{-2}$ of moderately ionised or `warm' gas."
" Importantly, the attenuation of the power law component below ~1 keV allows the spectral shape of the ‘soft excess’ to be resolved and quantified."," Importantly, the attenuation of the power law component below $\sim$ 1 keV allows the spectral shape of the `soft excess' to be resolved and quantified."
 In the first spectral fit (Model 1) the X-ray absorbing gas is cold and embedded dust could also be responsible for the optical reddening seen in this 2MASS QSO., In the first spectral fit (Model 1) the X-ray absorbing gas is cold and embedded dust could also be responsible for the optical reddening seen in this 2MASS QSO.
" While it is possible that the BLR has a flattened geometry and therefore suffers less absorption than the (centrally located) continuum X-ray source, the Type 1 optical classification of 2M23 suggests such ‘cold’ matter lies close to the continuum source."," While it is possible that the BLR has a flattened geometry and therefore suffers less absorption than the (centrally located) continuum X-ray source, the Type 1 optical classification of 2M23 suggests such `cold' matter lies close to the continuum source."
" If so, it must be relatively dense to survive."," If so, it must be relatively dense to survive."
 We note the growing evidence for ‘variable’ cold absorbing matter in AGN (eg Risaliti 22002)., We note the growing evidence for `variable' cold absorbing matter in AGN (eg Risaliti 2002).
 On the large scale this is seen in the tendency for some type 2 AGN switching between Compton-thick and Compton-thin X-ray absorption (eg Guainazzi 22004)., On the large scale this is seen in the tendency for some type 2 AGN switching between Compton-thick and Compton-thin X-ray absorption (eg Guainazzi 2004).
" Again, a substantial column density of near-neutral matter was seen to shrink as the continuum X-ray brightness increased strongly in the luminous type 1 Seyfert 1H0419-577 (Pounds 22004a)."," Again, a substantial column density of near-neutral matter was seen to shrink as the continuum X-ray brightness increased strongly in the luminous type 1 Seyfert 1H0419-577 (Pounds 2004a)."
" It is interesting to speculate that the ‘cold’ X-ray absorption in Type 1 AGN might be associated with the accretion process and related, for example, to the accreting clouds envisaged by Guilbert and Rees (1988) in their original paper predicting the subsequent discovery of X-ray ‘reflection’ (Pounds 11990)."," It is interesting to speculate that the `cold' X-ray absorption in Type 1 AGN might be associated with the accretion process and related, for example, to the accreting clouds envisaged by Guilbert and Rees (1988) in their original paper predicting the subsequent discovery of X-ray `reflection' (Pounds 1990)."
 Alternatively it may lie in dense structures at the base of an outflow or wind., Alternatively it may lie in dense structures at the base of an outflow or wind.
 In the alternative spectral fit to 2M23 (Model 2) the absorbing gas has the same ionisation parameter as that of the soft X-ray emission., In the alternative spectral fit to 2M23 (Model 2) the absorbing gas has the same ionisation parameter as that of the soft X-ray emission.
" However, in the context of the ‘red’ near-infrared colour of 2M23, it is noteworthy that the ionisation parameter (€~10) is lower than commonly found in Type 1 AGN."," However, in the context of the `red' near-infrared colour of 2M23, it is noteworthy that the ionisation parameter $\xi$$\sim$ 10) is lower than commonly found in Type 1 AGN."
" In this case the combination of a relatively low ionisation parameter and unusually high column density, Ny~1x10?? cm~?, may be causally linked with the designation of 2M23 as a 2MASS QSO."," In this case the combination of a relatively low ionisation parameter and unusually high column density, $_{H}$$\sim$ $\times 10^{22}$ $^{-2}$, may be causally linked with the designation of 2M23 as a 2MASS QSO."
 The main result from our spectral modelling of 2M23 is that the soft excess - revealed with unusual clarity due to the effective removal of the soft power law component - has a spectral form well matched by the blended line emission from an ionised gas., The main result from our spectral modelling of 2M23 is that the soft excess - revealed with unusual clarity due to the effective removal of the soft power law component - has a spectral form well matched by the blended line emission from an ionised gas.
" A similar finding was reported in the low flux state spectrum of the Seyfert 1 galaxy NGC 4051 (Pounds al.2004b), and proposed as the origin of a quasi-constant soft X-ray emission component in the luminous Seyfert 1 galaxy 1H0419-577 (Pounds 22004a)."," A similar finding was reported in the low flux state spectrum of the Seyfert 1 galaxy NGC 4051 (Pounds 2004b), and proposed as the origin of a quasi-constant soft X-ray emission component in the luminous Seyfert 1 galaxy 1H0419-577 (Pounds 2004a)."
" It was speculated in those papers that such an emission component might be a stronger analogue of the extended soft X-ray emission seen in Seyfert 2 galaxies such as NGC 1068 (Brinkman 22002, Kinkhabwala 22002) and Mkn 3 (Sako 22000, Pounds and Page 2005)."," It was speculated in those papers that such an emission component might be a stronger analogue of the extended soft X-ray emission seen in Seyfert 2 galaxies such as NGC 1068 (Brinkman 2002, Kinkhabwala 2002) and Mkn 3 (Sako 2000, Pounds and Page 2005)."
" In both those bright, nearby Seyfert 2 galaxies the high resolution soft X-ray spectrum is dominated by line emission from photoionised/ photoexcited gas, with the scattered continuum being weak."," In both those bright, nearby Seyfert 2 galaxies the high resolution soft X-ray spectrum is dominated by line emission from photoionised/ photoexcited gas, with the scattered continuum being weak."
 The shape of the soft emission in 2M23 suggests this is again the case., The shape of the soft emission in 2M23 suggests this is again the case.
" To quantify the comparative strength of the soft X-ray emission in these Type 1 AGN with an archetypal Type 2 AGN, we list in Table 1 their observed and derived soft X-ray luminosities, together with the intrinsic 2-10 keV luminosities as a proxy for the relative ionising fluxes."," To quantify the comparative strength of the soft X-ray emission in these Type 1 AGN with an archetypal Type 2 AGN, we list in Table 1 their observed and derived soft X-ray luminosities, together with the intrinsic 2-10 keV luminosities as a proxy for the relative ionising fluxes."
 From this admittedly small sample it is seen that the luminosity ratio of the soft X-ray emission to power law is up to an, From this admittedly small sample it is seen that the luminosity ratio of the soft X-ray emission to power law is up to an
in certain cells for the run with only sellgravity reach the same high values as the runs with feedback. but a much smaller fraction of the simulation volume has these hieh densities.,"in certain cells for the run with only self–gravity reach the same high values as the runs with feedback, but a much smaller fraction of the simulation volume has these high densities."
 Another blatant dillerence between the PDFs we [ind in our runs with feedback and the PDEs found by Wacla Norman (2001) and Wravisoy (2003) is that their runs do not show as high a peak at low densities., Another blatant difference between the PDFs we find in our runs with feedback and the PDFs found by Wada Norman (2001) and Kravtsov (2003) is that their runs do not show as high a peak at low densities.
 The smaller quantity of low density gas in their simulations is likely due to the much lower supernovac rates in Wada Norman's (2001) simulations (0.01 SN/vr as compared to0.5 43 SN/vr in our simulations) and in Ixravtsov's (2003) case. to the more realistic boundary conditions. which allow tenuous. hot gas Lo escape the disk.," The smaller quantity of low density gas in their simulations is likely due to the much lower supernovae rates in Wada Norman's (2001) simulations (0.01 SN/yr as compared to 0.5 – 4 SN/yr in our simulations) and in Kravtsov's (2003) case, to the more realistic boundary conditions, which allow tenuous, hot gas to escape the disk."
 Enerev spectra of the ISAL carry complementary information to that given by a study of its density structure., Energy spectra of the ISM carry complementary information to that given by a study of its density structure.
 With the, With the
rotation curve can be reproduced without the need for a dark matter halo.,rotation curve can be reproduced without the need for a dark matter halo.
" However, the resulting Misp/Lg ratios would be unrealistically high (~25 for Malin 1 and ~15 for NGC 7589), as it is usually found by maximizing the disks of “ordinary” LSB galaxies (e.g. ?))."," However, the resulting $M_{\rm{LSB}}/L_{\rm{R}}$ ratios would be unrealistically high $\sim$ 25 for Malin 1 and $\sim$ 15 for NGC 7589), as it is usually found by maximizing the disks of “ordinary” LSB galaxies (e.g. \citealt{deBlok2001}) )."
" Therefore, we estimated Misp/Lg=0.9 using the relations of ? and assuming B—V=0.44."," Therefore, we estimated $M_{\rm{LSB}}/L_{\rm{R}} = 0.9$ using the relations of \citet{Bell2003} and assuming $\rm{B - V} = 0.44$."
 The latter value is the mean colour found by ? for a sample of bulgeless LSB disks., The latter value is the mean colour found by \citet{McGaugh1994a} for a sample of bulgeless LSB disks.
 Theresults of the fits are listed in table 4 and shown in the bottom panels of figure 8.., Theresults of the fits are listed in table \ref{table:darkmatter} and shown in the bottom panels of figure \ref{fig:decoHL}.
 We also performed a standard bulge-disk decomposition for the inner HSB part and found a degeneracy of the mass-to-light ratios of bulge and disk., We also performed a standard bulge-disk decomposition for the inner HSB part and found a degeneracy of the mass-to-light ratios of bulge and disk.
" Different combinations of Myuge/Ln and Maisc/Lg provide comparable fits, as the rotation curves are not resolved in the inner regions."," Different combinations of $M_{\rm{bulge}}/L_{\rm{R}}$ and $M_{\rm{disk}}/L_{\rm{R}}$ provide comparable fits, as the rotation curves are not resolved in the inner regions."
 By maximizing both bulge and disk we found that their contributions to the first point of the rotation curve are approximately equal., By maximizing both bulge and disk we found that their contributions to the first point of the rotation curve are approximately equal.
" Therefore, in this “maximum light"" hypothesis, the inner steep rise is not explained only with the bulge, but also the disk is dynamically important."," Therefore, in this “maximum light” hypothesis, the inner steep rise is not explained only with the bulge, but also the disk is dynamically important."
 This further suggests that Malin 1 and NGC 7589 have an inner “normal” HSB disk (see also ?))., This further suggests that Malin 1 and NGC 7589 have an inner “normal” HSB disk (see also \citealt{Barth2007}) ).
" Following ?,, we used the derived values of M./Lg and My, to check the position of Malin 1 and NGC 7589 on the baryonic Tully-Fisher relation."," Following \citet{McGaugh2005}, we used the derived values of $M_{*}/L_{\rm{R}}$ and $M_{\rm{\hi}}$ to check the position of Malin 1 and NGC 7589 on the baryonic Tully-Fisher relation."
 We found that both galaxies follow the relation within the observed scatter., We found that both galaxies follow the relation within the observed scatter.
 The modified newtonian dynamics (MOND) was proposed by ? as an alternative to dark matter., The modified newtonian dynamics (MOND) was proposed by \citet{Milgrom1983} as an alternative to dark matter.
 This theory modifies the force law at accelerations a lower than a critical value ay and naturally explains the flatness of rotation curves (see the review by ?))., This theory modifies the force law at accelerations $a$ lower than a critical value $a_{0}$ and naturally explains the flatness of rotation curves (see the review by \citealt{Sanders2002}) ).
 GLSB galaxies are ideal systems to test MOND., GLSB galaxies are ideal systems to test MOND.
" Firstly, they are very diffuse and extended, thus the test is possible down to very low values of a/ag (deep MONDian regime)."," Firstly, they are very diffuse and extended, thus the test is possible down to very low values of $a/a_{\rm{0}}$ (deep MONDian regime)."
" Secondly, they have massive gaseous disks."," Secondly, they have massive gaseous disks."
" In the MONDian framework, this gives an important contribution to the galaxy dynamics, which is not subject to uncertainties as that of stars (i.e. the value of M./L)."," In the MONDian framework, this gives an important contribution to the galaxy dynamics, which is not subject to uncertainties as that of stars (i.e. the value of $M_{*}/L$ )."
" For the MOND fit to the rotation curves, we followed ? and assumed ay=3000 km? s? kpc! (?).."," For the MOND fit to the rotation curves, we followed \citet{Begeman1991} and assumed $a_{0} = 3000$ $^{2}$ $^{-2}$ $^{-1}$ \citep{Bottema2002}."
 Two different M/L ratios for the inner and outer regions of the galaxies have been used as free parameters in the fit., Two different M/L ratios for the inner and outer regions of the galaxies have been used as free parameters in the fit.
" For NGC 7589, we obtained a good fit by using a bulge- decomposition."," For NGC 7589, we obtained a good fit by using a bulge-disk decomposition."
 The stellar mass-to-light ratiospredicted by MOND are acceptable: Mpuige/L=4.7 and Maii/L=1.3., The stellar mass-to-light ratiospredicted by MOND are acceptable: $M_{\rm{bulge}}/L = 4.7$ and $M_{\rm{disk}}/L = 1.3$.
" For Malin 1, it is preferable to use the HSB-LSB decomposition shown in figure 8,, but it is not possible to leave both Mysp/L and Misp/L free."," For Malin 1, it is preferable to use the HSB-LSB decomposition shown in figure \ref{fig:decoHL}, but it is not possible to leave both $M_{\rm{HSB}}/L$ and $M_{\rm{LSB}}/L$ free."
" Indeed, the fit would require a negative value for Misp/Lg."," Indeed, the fit would require a negative value for $M_{\rm{LSB}}/L_{\rm{R}}$ ."
" This suggests that the presence of the extended LSB disk, combined with the large amount of gas, creates problems for MOND."," This suggests that the presence of the extended LSB disk, combined with the large amount of gas, creates problems for MOND."
 We fixed Misp/Lg and used only Mysp/Lp as a free parameter., We fixed $M_{\rm{LSB}}/L_{\rm{R}}$ and used only $M_{\rm{HSB}}/L_{\rm{R}}$ as a free parameter.
 Figure 9 shows the result for Misp/Lg= 0.5., Figure \ref{fig:MOND2c} shows the result for $M_{\rm{LSB}}/L_{\rm{R}} = 0.5$ .
" Using the relations by ?,, this corresponds to an extremely blue colour for the outer disk, i.e. B-V = 0.2."," Using the relations by \citet{Bell2003},, this corresponds to an extremely blue colour for the outer disk, i.e. B-V = 0.2."
The position errors of the source on the X-ray images was measured using the LIRAF taskincentroid.,The position errors of the source on the X-ray images was measured using the IRAF task.
" These errors were all <0.1"" ancl are also provided in Table 2..", These errors were all $<$ $''$ and are also provided in Table \ref{xpos}.
 The final position ancl errors on (he position were calculated by adding the alignment. ancl position errors for each measurement in quadrature and (hen taking the weighted mean of the three independent position measurements., The final position and errors on the position were calculated by adding the alignment and position errors for each measurement in quadrature and then taking the weighted mean of the three independent position measurements.
 Finally. the spectra of 01-36 were extracted using5 the CLAO task binning (he spectrum in energy so that each bin contained. < 20 counts.," Finally, the spectra of r1-36 were extracted using the CIAO task binning the spectrum in energy so that each bin contained $\gap$ 20 counts."
 We then [it each spectrum using CIAO 3.1/Sherpa. trving both a power-law with absorption and a disk blackbody model with absorption.," We then fit each spectrum using CIAO 3.1/Sherpa, trying both a power-law with absorption and a disk blackbody model with absorption."
 All of the errors reported are lo unless otherwise specilied., All of the errors reported are $\sigma$ unless otherwise specified.
 [Errors in (he spectral fits were obtained with the routine in Sherpa., Errors in the spectral fits were obtained with the routine in Sherpa.
 Luminosity errors take into account the normalization errors in the spectral fitting. but count rate errors do not.," Luminosity errors take into account the normalization errors in the spectral fitting, but count rate errors do not."
" Results are discussed in ο,", Results are discussed in \ref{results}.
 We obtained three sets of images containing the position of 11-36 with the Advanced Camera [or Survevs (ACS) aboardUST., We obtained three sets of images containing the position of r1-36 with the Advanced Camera for Surveys (ACS) aboard.
 The first data set. pointed at I.À.—00:42:43.36. Dec.—41:16:30.1 with an orientation of 55 degrees. was taken on 23-Jan-2004 in an attempt to detect a dillerent transient source.," The first data set, pointed at R.A.=00:42:43.86, Dec.=41:16:30.1 with an orientation of 55 degrees, was taken on 23-Jan-2004 in an attempt to detect a different transient source."
 This dataset fortuitously. contained the position of before it became active., This dataset fortuitously contained the position of r1-36 before it became active.
 The second. pointed at R.A.=00:42:41.50. Dec.—41:17:00.0 with an orientation of 220 degrees. was taken 15-Aug-2004 while the X-ray source was active.," The second, pointed at R.A.=00:42:41.50, Dec.=41:17:00.0 with an orientation of 220 degrees, was taken 15-Aug-2004 while the X-ray source was active."
 The final observations. pointed at R.A.200:42:41.50. Dec.—41:17:05.0 with an orientation of 116 degrees. was taken 22-Nov-2004 aller the X-ray. source [aded," The final observations, pointed at R.A.=00:42:41.50, Dec.=41:17:05.0 with an orientation of 116 degrees, was taken 22-Nov-2004 after the X-ray source faded."
 All three data sets were obtained through the FA85W (D equivalent) filter using the ACS J4-point box dither pattern in order to allow us to recover the best possible spatial resolution., All three data sets were obtained through the F435W $B$ equivalent) filter using the ACS 4-point box dither pattern in order to allow us to recover the best possible spatial resolution.
 The total exposure time for each set was 2200 s. Each epoch was combined in an identical manner using the task which has been optimized to process ACS imaging data., The total exposure time for each set was 2200 s. Each epoch was combined in an identical manner using the task which has been optimized to process ACS imaging data.
 The task removes cosmic Ταν events and geometric distortions. and il combines (he dithered frames together into one final photometric image in units of counts per second.," The task removes cosmic ray events and geometric distortions, and it combines the dithered frames together into one final photometric image in units of counts per second."
 We utilized the capabilities of this task (ο improve the sampling of the final, We utilized the capabilities of this task to improve the sampling of the final
there are two possible causes for the deflection: 1) the side reconnection with the null point which happcus a the east side of the CAIE and ii) the effect of the Lorentz force.,there are two possible causes for the deflection: i) the side reconnection with the null point which happens at the east side of the CME and ii) the effect of the Lorentz force.
 The flux rope is created by a southward line curren and the elobal maguetic field above AR 10798 is of uuipolay positive polavity., The flux rope is created by a southward line current and the global magnetic field above AR 10798 is of unipolar positive polarity.
 The Lorentz force acting on the flux rope is. therefore. alinost purely castward.," The Lorentz force acting on the flux rope is, therefore, almost purely eastward."
 The Loreutz force here is simular to that of the models of Chen(1996). and Iseuberg&Forbes(2007) but iu these studies. it results in acceleration and rotation. respectively duc to different magnetic topologics.," The Lorentz force here is similar to that of the models of \citet{Chen:1996} and \citet{Isenberg:2007} but in these studies, it results in acceleration and rotation, respectively due to different magnetic topologies."
 Tere. due to the larecly unipolar polarity. the Lorentz force can deflect the CME.," Here, due to the largely unipolar polarity, the Lorentz force can deflect the CME."
 Side reconnection resulting iu CME deflection cau be seen iu the work of Chen&Shi-vata(2000) for example., Side reconnection resulting in CME deflection can be seen in the work of \citet{Chen:2000} for example.
 It is not straight-forward to determine the cause of the CAIE deflection., It is not straight-forward to determine the cause of the CME deflection.
 However. during the first phase of the reconnection (first six winutes). the main part of the flux rope appears to be relatively unaffected by the reconnection as its negative ootpoiut reconuects (see top left panel of Figure 7. for exainple).," However, during the first phase of the reconnection (first six minutes), the main part of the flux rope appears to be relatively unaffected by the reconnection as its negative footpoint reconnects (see top left panel of Figure 7, for example)."
 Iu addition. the recounection is a breakout ype reconnection (above the flux rope) aud not a side reconnection.," In addition, the reconnection is a breakout type reconnection (above the flux rope) and not a side reconnection."
 No deflection has vet been reported for asviuuetrce breakout simulations (e.g..seeLyuchetal. 2009).," No deflection has yet been reported for asymmetric breakout simulations \citep[e.g., see][]{Lynch:2009}."
. Iu addition. the deflection rate of the null point does not correspond to that of the apex of the flux rope. as shown in the right pauel of Figure 7.," In addition, the deflection rate of the null point does not correspond to that of the apex of the flux rope, as shown in the right panel of Figure 7."
 During the second phase of the reconnection. the side reconnection results in an opening of the maeuetic field of the flux rope and it is unclear whether it affects or not the CALE direction.," During the second phase of the reconnection, the side reconnection results in an opening of the magnetic field of the flux rope and it is unclear whether it affects or not the CME direction."
 For these reasons. we believe that it is Likely that the Loreutz force is the main contributor to the deflection of the flux rope.," For these reasons, we believe that it is likely that the Lorentz force is the main contributor to the deflection of the flux rope."
 The Loreutz force is expected to decrease as the CATE rises since both the magnetic field streugth aud the curent decrease with height., The Lorentz force is expected to decrease as the CME rises since both the magnetic field strength and the current decrease with height.
 We fud that the deflection rate decreases with height aud the CME only deflects by about 2? in the second 30 minutes of its propagation., We find that the deflection rate decreases with height and the CME only deflects by about $^\circ$ in the second 30 minutes of its propagation.
 It is consistent with the Lorentz force being the cause of the deflection., It is consistent with the Lorentz force being the cause of the deflection.
 The deflection of the actual filamceut at the Sun could be exeater than what is found in the siuulatiou because its speed iu the low corona. where the Loreutz force is stronger. was certainly lower than that in our simulation.," The deflection of the actual filament at the Sun could be greater than what is found in the simulation because its speed in the low corona, where the Lorentz force is stronger, was certainly lower than that in our simulation."
 As noted above. the simulation does not reproduce the slow rise of the CADE for the LS nünutes frou the start of the SNR flare to the start of the TINR dare.," As noted above, the simulation does not reproduce the slow rise of the CME for the 18 minutes from the start of the SXR flare to the start of the HXR flare."
" It wieght create an additional deflection of about 10"" in the low corona. assmuine that the actual deflection rate is more or less consistent with what we find in our simulation."," It might create an additional deflection of about $^\circ$ in the low corona, assuming that the actual deflection rate is more or less consistent with what we find in our simulation."
 Another reason why our simulation does not reproduce the observed deflection is also due the initiation mechanism., Another reason why our simulation does not reproduce the observed deflection is also due the initiation mechanism.
 As the orce-free field of the flux rope is superiuposed onto he coronal magnetic field. if results in a magnetic configuration which is not force-free.," As the force-free field of the flux rope is superimposed onto the coronal magnetic field, it results in a magnetic configuration which is not force-free."
 Therefore. there are additional Loreutz forces on all directions. which xwtiallv compcusate the main castward Lorentz force discussecl before.," Therefore, there are additional Lorentz forces on all directions, which partially compensate the main eastward Lorentz force discussed before."
 In the real Sun. it is likely that the filament was ina force-free state before the eruption and it would oulv be subject to the main castward Lorentz orce. Which would vield a larger deflection.," In the real Sun, it is likely that the filament was in a force-free state before the eruption and it would only be subject to the main eastward Lorentz force, which would yield a larger deflection."
 Deflection of CNIES from the same active region curing the following Carrington rotation has also been reported2006)., Deflection of CMEs from the same active region during the following Carrington rotation has also been reported.
. We have investigated. using a numerical simulation. he first of two cruptions ou 2005 August 22 from anclmonue active region 10798.," We have investigated, using a numerical simulation, the first of two eruptions on 2005 August 22 from anemone active region 10798."
 We lave used a recently developed component of the SWAIF to include realistic hermocdyvnamics iuto the representation of the low corona (LC)., We have used a recently developed component of the SWMF to include realistic thermodynamics into the representation of the low corona (LC).
 Using the LC componcut has enabled us to xoduce EUV images of the solar corona which we cau compare to ETT observations., Using the LC component has enabled us to produce EUV images of the solar corona which we can compare to EIT observations.
 Ti addition. previous works iive shown that including thermodynamics is innportaut o reproduce the plasima aud magnetic properties of he eutire corona.," In addition, previous works have shown that including thermodynamics is important to reproduce the plasma and magnetic properties of the entire corona."
 We have investigated the pre-eveut opologv of the active region and found that it indeed developed iuto au anemione active region with a single ill] point separating the active region flux svstem to unipolar positive magnetic field., We have investigated the pre-event topology of the active region and found that it indeed developed into an anemone active region with a single null point separating the active region flux system to unipolar positive magnetic field.
 We have found that sole of these unipolar magnetic field lines are open while others are part of the streamer belt connected to a region rear disk ceuter., We have found that some of these unipolar magnetic field lines are open while others are part of the streamer belt connected to a region near disk center.
 Pre-eveut 171 and 195 images of he corona are in good agreement with simulated ones. validating our steady-state model.," Pre-event 171 and 195 images of the corona are in good agreement with simulated ones, validating our steady-state model."
 In agreement with observations of a Πα filament. which was well observed xior to the cruption. we have superposed a lightly wisted flux rope onto the steady-state solar corona at the sale location aud with the same orientation aud chirality (vielt-handed and soutlaward).," In agreement with observations of a $\alpha$ filament, which was well observed prior to the eruption, we have superposed a highly twisted flux rope onto the steady-state solar corona at the same location and with the same orientation and chirality (right-handed and southward)."
 Due to force imbalance. he flux rope iuuuediatelv erupts aud we have focused our study on its interaction with the adjacent magnetic Hux systems and ii particular with the open fiux from he surrounding coronal holes.," Due to force imbalance, the flux rope immediately erupts and we have focused our study on its interaction with the adjacent magnetic flux systems and in particular with the open flux from the surrounding coronal holes."
 The aneinone nature of the source region has two major consequences ou the evolution of the CATE: an castwared deflection due to the Loreutz force and a recounection of the negative footpoiut of the flux rope. which results in the appearance of a long«duration diuuuius region and a nux of open and closed feld lines in the CAIE.," The anemone nature of the source region has two major consequences on the evolution of the CME: an eastward deflection due to the Lorentz force and a reconnection of the negative footpoint of the flux rope, which results in the appearance of a long-duration dimming region and a mix of open and closed field lines in the CME."
 Since the flux rope was southward and right-hauded. here was a strong southward axial current.," Since the flux rope was southward and right-handed, there was a strong southward axial current."
 The Loreutz force generated by this current aud the unipolar positive uagnetic field in the vicinity of the active region is cast- and. in our simulation. it deflects the CATE by about 15° as the CME propagates to 5 Πιν," The Lorentz force generated by this current and the unipolar positive magnetic field in the vicinity of the active region is east-directed and, in our simulation, it deflects the CME by about $^\circ$ as the CME propagates to 8 $_\odot$."
 We have rot acdldress the reason why CMES from anemone active reeions tend to be fast CMEs (Lin2007)., We have not address the reason why CMEs from anemone active regions tend to be fast CMEs \citep[]{Liu:2007}.
. It. should (0 however noted that. in the simulation. the CATE encounters closed field Hines in the low corona. so the nagnetic topology cannot be fully responsible for the ‘ast speed.," It should be however noted that, in the simulation, the CME encounters closed field lines in the low corona, so the magnetic topology cannot be fully responsible for the fast speed."
 It is therefore likely that the fast speed of CAIEs from anemione active regions is at least partially due to the fact the CALE propagates into a low density. ‘ast solar wind. where the hydrodvnamical drag is low (Careill2001).," It is therefore likely that the fast speed of CMEs from anemone active regions is at least partially due to the fact the CME propagates into a low density, fast solar wind, where the hydrodynamical drag is low \citep[]{Cargill:2004}."
. The negative footpoiut of the flux rope quickly reconnects with the positive field part of the streamer belt., The negative footpoint of the flux rope quickly reconnects with the positive field part of the streamer belt.
 It leads to à CME with oue footpoiut (the positive one) that is line-tied. aud the second oue which couuects to a region of the Sun about 107 away from the active region.," It leads to a CME with one footpoint (the positive one) that is line-tied, and the second one which connects to a region of the Sun about $^\circ$ away from the active region."
 A dinunine region was found to develop in the reeion of the second footpoiut. both in simulated EUV nuages and in EIT ππασος aud to persist at tle Sun for as inuch as a dav after the eruption.," A dimming region was found to develop in the region of the second footpoint, both in simulated EUV images and in EIT images and to persist at the Sun for as much as a day after the eruption."
 Additioually.," Additionally,"
2000).,.
 The spectrum of the inverse-Compton photons off the jitter radiation would be similar to the jitter spectrum. which results ma harder spectrum than SSC at low frequency.," The spectrum of the inverse-Compton photons off the jitter radiation would be similar to the jitter spectrum, which results in a harder spectrum than SSC at low frequency."
" Therefore. the difference (F,,x» vs v?) might appear in the X-ray regime and could be tested."," Therefore, the difference $F_\nu \propto \nu$ vs $\nu^{1/3}$ ) might appear in the X-ray regime and could be tested."
 For photon energies =10 GeV. the 25.»e*e opacity against surrounding supernova photons becomes greater than unity (e.g..Steckeretal.1992).," For photon energies $\gtrsim 10^2$ GeV, the $\gamma\gamma \to e^+e^-$ opacity against surrounding supernova photons becomes greater than unity \citep[e.g.,][]{Stecker1992}."
. Thus. these high-energy photons cannot escape from the source. giving a sharp cutoff in the spectrum as shown in Figure |..," Thus, these high-energy photons cannot escape from the source, giving a sharp cutoff in the spectrum as shown in Figure \ref{fig:spect_e}."
 The resulting e* pairs. having Lorentz factor above —10?. then radiate synchrotron photons over a wide frequency range.," The resulting $e^\pm$ pairs, having Lorentz factor above $\sim$ $^5$, then radiate synchrotron photons over a wide frequency range."
 The inverse-Compton scattering of supernova photons by these pairs occurs mainly in the Klein-Nishina regime. where the scattering cross section is suppressed (Rybicki&Lightman1979).. and can be neglected here.," The inverse-Compton scattering of supernova photons by these pairs occurs mainly in the Klein-Nishina regime, where the scattering cross section is suppressed \citep{Rybicki1979}, and can be neglected here."
 The resulting pairs rapidly lose their energy through radiation. and thus they are in the strongly fast-cooling regime.," The resulting pairs rapidly lose their energy through radiation, and thus they are in the strongly fast-cooling regime."
 We computed the synchrotron spectrum from these pairs. and show the results in Figure 4. for the fiducial case and d=100 Mpe.," We computed the synchrotron spectrum from these pairs, and show the results in Figure \ref{fig:spect_reg} for the fiducial case and $d = 100$ Mpc."
 We find that this radiation component is significantly lower than both the SSC and EIC components for all the available energy ranges., We find that this radiation component is significantly lower than both the SSC and EIC components for all the available energy ranges.
" While we show here only the result for the fiducial case. the same conclusion would apply to other parameter sets. because the dependence of the regenerated emission on e, or eg 1s more or less the same as that of the SSC or EIC emission."," While we show here only the result for the fiducial case, the same conclusion would apply to other parameter sets, because the dependence of the regenerated emission on $\epsilon_e$ or $\epsilon_B$ is more or less the same as that of the SSC or EIC emission."
 In addition. if the EIC spectrum had any cutoff above 107 GeV. this would reduce even further the regenerated flux.," In addition, if the EIC spectrum had any cutoff above $10^2$ GeV, this would reduce even further the regenerated flux."
 Thus. the regeneration emission component can be safely ignored in the present discussion.," Thus, the regeneration emission component can be safely ignored in the present discussion."
 Recently. Wangetal.(2007) and Budniketal.(2008) suggested that hypernovae could account for some of the cosmic-ray flux that we measure today. and also investigated the possibility of the associated neutrino emission and detection in order to further test such a scenario.," Recently, \citet{Wang2007} and \citet{Budnik2008} suggested that hypernovae could account for some of the cosmic-ray flux that we measure today, and also investigated the possibility of the associated neutrino emission and detection in order to further test such a scenario."
 Asano&Mészáros(2008)— studied another consequence of this mechanism. the gamma-ray emission from secondary electrons produced by interactions. involving accelerated protons.," \citet{Asano2008} studied another consequence of this mechanism, the gamma-ray emission from secondary electrons produced by interactions involving accelerated protons."
 Their conelusion is that the synchrotron flux from these secondary electrons could be detected by modern X-ray telescopes., Their conclusion is that the synchrotron flux from these secondary electrons could be detected by modern X-ray telescopes.
 Here we further extend their study. comparing the results with that of the primary electron component. and investigate which emission components would dominate in different energy ranges. depending on the parameters.," Here we further extend their study, comparing the results with that of the primary electron component, and investigate which emission components would dominate in different energy ranges, depending on the parameters."
 As the origin of the secondary electrons. we consider decays of charged pions produced by interactions between protons and supernova photons. e.g.. p>»am? (labeled simply as p>). and pair production through the Bethe-Heitler process. p>>pete” (labeled as BH).," As the origin of the secondary electrons, we consider decays of charged pions produced by interactions between protons and supernova photons, e.g., $p \gamma \to n \pi^+$ (labeled simply as $p\gamma$ ), and pair production through the Bethe-Heitler process, $p \gamma \to p e^+
e^-$ (labeled as BH)."
 The BH process was not taken into account in the computation of Asano&Mészáros(2008)., The BH process was not taken into account in the computation of \citet{Asano2008}.
 Following Wangetal.(2007) and Asano&Mészáros(2008).. we assume that the total energy of cosmic-ray protons Is Een~Ex/6. and the energy distribution follows d - spectrum.," Following \citet{Wang2007} and \citet{Asano2008}, we assume that the total energy of cosmic-ray protons is $E_{\rm CR} \simeq E_K / 6$, and the energy distribution follows a $\varepsilon_p^{-2}$ spectrum."
" The maximum acceleration energy of the protons is assumed to be given by 2,44,=eBR02«105(e5/1077)? eV."," The maximum acceleration energy of the protons is assumed to be given by $\varepsilon_{p,{\rm max}} = e B R \beta \simeq 2 \times 10^{18}
(\epsilon_B / 10^{-2})^{1/2} ~ {\rm eV}$ ."
 We use again ey=107 as our fiducial value. to compare with our primary electron radiation.," We use again $\epsilon_B = 10^{-2}$ as our fiducial value, to compare with our primary electron radiation."
 Note that Asano&Mészáros(2008) used eg=0.1., Note that \citet{Asano2008} used $\epsilon_B = 0.1$.
 Due to this our estimate results in slightly smaller fluxes than those of Asano&Mészáros(2008): as we argue below. the main dependence on ey is expected to be weak.," Due to this our estimate results in slightly smaller fluxes than those of \citet{Asano2008}; as we argue below, the main dependence on $\epsilon_B$ is expected to be weak."
" The threshold energy for pion production through. A-resonance is roughly given by 2)"",20.3GeV/z.~10! eV; thus protons with energies near can produce charged pions through this mechanism."," The threshold energy for pion production through $\Delta$ -resonance is roughly given by $\varepsilon_{p,{\rm th}}^{p\gamma} = 0.3 ~ {\rm
GeV}^2 / \varepsilon_\gamma \simeq 10^{17}$ eV; thus protons with energies near $\varepsilon_{p,{\rm max}}$ can produce charged pions through this mechanism."
" The 7,,,4,energy-loss time scale of the protons is given by t.2(0.20,5.0)!c10° s, where o,.=5«107 em? (almost independent of energy) is the p cross section and the factor 0.2 represents the energy lost by the proton by each interaction."," The energy-loss time scale of the protons is given by $t_{p\gamma} =
(0.2 \sigma_{p\gamma} n_\gamma c)^{-1} \simeq 10^6$ s, where $\sigma_{p \gamma} = 5 \times 10^{-28}$ $^2$ (almost independent of energy) is the $p\gamma$ cross section and the factor 0.2 represents the energy lost by the proton by each interaction."
" The fraction of the total proton energy Fer taken away by this process may then be estimated by fpsShinfp,20.2."," The fraction of the total proton energy $E_{\rm CR}$ taken away by this process may then be estimated by $f_{p\gamma} = t_{\rm dyn} /
t_{p\gamma} \approx 0.2$."
 We assume that halfof the outcome of p> interaction is in the form of charged pions. and that they carry off of the parent proton energy. conserving the original € spectrum.," We assume that half of the outcome of $p\gamma$ interaction is in the form of charged pions, and that they carry off of the parent proton energy, conserving the original $\epsilon_\pi^{-2}$ spectrum."
" Secondary electrons or positrons are produced by pror decays and we assume that their energies are. 1/4 of that of the parent pions (z,=0.257, 0.057,). since a charged pion eventually decays into four light particles (one chargec lepton and three neutrinos)."," Secondary electrons or positrons are produced by pion decays and we assume that their energies are 1/4 of that of the parent pions $\varepsilon_e = 0.25 \varepsilon_\pi = 0.05 \varepsilon_p$ ), since a charged pion eventually decays into four light particles (one charged lepton and three neutrinos)."
" The electron spectrum is agai approximately represented by a power law with the same index. i.e.. for6«10?—-,/eV<10(eg107)""7."," The electron spectrum is again approximately represented by a power law with the same index, i.e., $\varepsilon_e^{-2}$ for $6 \times 10^{15} < \varepsilon_e / {\rm eV} 
< 10^{17} (\epsilon_B / 10^{-2})^{1/2}$."
 These electrons rapidly7 lose their energy radiatively./ and we obtain the final energy distribution by equating radiative energy loss and injection due to x production and decay.," These electrons rapidly lose their energy radiatively, and we obtain the final energy distribution by equating radiative energy loss and injection due to $\pi$ production and decay."
" The typical energy of the synchrotron photons corresponding to 2,= eV is ru~10% eV. Thus. the synchrotron spectrum from these secondary electrons is very hard and peaks in the multi-TeV region. where absorption by supernova photons is significant."," The typical energy of the synchrotron photons corresponding to $\varepsilon_e = 10^{16}$ eV is $\nu_{\rm syn} \sim 10^{13}$ eV. Thus, the synchrotron spectrum from these secondary electrons is very hard and peaks in the multi-TeV region, where absorption by supernova photons is significant."
 A third generation of electrons and positrons is then produced. and they again radiate synchrotron photons.," A third generation of electrons and positrons is then produced, and they again radiate synchrotron photons,"
Ever since the first discoveries of Jupiterauass plaucts around solar-type stars (Mavor Queloz 1995: Butler Alarev. 1996). the origins of hot Jupiters have inspired variety of ideas.,"Ever since the first discoveries of Jupiter-mass planets around solar-type stars (Mayor Queloz 1995; Butler Marcy 1996), the origins of hot Jupiters have inspired a variety of ideas."
 Formation scenarios invoked iun the T-—erature include orbital mieration. radiative strippiug. planet scattering and secular chaos (Ford Rasio 2008: Ida Lin 2010: Naoz et al," Formation scenarios invoked in the literature include orbital migration, radiative stripping, planet scattering ,and secular chaos (Ford Rasio 2008; Ida Lin 2010; Naoz et al."
 2011)., 2011).
 Each of these miechanisnis cau explain some of observational properties of hot Jupiters. but not all of them.," Each of these mechanisms can explain some of observational properties of hot Jupiters, but not all of them."
" The main propertics are summarized as follows: a) a pile up of orbital periods around three days: b) a wide range of obliquity: c) a lack of close conipauion stars. d) aimass function different thaw or other types of exoplancts. e) a Πίος, mass range of lost star masses. f) a steep rise in frequency for supeor-solar imetallicities. ο) a frequency insensitive to stellar uetallicity from solar down to very metal poor stars. and 1) a wide range of planetary sizes for a given planctary uass."," The main properties are summarized as follows: a) a pile up of orbital periods around three days; b) a wide range of obliquity; c) a lack of close companion stars, d) a mass function different than for other types of exoplanets, e) a limited mass range of host star masses, f) a steep rise in frequency for super-solar metallicities, g) a frequency insensitive to stellar metallicity from solar down to very metal poor stars, and h) a wide range of planetary sizes for a given planetary mass."
 Together. all these observational patterus constitute a substantial set of constrains that no single theory has con able to explain. though plauet population svuthesis (Mordasini et al.," Together, all these observational patterns constitute a substantial set of constrains that no single theory has been able to explain, though planet population synthesis (Mordasini et al."
 2009) las iet with success iu accounting or several of the observed features., 2009) has met with success in accounting for several of the observed features.
 It is possible that different patlisvavs exist to the formation of hot Jupiters., It is possible that different pathways exist to the formation of hot Jupiters.
 Topefully they would leave different signatures on the auet properties that could be used to ideutifv their origin., Hopefully they would leave different signatures on the planet properties that could be used to identify their origin.
 Iu this paper we hivpothesize a new scenario. where hey forma in excretion disks xoduced by low-inass binary nerecrs. which. we argue. could be the cause of the large sizes of sole ransitiug hot Jupiters.," In this paper we hypothesize a new scenario, where they form in excretion disks produced by low-mass binary mergers, which, we argue, could be the cause of the large sizes of some transiting hot Jupiters."
 The radius of brown dwarfs and eiut planets is determined by the interior temperatures which. for single objects. depend ou the total mass. age. aud the opacity.," The radius of brown dwarfs and giant planets is determined by the interior temperatures which, for single objects, depend on the total mass, age, and the opacity."
 At a given plauctary mass and composition. all planets are expected to converge fo a coustan radius (Ry) with time. which is determined by clectron degencracy pressure. Coulomb pressure aud exchange pressure in the interior (Zapolsky Salpeter 1969: Stevenson 1991).," At a given planetary mass and composition, all planets are expected to converge to a constant radius $R_0$ ) with time, which is determined by electron degeneracy pressure, Coulomb pressure and exchange pressure in the interior (Zapolsky Salpeter 1969; Stevenson 1991)."
 Since the stellar hosts of exoplaucts are usually on the main sequence. it is difficult to estimate their ages. but typically they are thought to be several Corr old. aud if their planets have the sane age. they should have had enough time to evolve to their degeneracy dominated confieuration where Ro»Ry.," Since the stellar hosts of exoplanets are usually on the main sequence, it is difficult to estimate their ages, but typically they are thought to be several Gyr old, and if their planets have the same age, they should have had enough time to evolve to their degeneracy dominated configuration where $R\rightarrow R_0$."
 The adveut of wide area transit survevs roni the eround (such as WASP) and in space (CoRoT aud Wepler} has brought about the discovery of over a hundred transiting exoplanets., The advent of wide area transit surveys from the ground (such as WASP) and in space (CoRoT and Kepler) has brought about the discovery of over a hundred transiting exoplanets.
 The colmbination of detailed photometric transit observations and high-precision radial velocity measurements have xovided ασ» and radius cdeternünations for those dlanets., The combination of detailed photometric transit observations and high-precision radial velocity measurements have provided mass and radius determinations for those planets.
 With al this wealth of data. it has become clear hat our understanding of the mass-racdius relationship in exoplanets is far from satisfactory.," With all this wealth of data, it has become clear that our understanding of the mass-radius relationship in exoplanets is far from satisfactory."
 In particular two xoperties of transiting exoplauets are specially troubling: (1) The large spread of planetary radi for masses around 1l Jupiter mass (from to as uch a factor 2): and (2) he large size of some exoplanets. which has been termed he radius anomaly of hot JipLI.ers (Bodenheiuer et 22001: Caullot et 22006).," In particular two properties of transiting exoplanets are specially troubling: (1) The large spread of planetary radii for masses around 1 Jupiter mass (from to as much a factor 2); and (2) the large size of some exoplanets, which has been termed the `radius anomaly' of hot Jupiters (Bodenheimer et 2001; Guillot et 2006)."
 Several iiechiuisuis have been invoked to account. for he radius anomaly. such as diffusion of stellar παπαοι into the planet interior. enhanced opacities. imefücieut convection. fida heating (Lecoute et 22009). Olunic wating (Batvein Stevenson 2010). burial of heat bv urbuleuce (Youdin Mitchell 2010). aud stroug winds (Chabrier. Lecoute Baraffe 2010).," Several mechanisms have been invoked to account for the radius anomaly, such as diffusion of stellar irradiation into the planet interior, enhanced opacities, inefficient convection, tidal heating (Leconte et 2009), Ohmic heating (Batygin Stevenson 2010), burial of heat by turbulence (Youdin Mitchell 2010), and strong winds (Chabrier, Leconte Baraffe 2010)."
 The wide spread iu radius excess (muore than a factor of 2 for planets with nasses around that of Jupiter) makes it unlikely that the cause of the radius anomaly lies in systematic errors dn he physics of interior structure (EOS. opacities) or other actors that are just functious of intrinsic properties of the udanet.," The wide spread in radius excess (more than a factor of 2 for planets with masses around that of Jupiter) makes it unlikely that the cause of the radius anomaly lies in systematic errors in the physics of interior structure (EOS, opacities) or other factors that are just functions of intrinsic properties of the planet."
 Recent Spitzer observations of the transiting planet WASP-12) have indicated that the atmospheric colmposition may be bhighlv uou-seolar (Madliusudliau et al, Recent Spitzer observations of the transiting planet WASP-12b have indicated that the atmospheric composition may be highly non-solar (Madhusudhan et al.
 2011)., 2011).
 Then it is possible that the scatter iu the amass-radius relatiouship is at least partly duc to variations in the atimospheric composition of hot Jupiters and in their cloud thickness (Burrows. Ποιο Nampaisarn 2011).," Then it is possible that the scatter in the mass-radius relationship is at least partly due to variations in the atmospheric composition of hot Jupiters and in their cloud thickness (Burrows, Heng Nampaisarn 2011)."
 However. to produce hot Jupiters as inflated as observed with cuhanced atimospheric opacitics a factor 10 above the solar musture is needed for cooling and contraction to be slowed down significantly (Caullot," However, to produce hot Jupiters as inflated as observed with enhanced atmospheric opacities a factor 10 above the solar mixture is needed for cooling and contraction to be slowed down significantly (Guillot"
which case different observers would see the peripheries as weak GRB. but with a laree rate per unit comoving volume. and the product of the rate density and apparent energy. of the weak GRB might conceivably dominate.,"which case different observers would see the peripheries as weak GRB, but with a large rate per unit comoving volume, and the product of the rate density and apparent energy of the weak GRB might conceivably dominate."
 In (his letter we examine to what extent weak. soft. wide angle and/or high redshift GRB contribute to the present cosmic GRD radiation densitv.," In this letter we examine to what extent weak, soft, wide angle and/or high redshift GRB contribute to the present cosmic GRB radiation density."
 In contrast to previous works on this matter. which appear to be model dependent and give a large range of numbers. our approach is nearly model independent - we compute the total [lux recorded by different GRB detectors. each with different ranges and sensitivities.," In contrast to previous works on this matter, which appear to be model dependent and give a large range of numbers, our approach is nearly model independent - we compute the total flux recorded by different GRB detectors, each with different ranges and sensitivities."
 We show that most of the GRB enerev output comes in hieh (hence GRBs in the spectral range 20-5.000 IxeV. independent of detector and trigger sensitivity. and such bursts could be detected by most major detectors.," We show that most of the GRB energy output comes in high fluence GRBs in the spectral range 20-5,000 KeV, independent of detector and trigger sensitivity, and such bursts could be detected by most major detectors."
" Moreover. these bright bursts would (wpically overwrite data that was being transmitted from BATSE. should thev have occured during the time such data transmission was taking place (considered ""dead. time” except when overwritten [Paciesas et al.."," Moreover, these bright bursts would typically overwrite data that was being transmitted from BATSE, should they have occurred during the time such data transmission was taking place (considered ""dead time"" except when overwritten [Paciesas et al.,"
 L999]). ancl would not have been overlooked by BATSE (see below).," 1999]), and would not have been overlooked by BATSE (see below)."
 We conclude from this that. in the context of total energetics. GRBs respect a law of excluded middle: There appears to be very little energv in. semi-GRDs'. compromised GRBs. ragged GRB peripheries. and the like.," We conclude from this that, in the context of total energetics, GRBs respect a law of excluded middle: There appears to be very little energy in ""semi-GRBs"", compromised GRBs, ragged GRB peripheries, and the like."
 Theoretical models of GRB should respect (his principle of the excluded middle at a quantitative level., Theoretical models of GRB should respect this principle of the excluded middle at a quantitative level.
 In other words. they should not predict that half or more of the total photon energy is spraved or otherwise diluted in directions [rom which the burst would be perceived as à marginal. weak. or semi- GRD.," In other words, they should not predict that half or more of the total photon energy is sprayed or otherwise diluted in directions from which the burst would be perceived as a marginal, weak, or ""semi-"" GRB."
 While the detectability of anv given weak GRB is a function of detector sensitivitv. the measured allskv flux in GRB is insensitive to it.," While the detectability of any given weak GRB is a function of detector sensitivity, the measured allsky flux in GRB is insensitive to it."
 We [ind from the BATSE 4D catalog that the total allskv flux in GRB photons £4 is about 5.3x10.7 erg/cnivr., We find from the BATSE 4B catalog that the total allsky flux in GRB photons $F_{\gamma}$ is about $5.3\times 10^{-3}$ $^2$ yr.
 Specilically. we have summed the fences of all the GRBs with reported [Iuences. divided bv the number of vears (5.3) and. assuming an effective fiekl of view (FOV) of 27 sr. multiplied by 2 in order to obtain an average allskv. (47) flux.," Specifically, we have summed the fluences of all the GRBs with reported fluences, divided by the number of years (5.3) and, assuming an effective field of view (FOV) of $2\pi$ sr, multiplied by 2 in order to obtain an average allsky $4\pi$ ) flux."
" This flux corresponds {ο a cosmic average enerev input. ΕΠο1.24xI0 erg/Mpce?v. The present day. local input is about a factor of 2 lower due to the [act that the star formation rate. ancl hence the inferred long GRD rate. was considerably higher in (lie past (e.g. Eichler, Guetta and Pohl 2010). i.e. enough to override the average redshift factor by a factor of 2."," This flux corresponds to a cosmic average energy input $F_{\gamma}H/c = 1.24 \times 10^{43}$ $^3$ y. The present day, local input is about a factor of 2 lower due to the fact that the star formation rate, and hence the inferred long GRB rate, was considerably higher in the past (e.g. Eichler, Guetta and Pohl 2010), i.e. enough to override the average redshift factor by a factor of 2."
 These values include short bursts. and for long bursts would probably be lower by several percent OY SO.," These values include short bursts, and for long bursts would probably be lower by several percent or so."
" Previously (Eichler, Guetta and Pohl. 2010). we summed (he total {hence in GDM data [rom August 2009 through February 2010. an 18 month interval. and found an average flux in GRD of 4x10 ""erg ty!."," Previously (Eichler, Guetta and Pohl, 2010), we summed the total fluence in GBM data from August 2009 through February 2010, an 18 month interval, and found an average flux in GRB of $4\times 10^{-3}$ erg $^{-1} \rm y^{-1}$."
 The GBA catalogue contains 212 long burst over this period. or about 283 GRBs per vear per ἐπ sr.," The GBM catalogue contains 212 long burst over this period, or about 283 GRBs per year per $4\pi$ sr."
"that, for acceptable solar pairs of awr-9ML; larger values of the mixing length imply narrower and more intense superadiabatic regions to drive the convection.","that, for acceptable solar pairs of $\alpha_{ML}$ $g_{ML}$, larger values of the mixing length imply narrower and more intense superadiabatic regions to drive the convection."
" Larger values of mixing length parameter ayr, imply larger velocity scale (larger Um) as Table 2. shows.", Larger values of mixing length parameter $\alpha_{ML}$ imply larger velocity scale (larger $v_m$ ) as Table \ref{tablep} shows.
" Thus, models A-E are a sequence having increasingly vigorous and narrowly restricted regions of convective driving (acceleration)."," Thus, models A–E are a sequence having increasingly vigorous and narrowly restricted regions of convective driving (acceleration)."
 Figure 3 shows the structure of the SAR for models A through E. This may be compared to Fig., Figure \ref{fignabt} shows the structure of the SAR for models A through E. This may be compared to Fig.
 2., 2.
" of Schlattl,Weiss&Ludwig(1997).", of \cite{swl97}.
". Again model A resembles their *Eddington-approximation"" curve, and models C, D and E are similar to their “2D-hydro-model and “1D-model-atmosphere” curves."," Again model A resembles their “Eddington-approximation” curve, and models C, D and E are similar to their ``2D-hydro-model'' and “1D-model-atmosphere” curves."
" Here AV=V—V, is plotted against logarithm of pressure (dynes/cm?).", Here $\Delta \nabla \equiv \nabla - \nabla_a$ is plotted against logarithm of pressure $\rm dynes/ cm^2$ ).
" Above the horizontal line AV=0, buoyant forces accelerate the turbulent flow, while below the line we have buoyancy damping (deceleration; this region is barely visible at the left edge of the curve)."," Above the horizontal line $\Delta \nabla = 0$, buoyant forces accelerate the turbulent flow, while below the line we have buoyancy damping (deceleration; this region is barely visible at the left edge of the curve)."
" According to MLT with the Schwarzschild criterion for convection, there should be no flow for AV< 0."," According to MLT with the Schwarzschild criterion for convection, there should be no flow for $\Delta \nabla \le 0$ ."
 The area above (under) the curve gives the net buoyant acceleration (deceleration)., The area above (under) the curve gives the net buoyant acceleration (deceleration).
" Clearly the deceleration, seen as the small depression near logP=5, is overcome by the much larger region of acceleration around logP=5.2, so that the Schwarzschild criterion gives incorrect results here."," Clearly the deceleration, seen as the small depression near $\log P = 5$, is overcome by the much larger region of acceleration around $\log P \approx 5.2 $, so that the Schwarzschild criterion gives incorrect results here."
" The area argument implied in Figure 3 is essentially the bulk Richardson number criterion (Fernando1991;Meakin&Arnett 2007b),, and is nonlocal."," The area argument implied in Figure \ref{fignabt} is essentially the bulk Richardson number criterion \citep{fern91,ma07b}, , and is nonlocal."
" Therefore, because the pathological deceleration implied by use of the Schwarzschild criterion is incorrect, the velocities vm, given in Table 2 are directly related to those which produce solar line broadening."," Therefore, because the pathological deceleration implied by use of the Schwarzschild criterion is incorrect, the velocities $v_m$ given in Table \ref{tablep} are directly related to those which produce solar line broadening."
 Figure4 shows the run of turbulent velocity as a function of optical depth for, Figure\ref{figasp} shows the run of turbulent velocity as a function of optical depth for
The analvsis is based on two sets of data: 1) a 9.7-day run. May 2 through Alay 12. 2009. during the conuuission period to measure (he initial photometric performance of the instrument. and 2) (he first 33.5 davs of science operations. May. 13 through June 17. 2009. the end of which is marked bv a 90° quarterly roll of the spacecraft about the optical axis (IIlaasetal.2010).,"The analysis is based on two sets of data: 1) a 9.7-day run, May 2 through May 12, 2009, during the commission period to measure the initial photometric performance of the instrument, and 2) the first 33.5 days of science operations, May 13 through June 17, 2009, the end of which is marked by a $90^\circ$ quarterly roll of the spacecraft about the optical axis \citep{haas}."
. The former data set is referred (to as Quarter 0 (QU). while the latter is relerred (o as Quarter 1 (QI).," The former data set is referred to as Quarter 0 (Q0), while the latter is referred to as Quarter 1 (Q1)."
 All relatively uncrowded stars brighter than A'p=13.6 (where Apis (he apparent magnitude in (heAepler passband) make up a list of 52.496 targets observed during Q0.," All relatively uncrowded stars brighter than $Kp=13.6$ (where $Kp$ is the apparent magnitude in the passband) make up a list of 52,496 targets observed during Q0."
" The Q1 target list contains 156.097 stars. 93% of which were selected based on metrics (hat address the detectability of terrestrial-size planets (Datalba2010) metrics dependent on stellar properties (e.g. effective temperature. surface eravily, and apparent maenitude)."," The Q1 target list contains 156,097 stars, $93\%$ of which were selected based on metrics that address the detectability of terrestrial-size planets \citep{targetselection} – metrics dependent on stellar properties (e.g. effective temperature, surface gravity, and apparent magnitude)."
 The intersection of the two lists contains 45.742 targets.," The intersection of the two lists contains 45,742 targets."
 The aperture photometry. data conditioning. and transit search al@orithiuis are described in Jenkinsetal.(2010a).," The aperture photometry, data conditioning, and transit search algorithms are described in \citet{pipeline}."
. A diseussion of the resulting flux timeseries is given in (2010b).., A discussion of the resulting flux timeseries is given in \citet{longcadence}.
 The (ransitüng planet search applied (ο all stars vielded several thousand TCEs. some of which were triggered by artifacts in the light curves and single strong outliers.," The transiting planet search applied to all stars yielded several thousand TCEs, some of which were triggered by artifacts in the light curves and single strong outliers."
 This was mitigated by comparing the 47 statistic of a (transit model fit to the folded light curve against the 47 statistic of a linear fit to the same., This was mitigated by comparing the $\chi^2$ statistic of a transit model fit to the folded light curve against the $\chi^2$ statistic of a linear fit to the same.
 If the two are statistically indistinguishable. the TCE is disregarded.," If the two are statistically indistinguishable, the TCE is disregarded."
 A constraint on the transit duration (I-hour) is used to isolate TCEs triggered by strong outliers in the [hix timeseries., A constraint on the transit duration (1-hour) is used to isolate TCEs triggered by strong outliers in the flux timeseries.
 The transit model provides an estimate of the companion radius., The transit model provides an estimate of the companion radius.
 TCEs with a companion radius larger (han 2/2; are also disregarded., TCEs with a companion radius larger than $2R_J$ are also disregarded.
 Such laree radii are likely associated with late \l-dwarls ancl are not considered high priority planetary candidates., Such large radii are likely associated with late M-dwarfs and are not considered high priority planetary candidates.
 This leaves approximately 107 TCEx., This leaves approximately $10^3$ TCEs.
results between SPI aud IBIS led to confusion again. and the existence of x-ray polarization is still iu debate.,"results between SPI and IBIS led to confusion again, and the existence of $\gamma$ -ray polarization is still in debate."
 As the authors themselves pointed out in their reports. these results are with low statistics of ~26 level. aud may be stronely affected by the iustruucental svstematics uncertainties (valemeietal.2007:MeClvunGotzetal 2009).," As the authors themselves pointed out in their reports, these results are with low statistics of $\sim 2~\sigma$ level, and may be strongly affected by the instrumental systematics uncertainties \citep{kalemci2007, mcglynn2007, gotz2009}."
. These coutroversies aud conflicts indicate difficulties in detecting the 5-rav. polarization., These controversies and conflicts indicate difficulties in detecting the $\gamma$ -ray polarization.
 Iu this Letter. we report the detection of οταν polarization aud also the chauge of PA for extremely bright CRB LoOs26A using the newly developed CRB polarimeter.," In this Letter, we report the detection of $\gamma$ -ray polarization and also the change of PA for extremely bright GRB 100826A using the newly developed GRB polarimeter."
 Our polarimeter is completely designed for the polarization measurciuent of prompt CRBs. and well calibrated during developing phase before launch.," Our polarimeter is completely designed for the polarization measurement of prompt GRBs, and well calibrated during developing phase before launch."
 Especially. using a detector of prote-flight model. we experimentally aud wmuerically uuderstaud the response for polarized 5-rav with low systematic uncertainties.," Especially, using a detector of proto-flight model, we experimentally and numerically understand the response for polarized $\gamma$ -ray with low systematic uncertainties."
 Iu the following sections. we show the observation (section 2). data analysis (section 3) and discussion for the emission mechanisin of the prompt GRBs (section 1).," In the following sections, we show the observation (section 2), data analysis (section 3) and discussion for the emission mechanism of the prompt GRBs (section 4)."
 TSAROS (Ixawaeuchnietal.2008:Morict2010) (uterplauctary Wite-craft Accelerated by Radiation Of the Sun} is a small solar-power-sail demonstrator aux was successfully laauched ou 21 May 2010.," IKAROS \citep{kawaguchi2008, mori2010} (Interplanetary Kite-craft Accelerated by Radiation Of the Sun) is a small solar-power-sail demonstrator and was successfully launched on 21 May 2010."
 TXAROS las a laree polvimide membrane of 20 1i in diameter. aux this translates the solar radiation pressure to the thrust of the spacecraft.," IKAROS has a large polyimide membrane of 20 m in diameter, and this translates the solar radiation pressure to the thrust of the spacecraft."
 Since the deplovinent of the sail ou 9 June 2010. INAROS started the solar-sailing tow Venus.," Since the deployment of the sail on 9 June 2010, IKAROS started the solar-sailing toward Venus."
 The Ganuna-Rav Burst Polarimeter (GAP) (Yonetola aboard TXAROS is fully designed to measure a deeree of linear polarization iu the prompt enission of CRBs in the enerev range of TO300 keV. The detection principle Is an anisotropy of Compton scattered photons.," The Gamma-Ray Burst Polarimeter (GAP) \citep{yonetoku2011, murakami2010, yonetoku2006} aboard IKAROS is fully designed to measure a degree of linear polarization in the prompt emission of GRBs in the energy range of 70–300 keV. The detection principle is an anisotropy of Compton scattered photons."
 If the incident 5-ravs are linearly polarized. the distribution of scattered photons is due to Wlhein-Nishina effect which approximately shows sin?o curves. where o is the scattering angle.," If the incident $\gamma$ -rays are linearly polarized, the distribution of scattered photons is due to Klein-Nishina effect which approximately shows $\sin^{2} \phi$ curves, where $\phi$ is the scattering angle."
 The GAP cousists of a ceutral plastic scatterer of LF cim iu diunueter aud 6 cii in thickness. and the surrounding 12 CsICTI) scintillators with 5 nuu in thickness.," The GAP consists of a central plastic scatterer of 17 cm in diameter and 6 cm in thickness, and the surrounding 12 CsI(Tl) scintillators with 5 mm in thickness."
 The coincidence events within a eate time of 5ys between the signal from anv Csl aud that from the plastic are accepted for the polarization analysis., The coincidence events within a gate time of $5~{\rm \mu s}$ between the signal from any CsI and that from the plastic are accepted for the polarization analysis.
 The GAP’s high axial sviinietrv in shape aud the high eain wnitormity are the kevs for reliable imueasureimeut of polarizations. and to avoid a fake modulation due to background 5-ravs.," The GAP's high axial symmetry in shape and the high gain uniformity are the keys for reliable measurement of polarizations, and to avoid a fake modulation due to background $\gamma$ -rays."
 There are no external parts of spacecraft inside the GAP field of view., There are no external parts of spacecraft inside the GAP field of view.
 Moreover the detector cases (chassis) except for the detector top. are covered by thin lead sheets with 0.5 nua thickuess.," Moreover the detector cases (chassis), except for the detector top, are covered by thin lead sheets with 0.5 mm thickness."
 Therefore the effect of background 5-ravs scattered by the spacecraft body are negligible., Therefore the effect of background $\gamma$ -rays scattered by the spacecraft body are negligible.
 The GAP detected CRB 100826À on 26 August 2010 at 22:57:20.8 (UT) on the wav to Venus at about 0.21 AU apart from the Earth., The GAP detected GRB 100826A on 26 August 2010 at 22:57:20.8 (UT) on the way to Venus at about 0.21 AU apart from the Earth.
 The liehteurve of the prompt enuission is shown in Figure [. , The lightcurve of the prompt emission is shown in Figure \ref{fig1}. .
This burst was also detected by several satellites. auc was localized hy an iuterplanctary network (IPN) (Ihulewetal. 2010)..," This burst was also detected by several satellites, and was localized by an interplanetary network (IPN) \citep{hurley2010}. ."
 Combining the GAP data with the published IPN information. the direction of this burst is derived as (n.8)=(279.640.2.22.340.5). which corresponds to 20.0 deeree off-axis from the center of the GAP field of view.," Combining the GAP data with the published IPN information, the direction of this burst is derived as $(\alpha, \delta) = (279.6 \pm 0.2, -22.3 \pm 0.5)$, which corresponds to 20.0 degree off-axis from the center of the GAP field of view."
 An cucrey flucuce of this burst is (3.0d0.3)«10!erecin7 in 20 keV10 MeV band by KONUS (Colenetskiüetal.2010).. which is the top 1 of the brightest events listed in DATSE catalog.," An energy fluence of this burst is $(3.0 \pm 0.3) \times 10^{-4}~{\rm erg~cm^{-2}}$ in 20 keV–10 MeV band by KONUS \citep{golenetskii2010}, , which is the top 1 of the brightest events listed in BATSE catalog."
" The and high-energv photon indices are reported as ap=L1009 and Jp—21'S. respectively. and the PE, peak euerey as E,=606!iu)keV (Colenetskiietal.2010)."," The low- and high-energy photon indices are reported as $\alpha_B = -1.31^{+0.06}_{-0.05}$ and $\beta_B = -2.1^{+0.1}_{-0.2}$, respectively, and the $\nu F_{\nu}$ peak energy as $E_p = 606^{+134}_{-109}\;$ keV \citep{golenetskii2010}."
. An optical afterglow of this GRB was not reported. so its redshift is uukuownu.," An optical afterglow of this GRB was not reported, so its redshift is unknown."
 We divided the cutive data iuto two time intervals as abeled Tnterval-l aud -2 in a lighteurve of Figure L.., We divided the entire data into two time intervals as labeled Interval-1 and -2 in a lightcurve of Figure \ref{fig1}.
 The total nunbers of 5-rav photons after subtracting vackeround are 32.921 and 19.007 photous for each interval. respectively.," The total numbers of $\gamma$ -ray photons after subtracting background are 32,924 and 19,007 photons for each interval, respectively."
 The first part of this burst shows a large fare lasting [7 seconds since the trigeer. aud a ‘ollowing 53 seconds cousists of multiple spikes.," The first part of this burst shows a large flare lasting 47 seconds since the trigger, and a following 53 seconds consists of multiple spikes."
 Although he Iuterval-2 has several spikes. we combined all of them. ο keep photon statistics.," Although the Interval-2 has several spikes, we combined all of them to keep photon statistics."
 We used the backeround iodulatiou curve obtained w the 36.7 hours inteeration just before aud after the CRB trigger time., We used the background modulation curve obtained by the 36.7 hours integration just before and after the GRB trigger time.
 Iu this case. the averaged coincidence jackeround rate is quite stable as 5.6countss1C.|. and its modulation curve can be described as coustaut with the standard deviation of 0.1 level," In this case, the averaged coincidence background rate is quite stable as $5.6~{\rm counts~s^{-1}CsI^{-1}}$, and its modulation curve can be described as constant with the standard deviation of 0.1 level."
 After subtracting the backeround imeodiulation. the total muubers of the coimceideuce 5-xavs (polarization signals) are LAS21 and 2.733 for IDuterval-1 aud -2. respectively.," After subtracting the background modulation, the total numbers of the coincidence $\gamma$ -rays (polarization signals) are 4,821 and 2,733 for Interval-1 and -2, respectively."
 When we perform the polarization analvsis. we calculated the model function (or the detector respouse) for the polarized eamunaravs with Cant Lt Moute- simulator.," When we perform the polarization analysis, we calculated the model function (or the detector response) for the polarized gamma-rays with Geant 4 Monte-Carlo simulator."
 Ao ecometrical mass model of GAP was introduced in Youetokuetal.(2011)., A geometrical mass model of GAP was introduced in \citet{yonetoku2011}.
.. We can set all the possible situations with different off-axis angles. azimuthal phase angles. spectral parameters. ax well as polarization degrees aud angles of incoming 5-ravs.," We can set all the possible situations with different off-axis angles, azimuthal phase angles, spectral parameters, as well as polarization degrees and angles of incoming $\gamma$ -rays."
 In this analysis. the free paraincters are the polarization degrees and angles since the other parameters are constrained as shown in section 2.," In this analysis, the free parameters are the polarization degrees and angles, since the other parameters are constrained as shown in section 2."
 Suuulatius the interactions between the inconius eanuniravs and the ecometrical mass model with Monte-Carlo method. we created the model functions of the modulation curves.," Simulating the interactions between the incoming gamma-rays and the geometrical mass model with Monte-Carlo method, we created the model functions of the modulation curves."
 Then we selected the coimcidence events between the plastic aud one C8I scintillator as the polarization signals., Then we selected the coincidence events between the plastic and one CsI scintillator as the polarization signals.
 We simulated the model function with the step resolutions of 5 for polarization degrees and 5 degrees for phase angles., We simulated the model function with the step resolutions of 5 for polarization degrees and 5 degrees for phase angles.
% We will fit the observed modulation curves with the model function by a least squares moethod iu the following subsections., We will fit the observed modulation curves with the model function by a least squares method in the following subsections.
 We should treat not only the statistical error but also the systematic uncertainties correctly., We should treat not only the statistical error but also the systematic uncertainties correctly.
 Tere. the systematic uncertainties can any unexplainable difference between the observedmodulation curve aud the zinulated onc.," Here, the systematic uncertainties mean any unexplainable difference between the observedmodulation curve and the simulated one."
 huperfect tunimes of parameters inthe eround aud m- calibrations for the incident radiation. especiallyfrom off-axis direction. cause more miportaut svsteniatic," Imperfect tunings of parameters inthe ground and in-orbit calibrations for the incident radiation, especiallyfrom off-axis direction, cause more important systematic"
about 0.35rg.,about $0.35 r_{\rm H}$.
 They attributed this value of the dise radius as following from the estimate of Quillen&‘Trilling(1998).. for which the radius is determined by the angular momentum of the accreting gas.," They attributed this value of the disc radius as following from the estimate of \cite{quillen98}, for which the radius is determined by the angular momentum of the accreting gas."
 AMachida(2009) also performed three-dimensional calculations of cireumplanetary dises., \cite{machida09} also performed three-dimensional calculations of circumplanetary discs.
 For a Jupiter mass planet. he ound a peak in the surface density at about 21 protoplanctary racii (or 0.028rg) from the protoplanet centre.," For a Jupiter mass planet, he found a peak in the surface density at about 21 protoplanetary radii (or $0.028\,r_{\rm H}$ ) from the protoplanet centre."
 According o these results. the peak is due to the balance between the centrifugal force and the gravity of the protoplanet.," According to these results, the peak is due to the balance between the centrifugal force and the gravity of the protoplanet."
 Such small scale features might then correspond to the formation sites of the regular satellites around Jupiter ancl Saturn., Such small scale features might then correspond to the formation sites of the regular satellites around Jupiter and Saturn.
 However. his peak in the surface density was not reproduced by Avlille&Bate(2009) who found that the surface density decreases monotonically with radius from the protoplanet.," However, this peak in the surface density was not reproduced by \cite{ayliffe09} who found that the surface density decreases monotonically with radius from the protoplanet."
 We analyse the general properties of cireumplanetary disces by the use of some simplified models. in order to understand he physical processes that alfect dise structure.," We analyse the general properties of circumplanetary discs by the use of some simplified models, in order to understand the physical processes that affect disc structure."
 In Section 2.. we estimate the scale height and temperature of the disc.," In Section \ref{properties}, we estimate the scale height and temperature of the disc."
 In Section 3. we consider a ballistic disc model consisting of periodic particle orbits around a planet., In Section \ref{test} we consider a ballistic disc model consisting of periodic particle orbits around a planet.
 In Section 4. we examine he cireumplanetary. [uid dise evolution equations in the Lill approximation and cliscuss their relationship to the binary star case., In Section \ref{eqs} we examine the circumplanetary fluid disc evolution equations in the Hill approximation and discuss their relationship to the binary star case.
 In Section 5. we consider whether resonances could lie within a cireumplanetary disc., In Section \ref{res} we consider whether resonances could lie within a circumplanetary disc.
 In Section 6.. we analyse the effects ofa tidal torque on a purely viscous disc.," In Section \ref{tidal}, we analyse the effects of a tidal torque on a purely viscous disc."
 In Section T we describe results of a model that includes the cllects of eas pressure ov means of some SPL simulations., In Section \ref{sph} we describe results of a model that includes the effects of gas pressure by means of some SPH simulations.
 Section & contains a discussion and the conclusions., Section \ref{concs} contains a discussion and the conclusions.
" We consider a system with a planet of mass Ad, in à circular orbit around a star of mass M, at à separation ea.", We consider a system with a planet of mass $M_{\rm p}$ in a circular orbit around a star of mass $M_{\rm s}$ at a separation $a$.
 The Lill sphere is the approximate region wherethe planets gravity dominates that of the star and its radius is given by where & , The Hill sphere is the approximate region wherethe planet's gravity dominates that of the star and its radius is given by where $\mu=M_{\rm p}/(M_{\rm p}+M_{\rm s}) \ll 1$.
is the Boltzmann constant. Z is the temperature. say is the mean molecular weight and mg is the mass of a hydrogenoO atom.," The sound speed in an ideal gas is where $k$ is the Boltzmann constant, $T$ is the temperature, $\mu_{\rm
 m}$ is the mean molecular weight and $m_{\rm H}$ is the mass of a hydrogen atom."
 The surface temperature of a steady state accretion disc around an object of mass AJ at a radius r is ogiven by (Pringle1981). where σ is the Stefan-Doltzmann constant., The surface temperature of a steady state accretion disc around an object of mass $M$ at a radius $r$ is given by \citep{pringle81} where $\sigma$ is the Stefan-Boltzmann constant.
 We estimate the temperature of the accreting eas with the disc., We estimate the temperature of the accreting gas with the disc.
 Similar approaches were taken by (2002)., Similar approaches were taken by \cite{canup02}.
. Consider a disc orbiting a solar niass star that is accreting gas at à rate of LO7Mve toa typical value inferred from observations of T ‘Tauri stars.," Consider a disc orbiting a solar mass star that is accreting gas at a rate of $10^{-8}\,\rm
M_\odot \, yr^{-1}$, a typical value inferred from observations of T Tauri stars."
 Equation 3. implies that at a distance from the central star of 5XU. the disc has a temperature of about. 26 Ix. ignoring effects of stellar heating.," Equation \ref{temp} implies that at a distance from the central star of $5\,\rm AU,$ the disc has a temperature of about $26\,\rm K$ , ignoring effects of stellar heating."
 This temperature corresponds to a disce aspect ratio Lffrx0.04., This temperature corresponds to a disc aspect ratio $H/r \simeq 0.04$.
 The aspect ratio of the cireumplanetary disc taken with respect to the planet is where Af) is the mass of Jupiter and + is the distance from the planet. (, The aspect ratio of the circumplanetary disc taken with respect to the planet is where $M_{\rm J}$ is the mass of Jupiter and $r$ is the distance from the planet. (
Phere is a small correction due to the vertical eravity of the star that we ignore.),There is a small correction due to the vertical gravity of the star that we ignore.)
 This aspect ratio is much larger than the value for the cireumstellar disc. (14/7).2 0.04.," This aspect ratio is much larger than the value for the circumstellar disc, $(H/r)_{\rm s} \simeq 0.04$ ."
 We have taken the disc sound speed ce; to be given by its value at the dise surface and have ignored any increase at the midplane with optical depth., We have taken the disc sound speed $c_{\rm s}$ to be given by its value at the disc surface and have ignored any increase at the midplane with optical depth.
" The relatively high value of (44/7), is due to the higher gas temperatures and the typically. smaller Or values in the cireumplanetary disc case.", The relatively high value of $(H/r)_{\rm p}$ is due to the higher gas temperatures and the typically smaller $\Omega r$ values in the circumplanetary disc case.
 The value of the disc aspect ratio from equation (4)) is similar to the values obtained in the simulations by Avtille&Bate(2009)., The value of the disc aspect ratio from equation \ref{Hr}) ) is similar to the values obtained in the simulations by \cite{ayliffe09}.
. They found no evidence for cireumplanetary clises in their simulations of low mass planets (less than 100 Earth masses)., They found no evidence for circumplanetary discs in their simulations of low mass planets (less than 100 Earth masses).
 They attributed this result to the weakness of the planet's vertical gravity compared to pressure. as can be seen in equation (4)).," They attributed this result to the weakness of the planet's vertical gravity compared to pressure, as can be seen in equation \ref{Hr}) )."
" For a sulliciently low mass planet. (14/7), can be large enough that the concept of a dise breaks down and the gas around the planet is essentially spherical."," For a sufficiently low mass planet, $(H/r)_p$ can be large enough that the concept of a disc breaks down and the gas around the planet is essentially spherical."
 As is well known. the cireumplanetary disc temperatures at these accretion rates are too high to explain the existence of the icv satellites of Jupiter and Saturn (Canup&Ward2002).," As is well known, the circumplanetary disc temperatures at these accretion rates are too high to explain the existence of the icy satellites of Jupiter and Saturn \citep{canup02}."
. However. in a slower accretion phase as proposed bv Canup&Ward(2002). and further investigated by Barr&Canup(2008).. the temperature would be lower and £/r would be smaller. /r0.1.," However, in a slower accretion phase as proposed by \cite{canup02} and further investigated by \cite{barr08}, the temperature would be lower and $H/r$ would be smaller, $H/r \sim
0.1$."
 We discuss this further in Section &.., We discuss this further in Section \ref{concs}.
 We compare the surface density in the circumplanctary disc. X5. to the surface density in the cireumstellar dise. Sl.," We compare the surface density in the circumplanetary disc, $\Sigma_{\rm p}$, to the surface density in the circumstellar disc, $\Sigma_{\rm s}$ ."
 We assume that most of the mass being accreted through the circumstellardisc outside the orbitof the planet is acereted by it., We assume that most of the mass being accreted through the circumstellardisc outside the orbitof the planet is accreted by it.
 As discussed in the Introduction. this ellicient gas capture has been found in simulations of Jupiter mass planets that orbit," As discussed in the Introduction, this efficient gas capture has been found in simulations of Jupiter mass planets that orbit"
"All of the tests in this sectiou were performed ou the samples of mock galaxies with deusitv simular to that of M,20 galaxies.",All of the tests in this section were performed on the samples of mock galaxies with density similar to that of $M_r<-20$ galaxies.
" First. we removed galaxy clusters from the mock galaxy distribution using au isotropic frieuds-of-frieuds algorithm with hj=bj,=02 and nüniuun eroup size of Niji,—5."," First, we removed galaxy clusters from the mock galaxy distribution using an isotropic friends-of-friends algorithm with $b_{||}=b_{\perp}=0.2$ and a minimum group size of $N_{min}=5$."
 Foy |!=10Mypc. the filament incompleteness aud coutzuuimation rates for the cluster-free fibuneut distribution (33 and 39 per ceut. respectively) are much larger than those iu real space (16 and 25 per cent). sugecsting that overdensitics ou mceaparsec scales are plaviug an iuportaut role iu defining filaments on 10h sscales.," For $l=10$, the filament incompleteness and contamination rates for the cluster-free filament distribution $33$ and $39$ per cent, respectively) are much larger than those in real space $16$ and $25$ per cent), suggesting that overdensities on megaparsec scales are playing an important role in defining filaments on $10$ scales."
 Similar results are obtained when clusters are found and removed in redshift space using the approac[um of b..., Similar results are obtained when clusters are found and removed in redshift space using the approach of \ref{subsec:ClusterFind}.
 Tfς we instead the fingers-ofeod preseuted by ealaxy clusters. we can remove most of the contamination without having to remove the clusters themselves.," If we instead the fingers-of-god presented by galaxy clusters, we can remove most of the contamination without having to remove the clusters themselves."
 For this study. we will take the very simple approach of moving all members of a particular cluster to their ean position — that is. we will collapse the fiugers-of-eod to a point weighted by umber of galaxies iuthe cluster.," For this study, we will take the very simple approach of moving all members of a particular cluster to their mean position – that is, we will collapse the fingers-of-god to a point weighted by number of galaxies inthe cluster."
 If we follow this procedure. the incompleteness and contamination are smaller (19 aud 26 per cent) than the cluster-tree mock galaxy distributions aud a mareinal iuprovenieut over the redshift space distribution with uo special treatiueut of clusters (20 aud27 per cent).," If we follow this procedure, the incompleteness and contamination are smaller $19$ and $26$ per cent) than the cluster-free mock galaxy distributions and a marginal improvement over the redshift space distribution with no special treatment of clusters $20$ and $27$ per cent)."
 We repeated this exercise for filaments found on a 5 sincothing scale., We repeated this exercise for filaments found on a $5$ smoothing scale.
 Collapsing the fingers-of-eod docs lead to a marginal improvement. but filauncuts are still very poorly defined iu redshift space at these densities. with ~10 per cent contamination rates.," Collapsing the fingers-of-god does lead to a marginal improvement, but filaments are still very poorly defined in redshift space at these densities, with $\sim 40$ per cent contamination rates."
 A more sophisticated treatment of the clusters may be needed. but is bevoud the scope of this paper.," A more sophisticated treatment of the clusters may be needed, but is beyond the scope of this paper."
 In the section that follows we will discuss theapplication of the filament fiuder to real SDSS data., In the section that follows we will discuss theapplication of the filament finder to real SDSS data.
 To minimize contanination. we will be working only with filameuts fonnd on 10 aaud 157 sscales auc ouly after collapsing fiugers-of-god.," To minimize contamination, we will be working only with filaments found on $10$ and $15$ scales and only after collapsing fingers-of-god."
" The Sloan Digital Sky Survey has imaged a quarter of the skv in five wavebands. ranging frou, 3000 to 10000À.. to a depth of k&~22.5 (?).."," The Sloan Digital Sky Survey has imaged a quarter of the sky in five wavebands, ranging from $3000$ to $10000$, to a depth of $r \sim 22.5$ \citep{SDSS}. ."
 As of Data Release 6 νε spectra had been taken of ~800.000. ealaxics. covering 9583 square degrees aud extending to Petrosiun pocαντ (2).," As of Data Release $6$ \citep{DR6}, , spectra had been taken of $\sim 800,000$ galaxies, covering $9583$ square degrees and extending to Petrosian $r \sim 17.7$ \citep{Spectro}."
 Galaxy redshifts are typically accurate to ~30 Jun +. πα]ιο it ideal for studies of lareescale structure.," Galaxy redshifts are typically accurate to $\sim 30$ km $^{-1}$, making it ideal for studies of large--scale structure."
 For this study. we need a portion of skv with relatively few coverage gaps to minimize the effect of the window function on the A space distributions.," For this study, we need a portion of sky with relatively few coverage gaps to minimize the effect of the window function on the $\lambda$ –space distributions."
" With this in nünd. we construct two volume subsamples from the northern portion (8<a<16 Galshteh and 25<ὃ 60) of the NYU Valuc-Added Catalog(NYU-VACGC.?.throughDRG . the first τος140 «31041 ) iu size with A,Mpe6).20.5 (M r205) and the 170«100 (5.i Y in size with AL,<"" ME (Afr21)."," With this in mind, we construct two volume–limited subsamples from the northern portion $8<\alpha<16$ h and $25<\delta<60$ ) of the NYU Value-Added Galaxy Catalog \citep[NYU-VAGC, ][through DR$6$]{VAGC}, , the first $140 \times 140 \times 340$ $h^{-1}$ $^3$ in size with $M_r<-20.5$ $Mr205$ ) and the second $170 \times 170 \times 400$ $h^{-1}$ $^3$ in size with $M_r<-21$ $Mr21$ )."
 The samples extend to maxima redshifts 2=0.12 and +=0.15. respectively. aud are plotted in redshift space in Fie. 8..," The samples extend to maximum redshifts of $z=0.12$ and $z=0.15$, respectively, and are plotted in redshift space in Fig. \ref{fig:DataBoxes}. ."
 Absolute magnitudes were colmputed with hkearrect (2?) using SDSS ;-baudPetrosian magnitudes shifted to τί=(0.1(and using hl)., Absolute magnitudes were computed with \citep{Blanton03} using SDSS $r$ -bandPetrosian magnitudes shifted to $z=0.1$(and using $h=1$ ).
 We described the compilation and processing of the SDSS subsamples audthem mock counterparts im Paper 1., We described the compilation and processing of the SDSS subsamples andtheir mock counterparts in Paper $1$ .
 Before geeneratiug the filament distributions. we identify andcollapse the flugers-ofeod as described," Before generating the filament distributions, we identify andcollapse the fingers-of-god as described"
spectrum of the star is remarkably hard. a feature that is often associated with interacting wind binaries (Nazé et citeNRM)).,"spectrum of the star is remarkably hard, a feature that is often associated with interacting wind binaries (Nazé et \\cite{NRM}) )."
" To address the issue of multiplicity. we measured the RVs of a number of stellar emission lines (H8. Πα. ο 55412. 6406. 6683. UU55204. 5737. 6383. and 244945, see refRV20b)). as well as interstellar and nebular features such as the diffuse interstellar bands (DIBs) at 11952609. 5705. 5850. 6010. 6614 (see Herbig 1995)). the Nat,U1 53890, 5896 doublet. and the [O [t 55007 line."," To address the issue of multiplicity, we measured the RVs of a number of stellar emission lines $\beta$, $\alpha$ , $\lambda\lambda$ 5412, 6406, 6683, $\lambda\lambda$ 5204, 5737, 6383, and $\lambda$ 4945, see \\ref{RV20b}) ), as well as interstellar and nebular features such as the diffuse interstellar bands (DIBs) at $\lambda\lambda$ 5363, 5705, 5850, 6010, 6614 (see Herbig \cite{Herbig}) ), the $\lambda\lambda$ 5890, 5896 doublet, and the $[$ $]$ $\lambda$ 5007 line."
 The interstellar and nebular features allow us to check the stability of the wavelength calibration: their 1—c RV dispersion is ss! for the narrow absorptions. indicating a very stable calibration.," The interstellar and nebular features allow us to check the stability of the wavelength calibration: their $1-\sigma$ RV dispersion is $^{-1}$ for the narrow absorptions, indicating a very stable calibration."
 For the broad DIBs. the dispersion ts larger. of the order of 4kkmss7!.," For the broad DIBs, the dispersion is larger, of the order of $^{-1}$."
 For the majority of the stellar emission lines. the RV dispersions are significantly larger than these values (24.8. 17.7. and 6.Skkmss7! for theHt. and + lines. respectively).," For the majority of the stellar emission lines, the RV dispersions are significantly larger than these values (24.8, 17.7, and $^{-1}$ for the, and + lines, respectively)."
" However. as shown in refWR20bRV.. the variations in the different sets of lines are uncorrelated,"," However, as shown in \\ref{WR20bRV}, the variations in the different sets of lines are uncorrelated."
 To quantify the spectral variability of the star. we computed the temporal variance spectrum (TVS. see Fullerton. Gies Bolton 1996)) of our time series.," To quantify the spectral variability of the star, we computed the temporal variance spectrum (TVS, see Fullerton, Gies Bolton \cite{FGB}) ) of our time series."
 This method provides a quantitative assessment of the significance of the spectral line profile variability properly accounting for the noise level of the individual spectra., This method provides a quantitative assessment of the significance of the spectral line profile variability properly accounting for the noise level of the individual spectra.
 The results are shown in ref WR20bTVS.., The results are shown in \\ref{WR20bTVS}.
 Significant variations are detected in the strongest emission lines. but not in weaker ones such as the iv;Lt 66212-20 blend.," Significant variations are detected in the strongest emission lines, but not in weaker ones such as the $\lambda\lambda$ 6212-20 blend."
 The TVS of the strongest lines is quite complex. with multiple (up to four) peaks. indicating that the RV variability of these lines reflects their line profile variability.," The TVS of the strongest lines is quite complex, with multiple (up to four) peaks, indicating that the RV variability of these lines reflects their line profile variability."
 A more detailed inspection of the profiles of the hydrogen lines in different observations (see ref WR20bHalpha)) actually indicates. that the apparent RV variations are mainly due to changes in the upper part of the line profiles., A more detailed inspection of the profiles of the hydrogen lines in different observations (see \\ref{WR20bHalpha}) ) actually indicates that the apparent RV variations are mainly due to changes in the upper part of the line profiles.
 The results presented in this section suggest that 220b is indeed a spectroscopically variable object. but this variability is more likely to be related to (large-scale) structures in the stellar wind. than caused by genuine RV variations of the star.," The results presented in this section suggest that 20b is indeed a spectroscopically variable object, but this variability is more likely to be related to (large-scale) structures in the stellar wind, than caused by genuine RV variations of the star."
 Therefore. we conclude that there is at present no firm evidence that 220b is a short period (less than two week period) spectroscopic binary system (see also the discussion in refdiscussion)).," Therefore, we conclude that there is at present no firm evidence that 20b is a short period (less than two week period) spectroscopic binary system (see also the discussion in \\ref{discussion}) )."
" Following the classification. criteria of Smith. Shara Moffat (1996)). our spectra correspond most closely to a WNS5-6 spectral type (ratio of the 255412 to 255876 intensities above continuum = 3.6) with narrow lines (equivalent width. EW. of 455412 2 -5.9 mE finitesignatureo fhydrogenemissionrevealedbytheoscillatioi n,155412. and Hf on refWR20bTVS))."," Following the classification criteria of Smith, Shara Moffat \cite{SSM}) ), our spectra correspond most closely to a WN5-6 spectral type (ratio of the $\lambda$ 5412 to $\lambda$ 5876 intensities above continuum = 3.6) with narrow lines (equivalent width, EW, of $\lambda$ 5412 = $-5.9$ $> -40$ ) and a definite signature of hydrogen emission revealed by the oscillation of the Pickering decrement (compare the intensities of $\alpha$, $\lambda$ 5412, and $\beta$ on \\ref{WR20bTVS}) )."
 The only obvious stellar absorption feature in our spectra is the P-Cygni-type absorption trough of the 455876 line. which is obviously associated with the Wolf-Rayet star.," The only obvious stellar absorption feature in our spectra is the P-Cygni-type absorption trough of the $\lambda$ 5876 line, which is obviously associated with the Wolf-Rayet star."
 Therefore. our data are consistent with the WNoGha classification given by van der Hucht (2001)).," Therefore, our data are consistent with the WN6ha classification given by van der Hucht \cite{vdH}) )."
 On the basis of low-resolution spectroscopy. Moffat et ((1991)) classified 118 as an O7:V star.," On the basis of low-resolution spectroscopy, Moffat et \cite{MSP}) ) classified 18 as an O7:V star."
 This classification was revised to VV((f)) by Uzpen et ((2005)) and more recently to VV-HICf)) based on higher resolution spectra (Paper D., This classification was revised to V((f)) by Uzpen et \cite{Uzpen}) ) and more recently to V-III((f)) based on higher resolution spectra (Paper I).
 Both its optical brightness (Paper D and X-ray emission properties (brightness. spectral hardness and slight flux variability: see Nazé et citeNRM)) suggest that this star could be an interacting-wind binary system.," Both its optical brightness (Paper I) and X-ray emission properties (brightness, spectral hardness and slight flux variability; see Nazé et \\cite{NRM}) ) suggest that this star could be an interacting-wind binary system."
" We fitted Gaussian line profiles to measure the RVs of a number of stellar absorption lines (Hf. Πα. 1952576. ;,1 55412. 6683). as well as two faint emission lines 21155696. 6729-6731; Walborn 2001)."," We fitted Gaussian line profiles to measure the RVs of a number of stellar absorption lines $\beta$, $\alpha$, $\lambda$ 5876, $\lambda\lambda$ 5412, 6683), as well as two faint emission lines $\lambda\lambda$ 5696, 6729-6731; Walborn \cite{Walborn}) )."
 The results are listed in refRVMSPI8.., The results are listed in \\ref{RVMSP18}.
 We also measured the same interstellar and nebular features as for 220b às well as the nebular Hf and Ha emission lines., We also measured the same interstellar and nebular features as for 20b as well as the nebular $\beta$ and $\alpha$ emission lines.
 The |—c RV dispersions of the stellar absorption lines are between 1.8 and 5.9kkmss7!. which ts quite comparable to the corresponding values of the prominent interstellar features.," The $1-\sigma$ RV dispersions of the stellar absorption lines are between 1.8 and $^{-1}$, which is quite comparable to the corresponding values of the prominent interstellar features."
 The stellar emission lines have slightly larger RV dispersions. but these emission lines are actually extremely weak(normalized intensity of 0.02 above the continuum for 19526096) and their RVs are thus subject to larger uncertainties.," The stellar emission lines have slightly larger RV dispersions, but these emission lines are actually extremely weak(normalized intensity of 0.02 above the continuum for $\lambda$ 5696) and their RVs are thus subject to larger uncertainties."
 In addition. the TVS does not reveal any significant line profile variability.," In addition, the TVS does not reveal any significant line profile variability."
 We thus conclude that 118 is unlikely to be a short period (shorter than two week) spectroscopic binary (see also the discussion in refdiscussion))., We thus conclude that 18 is unlikely to be a short period (shorter than two week) spectroscopic binary (see also the discussion in \\ref{discussion}) ).
complemented a fully numerical seismology investigation by Arreguietal.(2007).,complemented a fully numerical seismology investigation by \cite{Arregui2007}.
 The eigenfunctions in the thin dissipative layer can be described by the functions F(r) and G(r) defined by Goossensetal.(1995) for the driven problem and the functions F(7) and G(r) defined by Rudemanetal.(1995) for the incompressible eigenvalue problem and by Tirry&Goossens(1996). for the compressible eigenvalue problem.," The eigenfunctions in the thin dissipative layer can be described by the functions $F(\tau)$ and $G(\tau)$ defined by \cite{Goossens1995a} for the driven problem and the functions $\tilde{F}(\tau)$ and $\tilde{G}(\tau)$ defined by \cite{Ruderman1995}
 for the incompressible eigenvalue problem and by \cite{Tirry1996}
 for the compressible eigenvalue problem."
 In the dissipative layer the MHD kink waves are highly Alfvénnic., In the dissipative layer the MHD kink waves are highly Alfvénnic.
 This can be understood as follows., This can be understood as follows.
" From the analysis by Sakuraietal.(1991a).. Goossensetal.(1995) and Tirry&Goossens(1996) it follows that in the dissipative layer. the Eulerian perturbation of total pressure P"" is constant and that |£|<<]&.|."," From the analysis by \cite{Sakurai1991a}, \cite{Goossens1995a} and \cite{Tirry1996} it follows that in the dissipative layer, the Eulerian perturbation of total pressure $P'$ is constant and that $\mid \xi_r \mid << \mid \xi_{\varphi} \mid $."
 We do not have to worry about £- since it is zero for a pressureless plasma., We do not have to worry about $\xi_z$ since it is zero for a pressureless plasma.
" In addition Here 7=(r—ra)/O, 15 the stretched independent variable used in the dissipative layer which has a typical width [559.959]."," In addition Here $\tau = \left(r - r_{\mathrm A}\right)/ \delta_{\mathrm A}$ is the stretched independent variable used in the dissipative layer which has a typical width $[-5\delta_{\mathrm A}, 5\delta_{\mathrm A}]$."
 F(t) is a complex function and £.(7) does not suffer a discontinuous jump as in the case of a uniform plasma., $\tilde{F}(\tau)$ is a complex function and $\xi_{\varphi}(\tau)$ does not suffer a discontinuous jump as in the case of a uniform plasma.
 It is characterised by rapid spatial variation in the dissipative layer., It is characterised by rapid spatial variation in the dissipative layer.
 From (34)) itfollows that the ratio of the g component of magnetic tension to that of pressure gradient is where we have used me|s|=£f., From \ref{xiDL}) ) itfollows that the ratio of the $\varphi$ component of magnetic tension to that of pressure gradient is where we have used $ \omega^2_{\mathrm A} /\mid \Delta_{\mathrm A} \mid \approx l$.
 Since F(t) is of order unity or bigger it follows that In thera MPdissipative layer the magnetic tension force is far bigger than the pressure gradient force., Since $\tilde{F}(\tau)$ is of order unity or bigger it follows that In the dissipative layer the magnetic tension force is far bigger than the pressure gradient force.
 Hence the MHD kink wave is highly Alfvénnie in the dissipative layer., Hence the MHD kink wave is highly Alfvénnic in the dissipative layer.
 In. this. subsection we go beyond the thin boundary approximation and we consider thick non-uniform layers with I/R.=0.2 and O.4., In this subsection we go beyond the thin boundary approximation and we consider thick non-uniform layers with $l/R = 0.2$ and $0.4$ .
 The eigenfunctions are. numerically calculated by solving the full set of linear. resistive MHD equations described in Terradasetal.(2006) and using the PDE2D code (Sewell.1995).," The eigenfunctions are numerically calculated by solving the full set of linear, resistive MHD equations described in \citet{Terradas2006} and using the PDE2D code \citep{Sewell2005}."
. Figure 2. displays the obtained results., Figure \ref{eigennonunif} displays the obtained results.
 It is clear that the two inhomogeneous solutions are almost identical to the homogeneous solution. except in the non-uniform layer. where large displacements are found.," It is clear that the two inhomogeneous solutions are almost identical to the homogeneous solution, except in the non-uniform layer, where large displacements are found."
 This is the location where the resonance takes place., This is the location where the resonance takes place.
 In Figure 3. we see that both the radial and azimuthal components of the Lorentz force are dominated by magnetic tension., In Figure \ref{forces} we see that both the radial and azimuthal components of the Lorentz force are dominated by magnetic tension.
 The magnitude of the magnetic tension and pressure Is of the same order (with the tension about twice as important as the magnetic pressure force such as corresponds to a density contrast of3) except in the inhomogeneous layer where the tension in the azimuthal direction is clearly dominant. reflecting the strong Alfvénnic nature of the solution in the dissipative layer.," The magnitude of the magnetic tension and pressure is of the same order (with the tension about twice as important as the magnetic pressure force such as corresponds to a density contrast of 3) except in the inhomogeneous layer where the tension in the azimuthal direction is clearly dominant, reflecting the strong Alfvénnic nature of the solution in the dissipative layer."
 The equations for incompressible MHD waves can be found in Goossensetal.(1992)., The equations for incompressible MHD waves can be found in \cite{Goossens1992}.
. Incompressibility means that we take the limit vs—co and enforce V-£=0. Note also that in the incompressible ease we=wa.," Incompressibility means that we take the limit $v_{\mathrm S} \rightarrow \infty$ and enforce $\nabla \cdot \vec{\xi}
= 0.$ Note also that in the incompressible case $\omega_{\mathrm C} =
\omega_{\mathrm A}$."
 The relevant equations are As before we concentrate on solutions with P#0 (although now we assume incompressibility) and rewrite (37)) as a second order ordinary differential equation for P’ (eliminating the displacements): Here we do not need to worry about propagating and/or evanescent behaviour of the solutions and the local radial wave number., The relevant equations are As before we concentrate on solutions with $P'\neq 0$ (although now we assume incompressibility) and rewrite \ref{MHDwavesInCom1}) ) as a second order ordinary differential equation for $P'$ (eliminating the displacements): Here we do not need to worry about propagating and/or evanescent behaviour of the solutions and the local radial wave number.
 The solutions are always evanescent or surface waves and the local radial wave number ts &-., The solutions are always evanescent or surface waves and the local radial wave number is $k_z$ .
 Note that in our notation El , Note that in our notation $\Gamma(\omega^2) = - k_z^2$.
"For a uniform plasma we can rewrite (38)) às Equation (39)) can then be solved in terms of Bessel functions Z,;€x) (x2 Κ.Π) in the internal part of the flux tube and K,,Cv) in the exterior region.", For a uniform plasma we can rewrite \ref{PInC1}) ) as Equation \ref{PInC2}) ) can then be solved in terms of Bessel functions $I_m(x)$ $x = k_z r$ ) in the internal part of the flux tube and $K_m(x)$ in the exterior region.
 Continuity of total pressure and the radial component of the Lagrangian displacement leads to the dispersion relation: The incompressible version of F is ↕∖∫⋯∖⇁⇁∖⋅∪∶∣∖⊳∶∣, Continuity of total pressure and the radial component of the Lagrangian displacement leads to the dispersion relation: The incompressible version of $F$ is Now $x_0 = k_z R$ .
∖↗↜⊤∣↴⊜↳∥⋋∣⊃⊜∣⋪⋋≣∪∏∣⋪⊖∣∐⊓∪∏⋖−⊔⊔∁⋅≏⋯⇂↴⊜↭∣∖⇁⊜∐ numerically., The dispersion relation \ref{DRInC1}) ) can be solved numerically.
 However. if we considerthe TT approximation. the Bessel functionsμία) and Αμ) in (41)) are replaced with their first order asymptotic expansions. and the dispersion," However, if we considerthe TT approximation, the Bessel functions$I_m(x)$ and $K_m(x)$ in \ref{DRInC1}) ) are replaced with their first order asymptotic expansions, and the dispersion"
without DENIS data. aud to calculate the star couuts of the section 83...,"without DENIS data, and to calculate the star counts of the section \ref{.counts}."
" With regard to the extinction shown iu 8l. the caleulatious were carried out frou, DENIS data by Schultheis et al. ("," With regard to the extinction shown in \ref{.ext}, the calculations were carried out from DENIS data by Schultheis et al. ("
private commmuication).,private communication).
 The images were reduced at the Paris Data Analvsis Centre (PDAC) of DENIS. and homoegencous criteria of reduction were applied to the whole special sample of the bulge. which is used here.," The images were reduced at the Paris Data Analysis Centre (PDAC) of DENIS, and homogeneous criteria of reduction were applied to the whole special sample of the bulge, which is used here."
 They were checked and calibrated (Borsenberecr 1997: Ruphy ct al., They were checked and calibrated (Borsenberger 1997; Ruphy et al.
 L997) at PDAC., 1997) at PDAC.
 The point source extraction was also carried out at PDAC by fitting PSEs optimized for crowded fields., The point source extraction was also carried out at PDAC by fitting PSFs optimized for crowded fields.
 For the astromoetrv. the iudividual DENIS frames were cross-correlated with the PNM catalogue (USNO-A2.0) with a precision of better than oue aresecond.," For the astrometry, the individual DENIS frames were cross-correlated with the PMM catalogue (USNO-A2.0) with a precision of better than one arcsecond."
 Star counts provide a powerful tool when scarching for asvuuetries inthe stellar distribution iu the immer Galaxy., Star counts provide a powerful tool when searching for asymmetries in the stellar distribution in the inner Galaxy.
 Other researchers have already used this method (sec introduction): however. thev have mostly observed regions where the bulee is predominant.," Other researchers have already used this method (see introduction); however, they have mostly observed regions where the bulge is predominant."
 While the bulge is observed iu offplane regions up to b~10 degrees and in ||<15 degrees. the long. thin bar is visible oulv in the in-plane regions. |b]<2 degrees aud up to [/|=27° at positive longitudes. aud somewhat less at negative longitudes.," While the bulge is observed in off-plane regions up to $b\sim 10$ degrees and in $|l|< 15$ degrees, the long, thin bar is visible only in the in-plane regions, $|b|<2$ degrees and up to $|l|=27^\circ $ at positive longitudes, and somewhat less at negative longitudes."
 It is now well known that the star counts iu the bulge are non-axisviunetrc (e.g Lóppez-Corredoira et al., It is now well known that the star counts in the bulge are non-axisymmetric (e.g Lóppez-Corredoira et al.
 2000): however. the larger-scale in-plane asvnunetries are less well investigated.," 2000); however, the larger-scale in-plane asymmetries are less well investigated."
 The most appropriate DENIS filter for probing the inner Galaxy near plane is A. since the effect of extinction is much lower than in the other filters.," The most appropriate DENIS filter for probing the inner Galaxy near plane is $K_{\rm s}$, since the effect of extinction is much lower than in the other filters."
 Wowever. the other filters cau be used to provide iuformation about the population of the stars aud the extinction.," However, the other filters can be used to provide information about the population of the stars and the extinction."
 Although the nominal A. limiting magnitude of DENIS is around 11.0 (Epchteiu 1998): the presence. of confusion implies iat in the more crowded regions the conrpleteness can be two or three magnitudes brighter. particularly owards the inner bulge (Unavaue et al.," Although the nominal $K_{\rm s}$ limiting magnitude of DENIS is around 14.0 (Epchtein 1998); the presence of confusion implies that in the more crowded regions the completeness can be two or three magnitudes brighter, particularly towards the inner bulge (Unavane et al."
 1995)., 1998).
 For this studs. however. only the brighter stars will be used. so confusion is not an issue.," For this study, however, only the brighter stars will be used, so confusion is not an issue."
 Whilst it is often argued that the fainter the star counts the better. this is not necessarily so when observing the immer Galaxy.," Whilst it is often argued that the fainter the star counts the better, this is not necessarily so when observing the inner Galaxy."
 The ratio of inner Galaxy to local disc stars is more important., The ratio of inner Galaxy to local disc stars is more important.
 At magnitude 13 or Ll there will be more inner Calaxy sources than at magnitudes 8 to 10. but the nuubers of οσα]. disc sources rises even nore rapidly.," At magnitude 13 or 14 there will be more inner Galaxy sources than at magnitudes 8 to 10, but the numbers of local disc sources rises even more rapidly."
 Cóuzónunu ot al. (, Garzónn et al. (
1993) showed tha the best contrast when lookiug at the bulee was at near mg 15. and that at fainter uaeuitudes the contrast became poorer.,"1993) showed that the best contrast when looking at the bulge was at near $m_{\rm K}$ = +8, and that at fainter magnitudes the contrast became poorer."
 However. the iuuber of stars detected falls rapidly mich brighter than lis (few if auv bulge sources are brighter than my = 6. 1100) aud the statistical errors become important.," However, the number of stars detected falls rapidly much brighter than this (few if any bulge sources are brighter than $m_{\rm K}$ = 6, L00) and the statistical errors become important."
 For this oper star counts up fo mg=9 will be used. as the inner-CGalaxy-to-dise contrast is high. but there are still sufficient sources to provide good statistics.," For this paper star counts up to $m_{K_{\rm s}}=9$ will be used, as the inner-Galaxy-to-disc contrast is high, but there are still sufficient sources to provide good statistics."
" The star counts down to mg, = 9.0 are shown iu Fiewre L..", The star counts down to $m_{K_{\rm s}}$ = 9.0 are shown in Figure \ref{Fig:DENIScounts9}.
 The DENIS star counts show the average of th available counts in a region of width Al=5 ceutred on he position marked., The DENIS star counts show the average of the available counts in a region of width $\Delta l=5^\circ $ centred on the position marked.
 A large bin size was adopted. as thiρα averages out the fluctuation due to the varving extiuctio- and the asvuuuetrics beiug sorted are distributed over nay tens of degrees.," A large bin size was adopted, as this averages out the fluctuation due to the varying extinction and the asymmetries being sorted are distributed over many tens of degrees."
 To complete the longitude coverage. he star counts from TMGS are used (see. Tamunerslev et al.," To complete the longitude coverage, the star counts from TMGS are used (see Hammersley et al."
 1999 for more details ou the data)., 1999 for more details on the data).
 The TMCS data. however. are plotted at the average longitude where he strip crossed the plane.," The TMGS data, however, are plotted at the average longitude where the strip crossed the plane."
 In both cases the areas of skv covered were large so that the Poissonian error is welicible., In both cases the areas of sky covered were large so that the Poissonian error is negligible.
 There is a difference between the standard A Glter of the TAICS aud the DENIS AL: however. Coheu (1997) has shown that the difference between he two for any star ds less than 0.03 mae. so the effect on the counts of using one or other should be negligible.," There is a difference between the standard $K$ filter of the TMGS and the DENIS $K_{\rm s}$; however, Cohen (1997) has shown that the difference between the two for any star is less than 0.03 mag, so the effect on the counts of using one or other should be negligible."
 Fig., Fig.
" L shows five slices across the plane at 2°>51575.1»b»05.025""7»b» 025°. OS>bo1 alc 1ο>be2."," \ref{Fig:DENIScounts9} shows five slices across the plane at $2^\circ>b>1.5^\circ$ , $ 1^\circ>b> 0.5^\circ$, $ 0.25^\circ>b>-0.25^\circ$, $-0.5^\circ>b>-1^\circ$ and $-1.5^\circ>b>-2^\circ$."
 Therefore. at the distance of the Calacic Centre (8 kpc) the average heights over the plane are O pc. X100 pc aud + 210 pe.," Therefore, at the distance of the Galactic Centre (8 kpc) the average heights over the plane are 0 pc, $\pm 100$ pc and $\pm$ 240 pc."
 It is noticeable that he cuts furthest from the plane are substantially cliffercut roni those in the plane., It is noticeable that the cuts furthest from the plane are substantially different from those in the plane.
 It can further be seen hat the counts are more or less svauuetrical in latitude., It can further be seen that the counts are more or less symmetrical in latitude.
 The slight differences that there are can be attributed to the Sun cine about L1 pe above the Galactic plane (Hanunerslev ot al., The slight differences that there are can be attributed to the Sun being about 14 pc above the Galactic plane (Hammersley et al.
 1995). which would eive more counts at negative hana positive latitudes. aud differcuces iu the extinction above and below the plane.," 1995), which would give more counts at negative than at positive latitudes, and differences in the extinction above and below the plane."
 The simall differences that here are between the TAIGS aud. DENIS counts cau be attributed to he fact that they are ceutred ou differeut oneitudes a cover different areas of skv., The small differences that there are between the TMGS and DENIS counts can be attributed to the fact that they are centred on different longitudes and cover different areas of sky.
 Before attempting to diseutaugle the b=O° slice. it is Huportant to understaud the shape of the disc. as this will be an important feature in the plane.," Before attempting to disentangle the $b=0^\circ$ slice, it is important to understand the shape of the disc, as this will be an important feature in the plane."
 The slices at |b)=L.75° should be very scusitive to the shape of the disc., The slices at $|b|=1.75^\circ$ should be very sensitive to the shape of the disc.
 They are fay enoush frou the plane that the extinction will be low. typically 2 to 3 magnitudes in the visible or 0.2 to 0.3 mae at AL.," They are far enough from the plane that the extinction will be low, typically 2 to 3 magnitudes in the visible or 0.2 to 0.3 mag at $K_{\rm s}$."
 Furthermore. the lines of seht run sufficiently far above aud below the plane that the voung components close to the plane will not be a major colpoucnt in the counts and only inner Calaxy features with scale heights iun the huudreds of parsecs should provide a significant contribution.," Furthermore, the lines of sight run sufficiently far above and below the plane that the young components close to the plane will not be a major component in the counts and only inner Galaxy features with scale heights in the hundreds of parsecs should provide a significant contribution."
 Fie., Fig.
" 1 shows hat he slices for |b]=1.75"" are more or less sviuuetric iu longitude.", \ref{Fig:DENIScounts9} shows that the slices for $|b|=1.75^\circ$ are more or less symmetric in longitude.
 Initially. the counts rise with decreasing lougitude but then for 30°>|7|12° the counts are basically flat before rising steeply for [/|<107 as the bulee counts become increasingly important.," Initially, the counts rise with decreasing longitude but then for $30^\circ>|l|>12^\circ$ the counts are basically flat before rising steeply for $|l|<10^\circ$ as the bulge counts become increasingly important."
 These plateaux are not expected for a purely exponential disc., These plateaux are not expected for a purely exponential disc.
limit SPIRE is expected to resolve only a small fraction of the CFIRB (Lagacheetal.2003;Fernandez-Conde2008).,"limit SPIRE is expected to resolve only a small fraction of the CFIRB \citep{lag03, fer08}."
". A recent source extraction-based analysis of the SPIRE Science Demonstration Phase (SDP) data — the same data used in this paper — directly resolved 15, 10, and 6 per cent of the CFIRB at 250, 350, and 500um, respectively (Oliveretal.2010).."," A recent source extraction-based analysis of the SPIRE Science Demonstration Phase (SDP) data – the same data used in this paper – directly resolved $15$, $10$, and $6$ per cent of the CFIRB at 250, 350, and 500, respectively \citep{Oliver:2010b}."
" At shorter wavelengths, Berta(2010) directly resolved 52 and 45 per cent of the CFIRB at 100 and 160 uusing Herschel-PACS SDP data."," At shorter wavelengths, \citet{Berta:2010} directly resolved 52 and 45 per cent of the CFIRB at 100 and 160 using -PACS SDP data."
" The observations used in this analysis were obtained with the SPIRE instrument (Griffinetal.2010) on the Space Observatory (Pilbrattetal.2010) as part of the HerMES (Oliver et 22010, in prep) during the SDP."," The observations used in this analysis were obtained with the SPIRE instrument \citep{Griffin:2010} on the Space Observatory \citep{Pilbratt:2010} as part of the HerMES (Oliver et 2010, in prep) during the SDP."
" SPIRE observes simultaneously in three passbands: 250, 350 and 500um."," SPIRE observes simultaneously in three passbands: 250, 350 and 500."
" The on-orbit beam sizes, including the effects of the scanning strategy, are 18.1, 25.2, and 36.6 arcseconds, respectively, with mean ellipticities of 7 -1296."," The on-orbit beam sizes, including the effects of the scanning strategy, are 18.1, 25.2, and 36.6 arcseconds, respectively, with mean ellipticities of 7 -."
" The calibration is based on observations of Neptune, and is described in Swinyardetal. (2010)."," The calibration is based on observations of Neptune, and is described in \citet{Swinyard:2010}."
". Observations of five fields were obtained during SDP, but only three are used in this analysis: GOODS-N, Lockman-North, and Lockman-SWIRE."," Observations of five fields were obtained during SDP, but only three are used in this analysis: GOODS-N, Lockman-North, and Lockman-SWIRE."
 Their properties are summarized in table 1.., Their properties are summarized in table \ref{tbl:data}.
 The Lockman-North region are contained within the shallower Lockman-SWIRE field., The Lockman-North region are contained within the shallower Lockman-SWIRE field.
" The HerMES SDP fields omitted from this analysis are: FLS, which was left out because it is the same depth as the much-larger Lockman-SWIRE field and is significantly contaminated by infrared cirrus, and Abell 2218, because the strong lensing in this field complicates the interpretation of the background number counts."," The HerMES SDP fields omitted from this analysis are: FLS, which was left out because it is the same depth as the much-larger Lockman-SWIRE field and is significantly contaminated by infrared cirrus, and Abell 2218, because the strong lensing in this field complicates the interpretation of the background number counts."
" The detector (bolometer) timelines were processed using the standard SPIRE pipeline, which detects cosmic rays and removes instrumental signatures and temperature drifts (Dowelletal. 2010)."," The detector (bolometer) timelines were processed using the standard SPIRE pipeline, which detects cosmic rays and removes instrumental signatures and temperature drifts \citep{Dowell:2010}."
". The maps were produced using the SMAP package (Levenson et 22010, in prep) using 1/3 beam-FWHM pixels (6, 8 1/3, and 12 arcsec); this is a compromise between adequately sampling the beam and maintaining even coverage over the map."," The maps were produced using the SMAP package (Levenson et 2010, in prep) using $1/3$ beam-FWHM pixels (6, 8 1/3, and 12 arcsec); this is a compromise between adequately sampling the beam and maintaining even coverage over the map."
 Samples flagged as contaminated by cosmic-rays were excluded., Samples flagged as contaminated by cosmic-rays were excluded.
 Each map was masked to form an even-coverage region and was mean subtracted., Each map was masked to form an even-coverage region and was mean subtracted.
" In addition to the even-coverage mask, a small amount of additional masking was required, as there are five resolved sources in the Lockman-SWIRE field."," In addition to the even-coverage mask, a small amount of additional masking was required, as there are five resolved sources in the Lockman-SWIRE field."
" These sources are relatively bright, but are not the brightest in the field."," These sources are relatively bright, but are not the brightest in the field."
" Since the fformalism is based on unresolved point sources, we mask these objects with a 2 arcmin circular mask, and then correct our final number counts using the measured flux of each excluded source."," Since the formalism is based on unresolved point sources, we mask these objects with a 2 arcmin circular mask, and then correct our final number counts using the measured flux of each excluded source."
" Our instrumental noise estimates are based on the technique of Nguyenetal.(2010), and are assumed to have uncertainty, which represents only the uncertainty for a fixed calibration."," Our instrumental noise estimates are based on the technique of \citet{Nguyen:2010}, and are assumed to have uncertainty, which represents only the uncertainty for a fixed calibration."
 The overall SPIRE calibration error is discussed in 5.1.., The overall SPIRE calibration error is discussed in \ref{subsec:systematics}.
 The resulting pixel flux density histograms are shown in figure 1.., The resulting pixel flux density histograms are shown in figure \ref{fig:histos}.
 Smoothing the maps by the beam (via cross-correlation) is beneficial for finding isolated point sources., Smoothing the maps by the beam (via cross-correlation) is beneficial for finding isolated point sources.
" For BLAST observations, Chapinetal.(2010) show that the standard point-source-optimized filter should be modified in the presence of confusion noise."," For BLAST observations, \citet{Chapin:2010} show that the standard point-source-optimized filter should be modified in the presence of confusion noise."
" However, there is no guarantee this will benefit a aanalysis."," However, there is no guarantee this will benefit a analysis."
" Smoothing the map has the effect of broadening the effective beam, which decreases the depth that the ccan probe, while also reducing (but correlating) the instrumental noise."," Smoothing the map has the effect of broadening the effective beam, which decreases the depth that the can probe, while also reducing (but correlating) the instrumental noise."
" We empirically determined if smoothing is beneficial for our analysis by fitting simulated data with and without smoothing and comparing the scatter in the recovered model parameters to the error estimates, and found that it helps for all but our deepest map (GOODS-N); note that our GOODS-N data are several times deeper (relative to the confusion noise) than the deepest BLAST observations."," We empirically determined if smoothing is beneficial for our analysis by fitting simulated data with and without smoothing and comparing the scatter in the recovered model parameters to the error estimates, and found that it helps for all but our deepest map (GOODS-N); note that our GOODS-N data are several times deeper (relative to the confusion noise) than the deepest BLAST observations."
" It is likely that some amount of spatial de-convolution would be beneficial for the GOODS-N field, but since this would significantly complicate the instrumental noise properties of our maps, and hence require extensive testing, we do not pursue de-convolution here, but do beam-smooth the two shallower fields."," It is likely that some amount of spatial de-convolution would be beneficial for the GOODS-N field, but since this would significantly complicate the instrumental noise properties of our maps, and hence require extensive testing, we do not pursue de-convolution here, but do beam-smooth the two shallower fields."
" In addition, we have applied a 6 arcmin high-pass spatial filter to our maps to reduce the effects of clustering on our results."," In addition, we have applied a 6 arcmin high-pass spatial filter to our maps to reduce the effects of clustering on our results."
 The motivation for this is discussed in 3.3.., The motivation for this is discussed in \ref{subsec:filtering}.
" A aanalysis is critically dependent on measurements of the effective beam (which includes the effects of the map making and reduction, as well as any smoothing applied)."," A analysis is critically dependent on measurements of the effective beam (which includes the effects of the map making and reduction, as well as any smoothing applied)."
 Our beam map is based on in-flight fine-scan (highly oversampled) observations of Neptune with a large number of repeats and small offsets between each scan., Our beam map is based on in-flight fine-scan (highly oversampled) observations of Neptune with a large number of repeats and small offsets between each scan.
" These observations allowed us to measure the beam with finer resolution than our data maps to properly match the SPIRE calibration (which is timestream rather than map based), and to build beam maps for the individual bolometers."," These observations allowed us to measure the beam with finer resolution than our data maps to properly match the SPIRE calibration (which is timestream rather than map based), and to build beam maps for the individual bolometers."
 We corrected these observations for the relative motion of Neptune during the scans using the ephemeris computation service at the orbit ofHerschel., We corrected these observations for the relative motion of Neptune during the scans using the ephemeris computation service at the orbit of.
 The Neptune observations are deep enough that the third Airy ring is clearly detected for the array-averaged beams., The Neptune observations are deep enough that the third Airy ring is clearly detected for the array-averaged beams.
" As discussed later 5.1)), beam uncertainties are our largest identified systematic."," As discussed later \ref{subsec:systematics}) ), beam uncertainties are our largest identified systematic."
]envon&Luu(1998.19990) and Wenvon&Bromley(2001.2002a.2004) describe algoritlinis and tests of our multiannulus planet formation code.,"\citet[][1999]{kl98} and \citet[][2002a, 2004]{kb01} describe algorithms and tests of our multiannulus planet formation code."
 Here. we describe an update io the fragmentation algorithm.," Here, we describe an update to the fragmentation algorithm."
 In previous calculations. we used fragmentation prescriplions summarized bv Davis(1985) and Wetherill&Stewart.(1993).," In previous calculations, we used fragmentation prescriptions summarized by \citet{dav85} and \citet{ws93}."
". Both methods write the strength 5 of a pair ol colliding objects as the sum of a constant bulk strength.55. ancl the gravitational binding energy. Ly. For E,x(n;+mj)/ri. the strength is S = Sy -υ where 51 is a constant and rj; is the radius of an object with mass m;-mj."," Both methods write the strength $S$ of a pair of colliding objects as the sum of a constant bulk strength,$S_0$, and the gravitational binding energy, $E_g$ , For $E_g \propto (m_i + m_j) / r_{ij}$, the strength is $S$ = $S_0$ + $S_1 r_{ij}^2$, where $S_1$ is a constant and $r_{ij}$ is the radius of an object with mass $m_i + m_j$."
" When the collision energy. Qj. exceeds 5. Wetherill&Stewart.(1993) derive the mass lost by catastrophic disruption as (he ratio of (he impact energy (to a crushing energy Q,."," When the collision energy, $Q_I$, exceeds $S$, \citet{ws93} derive the mass lost by catastrophic disruption as the ratio of the impact energy to a crushing energy $Q_c$."
 Davisetal.(1985) assume Chat a fixed fraction. fre. ol the impact kinetic energy is (rauslerrecd (to ejected material and derive the lraction of the combined mass lost to disruption.," \citet{dav85}
 assume that a fixed fraction, $f_{KE}$, of the impact kinetic energy is transferred to ejected material and derive the fraction of the combined mass lost to disruption."
 In most cases. the Davisetal.(1985). algorithm vields less debris than the Wetherill&Stewart.(1993) algorithm.," In most cases, the \citet{dav85} algorithm yields less debris than the \citet{ws93} algorithm."
" To take advantage of recent advances in numerical simulations of collisions (e.g..BenzAsphaug1999:Micheletal.2001.2002. 2003).. we now celine a disruption energv (Qu, required to eject of the combined mass of two colliding bodies (equation 3 in the main text)."," To take advantage of recent advances in numerical simulations of collisions \citep[e.g.,][2003]{ben99,mic01,mic02}, we now define a disruption energy $Q_d$ required to eject of the combined mass of two colliding bodies (equation 3 in the main text)."
" For collision energy Qj. the mass ejected in a catastrophic collision is m,; ."," For collision energy $Q_I$, the mass ejected in a catastrophic collision is $m_{ej} = 0.5 (m_1 + m_2) (Q_I/Q_d)^{\beta_e}$ ."
" For most applications we set 3, = 1.125 (e.g..Davisοἱal.1985)."," For most applications we set $\beta_e$ = 1.125 \citep[e.g.,][]{dav85}."
".. When m,;<LO (midm;). we follow Wetherill&Stewart.(1993). and set m,; = 0."," When $m_{ej} < 10^{-8}$ $(m_i + m_j)$, we follow \citet{ws93}
 and set $m_{ej}$ = 0."
 To derive the size and velocity. distribution of the ejected material. we adopt a simple procedure for all collisions.," To derive the size and velocity distribution of the ejected material, we adopt a simple procedure for all collisions."
" We deline the remnant mass. The mass of the largest ejected body is We- adopt a cumulative. size. distribution. forB the remaining. ejected. bodies.. Πρ)x2xm.l"". with 6 = 0.8 (Dohnanvi1969:Willams&Wetherill 1994).. and require (hat themass integrated over the size distribution equal m,;."," We define the remnant mass, The mass of the largest ejected body is We adopt a cumulative size distribution for the remaining ejected bodies, $n_c(m) \propto m^{-b}$, with $b$ = 0.8 \citep{doh69,wil94}, and require that themass integrated over the size distribution equal $m_{ej}$ ."
 We assume that all bodies receive the same kinetic energv per unit mass. given by the initial relative velocities of the two boclies.," We assume that all bodies receive the same kinetic energy per unit mass, given by the initial relative velocities of the two bodies,"
"B,=J/B2+B2 after the flare (cf.",$B_h=\sqrt{B_x^2+B_y^2}$ after the flare (cf.
 Figures 1. (a) and (b))., Figures \ref{f1} (a) and (b)).
" Close temporal association of this field change with flare emissions and its permanence relative to the flare duration are demonstrated in Figure 2((a) covering a period spanning 10 hrs, in which we find that (Bj) at the region R increases by ~30% from ~1300 G to ~1700 G in ~330 minutes."," Close temporal association of this field change with flare emissions and its permanence relative to the flare duration are demonstrated in Figure \ref{f2}( (a) covering a period spanning 10 hrs, in which we find that $\langle B_h \rangle$ at the region R increases by $\sim$ from $\sim$ 1300 G to $\sim$ 1700 G in 30 minutes."
" This rapid evolution ensues from the beginning of the flare at 01:44 UT, with the change-over time cotemporal with the rapid rising of soft X-ray flux and peaking of HXR emissions."," This rapid evolution ensues from the beginning of the flare at 01:44 UT, with the change-over time cotemporal with the rapid rising of soft X-ray flux and peaking of HXR emissions."
" To further characterize the properties of magnetic field, we calculate magnetic shear S, weighted shear angle $, and magnetic inclination angle vy."," To further characterize the properties of magnetic field, we calculate magnetic shear $\tilde{S}$, weighted shear angle $\mathring{S}$, and magnetic inclination angle $\varphi$."
" Here S is defined as the product of field strength and shear angle $=B-6 (Wangetal.1994,2006;Jing2008),, where B=|B|, 0=cos'(B-B,)/(BB,), and the subscript p represents the potential field."," Here $\tilde{S}$ is defined as the product of field strength and shear angle $\tilde{S}=B \cdot \theta$ \citep{wang94,wang06shear,jing08}, where $B=|\mathbi{B}|$, $\theta={\rm cos}^{-1}(\mathbi{B} \cdot \mathbi{B}_p)/(BB_p)$, and the subscript $p$ represents the potential field."
" The weighted shear angle of a region of interest with n pixels is then $=»;S;/V3;Bi, where i=1,2,...n. The inclination angle q relative to the horizontal plane is o=tan!(|B,|/(B2+Bj) 2), "," The weighted shear angle of a region of interest with n pixels is then $\mathring{S}=\sum_{i} \tilde{S}_i / \sum_i B_i$, where i=1,2,...n. The inclination angle $\varphi$ relative to the horizontal plane is $\varphi = {\rm tan}^{-1} (|B_v| / (B_x^2+B_y^2)^{1/2})$ ."
"From the results shown in Figures 2((b)-(d), it can be clearly seen that all of (S), S, and (y) exhibit an abrupt change in the field strength, inclination angle, and azimuthal angle by about 400 G radian, 7°, and —10°, respectively, within the same transition time as (Bj) upon the flare occurrence."," From the results shown in Figures \ref{f2}( (b)–(d), it can be clearly seen that all of $\langle \tilde{S} \rangle$ , $\mathring{S}$, and $\langle\varphi\rangle$ exhibit an abrupt change in the field strength, inclination angle, and azimuthal angle by about 400 G radian, $^{\circ}$, and $-10^{\circ}$, respectively, within the same transition time as $\langle B_h \rangle$ upon the flare occurrence."
" Please note that in order to demonstrate that the rapid changes are very significant compared to variations seen in the long-term evolution, we plot the 30pre (30pos) as error bars in Figure 2,, where σργε {σγος) is derived from the linear fit of the temporal evolution of each quantity in the preflare (postflare) state."," Please note that in order to demonstrate that the rapid changes are very significant compared to variations seen in the long-term evolution, we plot the $3\sigma_{pre}$ $3\sigma_{pos}$ ) as error bars in Figure \ref{f2}, where $\sigma_{pre}$ $\sigma_{pos}$ ) is derived from the linear fit of the temporal evolution of each quantity in the preflare (postflare) state."
" Corroborating our previous studies (Wang&Liu2010,andreferencestherein), the above rapid developments evidenced by the unambiguous HMI observation strongly suggest a more horizontal and sheared state of the photospheric magnetic field at the region R after the flare."," Corroborating our previous studies \citep[][and references therein]{wang10}, the above rapid developments evidenced by the unambiguous HMI observation strongly suggest a more horizontal and sheared state of the photospheric magnetic field at the region R after the flare."
" We note that although the increase of S and $ seems contrary to the relaxation of nonpotentiality as required to energize eruptions, it has been demonstrated using field extrapolations that the increase is localized and both $ and $ decrease above a certain height (Jingetal. 2011a).."," We note that although the increase of $\tilde{S}$ and $\mathring{S}$ seems contrary to the relaxation of nonpotentiality as required to energize eruptions, it has been demonstrated using field extrapolations that the increase is localized and both $\tilde{S}$ and $\mathring{S}$ decrease above a certain height \citep{jing08,sun11,chang11}."
 The magnetic free energy in the 3-D volume is reduced after the flare., The magnetic free energy in the 3-D volume is reduced after the flare.
" On the spatial relationship between the field change and flare emissions, the region R lies between the earliest conjugate HXR footpoint sources at ~001:49 UT (Figure 1((c)) and the ends of the double J-shaped flare ribbons (Figure 1((d))."," On the spatial relationship between the field change and flare emissions, the region R lies between the earliest conjugate HXR footpoint sources at 01:49 UT (Figure \ref{f1}( (c)) and the ends of the double J-shaped flare ribbons (Figure \ref{f1}( (d))."
" Intriguingly, AIA 94 iimages show two extra footpoint-like flare brightenings FP3 and FP4 at the two far ends of the flaring PIL (Figure 1((e))."," Intriguingly, AIA 94 images show two extra footpoint-like flare brightenings FP3 and FP4 at the two far ends of the flaring PIL (Figure \ref{f1}( (e))."
" Co-spatial HXR emissions at FP3 and FP4 were observed few minutes later (Figure 1((f)), and the motion of HXR footpoints as well as the evolution of chromospheric ribbons generally proceed in such a manner as to reduce the magnetic shear, along the PIL (illustrated by the arrows in Figure 1((f)) as reported in eruptive sigmoids(e.g.,Jietal.2008)..These lead us to a picture as we schematically illustrate in Figure 3,, where the flare could be triggered by the tether-cutting reconnection"," Co-spatial HXR emissions at FP3 and FP4 were observed few minutes later (Figure \ref{f1}( (f)), and the motion of HXR footpoints as well as the evolution of chromospheric ribbons generally proceed in such a manner as to reduce the magnetic shear, along the PIL (illustrated by the arrows in Figure \ref{f1}( (f)) as reported in eruptive sigmoids\citep[e.g.,][]{ji08}..These lead us to a picture as we schematically illustrate in Figure \ref{f3}, , where the flare could be triggered by the tether-cutting reconnection"
"(see Watson.Stanger&Griffiths (1984):;Bregman.Schulman&Tomisaka(1995) ancl &Griffiths(1999) among many other papers too numerous to mention) lies 9"" (~160pc) away [rom the dvnamical center of the galaxy.",(see \citet{watson84}; ;\citet{bregman95} and \citet{ptak99} among many other papers too numerous to mention) lies $9\arcsec$ $\sim 160 \pc$ ) away from the dynamical center of the galaxy.
 The nature and origin of these objects remains unknown., The nature and origin of these objects remains unknown.
 The variability argues against models of supernova remnants evolving in dense surroundings (Francoetal.1993).. where a monotonic decline in X-ray Iuminositwv might naively be expected (but note this is nol necessarily always (rue (CidPernandesetal. 1996))).," The variability argues against models of supernova remnants evolving in dense surroundings \citep{franco1993}, where a monotonic decline in X-ray luminosity might naively be expected (but note this is not necessarily always true \citep{cid96}) )."
 The possible association between ihe luminous X-ray source in M82 and the known vounug supernova remnant 41.954-57.5 (Stevens.Strickland&Wills|1999).. advanced based on the source position derived [romROSAT URI data. appears incorrect. based on the more accurate. astrometry.," The possible association between the luminous X-ray source in M82 and the known young supernova remnant 41.95+57.5 \citep{stevens99}, advanced based on the source position derived from HRI data, appears incorrect based on the more accurate astrometry."
 No persistent counterpart to this source at other wavelengths has been confirmed. although the transient radio source 41.4+59.7 (INronberg&Sramek1985).. seen in one observation in 1981 but not subsequently. lies within the positional error circle of the M82 and may be associated with the INO.," No persistent counterpart to this source at other wavelengths has been confirmed, although the transient radio source 41.4+59.7 \citep{kronberg85}, seen in one observation in 1981 but not subsequently, lies within the positional error circle of the M82 \citep{kaaret2001} and may be associated with the IXO."
 The currently favored interpretation is that these might be black holes of mass 10? 104. intermediate between the SMBIIs (pyc10) LO?AL.) found. in traditional AGN such as Sevfert galaxies (formed in some as vel unknown manner. which is possibly related to the growth of stellar bulges (Gebhardtetal.(2000):: (2000)))) and black holes of a few solar masses (formed naturally as a consequence of the stellar evolution of the most massive stars (Brown.Weingartner&Wijers 1996))).," The currently favored interpretation is that these might be black holes of mass $10^{2} \la M_{\rm BH} (\Msol) \la 10^{4}$ , intermediate between the SMBHs $M_{\rm BH} \sim 10^{6}$ – $10^{9}\Msol$ ) found in traditional AGN such as Seyfert galaxies (formed in some as yet unknown manner, which is possibly related to the growth of stellar bulges \citet{gebhardt2000}; \citet{ferrarese2000}) )) and black holes of a few solar masses (formed naturally as a consequence of the stellar evolution of the most massive stars \citep{brown96}) )."
 However. alternative explanations not involving intermediate mass black holes do exist.," However, alternative explanations not involving intermediate mass black holes do exist."
 Ning(2001) argue that these N-rav sources are conventional high-mass X-ray. binaries in which the X-ray emission happens to be beamed in our direction., \citet{king2001} argue that these X-ray sources are conventional high-mass X-ray binaries in which the X-ray emission happens to be beamed in our direction.
 Robertsοἱal.(2001) report a possible optical blue continuum counterpart to an INO in NGC 5204. which might be an extremely luminous O star or a cluster of voung stars.," \citet{roberts2001} report a possible optical blue continuum counterpart to an IXO in NGC 5204, which might be an extremely luminous O star or a cluster of young stars."
" llere we report on the re-detection of a stronely variable A-ray luminous point source (peak Lx~5xI0""eres. 1) located at least 900 pe from the nucleus of the nearby starburst ealaxv NGC 3628 (D = 10 Apc. Soifer (1987))."," Here we report on the re-detection of a strongly variable X-ray luminous point source (peak $L_{\rm X} \sim 5 \times 10^{40} \ergps$ ) located at least 900 pc from the nucleus of the nearby starburst galaxy NGC 3628 (D = 10 Mpc, Soifer (1987))."
 This source. discovered to be variable by Dahlem (1995). had [aded by a factor Z27 between twoROSAT IRI observations taken in late 1991 and early 1994 (when it was too [nint to be detected).," This source, discovered to be variable by Dahlem (1995), had faded by a factor $\ga 27$ between two HRI observations taken in late 1991 and early 1994 (when it was too faint to be detected)."
 This source is the brightest N-rav. source in a recent. 52 ks-Jong Chandra ACIS-S observation we obtainedof NGC 3623., This source is the brightest X-ray source in a recent 52 ks-long Chandra ACIS-S observation we obtainedof NGC 3628.
 With the arcsecond spatial resolution of we have been able to measure the, With the arcsecond spatial resolution of we have been able to measure the
are generally inferred from. seeing-limited photometric or spectroscopic fits to determine the spectral types of the bodies making up the binary.,are generally inferred from seeing-limited photometric or spectroscopic fits to determine the spectral types of the bodies making up the binary.
 The discovery of closer companions can thus dramatically alter the inferred. total masses of the systems. and therefore the expected applicable separation limits.," The discovery of closer companions can thus dramatically alter the inferred total masses of the systems, and therefore the expected applicable separation limits."
 If a significant fraction of apparently wide binaries are actually much higher mass high-order multiple systems. the inferences on their formation environments must be correspondingly altered.," If a significant fraction of apparently wide binaries are actually much higher mass high-order multiple systems, the inferences on their formation environments must be correspondingly altered."
 In addition. wide binary systems which survive the formation process may have experienced smaller dynamic influences on their evolution.," In addition, wide binary systems which survive the formation process may have experienced smaller dynamic influences on their evolution."
 It is concetvable that lower-mass substellar or even planetary companions may be more common in such systems., It is conceivable that lower-mass substellar or even planetary companions may be more common in such systems.
 New proper motion surveys based on the Sloan Digital Sky Survey (SDSS:?) are finding large samples of M-dwarf binaries which apparently violate the 2? ?/— empirical limits (e.g. SLoWPoKES: ?))., New proper motion surveys based on the Sloan Digital Sky Survey \citep[SDSS; ][]{York2000} are finding large samples of M-dwarf binaries which apparently violate the \citet{Reid2001} \citet{Burgasser2003} empirical limits (e.g. SLoWPoKES; \citealt{Dhital2010}) ).
" For the first time. these new samples allow us to undertake detailed statistical studies of the systems"" multiplicity properties."," For the first time, these new samples allow us to undertake detailed statistical studies of the systems' multiplicity properties."