ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2. Llowever. it is suspected that the Haring mechanism is the same for both the objects with very lower 5m and those with higher m.," However, it is suspected that the flaring mechanism is the same for both the objects with very lower $\dot{m}$ and those with higher $\dot{m}$."3 In our sample. Hi must have a broad cistribution: the LLACGNs accrete material at a very lower rate and the luminous QSOs may accrete material at high rate.," In our sample, $\dot{m}$ must have a broad distribution: the LLAGNs accrete material at a very lower rate and the luminous QSOs may accrete material at high rate."4 Probably the X-ray emission mechanism for LLACNs is dillerent from. normal Sevfert 1 galaxies and QSOs., Probably the X-ray emission mechanism for LLAGNs is different from normal Seyfert 1 galaxies and QSOs.5 For example. the X-rav may be emitted in an inner hot ADAE. for LLAGNs. such as MIST. NCC4258 and NGC4579 (Ganunic.Naravan&Dlandford1999:Quataertetal. 1999).. while it is possibly emitted. from a corona for. normal Sevfert. galaxies. and QSOs.," For example, the X-ray may be emitted in an inner hot ADAF for LLAGNs, such as M81, NGC4258 and NGC4579 \cite{gnb,quataert}, while it is possibly emitted from a corona for normal Seyfert galaxies and QSOs."6" Another possibility is à toy model of obscurational variability in AGNs. which can also give a relationship otween ez, and the size of central black hole. but in a more complicated wav ancl requiring fine tuning of several xwameters. such as the covering factor. the number of the clouds. ete. ("," Another possibility is a toy model of obscurational variability in AGNs, which can also give a relationship between $\sigma^2_{\rm rms}$ and the size of central black hole, but in a more complicated way and requiring fine tuning of several parameters, such as the covering factor, the number of the clouds, etc. ("7see Abrassart Czerny 2000 Lor details).,see Abrassart Czerny 2000 for details).8 The X-ray spectra of raclio-loucl objects are generally latter than those of radio-quiet objects., The X-ray spectra of radio-loud objects are generally flatter than those of radio-quiet objects.9 It is believed. that he X-ray is emitted in the jet rather than the disk-corona or radio-oud objects., It is believed that the X-ray is emitted in the jet rather than the disk-corona for radio-loud objects.10 So. the N-rav variability should. be enhanced. by relativistic ellects.," So, the X-ray variability should be enhanced by relativistic effects."11 Three radio-Ioud. objects. 3C120. 8C390.8 and 80373. are included. in our sample.," Three radio-loud objects, 3C120, 3C390.3 and 3C373, are included in our sample."12 lHlowever. they donot show significantly. enhanced: excess variability as shown in Figure 1l..," However, they donot show significantly enhanced excess variability as shown in Figure \ref{fig-1}. ."13 Phis may be caused. by that those objects have hieh inclinations of the disk., This may be caused by that those objects have high inclinations of the disk.14 Indeed.," Indeed,"15 aud tlC Werecy evolutiou of quasars aud red ealaxiesi, and the merger--driven evolution of quasars and red galaxies.16on The prevalence of this process is further su»ported by direct observational evidence for outflows iu both Sevfert ealaxies and amore Iunuinuous qiasars driven by stroug thermal aud radiative feedback originating from the ACN?)., The prevalence of this process is further supported by direct observational evidence for outflows in both Seyfert galaxies and more luminous quasars driven by strong thermal and radiative feedback originating from the AGN.17. Despite these observational aud. theorctica acivauces. he fundamental character of these SMDIT correlatious remains poorly understood.," Despite these observational and theoretical advances, the fundamental character of these SMBH correlations remains poorly understood."18 It appears clear the SAIBIIs are closely linked with the structiral properties of their ost ealaxyw’s bulge?2)., It appears clear the SMBHs are closely linked with the structural properties of their host galaxy's bulge.19". Ioxcever. an important question has not been adequately addressed: to he extent that the erowth of SMDIIs is self.regulated. wlat xopertv οἱ properties — of this bulee docs the SMDII ""see"""," However, an important question has not been adequately addressed: to the extent that the growth of SMBHs is self–regulated, what property – or properties – of this bulge does the SMBH “see?"""20" The more ""fuudanieutal the scaling relation. he nore it reflects the plivsical moechauisin driving the CO-evolution of SMDIIs aud bulges."," The more “fundamental"" the scaling relation, the more it reflects the physical mechanism driving the co-evolution of SMBHs and bulges."21 suggest that he final mass of the SMDII is set by the depth of he ocal potential: the SAIBIT erows uutil feedback uubiids he local gas supply. abruptv ternunating its crowth.," suggest that the final mass of the SMBH is set by the depth of the local potential; the SMBH grows until feedback unbinds the local gas supply, abruptly terminating its growth."22 However. this hypothesis. while prouising. has not vet COLL systematically tested.," However, this hypothesis, while promising, has not yet been systematically tested."23 Tn this paper. we use hydrodynamical simulations to exanune the selfregulated erowth of SMDIIs in three different fucling modes: major niergers. niünor niergers. and disk instabilities.," In this paper, we use hydrodynamical simulations to examine the self–regulated growth of SMBHs in three different fueling modes: major mergers, minor mergers, and disk instabilities."24 Our aim is to investigate the physical mechanism that determines the final mass of the SMDIT., Our aim is to investigate the physical mechanism that determines the final mass of the SMBH.25 It is orgauized as follows: in 2 we review our methodology. in 3 we describe the simulations. ins) [l we present our results. and iu 5 we CUSCTISS implications.," It is organized as follows: in \ref{sec:methods} we review our methodology, in \ref{sec:sims} we describe the simulations, in \ref{sec:results} we present our results, and in \ref{sec:discussion} we discuss implications."26values given by Harju et al. (,values given by Harju et al. (271993).,1993).28 They find a mean value of \(NET)~2.6:10.* in asample of 22 ammonia clumps in Orton., They find a mean value of $\chi ( NH_{3}) \sim 2.6 \cdot 10^{-8}$ in a sample of 22 ammonia clumps in Orion.29 Ammonia is considered as a molecule characterizing later stages of chemical evolution., Ammonia is considered as a molecule characterizing later stages of chemical evolution.30" Myers Benson (1993) found NJ, to be more abundant in older cores. where stars have already formed."," Myers Benson (1993) found $NH_{3}$ to be more abundant in older cores, where stars have already formed."31 This may indicate that [5055 J 20298+3559 is a rather young object., This may indicate that ISOSS J 20298+3559 is a rather young object.32 Tab., Tab.33 3. summarizes our ammonia results., \ref{ammonia_results} summarizes our ammonia results.34 ISOSS ] 2029843559 is located towards the southern outskirts of the Cygnus X Giant Molecular Cloud (Dame, ISOSS J 20298+3559 is located towards the southern outskirts of the Cygnus X Giant Molecular Cloud (Dame35"The results are that the pipeline certainly finds all targets with projected separations between 3” and 5"".",The results are that the pipeline certainly finds all targets with projected separations between $3''$ and $5''$.36 For smaller separations the pipeline is more erratic. especially for fainter systems.," For smaller separations the pipeline is more erratic, especially for fainter systems."37 We therefore exclude an area of 3” around each target from our search for companions and correct for the missed region via the (0 weight. where we assume that the distribution of pairs is uniform with projected separation over these small scales.," We therefore exclude an area of $3''$ around each target from our search for companions and correct for the missed region via the $w_b$ weight, where we assume that the distribution of pairs is uniform with projected separation over these small scales."38 The size of the added factor lies between and depending on redshift., The size of the added factor lies between and depending on redshift.39 The SDSS pipeline also tends to ‘double count’ light for close pairs. making these galaxies brighter than they would otherwise be.," The SDSS pipeline also tends to `double count' light for close pairs, making these galaxies brighter than they would otherwise be."40 Because fainter galaxies have more mergers and minor mergers are more common (Patton&Atfield2008).. this has the effect of artificially raising the pair fraction.," Because fainter galaxies have more mergers and minor mergers are more common \citep{patton08}, this has the effect of artificially raising the pair fraction."41" Masjedietal.(2006) estimate that this should decrease the actual merger rate for galaxies with c,<5’ by a factor of about 5.", \cite{masjedi06} estimate that this should decrease the actual merger rate for galaxies with $r_p < 5''$ by a factor of about 5.42 Since this is an uncertain correction and we are interested in an upper limit to the merger rate. we do not consider this effect here (the effect cannot be well modelled for our sample. as the SDSS pipeline cannot be run locally).," Since this is an uncertain correction and we are interested in an upper limit to the merger rate, we do not consider this effect here (the effect cannot be well modelled for our sample, as the SDSS pipeline cannot be run locally)."43 We also correct for possible companions missed by the redshift cuts' we apply to the survey., We also correct for possible companions missed by the redshift `cuts' we apply to the survey.44" If the primary lies within 500 km ! of the redshift boundaries. we ignore all companions between the primary and the borders and apply a weight ο=2 to all other companions in the opposite ""direction."," If the primary lies within 500 km $^{-1}$ of the redshift boundaries, we ignore all companions between the primary and the borders and apply a weight $w_{v_2}=2$ to all other companions in the opposite `direction'."45 We also need to apply a similar weight to account for potential primaries Jost in the redshift boundary., We also need to apply a similar weight to account for potential primaries lost in the redshift boundary.46" This weight is the reciprocal of the weight applied to the secondaries. Le. ut,=1/2 As all redshift surveys. 2SLAQ is not complete to its flux limit."," This weight is the reciprocal of the weight applied to the secondaries, i.e., $w_{v_1}=1/2$ As all redshift surveys, 2SLAQ is not complete to its flux limit."47 Our success rate is 574. for all galaxies within Sample 8 of C06 but the spectroscopic incompleteness may be dependent on the separation between galaxies., Our success rate is $87\%$ for all galaxies within Sample 8 of C06 but the spectroscopic incompleteness may be dependent on the separation between galaxies.48" Fibers in the 2dF positioner cannot be placed closer than about 25"" from each other in every single configuration.", Fibers in the 2dF positioner cannot be placed closer than about $25''$ from each other in every single configuration.49 However. this effect is compensated by the overlap between individual 28LAQ tiles and by the fact that fiber configurations were generally retouched halfway through each (typically 4 hours) exposure to place fibers assigned to galaxies for which a reliable redshift had already been obtained on to a nearby target (see C06 for a description of the observations).," However, this effect is compensated by the overlap between individual 2SLAQ tiles and by the fact that fiber configurations were generally retouched halfway through each (typically 4 hours) exposure to place fibers assigned to galaxies for which a reliable redshift had already been obtained on to a nearby target (see C06 for a description of the observations)."50 In order to correct for this source of bias we need to estimate the relative incompleteness for close pairs over the range of separations of interest (corresponding to 207;! kpe) and compare it with the incompleteness at large separations. where fiber interactions are not important.," In order to correct for this source of bias we need to estimate the relative incompleteness for close pairs over the range of separations of interest (corresponding to 20 $h^{-1}$ kpc) and compare it with the incompleteness at large separations, where fiber interactions are not important."51" We follow PO2 and estimate this weight bycomputing the ratio between the number of pairs between galaxies with redshift information CN.) and the number of pairs in the input photometric sample (.V,,,,). which is by definition complete. as a function of angular separation 0 The weight to be applied i5 the ratio between the spectroscopic and photometric pair completeness at each separation normalized to the value at large separations."," We follow P02 and estimate this weight bycomputing the ratio between the number of pairs between galaxies with redshift information $N_{zz}$ ) and the number of pairs in the input photometric sample $N_{pp}$ ), which is by definition complete, as a function of angular separation $\theta$ The weight to be applied is the ratio between the spectroscopic and photometric pair completeness at each separation normalized to the value at large separations."52 In 2SLAQ we have higher overall completeness than CNOC2? (by about a factor of 2) but sample a smaller range of projected separations (because of our higher mean redshift)., In 2SLAQ we have higher overall completeness than CNOC2 (by about a factor of 2) but sample a smaller range of projected separations (because of our higher mean redshift).53" We plot N..(0)/N,,(0) vs. 0 in Fig."," We plot $N_{zz} (\theta) /54N_{pp} (\theta)$ vs. $\theta$ in Fig."55 2 for Sample 8 targets: the error bars are assumed to be Poissonian., 2 for Sample 8 targets; the error bars are assumed to be Poissonian.56" Although the data are noisy (because there are relatively few potential pairs to start with). the value of at 0-<G"" is consistent with the value of this ratio at V../.V,,,0>100"", arguing that we are not systematically more incomplete at small separations than at large ones where we are not affected by fiber collisions."," Although the data are noisy (because there are relatively few potential pairs to start with), the value of $N_{zz}/N_{pp}$ at $\theta < 6''$ is consistent with the value of this ratio at $\theta > 100''$, arguing that we are not systematically more incomplete at small separations than at large ones where we are not affected by fiber collisions."57" P02 model the incompleteness in the CNOC2 survey by fitting a polynomial to the ratio δν.(0)/.N,,(0) as a function of projected separation 0.", P02 model the incompleteness in the CNOC2 survey by fitting a polynomial to the ratio $N_{zz}(\theta)/N_{pp}(\theta)$ as a function of projected separation $\theta$.58 Given the small number statistics and noisier nature of our data. it is not fully justified to model the completeness as a function of 0 with a polynomial as done in PO2.," Given the small number statistics and noisier nature of our data, it is not fully justified to model the completeness as a function of $\theta$ with a polynomial as done in P02."59 We therefore adopt an uniform weight of | for 2SLAQ galaxies over the separation of interest. as we do not appear to be systematically more incomplete than at larger separations.," We therefore adopt an uniform weight of 1 for 2SLAQ galaxies over the separation of interest, as we do not appear to be systematically more incomplete than at larger separations."60" Although pair completeness drops at 30""—60"" separations. because of fiber collisions. these large separations are not relevant to our study (as pairs with +,>50h + kpe or Av<1000 km s.1 are largely spurious — POO. De Propris et al."," Although pair completeness drops at $30''$ $60''$ separations, because of fiber collisions, these large separations are not relevant to our study (as pairs with $r_p > 50$ $h^{-1}$ kpc or $\Delta v 61< 1000$ km $s^{-1}$ are largely spurious – P00, De Propris et al."62 2007)., 2007).63 Given the small number statistics. and the fact we are ultimately deriving an upper limit. we eventually decided to follow the approach by DeProprisetal.(2005). to estimate the contribution from close pairs missed because of fiber collisions.," Given the small number statistics, and the fact we are ultimately deriving an upper limit, we eventually decided to follow the approach by \cite{depropris05} to estimate the contribution from close pairs missed because of fiber collisions."64" We search around each of our main targets (Sample 8 galaxies with 23.0<Af(r)21.5 and 0.15 0.65) for a companion (ying within— r,,) in the sample of galaxies (from the input photometric sample) for which we did not obtain a valid redshift."," We search around each of our main targets (Sample 8 galaxies with $-23.0 < M(r) 65< -21.5$ and $0.45 < z < 0.65$ ) for a companion (lying within $r_p$ ) in the sample of galaxies (from the input photometric sample) for which we did not obtain a valid redshift."66 If such a companion exists. we assign to it the same redshift as the primary galaxy and require that the companion lies within the selection lines in Fig.," If such a companion exists, we assign to it the same redshift as the primary galaxy and require that the companion lies within the selection lines in Fig."67 I., 1.68 This identifies all pairs (a total of 3) which are potentially missed because of redshift incompleteness., This identifies all pairs (a total of 3) which are potentially missed because of redshift incompleteness.69 Pairs where both members are missed by the spectroscopic survey. will share in the general 2SLAQ incompleteness. without a bias for incompleteness at small angular separations.," Pairs where both members are missed by the spectroscopic survey, will share in the general 2SLAQ incompleteness, without a bias for incompleteness at small angular separations."70 We can assume that all these 3 ‘extra’ pairs are real and treat them as a source of systematic error on our determination of the pair fraction and the merger rate., We can assume that all these 3 `extra' pairs are real and treat them as a source of systematic error on our determination of the pair fraction and the merger rate.71 Figure 3 shows postage stamp images of the two dynamical pairs we find and of the three possible pairs., Figure 3 shows postage stamp images of the two dynamical pairs we find and of the three possible pairs.72" PO2 also use à weight «, which accounts for the local magnitude incompleteness around each galaxy. a geometric effect due to shit placement and limiting filters in the CNOC2 survey. a color term and a term that depends on the evolution of the galaxy luminosity function over the redshifts covered by the survey."," P02 also use a weight $w_s$ which accounts for the local magnitude incompleteness around each galaxy, a geometric effect due to slit placement and limiting filters in the CNOC2 survey, a color term and a term that depends on the evolution of the galaxy luminosity function over the redshifts covered by the survey."73 This weight is specific to the methods employed by the CNOC? survey (Yeeetal. 1996).., This weight is specific to the methods employed by the CNOC2 survey \citep{yee96}. .74 In our case. we have a very homogeneous sample. there is no geometric effect (other than the one corrected by «'y). all galaxies have similar," In our case, we have a very homogeneous sample, there is no geometric effect (other than the one corrected by $w_{\theta}$ ), all galaxies have similar"75The properties of galaxies in our catalogues are generated using the semi-analytic galaxy formation code. (2222)..,"The properties of galaxies in our catalogues are generated using the semi-analytic galaxy formation code, \citep{COL00,BEN02,BEN03,BAU05}."76 For the purposes of this paper. we may consider a semi-analytic model as being a means of predicting. given some dark matter halo at a redshift of interest. the galaxy population of that halo.," For the purposes of this paper, we may consider a semi-analytic model as being a means of predicting, given some dark matter halo at a redshift of interest, the galaxy population of that halo."77 Having that information. we can construct galaxy luminosity functions. correlation functions. tthat might be considered the results or predictions of the model.," Having that information, we can construct galaxy luminosity functions, correlation functions, that might be considered the results or predictions of the model."78 The tirst step in predicting the galaxy population of a halo is calculating the merger history of the halo., The first step in predicting the galaxy population of a halo is calculating the merger history of the halo.79 In simulations of sufficiently high resolution and with a sufficiently large number of outputs. this can be extracted from the V-body data.," In simulations of sufficiently high resolution and with a sufficiently large number of outputs, this can be extracted from the $N$ -body data."80 In the ease ofGALFORM.. this has been done recently by ὁ with the Millennium Simulation (2): the same simulation has alsobeen used by ? and ? to generate catalogues using a different semi-analytic code.," In the case of, this has been done recently by \citet{BOW06} with the Millennium Simulation \citep{SPR05b}; the same simulation has alsobeen used by \citet{CRO06} and \citet{LUC06} to generate catalogues using a different semi-analytic code."81 The simulations we describe above. by contrast. do not have sufticien resolution for us to extract reliable merger trees for the haloes of interest.," The simulations we describe above, by contrast, do not have sufficient resolution for us to extract reliable merger trees for the haloes of interest."82 A Monte Carlo scheme based on the work of ? and using the algorithm described by ?. is employed instead. therefore.," A Monte Carlo scheme based on the work of \citet{LAC93} and using the algorithm described by \citet{COL00} is employed instead, therefore."83 This generates a merger tree for a halo based only on the halo mass. the cosmology and the initial power spectrum. and does not use other data from the simulation.," This generates a merger tree for a halo based only on the halo mass, the cosmology and the initial power spectrum, and does not use other data from the simulation."84 This scheme does not. therefore. provide galaxy positions: our methodof placing galaxies is given. instead. in Section ??..," This scheme does not, therefore, provide galaxy positions; our methodof placing galaxies is given, instead, in Section \ref{subsec:galplace}."85 Unfortunately. the statistical properties of merger histories generated by this algorithm are not identical to histories extracted directly from an N-body simulation (2)..," Unfortunately, the statistical properties of merger histories generated by this algorithm are not identical to histories extracted directly from an $N$ -body simulation \citep{COL07}."86 2. and ?. have devised empirically motivated modifications to the algorithm to allow Tonte Carlo trees to fit the simulation data better., \citet*{PAR07} and \citet{NEI07} have devised empirically motivated modifications to the algorithm to allow Monte Carlo trees to fit the simulation data better.87 A detailed analysis of the effect of such a modification on semi-analytie galaxy oroperties is beyond the scope of this paper., A detailed analysis of the effect of such a modification on semi-analytic galaxy properties is beyond the scope of this paper.88 We have. though. tested some of our results using the new algorithm of 2.. and find that for our purposes the new trees make little difference.," We have, though, tested some of our results using the new algorithm of \citet{PAR07}, and find that for our purposes the new trees make little difference."89 Given the merger history of a halo. the model computes he evolution of the baryonic content of the halo using a variety of analytic prescriptions.," Given the merger history of a halo, the model computes the evolution of the baryonic content of the halo using a variety of analytic prescriptions."90 Many of the equations governing the shysical processes modelled by contain parameters which may be adjusted., Many of the equations governing the physical processes modelled by contain parameters which may be adjusted.91" Some of these (for example. the form ‘uetor fori/e which governs the size of merger remnants) dave a ""natural value determined by the physics: others (those governing the angular momentum distribution of infalling haloes. say) are derived by comparison to more detailed simulations."," Some of these (for example, the form factor $f_\mathrm{orbit}/c$ which governs the size of merger remnants) have a `natural' value determined by the physics; others (those governing the angular momentum distribution of infalling haloes, say) are derived by comparison to more detailed simulations."92 The function of allowing these parameters to change. then. is to allow investigation into the magnitude of the effect of different physical processes on the resulting galaxy properties in the model.," The function of allowing these parameters to change, then, is to allow investigation into the magnitude of the effect of different physical processes on the resulting galaxy properties in the model."93 Other parameters have no natural value. and can only be fixed by requiring that they take values which allow the model to fit observations.," Other parameters have no natural value, and can only be fixed by requiring that they take values which allow the model to fit observations."94 Much of the time. if we are able to fit some set of observations satisfactorily by choosing the parameters of the model judiciously. the same set of observations could also. be fit reasonably well by some very different choice of parameters.," Much of the time, if we are able to fit some set of observations satisfactorily by choosing the parameters of the model judiciously, the same set of observations could also be fit reasonably well by some very different choice of parameters."95 Therefore. within the framework. we have different models using different physics which are equally good at matching the observations (though this may not. of course. be the case if we were to choose a different set of observations to constrain the model).," Therefore, within the framework, we have different models using different physics which are equally good at matching the observations (though this may not, of course, be the case if we were to choose a different set of observations to constrain the model)."96 Our aim here is to try to constrain cosmological parameters by comparing clustering statistics from a simulation populated with semi-analytic galaxies to the corresponding measurements in an observational survey., Our aim here is to try to constrain cosmological parameters by comparing clustering statistics from a simulation populated with semi-analytic galaxies to the corresponding measurements in an observational survey.97 We would hope that our constraints are insensitive to the precise semi-analytic model used. and we would like to test whether this is the case.," We would hope that our constraints are insensitive to the precise semi-analytic model used, and we would like to test whether this is the case."98 Therefore. although we use only one code.GALFORM.. we use three different “models”. in the sense of different combinations of the physics we attempt to model and the parameters governing that physics.," Therefore, although we use only one code, we use three different `models', in the sense of different combinations of the physics we attempt to model and the parameters governing that physics."99 In the remainder of this section of the paper. we discuss the technical differences between the three models before briefly describing how galaxies are placed in the simulations in Section ??..," In the remainder of this section of the paper, we discuss the technical differences between the three models before briefly describing how galaxies are placed in the simulations in Section \ref{subsec:galplace}. ."100 A reader uninterested in the details of the models may therefore wish to skip to ??.. or toour results in Section ??..," A reader uninterested in the details of the models may therefore wish to skip to \ref{subsec:galplace}, , or toour results in Section \ref{sec:res}. ."101 The three models are as follows:, The three models are as follows:102The color magnitude diagrams (CMDs) of the pbulges are displayed in Fig.,The color magnitude diagrams (CMDs) of the pbulges are displayed in Fig.103 1., 1.104 To equalize the samples. the cluster pbulges have been segregated by their distance from the cluster center at J2000 10:56:59.93 -03:37:36.6 of galaxy 1484 in van Dokkum (2000).," To equalize the samples, the cluster pbulges have been segregated by their distance from the cluster center at J2000 10:56:59.93 -03:37:36.6 of galaxy 1484 in van Dokkum (2000)."105" The virial radius or σου measured as 241"" by Tran (1999) or 1.83 Mpe has been revised using a larger sample of 130 galaxies and is 216"" or 1.64 Mpe.", The virial radius or $R_{200}$ measured as $''$ by Tran (1999) or 1.83 Mpc has been revised using a larger sample of 130 galaxies and is $''$ or 1.64 Mpc.106 All the cluster galaxies 1n our sample lie well within Rago., All the cluster galaxies in our sample lie well within $_{200}$.107 The main result is clear from Fig.1: the slope. color. and dispersion of the CMDs are similar among the four samples.," The main result is clear from Fig.1: the slope, color, and dispersion of the CMDs are similar among the four samples."108 Relative to the solid line showing the median color of UL—Bz0.438 from the 71 cluster pbulges. the field pbulges appear slightly redder.," Relative to the solid line showing the median color of $U-B = 0.438$ from the 71 cluster pbulges, the field pbulges appear slightly redder."109 We also see from Fig., We also see from Fig.110" | that the pbulge colors are red. independent of the distance from the cluster center (δι, average density). pbulge to total (pB/T) ratio. and the color of the pdisk."," 1 that the pbulge colors are red, independent of the distance from the cluster center (i.e., average density), pbulge to total (pB/T) ratio, and the color of the pdisk."111 We also find no support for claims or findings (see references in introduction) of to fractions of bulges or spheroidals being very blue Q) in the field at this redshift or within a broader range of redshifts of 0.7 to 1.0 (GSS8)., We also find no support for claims or findings (see references in introduction) of to fractions of bulges or spheroidals being very blue $U-B \leq 0$ ) in the field at this redshift or within a broader range of redshifts of 0.7 to 1.0 (GSS8).112 Finally. we find that almost all pbulges are as red or redder than pdisks in both field and cluster.," Finally, we find that almost all pbulges are as red or redder than pdisks in both field and cluster."113 Before continuing. a few caveats and possible weaknesses of the survey are noted.," Before continuing, a few caveats and possible weaknesses of the survey are noted."114 The results here apply only to 1) bulges that are luminous (Mp= —19.5) and thus opens the possibility that lower luminosity bulges may have experienced a different formation history: 2) galaxies that are dominatd by one or two components. with the bulge having an r' profile while the disk has an exponential: 3) galaxies that do not have large enough color gradients in either subcomponent to change the results: and 4) bulges that do not belong to obviously merging systems.," The results here apply only to 1) bulges that are luminous $M_{B} \lta -19.5$ ) and thus opens the possibility that lower luminosity bulges may have experienced a different formation history; 2) galaxies that are dominatd by one or two components, with the bulge having an $r^{1/4}$ profile while the disk has an exponential; 3) galaxies that do not have large enough color gradients in either subcomponent to change the results; and 4) bulges that do not belong to obviously merging systems."115 Our field sample with only 21 galaxies is small. so that different properties might apply to minorities.," Our field sample with only 21 galaxies is small, so that different properties might apply to minorities."116 To quantify the CMD. Table 1 gives the results of fits to the data shown in Fig.," To quantify the CMD, Table 1 gives the results of fits to the data shown in Fig."117 |. using the biweight method of Beers.Flynn.&Gebhardt (1990).," 1, using the biweight method of \cite{beers90}."118. Since we found no luminosity dependence of the colors. we locked the slope to O and measured only the colors and dispersions.," Since we found no luminosity dependence of the colors, we locked the slope to 0 and measured only the colors and dispersions."119 We confirm what is seen by eye. namely. that the cluster and field pbulges show very similar colors. with a hint that the inner cluster pbulges may be redder than the middle or outer cluster pbulges and that field pbulges may be slightly redder than cluster pbulges on average. but none of these differences are statistically significant. confidence limit or better. 1.e.. twice the RMS from Table 1).," We confirm what is seen by eye, namely, that the cluster and field pbulges show very similar colors, with a hint that the inner cluster pbulges may be redder than the middle or outer cluster pbulges and that field pbulges may be slightly redder than cluster pbulges on average, but none of these differences are statistically significant confidence limit or better, i.e., twice the RMS from Table 1)."120 The color dispersions are also similar., The color dispersions are also similar.121 Since the measured dispersions from the fits are smaller than that expected from the GIM2D color measurement errors. no reliable intrinsic dispersions can be estimated. though the implication is that they must be small.," Since the measured dispersions from the fits are smaller than that expected from the GIM2D color measurement errors, no reliable intrinsic dispersions can be estimated, though the implication is that they must be small."122 We estimate intrinsic color dispersions of less than 0.05 mag based on the variation of the measured values and GSS8 results., We estimate intrinsic color dispersions of less than 0.05 mag based on the variation of the measured values and GSS8 results.123 The main result is that the restframe colors. slope. and color dispersion of (Mp« —19.5) cluster and field pbulges at redshift z—0.8 are nearly the same.," The main result is that the restframe colors, slope, and color dispersion of $M_B < -19.5$ ) cluster and field pbulges at redshift $z \sim 0.8$ are nearly the same."124 Making the simplifying assumption that colors are a surrogate for ages. we thus find that field and cluster pbulges are universally old (>1.5 Gyr) with no evidence that the ages of luminous pbulges are dependent on radial distance from the cluster center or whether they reside in the cluster or in the field. Le. environment.," Making the simplifying assumption that colors are a surrogate for ages, we thus find that field and cluster pbulges are universally old $ > 1.5$ Gyr) with no evidence that the ages of luminous pbulges are dependent on radial distance from the cluster center or whether they reside in the cluster or in the field, i.e. environment."125 Moreover. while we do find slightly higher proportions of galaxies with large pbulge fractions (pB/T>0.4) and with red pdisks (U—B70.25) in cluster MS1054 than in the field. we find that the very red pbulge colors in both the field and cluster are independent of pbulge luminosity. pB/T. as well as pdisk colors.," Moreover, while we do find slightly higher proportions of galaxies with large pbulge fractions $pB/T > 0.4$ ) and with red pdisks $U-B >1260.25$ ) in cluster MS1054 than in the field, we find that the very red pbulge colors in both the field and cluster are independent of pbulge luminosity, pB/T, as well as pdisk colors."127" The universally red pbulge colors with small dispersion found here can be contrasted with prior studies (see references in introduction) claiming that large fractions (30%—-50%)) of distant early-type galaxies have blue colors or large internal color dispersions that suggest active or recent star formation,", The universally red pbulge colors with small dispersion found here can be contrasted with prior studies (see references in introduction) claiming that large fractions ) of distant early-type galaxies have blue colors or large internal color dispersions that suggest active or recent star formation.128 These other studies differ from that undertaken here. either by including the underlying disk light in. the photometry of the bulge or by including lower luminosity bulges.," These other studies differ from that undertaken here, either by including the underlying disk light in the photometry of the bulge or by including lower luminosity bulges."129 Our results also disagree with models that predict cluster bulges are older than field bulges (e.g.. Kauffmann 1996) or that bulge ages depend on bulge fraction (e.g.. Kauffmann 1996) or on cluster-centric distance (e.g.. Diaferio 2001).," Our results also disagree with models that predict cluster bulges are older than field bulges (e.g., Kauffmann 1996) or that bulge ages depend on bulge fraction (e.g., Kauffmann 1996) or on cluster-centric distance (e.g., Diaferio 2001)."130 Our findings for distant galaxies. are. however. fully consistent with the conclusions of Peletierefαἱ.(1999) for a local sample of bulges in early-type galaxies.," Our findings for distant galaxies are, however, fully consistent with the conclusions of \citet{pel99} for a local sample of bulges in early-type galaxies."131 They derived from HST ΜΕΡΟΣ and NICMOS images the half-light colors of bulges in groups and the field., They derived from $HST$ WFPC2 and NICMOS images the half-light colors of bulges in groups and the field.132 The colors were similar to that of elliptical galaxies in the Coma cluster and implied ar age of 10 Gyr. corresponding to redshifts z~2.," The colors were similar to that of elliptical galaxies in the Coma cluster and implied an age of 10 Gyr, corresponding to redshifts $z \sim 2$."133 The color spread within an effective radius where dust extinction. was negligible was also found to be equally small (~ 0.10 mag i B-I. equivalent to 0.05 mag in U—B). implying an age spreac of at most 2 Gyr.," The color spread within an effective radius where dust extinction was negligible was also found to be equally small $\sim$ 0.10 mag in $B-I$, equivalent to 0.05 mag in $U-B$ ), implying an age spread of at most 2 Gyr."134 Finally. because the bulge ages were so similar and old. they concluded that secular evolution of disks would have difficulty in forming bulges in early-type galaxies.," Finally, because the bulge ages were so similar and old, they concluded that secular evolution of disks would have difficulty in forming bulges in early-type galaxies."135 Exponential bulges of many late-type spirals. however. appear to have colors and thus presumably ages that are correlated with their disk colors. and lending support to secular evolution scenarios for this class of bulges (Carollo 2001: Kormendy Kennicutt 2004).," Exponential bulges of many late-type spirals, however, appear to have colors and thus presumably ages that are correlated with their disk colors, and lending support to secular evolution scenarios for this class of bulges (Carollo 2001; Kormendy Kennicutt 2004)."136 This work places no constraint on such bulges if they are of lummosity fainter than our limit of My~—19.5.," This work places no constraint on such bulges if they are of luminosity fainter than our limit of $M_B137\sim -19.5$."138 The second key result is that the pbulge colors are very red with U—B~0.45.," The second key result is that the pbulge colors are very red with $U-B139\sim 0.45$."140 This color lies between the LU—5=0.52 at My2—20.5 of a fit to the CMD of 409 E-SO's from the RC3 (Schweizer Seitzer 1992) and the U—B=0.42 of 379 SU's (Table 2 of Fukugita et al 1995) or U—B=0.40 of 30 bulges in early-type spirals (Peletier Balcells 1996)., This color lies between the $U-B = 0.52$ at $M_B = -20.5$ of a fit to the CMD of 409 E-S0's from the RC3 (Schweizer Seitzer 1992) and the $U-B = 0.42$ of 379 S0's (Table 2 of Fukugita et al 1995) or $U-B = 0.40$ of 30 bulges in early-type spirals (Peletier Balcells 1996).141 The rough constancy of colors over the last 7 Gyr is a surprise for which we have no simple. compelling explanationconsidered.," The rough constancy of colors over the last 7 Gyr is a surprise for which we have no simple, compelling explanation."142 Since bulges should be bluer by 0.2 mag or more in LU—B at lookback times corresponding to redshifts zo0.83. we expect to observe typical colors nearer U—B 0.20 - 0.32. depending on the formation epoch of bulges.," Since bulges should be bluer by 0.2 mag or more in $U-B$ at lookback times corresponding to redshifts $z \sim 0.83$, we expect to observe typical colors nearer $U-B \sim $ 0.20 - 0.32, depending on the formation epoch of bulges."143 If. however. an old. metal-rich stellar population were to be combined with small amounts of additional star-formation (or equally old but bluer stars of lower metallicity). significant changes in luminosity from passive evolution of the stellar population can be accompanied by nearly constant restframe U—B colors (GSS8. GSS9: Schiavon in preparation: Harker in preparation).," If, however, an old, metal-rich stellar population were to be combined with small amounts of additional star-formation (or equally old but bluer stars of lower metallicity), significant changes in luminosity from passive evolution of the stellar population can be accompanied by nearly constant restframe $U-B$ colors (GSS8, GSS9; Schiavon in preparation; Harker in preparation)."144" Such “drizzling” or ""frosting"" pictures find support ranging from detailed studies of local old stellar populations (Trager 2000. Yi 2005) as well as observations that distant early-type galaxies sometimes exhibit ΠΟΠ emission. lines that may indicate small"," Such ""drizzling"" or ""frosting"" pictures find support ranging from detailed studies of local old stellar populations (Trager 2000, Yi 2005) as well as observations that distant early-type galaxies sometimes exhibit [OII] emission lines that may indicate small"145(pointing errors. attitude control. receiver. e(c.).,"(pointing errors, attitude control, receiver, etc.),"146 along wilh a minor dispersive contribution from the interstellar media. asteroil belt. and Ixuiper belt. etc.," along with a minor dispersive contribution from the interstellar media, asteroid belt, and Kuiper belt, etc."147 The plasma contribution 65; to the total deflection angle is related to the change in the optical path where e is the electrons charge. i its mass. and :/N(6) is the total columnar electron content along the beam. ἂν={n.dt.," The plasma contribution $\theta_{\tt pl}$ to the total deflection angle is related to the change in the optical path where $e$ is the electron's charge, $m_e$ it's mass, and $N_e(\ell)$ is the total columnar electron content along the beam, $N_e=\int n_e d\ell$."148 Therefore. in order to calibrate the plasma term. we should know the electron densitv along the path.," Therefore, in order to calibrate the plasma term, we should know the electron density along the path."149" We start by decomposing the electron clensily 5, in static. spherically svinmetrie part (ή) plus a fluctuation à». i.e. The steady-state behavior is reasonably well known. and we can use one of the several plasma models found in the literature (Tvleretal.1977:Muhbleman.Esposito.aud1977:Muhlemanand.Anderson 1981)."," We start by decomposing the electron density $n_e$ in static, spherically symmetric part $\overline{n}_e(r)$ plus a fluctuation $\delta n_e$, i.e. The steady-state behavior is reasonably well known, and we can use one of the several plasma models found in the literature \citep{tyl77,muh77,muh81}."150. To be more explicit. we will reler to one particular model. namely (we do not consider here a correction factor due to the heliographic latitude): where jj=r/R...," To be more explicit, we will refer to one particular model, namely (we do not consider here a correction factor due to the heliographic latitude): where $\eta=r/{\cal R}_\odot $."151 M large distances this model gives the expected behavior x1/7? of the solar wind., At large distances this model gives the expected behavior $\propto 1/r^2$ of the solar wind.152 We will now determine the contribution to the total deflection angle due to the solar plasma for the model above., We will now determine the contribution to the total deflection angle due to the solar plasma for the model above.153 Following the usual method outlined in lessοἱal. (1999)... we obtain the corresponding deflection due to solar plasma 85 as follows: with 4)=6.32MIIz.," Following the usual method outlined in \cite{and97}, , we obtain the corresponding deflection due to solar plasma $\theta_{\tt pl}$ as follows: with $\nu_0=6.32 ~{\rm MHz}$."154 Comparing 85; with @g we notice the opposite sign - gravity bends the rav oulwards. plasma inwards απ the different. dependence on 5. plasma being steeper.," Comparing $\theta_{\tt pl}$ with $\theta_{\tt gr}$ we notice the opposite sign - gravity bends the ray outwards, plasma inwards – and the different dependence on $b$, plasma being steeper."155 The plasma deflection as a function of the solar offsetb is shown in the Figure 3.., The plasma deflection as a function of the solar offset$b$ is shown in the Figure \ref{fig:plasma}. .1561n Figure Lowe show the time evolution for run 2. the results for run 1 are similar.,"In Figure 1 we show the time evolution for run 2, the results for run 1 are similar."157 We shall therefore. illustrate our discussion. by considering only run 2 and return to compare the results of run 1 and run 2 when we examine the evolution of the cooling Low., We shall therefore illustrate our discussion by considering only run 2 and return to compare the results of run 1 and run 2 when we examine the evolution of the cooling flow.158 Shown in the figure are the density. internal energy. (pressure) and velocity structure in the cluster gas and radio source at. three dilferent. epochs: 401 Myr. while radio source is still active: 518. Myr approximately-LOO nMyr after thejets have been turned olf: 1.89 Gyr an since thejets were switched olf which is approximately 25 times the mtime of the radio source.," Shown in the figure are the density, internal energy (pressure) and velocity structure in the cluster gas and radio source at three different epochs: 401 Myr while the radio source is still active; 518 Myr approximately 100 Myr after the jets have been turned off; 1.89 Gyr an epoch since the jets were switched off which is approximately 25 times the lifetime of the radio source."159 The evolution of the system while the--jets are switched on follows very closely the results from. simulations (see for example RIIB)., The evolution of the system while the jets are switched on follows very closely the results from earlier simulations (see for example RHB).160 In Figure 1 the cocoon is clearly visible as the low density region: the cocoon expands supersonicallv and is surrounded: by à strong bow shock., In Figure 1 the cocoon is clearly visible as the low density region; the cocoon expands supersonically and is surrounded by a strong bow shock.161 After thejets are switched olf the over-pressured. cocoon continues. to expand and is preceded still by à bow shock., After the jets are switched off the over-pressured cocoon continues to expand and is preceded still by a bow shock.162 The expanding remnants soon reach pressure balance. however in this stage they are very. light compared to the external gas ancl rise buovantlv through the cluster.," The expanding remnants soon reach pressure balance, however in this stage they are very light compared to the external gas and rise buoyantly through the cluster."163 This general. behaviour is similar to recent studies by various authors (Churazovοἱal. Ol: RIB).," This general behaviour is similar to recent studies by various authors \citealt{Churazov et al,164BruggenKaiser, Quilis et al, Saxton et al}; RHB)."165 Were however as in RIB the initial conditions correspond. more realistically to those pertaining to a dead radio source since we have followed the initial evolution when thejets are active and as shown by Alexander(2002). this active phase results in à substantial amount of gas being dragged. from the centre of the cluster by the expanding source., Here however as in RHB the initial conditions correspond more realistically to those pertaining to a dead radio source since we have followed the initial evolution when the jets are active and as shown by \citet{Alexander} this active phase results in a substantial amount of gas being dragged from the centre of the cluster by the expanding source.166 The long-term. interaction between the buovantly rising remnants and the cluster is an interplay between the remnants continuing to crag eas out of the cluster and inflow driven both by the cluster σας infilline between the remnants and continued cooling within the cluster., The long-term interaction between the buoyantly rising remnants and the cluster is an interplay between the remnants continuing to drag gas out of the cluster and inflow driven both by the cluster gas infilling between the remnants and continued cooling within the cluster.167 In Figure 2 we plot the integrated X-ray Luminosity and temperature in the cluster gas for run 2 at two epochs after thejets have been switched olf., In Figure 2 we plot the integrated X-ray luminosity and temperature in the cluster gas for run 2 at two epochs after the jets have been switched off.168 We shall now consider the evolution of the buovant phase in a little more detail., We shall now consider the evolution of the buoyant phase in a little more detail.169 Initially the buovant remnants are surrounded by a weak bow shock., Initially the buoyant remnants are surrounded by a weak bow shock.170 The simulated. X-ray emission (Figure 2) is reminiscent of the images recently obtained by of radio sources in cooling Low clusters (e.g. Perscus Fabianctal. 2000: 2002: and .X2052/3€$17 lüzzaetal.2000:Blantonal. 20013).," The simulated X-ray emission (Figure 2) is reminiscent of the images recently obtained by of radio sources in cooling flow clusters (e.g. Perseus \citealt{Fabian et al2000}; ; \citealt{Fabian et al2002}; and A2052/3C317 \citealt{Rizza et al,171Blanton et al}) )."172 In. Figure 2b we show the emissivitv-weieghted temperature integrated along the line of sight., In Figure 2b we show the emissivity-weighted temperature integrated along the line of sight.173 Immediately behind the bow shock the temperature is higher than the surrounding gas. however the temperature of the gas decreases systematically from behind the shock towards the contact surface with the remnant radio lobe.," Immediately behind the bow shock the temperature is higher than the surrounding gas, however the temperature of the gas decreases systematically from behind the shock towards the contact surface with the remnant radio lobe."174 This cooling is cliscussecl in detail in Alexander(2002) and results from aciabatic cooling of the swept-up ICM as it is dragged out of the cluster core., This cooling is discussed in detail in \citet{Alexander} and results from adiabatic cooling of the swept-up ICM as it is dragged out of the cluster core.175 Simulations and analytical results presented bv Brighenti&Mathews(2002b) support the formation of cool gas by expansion., Simulations and analytical results presented by \citet{BrighentiMathewsII} support the formation of cool gas by expansion.176 The region occupied by the raclio-emitting plasma here appears saturated in the erev-scale., The region occupied by the radio-emitting plasma here appears saturated in the grey-scale.177 At long times (after approximately TOO Myr) two main shocks are established within the ICM. which we shall refer to as the leading and trailing shocks (see Figure 1).," At long times (after approximately 700 Myr) two main shocks are established within the ICM, which we shall refer to as the leading and trailing shocks (see Figure 1)."178 The leading shock is again a weak bow shock generated by the buovanthy rising remnants. while the trailing shock is produced: by infalling ICM.," The leading shock is again a weak bow shock generated by the buoyantly rising remnants, while the trailing shock is produced by infalling ICM."179 “Phe temperature structure is now particularly interesting., The temperature structure is now particularly interesting.180 The ICM. is heated: slightly by the weak bow shock but then cools until the trailing shock: the infalline eas is then heated by the trailing shock and within the cluster core there is little variation in temperature except in regions associated. with the radio source remnants themselves., The ICM is heated slightly by the weak bow shock but then cools until the trailing shock; the infalling gas is then heated by the trailing shock and within the cluster core there is little variation in temperature except in regions associated with the radio source remnants themselves.181 The buovant remnants undergo substantial mixing with the ICM as they rise through the cluster., The buoyant remnants undergo substantial mixing with the ICM as they rise through the cluster.182 A shell of mixed ICM/radio plasma forms with a core of cooler ICM. material (Figure | and Figure. 2d) which is being dragged by the remnants out of the centre of the cluster (sec the velocity structure in Figure 1)., A shell of mixed ICM/radio plasma forms with a core of cooler ICM material (Figure 1 and Figure 2d) which is being dragged by the remnants out of the centre of the cluster (see the velocity structure in Figure 1).183 The buovant remnant is unstable to large scale instabilities and in these simulations a large part of the remnant is shed at approximately 2 Cave., The buoyant remnant is unstable to large scale instabilities and in these simulations a large part of the remnant is shed at approximately 2 Gyr.184 The results presented here are in broad agreement with the studies of buovant radio plasma presented by a number of authors (Churazovetal.2001:Brüggen&IxaiserQuilisetal.2001Bohringer2€)02:Saxton 2001).. however our results differ in important. details which are due to the initial conditions we establish by following the evolution of a radio source through its active phase.," The results presented here are in broad agreement with the studies of buoyant radio plasma presented by a number of authors \citep{Churazov et al, BruggenKaiser,185Quilis et al, Bohringer et al, Saxton et al}, however our results differ in important details which are due to the initial conditions we establish by following the evolution of a radio source through its active phase."186 In this Section we consider in detail the effect. of the radio source on the cooling How., In this Section we consider in detail the effect of the radio source on the cooling flow.187 As pointed out by RIB approximately half the energy input by the radio source goes directly into heating the cluster gas., As pointed out by RHB approximately half the energy input by the radio source goes directly into heating the cluster gas.188 Phis is a consequence of the self-similar nature of the expansion., This is a consequence of the self-similar nature of the expansion.189" The stored energy of the radio source is just. proportional to p,Vi. while the work done by the radio source on the ICM is fpod. where subscripts e and s refer to the cocoon ancl shocked ICM. respectively anc we make use of the fact that the shocked. LOCAL is approximately atconstant pressure and equal to the pressure within the cocoon."," The stored energy of the radio source is just proportional to $p_c V_c$, while the work done by the radio source on the ICM is $\int p_c dV_s$, where subscripts $c$ and $s$ refer to the cocoon and shocked ICM respectively and we make use of the fact that the shocked ICM is approximately atconstant pressure and equal to the pressure within the cocoon."190" For the case of scl-similar expansion V,xV; and V; and p; have a power-aw dependence on time (Ixaiser&Alexander1997).", For the case of self-similar expansion $V_s \propto V_c$ and $V_c$ and $p_c$ have a power-law dependence on time \citep{KaiserAlexander1997}.191. Hence re stored energy in the cocoon and the energy input. to 1 ICM are proportional at all times and since the volume of the shocked. gas is comparable to that of the cocoon the constant of proportionality is approximately unity., Hence the stored energy in the cocoon and the energy input to the ICM are proportional at all times and since the volume of the shocked gas is comparable to that of the cocoon the constant of proportionality is approximately unity.192 Clearly wrefore over the lifetime of the source sullicient energy. is input to significantly allect the cluster as has been noted w many authors(e.g. Binney&Tabor1995:aLChurazoyet 2001:: RIB: xander 2002)).," Clearly therefore over the lifetime of the source sufficient energy is input to significantly affect the cluster as has been noted by many authors (e.g. \citealt{BinneyTabor,193Churazov et al,194BruggenKaiser,195Quilis et al}; ; RHB; \citealt{Alexander}) )."196 The results presented. here demonstrate hat it is not only the heatingbut also the redistribution of, The results presented here demonstrate that it is not only the heatingbut also the redistribution of197For iain sequence stars iuassive enough to show a convective core (AL211A... for solar-like metallicity). the associated mixed region plavs the role o [a reservolr for clear reactions.,"For main sequence stars massive enough to show a convective core $M\gtrsim1.1M_{\odot}$ for solar-like metallicity), the associated mixed region plays the role of a reservoir for nuclear reactions."198 The evolution pace of tlicse stars aud the time they speud on the main sequence depend ¢ivectlv O11 ie size of this reservoir., The evolution pace of these stars and the time they spend on the main sequence depend directly on the size of this reservoir.199 The imprecise kLowledee we have of the musing processes. particularly at t1ο boundary of the core. ecuerates large uncertainties on f1ο exteuskna | the mixed core aud subsequently ou the «telar age lass for a given set of surface parameers.," The imprecise knowledge we have of the mixing processes, particularly at the boundary of the core, generates large uncertainties on the extension of the mixed core and subsequently on the stellar age and mass for a given set of surface parameters."200 Aol18o xocesses of transport of chemucal elemeus that could ribute to the creation ofa mixed zoue bevoud t10 edge i6 convective core. overshooting is the oue mvoked LOS often.," Among the processes of transport of chemical elements that could contribute to the creation of a mixed zone beyond the edge of the convective core, overshooting is the one invoked most often."201 Iu the deep interior. convective clement:4 iiSC adiabatically.," In the deep interior, convective elements rise adiabatically."202 Thev are accelerated until trev reach the TOSion of coivective stability. 1.6. VidVad: Theli.," They are accelerated until they reach the position of convective stability, i.e. $\nabla\ind{rad}=\nabla\ind{ad}$."203 he buovancy forces cause a braking of liο οdies iu he radiative region.," Then, the buoyancy forces cause a braking of the eddies in the radiative region."204 It is. however. unlikely that thev shotId stop axuptlv at he boundary between he two reelues.," It is, however, unlikely that they should stop abruptly at the boundary between the two regimes."205" Thev nuelt penerate. over a dista1e dudi the rockuns of staημίν owing to them inertia. and generate a region of mining bevoid the edge of t1ο οοτο,"," They might penetrate, over a distance $d\ind{ov}$, in the regions of stability owing to their inertia, and generate a region of mixing beyond the edge of the core."206" This phe1011011Ο1ὰ las been investigated by several auhoTS (see Zahu1997 and references therein). but no satisAwl""m theoretical or nuuerical «escription have been prop.SCd."," This phenomenon has been investigated by several authors (see \cite{1991A&A...252..179Z} and references therein), but no satisfying theoretical or numerical description have been proposed."207" Ii practice. this region is modeled as an adiabatic aver above the core. whose thickuess is a fraction eo of t pressure scale height IT, (d. oovDHT,) aud where t eclucuts are mixed."," In practice, this region is modeled as an adiabatic layer above the core, whose thickness is a fraction $\aov$ of the pressure scale height $H_p$ $d\ind{ov}=\aov H_p$ ) and where the elements are mixed."208 While it is adnutted that Noy ilv à crude account for the complex processes of 1iixi18o the boundary of the convective core. it Is convenieut id usual mo the niocling of stellar interiors to adopt a representatiou of tlrese processes depending on this paraiucter alouc.," While it is admitted that $\aov$ is only a crude account for the complex processes of mixing at the boundary of the convective core, it is convenient and usual in the modeling of stellar interiors to adopt a representation of these processes depending on this parameter alone."209 Different studies have led to a wide range of Moy: betweei zero (Lauger 1986)) and. about 2 (Niong1985))., Different studies have led to a wide range of $\aov$: between zero \cite{1986A&A...164...45L}) ) and about 2 \cite{1985A&A...150..133X}) ).210 Iu. fact. it is currently adnütted that different values of ey unüght be needed to model stars of differen masses and ages (see Claret2007)).," In fact, it is currently admitted that different values of $\aov$ might be needed to model stars of different masses and ages (see \cite{2007A&A...475.1019C}) )."211 We therefore co not have precise knowledge of the amount of ΠΗΝΙΟ a the edge of the core. aud it is oneof the main goals of aseroseisuologv to coustrain it with observations (sce Lebreretal. 1995... Micheletal.200G6)).," We therefore do not have precise knowledge of the amount of mixing at the edge of the core, and it is oneof the main goals of asteroseismology to constrain it with observations (see \cite{1995IAUS..166..135L}, , \cite{2006ESASP1306...39M}) )."212 For iuterinediate-nunass aud high-uass stars. d is adiuitted to play au importa11 role (Maeder1976)).," For intermediate-mass and high-mass stars, $d\ind{ov}$ is admitted to play an important role \cite{1976A&A....47..389M}) )."213 The case of low-inass stars is nof as clear., The case of low-mass stars is not as clear.214 When they reach he ZAMS. these stars preseut a small convective core hat disappears alinost innidiately.," When they reach the ZAMS, these stars present a small convective core that disappears almost immediately."215 It has already beeu nentioned tlat an extra mixing at the edge of this carly convective core nüght increase its longevity. by providing frst more ?C. and then more “Te in the ceuter (Roxbureh 1955).," It has already been mentioned that an extra mixing at the edge of this early convective core might increase its longevity, by providing first more $^{12}$ C, and then more $^3$ He in the center (Roxburgh 1985)."216 Iu he specific case o: the Sun. core overshooting was added i1i the models. but it was concluded that it had 10 relevaur inupact ou the $1ws prescut structure. uuless we add an unreasonable aim.ount of extra mining.," In the specific case of the Sun, core overshooting was added in the models, but it was concluded that it had no relevant impact on the Sun's present structure, unless we add an unreasonable amount of extra mixing."217 Later O1. stellar nodels of low-1eiss stars sueseeosted that the overshooti1g at the edge of t1ο core could make it survive almost until the eud of the nain sequence. although the plenomenoji was not explained (Mowlavi 1993)).," Later on, stellar models of low-mass stars suggested that the overshooting at the edge of the core could make it survive almost until the end of the main sequence, although the phenomenon was not explained \cite{1993ASPC...40..454M}) )."218 Iu this article. we revisit this phenomenon in the case ο“TID203608... a ow-lnass FSV star that presents solar-liSO (cllatious.," In this article, we revisit this phenomenon in the case of, a low-mass F8V star that presents solar-like oscillations."219 It Was οSOLved with the high-resolution specroleter aat the ESO 3.6-1n telesco]ο in Απο 2006 (Mosser ot al., It was observed with the high-resolution spectrometer at the ESO 3.6-m telescope in August 2006 (Mosser et al.220 Re]XS. hereafter 3Os).," 2008, hereafter M08)."221 The authors analyzed the oscilation spectnI ac ideutified 15 (=0 aud {1 eieenmodes., The authors analyzed the oscillation spectrum and identified 15 $\ell=0$ and $\ell=1$ eigenmodes.222" They. fou a inodel that agrees with the physical parameters and all the seiunic parameters but one: the behavior of the snall spacing UV,14)/2 with frequency."," They found a model that agrees with the physical parameters and all the seismic parameters but one: the behavior of the small spacing $\zeroun=\nu_{n,0}-(\nu_{n,1}+\nu_{n-1,1})/2$ with frequency."223 Since Hs a low-inuass star (less tha ul AJ.) the effect of core overshooting was neglected in the preliminary modeling performed in MOS.," Since is a low-mass star (less than 1 $M_\odot$ ), the effect of core overshooting was neglected in the preliminary modeling performed in M08."224 For this range of mass. stars are not expected to have a convective core ou tle iain sequence. except for a small«ne which disappears shortly after the ZAAIS.," For this range of mass, stars are not expected to have a convective core on the main sequence, except for a small one which disappears shortly after the ZAMS."225 In Sect. 2.. ," In Sect. \ref{sect_models}, ,"226weshow that iu the case, weshow that in the case227jobs ofClaimAd; AM>,Recall that migrates $\lceil m/2\rceil$ jobs from machines in $B$.228 ofTj& (J;] are smallat time /* 1and their total processing time isαἱ most Proof /* 1)< O.75L;«," Hence, using the above Lemma \ref{lem:3r3}, we obtain that the total number of migrations is at most $4\lfloor m/2\rfloor + \lceil m/2\rceil\leq 2.5m$."229 jobsbythe choice of/*. (.Cj. Proofof, This finishesthe proof of Theorem \ref{th:4}. .230ty47:—22((6)92:—3.,	t_h.2313. Woosley et al. (, Woosley et al. (232"2004). using the KEPLER code (Weaveral. 1978).. estimate that convection ends when 7.z7.8« and p,z:2.6«10?&em™.","2004), using the KEPLER code \citep{wea78}, estimate that convection ends when $T_c\approx7.8\times10^8\ {\rm K}$ and $\rho_c\approx2.6\times10^9\ {\rm g\ cm^{-3}}$."233" This gives 4,—10s. roughly the convective overturn timescale as shown below."," This gives $t_h\sim10\ {\rm s}$, roughly the convective overturn timescale as shown below."234 Integrating over the burning region of the core. the total lummosity carried by convection is (Woosleyetal.2004) L.z:747:23((9)d.," Integrating over the burning region of the core, the total luminosity carried by convection is \citep{woo04}235 	L_c."2363ergs(10) The convective velocity at the largest scales is V.—(L/4ar-p) 7., The convective velocity at the largest scales is $V_c\approx (L/4\pi r^2\rho)^{1/3}$ .237" What is crucial for driving g-modes is the properties near the top of the convective zone. which has density and radius p, and +,. respectively."," What is crucial for driving -modes is the properties near the top of the convective zone, which has density and radius $\rho_t$ and $r_t$ , respectively."238 The velocity here is, The velocity here is.239"TIDL4em.(13) where ps=p,/10*eem? and 5a=75/105em."," where $\rho_{t,8}=\rho_t/10^8\ {\rm g\ cm^{-3}}$ and $r_{t,8}=r_t/10^8\ {\rm cm}$."240" The spectrum of g-modes excited by the convection is peaked at a frequency equal to the eddy turnover frequency uw.V./H,. where H, 1s the scaleheight at the top of the convection."," The spectrum of -modes excited by the convection is peaked at a frequency equal to the eddy turnover frequency $\omega_c\approx V_c/H_t$, where $H_t$ is the scaleheight at the top of the convection."241" Taking H,=~2«dOpisquEmcm. where87. results wtb(16) Note the convective timescale is f£,~H,/V.—IOs. roughly in agreement with when the convectionshould end. as discussed above."," Taking $H_t\approx2\times10^7\rho_{t,8}^{1/3}g_{10}^{-1}\ {\rm cm}$, where, results in _c. Note the convective timescale is $t_c\sim H_t/V_c\sim 10\ {\rm s}$, roughly in agreement with when the convectionshould end, as discussed above."242 These waves propagate in the non-convective WD surface layers if their frequency satisfies ayΝ. where N is the Brunt-Váusállà frequency.," These waves propagate in the non-convective WD surface layers if their frequency satisfies $\omega_c<N$, where $N$ is the Brunt-Väiisällä frequency."243 This is approximated asΕΕ”.1419) where p (no subscript) is the density at some position near the WD surface. ps=p/10gem™. ky is Boltzmann's constant. 7; is the temperature at the top of the convection. Z is the average charge per ton. and £j. is the Fermi energy for a degenerate. relativistic electron gas.," This is approximated as, where $\rho$ (no subscript) is the density at some position near the WD surface, $\rho_8=\rho/10^8\ {\rm g\ cm^{-3}}$, $k_{\rm B}$ is Boltzmann's constant, $T_t$ is the temperature at the top of the convection, $Z$ is the average charge per ion, and $E_{\rm F}$ is the Fermi energy for a degenerate, relativistic electron gas."244 I take Z=13.8. as Is appropriate for a mixture of equal parts carbon and oxygen. and ignore the scalings with composition to simplify the presentation.," I take $Z=13.8$, as is appropriate for a mixture of equal parts carbon and oxygen, and ignore the scalings with composition to simplify the presentation."245" Since v,«N at the convective boundary. and at shallower depths Nx7*7. the g-modes propagate freely toward the surface."," Since $\omega_c<N$ at the convective boundary, and at shallower depths $N\propto\rho^{-1/3}$, the -modes propagate freely toward the surface."246" The fraction of L, that can be put into g-modes is directly proportional to the Mach number of the convective eddies near the top of the convective zone (Goldreich&Kumar1990)negligible.", The fraction of $L_c$ that can be put into -modes is directly proportional to the Mach number of the convective eddies near the top of the convective zone \citep{gk90}.247".. For a soundspeed c,=(42/3p)7, Maz V,er""(21 1."," For a soundspeed $c_s=(4P/3\rho)^{1/2}$, 	Ma = ."248"4.(23) The g-mode luminosity is then L. Mal, SS24Ηfracp,.", The -mode luminosity is then 	L_g Ma L_c .249o2:5.7ergssl. Comparing the dependence on 7. in equations (26)) and (7)) shows that Ly*f+., Comparing the dependence on $T_c$ in equations \ref{eq:lg}) ) and \ref{eq:theat}) ) shows that $L_g\propto t_h^{-1.4}$.250" The total amount of energy put into up to any given time ts therefore E, = L.dt 22.5Lety 4d08p282327)fracp,.", The total amount of energy put into }-modes up to any given time is therefore 	E_g = L_g 2.5 L_g t_h .25192:2.4ergs. Depending the final Τ.. about goes into e-modes.," Depending the final $T_c$, about goes into }-modes."252" As the g-modes propagate into the non-convective surface layers. they satisfy the dispersion relation and have a group velocity of V,=.../k,. where &, and ky, are the radial and horizontal wavenumbers. respectively."," As the -modes propagate into the non-convective surface layers, they satisfy the dispersion relation and have a group velocity of $V_g=\omega_c/k_r$, where $k_r$ and $k_h$ are the radial and horizontal wavenumbers, respectively."253 This relation does not include rotational modifications., This relation does not include rotational modifications.254 The radial wavenumber is unaffected by the Coriolis force if the spin O is small in comparison to the buoyaney NCH ως sl.," The radial wavenumber is unaffected by the Coriolis force if the spin $\Omega$ is small in comparison to the buoyancy\citep{cl70,bg94}255 N^2 r ."256 In this limit the mode equations can be simplified using the “traditional approximation” to separate vertical and horizontal parts., In this limit the mode equations can be simplified using the “traditional approximation” to separate vertical and horizontal parts.257" Although &, remains the same. Kj can depend onthe angle with respect to the rotation axis. since the Coriolis"," Although $k_r$ remains the same, $k_h$ can depend onthe angle with respect to the rotation axis, since the Coriolis"258D stars.,B stars.259" Moreover. because runavavs are svsteimaticallv ejected at low velocities ίσιο,PortegiesZavart2000). the fastest runaways are those ejected in the direction of Calactic rotation and thus prefercutially found at low Calactic latitudes."," Moreover, because runaways are systematically ejected at low velocities \citep[e.g.,][]{portegies00}, the fastest runaways are those ejected in the direction of Galactic rotation and thus preferentially found at low Galactic latitudes."260 Thus the expected distribution of runaway longitudes aud. latitudes are coutrary to the observed distribution of IIVSs., Thus the expected distribution of runaway longitudes and latitudes are contrary to the observed distribution of HVSs.261 2pt The distribution of Local Croup dwarf galaxies ds anisotropic. possibly due to a tidal origin (e.5..Metzetal.2008).," 2pt The distribution of Local Group dwarf galaxies is anisotropic, possibly due to a tidal origin \citep[e.g.,][]{metz08}."262. À tidal debris origin appears supported by the clumping of IIVS travel times around 100-200 Nvr. however the travel times are simply a product of the IIVS's ~500 lau + velocities aud our iiaenitude-Iited survey depth of 50-100 kpc.," A tidal debris origin appears supported by the clumping of HVS travel times around 100-200 Myr, however the travel times are simply a product of the HVS's $\sim$ 500 km $^{-1}$ velocities and our magnitude-limited survey depth of 50-100 kpc."263 ITVS travel times are in fact problematic for a tidal debris origi because the times are a siguificant fraction of the stars madn sequence lifetimes. and multiple (gas-xich) tidal disruption eveuts would be required to explain the full 2«105 vr span of ITVS travel times.," HVS travel times are in fact problematic for a tidal debris origin because the times are a significant fraction of the stars' main sequence lifetimes, and multiple (gas-rich) tidal disruption events would be required to explain the full $2\times10^8$ yr span of HVS travel times."264 No dwart ealaxy in the Local Group travels with radial velocities comparable to the unbound IIVSs: known dwarf galaxy rocunants like the Ser stream (Ibatactal.1991) are bound., No dwarf galaxy in the Local Group travels with radial velocities comparable to the unbound HVSs; known dwarf galaxy remnants like the Sgr stream \citep{ibata94} are bound.265 We thus consider tidal debris aui uulikelv explanation for the observed set of IIVSs (however.seeAbadietal. 2008).," We thus consider tidal debris an unlikely explanation for the observed set of HVSs \citep[however,266see][]{abadi08}."267. 2pt Wile an equalanass binary MDBIT is ruled out in the Calactic Center (Reid&Druthaler 2001).. theorists speculate that the massive star clusters in the Galactic Center form intermediate mass black holes (IAIBITs) iu their cores.," 2pt While an equal-mass binary MBH is ruled out in the Galactic Center \citep{reid04}, theorists speculate that the massive star clusters in the Galactic Center form intermediate mass black holes (IMBHs) in their cores."268 If such IMDIIS exist. dvuamical friction causes them to in-piral iuto the central MDBIT. prefercutially expecteejecting IVSs from their orbital planes.," If such IMBHs exist, dynamical friction causes them to in-spiral into the central MBH, preferentially ejecting HVSs from their orbital planes."269 Thus the dsignature of à IMDIT nespiral is a rius of IIVSs around the skv (Gualandrisetal.2005:Levin2006:Sesaua 20063.," Thus the expected signature of a IMBH in-spiral is a ring of HVSs around the sky \citep{gualandris05, levin06, sesana06}."270. Daunugdtetal(2006) argue. however. that stellar iuteractions perturh the orbital plane of an in-spiraling IMDITI: the resulting IIVS distribution in this scenario may in fact be isotropic.," \citet{baumgardt06} argue, however, that stellar interactions perturb the orbital plane of an in-spiraling IMBH; the resulting HVS distribution in this scenario may in fact be isotropic."271 Moreover. a single IMDIT inspiral eveut happens ou timescales & shorter than the observed span of ITVS travel times: uniltiple INIT in-piral eveuts are required to explain the observed IIVSs.," Moreover, a single IMBH in-spiral event happens on timescales $\times$ shorter than the observed span of HVS travel times; multiple IMBH in-spiral events are required to explain the observed HVSs."272 2pt The Calactic ceuter contains many well-defined structures., 2pt The Galactic center contains many well-defined structures.273 As illustrated in Pauuudetal.(2006).. the molecular gas cireunruuclear disk and the ionized northern armi are roughly aligned with the plane of the Milkv Way.," As illustrated in \citet{paumard06}, the molecular gas circum-nuclear disk and the ionized northern arm are roughly aligned with the plane of the Milky Way."274 The gaseous nmüuispiral is perpendicular to the plane of the Milkv. Way., The gaseous minispiral is perpendicular to the plane of the Milky Way.275 Notably. the stellar disk 0.1 pe from the AIBIT i$ roughly perpendicular to thegaseous components (Lu2008).," Notably, the stellar disk 0.1 pc from the MBH is roughly perpendicular to thegaseous components \citep{lu08}."276. The stellar disk coutaius massive stars (Tameretal.2006:απατάct2006).. possibly formed im- from a eas accretion disk (Geuzeletal.2003:Levin&Belohorodov2003).," 	The stellar disk contains massive stars \citep{tanner06, paumard06}, possibly formed in-situ from a gas accretion disk \citep{genzel03, levin03}."277. Dynamical interactions between a pair of stellar disks may scatter stars in towards the MBI. explaining both the S-stars aud the IIVSs (Lockimannetal.2008:Perets2008)..," Dynamical interactions between a pair of stellar disks may scatter stars in towards the MBH, explaining both the S-stars and the HVSs \citep{lockmann08b, perets08c}."278 Clearly. the Galactic center contains nou-óisotropie distributious of stars and eas which may provide a natural source for the observed anisotropy of IIVSs cjyected from the Galactic ceuter.," Clearly, the Galactic center contains non-isotropic distributions of stars and gas which may provide a natural source for the observed anisotropy of HVSs ejected from the Galactic center."279 However. it is unclear if the observed structures can persist long enough to explain the anisotropic distribution of IIVSs.," However, it is unclear if the observed structures can persist long enough to explain the anisotropic distribution of HVSs."280 Unbound IIVSs are spatially anisotropic at the σ level., 	Unbound HVSs are spatially anisotropic at the $\sigma$ level.281 The anisotropy is most significant iu Calactic longitude. aud uot im latitude.," The anisotropy is most significant in Galactic longitude, and not in latitude."282 Lower velocity IIVSs are systematically more isotropic. and apparent close pairs of IIVSs are plivsically uurelated.," Lower velocity HVSs are systematically more isotropic, and apparent close pairs of HVSs are physically unrelated."283 The observed distribution of IIVSs is linked to the origin of the ITWSs., 	The observed distribution of HVSs is linked to the origin of the HVSs.284 Abadietal.(2008) propose a tidal debris explanation. although this appears difficult to recoucile with all the observations.," \citet{abadi08} propose a tidal debris explanation, although this appears difficult to reconcile with all the observations."285 We investigate other plivsical models for the anisotropy in a separate paper., We investigate other physical models for the anisotropy in a separate paper.286 Tn the future. easing the distribution of bound aud nubound IIVSs over the southern sky will allow us to better coustrain the anisotropy aud the origin of IIVSs.," In the future, measuring the distribution of bound and unbound HVSs over the southern sky will allow us to better constrain the anisotropy and the origin of HVSs."287 Thus work is based on observatious obtained at the MAIT Observatory. a joint facility of the Smithsonian Tustitution aud the University of Arizona.," 	This work is based on observations obtained at the MMT Observatory, a joint facility of the Smithsonian Institution and the University of Arizona."288 This rescarcli makes use of NASA's Astrophysics Data System Bibliographic Services., This research makes use of NASA's Astrophysics Data System Bibliographic Services.289 We thank the aunouviunous referee and Oleg Cuediu for helpful couuucuts., We thank the anonymous referee and Oleg Gnedin for helpful comments.290 This work was supported by the Siunithsouian Iustitutiou., This work was supported by the Smithsonian Institution.291to the outmost emission region.,to the outmost emission region.292 For the 5 CGllIz polarization image with high signal to noise ratio. the electric field information can be well extracted on a scale of more than 30 mas from the core. corresponding to a linear size of240 pc.," For the 5 GHz polarization image with high signal to noise ratio, the electric field information can be well extracted on a scale of more than 30 mas from the core, corresponding to a linear size of$240$ pc."293 At this [requeney. the electric vector is distributed with PLA. of about 50° near the core. and then bifurcate apparently at a distance of about 5 mas north to the core (see also the polarization image at S CGllz).," At this frequency, the electric vector is distributed with P.A. of about $50^\circ$ near the core, and then bifurcate apparently at a distance of about $5$ mas north to the core (see also the polarization image at 8 GHz)."294 When the jet reaches a distance of about 10 mas from the core. the electric field turns almost perpendicular to the overall jet direction with an overall PX of787.," When the jet reaches a distance of about $10$ mas from the core, the electric field turns almost perpendicular to the overall jet direction with an overall P.A of."295.. Ata distance of about 25 mas north to the core. the electric vector becomes roughly. parallel to the north direction. with an overall PX. of23.," At a distance of about $25$ mas north to the core, the electric vector becomes roughly parallel to the north direction with an overall P.A. of."296. Such an orientation change down the jet may be attributed to magnetic fields that are ordered. by local phenomena at various places in the jet. while Gabuzda.Murray&Cronin(2004) suggested that these alternating magnetic fields may inclicate oscillations or instabilities of a global jet magnetic Ποια.," Such an orientation change down the jet may be attributed to magnetic fields that are ordered by local phenomena at various places in the jet, while \citet{gab04} suggested that these alternating magnetic fields may indicate oscillations or instabilities of a global jet magnetic field."297 Faraday rotation reveals some physical condition along the line of sight., Faraday rotation reveals some physical condition along the line of sight.298" When the polarized. emission. propagates through a magnetized. plasma. the polarization plane will rotates with wavelength A according to the following expression (e.g. Tavlor 1998)). where 4x is polarization angle olfset due to Faraday rotation. N(s) is the electron number density in em.7. is magnetic field in mG. and the integral is taken over a passage in parsec along the line of sight from the source to the observer. RAL is the rotation measure in radm"""," When the polarized emission propagates through a magnetized plasma, the polarization plane will rotates with wavelength $\lambda$ according to the following expression (e.g. \citealt{tay98}) ), where $\Delta\chi$ is polarization angle offset due to Faraday rotation, $\rm{N(\textbf{s})}$ is the electron number density in $\rm{cm^{-3}}$, is magnetic field in mG, and the integral is taken over a passage in parsec along the line of sight from the source to the observer, RM is the rotation measure in $\rm{rad~m^{-2}}$."299 The RM distribution can be used to explore the intervening plasma and magnetic field along the line of sight., The RM distribution can be used to explore the intervening plasma and magnetic field along the line of sight.300 To ensure proper estimation of RAL. at least 3. band of ΕΟΝΑν are required.," To ensure proper estimation of RM, at least 3 band of EVPAs are required."301 Under the limitation of resolution ancl sensitivity for the current observational data. we cannot obtain a transverse RAL variation across several beamwidths.," Under the limitation of resolution and sensitivity for the current observational data, we cannot obtain a transverse RM variation across several beamwidths."302 Dased on the resultant EWDPAs between 15 and 43 Gilz from mocel fitting. the components and are fitted with equation (2).," Based on the resultant EVPAs between 15 and 43 GHz from model fitting, the components and are fitted with equation (2)."303 As a result. the component is badly fitted. while the EVPAs for the component obeved a A7 Earaday rotation law well. with the observed RALof 1061.9250.2racm7.," As a result, the component is badly fitted, while the EVPAs for the component obeyed a $\lambda^2$ Faraday rotation law well, with the observed RM of $-1061.9 \pm 0.2~ \rm{rad~m^{-2}}$."304 Ehe fitting result for both components is shown in Figure 6.. where the errors are estimated from the Stokes Q and C images with 5a of noise level adopted.," The fitting result for both components is shown in Figure \ref{fig6}, where the errors are estimated from the Stokes $Q$ and $U$ images with $5\sigma$ of noise level adopted."305 In view of that the absolute RAL due to our galaxy is usually no more than 200radm7 in any. direction (Pushkearey2001).. and the RAL in the rest frame of are higher by a factor of (1|2)7. the RAL in its rest frame will become quite large. with the absolute value of more than 3000 radm.27. which implies that most of the RAL should. arise from the source itself or plasma near the source. rather than the plasma in our galaxy.," In view of that the absolute RM due to our galaxy is usually no more than $200~ \rm{rad~m^{-2}}$ in any direction \citep{pus01}, and the RM in the rest frame of are higher by a factor of $(1 + z)^2$, the RM in its rest frame will become quite large, with the absolute value of more than 3000 $\rm{rad~m^{-2}}$, which implies that most of the RM should arise from the source itself or plasma near the source, rather than the plasma in our galaxy."306 Since there exists. quite a large Faraday. rotation for component£.. it naturally reminds one that for the Faraday rotation occurring within the jet. radiation emitted. at different depths is rotated through different angles. thus the net Dux will be depolarized. and the polarization level will fall correspondingly.," Since there exists quite a large Faraday rotation for component, it naturally reminds one that for the Faraday rotation occurring within the jet, radiation emitted at different depths is rotated through different angles, thus the net flux will be depolarized, and the polarization level will fall correspondingly."307 By contrast. external Faraday rotation is relatively hard. to result. in depolarization (e.g. Burn 1966)).," By contrast, external Faraday rotation is relatively hard to result in depolarization (e.g. \citealt{bur66}) )."308 For an opticallv-thin uniform source. the internal Faraday dispersion can be deseribed by an additional factor of sin()/d* where &=2.07 (Burn1966:Llomanctal. 2009).," For an optically-thin uniform source, the internal Faraday dispersion can be described by an additional factor of $\rm{sin(\Phi)/\Phi}$ where $\Phi =3092.0\Delta\chi$ \citep{bur66, hom09}."310. Based on the equation (2)). the additional factors are obtained with values of 0.895. 0.975. and 0.998 at the corresponcing frequency 15. 22. and 43 Cllz. respectively.," Based on the equation \ref{eq2}) ), the additional factors are obtained with values of 0.895, 0.975, and 0.998 at the corresponding frequency 15, 22, and 43 GHz, respectively."311 This means that even if all the observed: rotation is completely internal to the jet. this rotation is not large enough to cause all of the decrease in fractional polarization observed. assuming that the component. has comparable degree of polarization without. depolarization.," This means that even if all the observed rotation is completely internal to the jet, this rotation is not large enough to cause all of the decrease in fractional polarization observed, assuming that the component has comparable degree of polarization without depolarization."312 Llomanοἱal.(2009) argued. that the alternative cause inducing depolarization is the opacity by the mecdiums in the passage. by which the depolarization also gets worse with wavelength.," \citet{hom09} argued that the alternative cause inducing depolarization is the opacity by the mediums in the passage, by which the depolarization also gets worse with wavelength."313 This implies that both the opacity and. internal Faraday rotation contribute to the decrease in fractional polarization with wavelength. in which the opacity plavs quite a large part of role in the depolarization.," This implies that both the opacity and internal Faraday rotation contribute to the decrease in fractional polarization with wavelength, in which the opacity plays quite a large part of role in the depolarization."314 In general. the core component in blazars lies in the extreme end of the jet. which is usually optically thick. and highly Doppler boosted with dominant Εαν density. relatively hard spectrum and high brightness temperature.," In general, the core component in blazars lies in the extreme end of the jet, which is usually optically thick, and highly Doppler boosted with dominant flux density, relatively hard spectrum and high brightness temperature."315 ὃν comparing the physical quantities of component and in Table . one can find that component has higher brightness temperature and harder spectrum. both of which are indicators of radio core (c.g. Shenetal. 2005)).," By comparing the physical quantities of component and in Table \ref{tb2}, , one can find that component has higher brightness temperature and harder spectrum, both of which are indicators of radio core (e.g. \citealt{she05}) )."316 On the other hand. the component exhibits higher IHux density," On the other hand, the component exhibits higher flux density"317to the expectation of Sciama(2000a)} who suggested a spectral change above 1 GeV a the point where sclfshiclding becomes important.,to the expectation of \citet{Sci00a} who suggested a spectral change above 1 GeV at the point where self-shielding becomes important.318 For some purposes the cucreyv-iuteeratcd cluissivities are of more interest than their differeutial counterparts. so we present these quantities iu figures 6 and 9.. for cosmic rav protons aud electrons. respectively.," For some purposes the energy-integrated emissivities are of more interest than their differential counterparts, so we present these quantities in figures \ref{fig:f6}319 and \ref{fig:f7}, for cosmic ray protons and electrons, respectively."320" At low column deusities. where cascades arc selt-shiclding are wuninuportaut. there is very little variation of the integrated ciissivity with cloud column deusitv. aud the thin material luit can ↴⋝↸∖⋜∥↧≺∏≻↑↸∖≼↧↕∪↥⋅↸⊳∪↿∐⊔∐↴∖↴↕↸∖↴∖↴↴∖↴↑↕⋜⋯↓∩∶↴∙⊾↸⊳⋯−. ↽⋅≻ for. protons (1↽gcn ""nfor electrons)"," At low column densities, where cascades and self-shielding are unimportant, there is very little variation of the integrated emissivity with cloud column density, and the thin material limit can be adopted for columns less than $10\;{\rm g\,cm^{-2}}$ for protons $1\;{\rm g\,cm^{-2}}$ for electrons)."321 Above this point. however. the cuussivity drops rapidly with increasing cloud. cobluun-deusity.," Above this point, however, the emissivity drops rapidly with increasing cloud column-density."322" In order to eauge the seusitivitv of these calculations to the asse COSDUbC-YAY Spectra. we have computeck our results for two different ποσα cosimic-rav proton spectra. and three ifferent iucideut cosniüc-rav electron spectra. as shown iu figures 6 and 7. respectively,"," In order to gauge the sensitivity of these calculations to the assumed cosmic-ray spectra, we have computed our results for two different incident cosmic-ray proton spectra, and three different incident cosmic-ray electron spectra, as shown in figures 6 and 7, respectively."323 The «Mferenees aye seen to be small im comparison with the variation as a function of cloud colin density. but at a fixed colummu-density the systematic uncertaiuties associated with the input cosmic-ray spectra are nevertheless significant.," The differences are seen to be small in comparison with the variation as a function of cloud column density, but at a fixed column-density the systematic uncertainties associated with the input cosmic-ray spectra are nevertheless significant."324 Usiug the gamma-ray production spectra obtained in the previous section. we have calculated the diffuse οαατα cussion from the Galaxy as follows.," Using the gamma-ray production spectra obtained in the previous section, we have calculated the diffuse gamma-ray emission from the Galaxy as follows."325 The predicted gamma-ray spectrum for cach case is where αλας is the spectrum returned by the sinulation i units of protons/MeV/primary. for an individual cloud. apxopriate to the incident cosmic-ray spectruin.," The predicted gamma-ray spectrum for each case is where ${\rm d}N/{\rm d}E$ is the spectrum returned by the simulation in units of photons/MeV/primary, for an individual cloud, appropriate to the incident cosmic-ray spectrum."326 T1ο quantitv Ja) is the iutensitv of cosmic ravs at a distance 5 along the line of sight. in units of ∐∐∐∡∐↕↸↴∖↸↕∐∖ ⊐↴⊥↴⋅⊥. and p(s) is the mean ¢leuxitv iu gas clouds of column density X.," The quantity $J_{\rm cr}(s)$ is the intensity of cosmic rays at a distance $s$ along the line of sight, in units of primaries $^{-2}$ $^{-1}$ $^{-1}$, and $\rho(s)$ is the mean density in gas clouds of column density $\Sigma$."327 The Calactic variatio1 of the spectrum (5) is not well constrained by existing data. and consequently we adopt tl1¢ situplifving assumption that the shape of the ¢osuc-ray spectra (both," The Galactic variation of the spectrum $J_{cr}(s)$ is not well constrained by existing data, and consequently we adopt the simplifying assumption that the shape of the cosmic-ray spectra (both"328overlay of an undefined number of Gaussians we perform the same analysis as observers and fit straight lines to the inner and outer oxurts of the probability plot.,overlay of an undefined number of Gaussians we perform the same analysis as observers and fit straight lines to the inner and outer parts of the probability plot.329 The probability plot is well fitted with wo lines: low dispersion velocities («|Lo |). from the population of stars confined to the thin disc. and high dispersion velocities (|lo |i. where the population of the stars corresponds to the uck dise.," The probability plot is well fitted with two lines: low dispersion velocities $< \mid1330\sigma \mid$ ), from the population of stars confined to the thin disc, and high dispersion velocities $> \mid1 \sigma \mid$ ), where the population of the stars corresponds to the thick disc."331" In our example. the slope for the velocity dispersion for re stars of the thin dise is 4.01 kmss! and for the thick dise stars 20.5kmss +,"," In our example, the slope for the velocity dispersion for the stars of the thin disc is $4.01$ $^{-1}$ and for the thick disc stars $20.5$ $^{-1}$."332 Additionally we measure the full width half maximum of qe Velocity distribution and give these values as the last entry in Tab. 3.., Additionally we measure the full width half maximum of the velocity distribution and give these values as the last entry in Tab. \ref{tab:vel}.333 As with the 2. distribution we now investigate how our results change if we adopt an ICMF., As with the $z-$ distribution we now investigate how our results change if we adopt an ICMF.334 We combine our models the same way as done for the z-distribution and show the combined velocity distribution in the top panel of Fig., We combine our models the same way as done for the $z$ -distribution and show the combined velocity distribution in the top panel of Fig.335 4. (and also in the last line of Tab. 35)., \ref{fig:vel-icmf} (and also in the last line of Tab. \ref{tab:vel}) ).336 The derived probability plot to this distribution is shown in the lower panel of Fig. 4.., The derived probability plot to this distribution is shown in the lower panel of Fig. \ref{fig:vel-icmf}.337" Fitting straight lines to the inner and ""uter parts give velocity dispersions of 1.3 and SO + for the thin and thick dise component respectively.", Fitting straight lines to the inner and outer parts give velocity dispersions of $1.3$ and $80$ $^{-1}$ for the thin and thick disc component respectively.338 As we can see. the velocity dispersion is much larger in the thick dise than in the thin disc.," As we can see, the velocity dispersion is much larger in the thick disc than in the thin disc."339 Comparing our values of the solar neighborhood with published observations (Bensbyetal.2003:Vallenariet2006:Veltz 2008).. which report Coμοι=3S 1 and oythin16 + for the thick dise and thin dise respectively. shows that we underestimate the thin dise values quite substantially and overestimate the value for the thick disc.," Comparing our values of the solar neighborhood with published observations \citep{bens03,vall06,veltz08}, which report $\sigma_{\rm w,thick} =34038$ $^{-1}$ and $\sigma_{\rm w, thin} = 16$ $^{-1}$ for the thick disc and thin disc respectively, shows that we underestimate the thin disc values quite substantially and overestimate the value for the thick disc."341 An explanation could be that we do not take any other sscattering mechanisms into account. which would elevate our thin dise dispersion.," An explanation could be that we do not take any other scattering mechanisms into account, which would elevate our thin disc dispersion."342 For the thick dise it could probably be. that there never were any SCs with 10* M. formed in the history of the MW.," For the thick disc it could probably be, that there never were any SCs with $10^{7}$ $_{\odot}$ formed in the history of the MW."343 These high-mass clusters are mainly responsible for the elevated velocity dispersion in the thick disc component of our models., These high-mass clusters are mainly responsible for the elevated velocity dispersion in the thick disc component of our models.344 One interesting result can be observed in our simulation when we focus on the z-distribution of the stars at early stages ες1 Gyr) of the SC evolution., One interesting result can be observed in our simulation when we focus on the $z$ -distribution of the stars at early stages $< 1$ Gyr) of the SC evolution.345" For an example. in Fig 5. we show the z- of the stars. after 270 Myr. for the simulation with an initial mass of 2.2.10"" M. and a SFE of 0.2."," For an example, in Fig \ref{fig:chris2} we show the $z$ -distribution of the stars, after $270$ Myr, for the simulation with an initial mass of $2.2 \times 10^{7}$ $_{\odot}$ and a SFE of $0.2$."346 The shape of the z-distribution looks like a Christmas tree (this effect was seen in our simulations shortly before Christmas 2009)., The shape of the $z$ -distribution looks like a Christmas tree (this effect was seen in our simulations shortly before Christmas 2009).347 The origin of this shape can be explained by considering the dynamical processes during the SC evolution., The origin of this shape can be explained by considering the dynamical processes during the SC evolution.348 The star cluster is orbiting on a circular orbit around the Galactic Centre., The star cluster is orbiting on a circular orbit around the Galactic Centre.349 After gas expulsion. some of the stars acquire higher expanding velocities and form the thick disc.," After gas expulsion, some of the stars acquire higher expanding velocities and form the thick disc."350" We show in Fig 6 in the top left panel the 1 -velocity distribution as a function of the distance between 5 and 9 kpe from the Galactic Centre after 270 Myr νο, about one orbital period).", We show in Fig \ref{fig:fish-1} in the top left panel the $W$ -velocity distribution as a function of the distance between $8$ and $9$ kpc from the Galactic Centre after $270$ Myr (i.e. about one orbital period).351 The velocities show discrete layers., The velocities show discrete layers.352 A similar behaviour is also described in Küpperetal.(2010). where the tidal tails of SCs show density enhancements with regular spacings stemming, A similar behaviour is also described in \citet{kuepper10} where the tidal tails of SCs show density enhancements with regular spacings stemming353Canuna-ray. bursts (GRBs)MER are one of. the brightest.‘ phenomena in. the ⇁⊀⊽∙Universe.,Gamma-ray bursts (GRBs) are one of the brightest phenomena in the Universe.354 Although the majority of ⊲∙⋠⊽⋅⊲≽≻GRBs were. detected at energies-- ranging.: from⋅∙ hundreds.. of LVkeV to several MeV.⊳ they⊽ were. also observed∙⊽ at much higher:. energies. up to tens of ⋅⊲⊳GeV. (Hurleyetal1994).," Although the majority of GRBs were detected at energies ranging from hundreds of keV to several MeV, they were also observed at much higher energies up to tens of GeV \citep{Hurley1994}."355".. opPhe advent of the Fermi Large Area Telescope (LAT) with its aünprecedented sensitivity (Atwoodetal.2009:Band2009) has great . increasedcapability"" to study the fromM IxCiltDs⋅ (Omocdei: ouret.al.2009)..databaseo"," The advent of the Fermi Large Area Telescope (LAT) with its unprecedented sensitivity \citep{Atwood2009,Band2009} has greatly increased our capability to study the high-energy emission from GRBs \citep{Omodei2009}."356es ‘The high-energy (LIE) emission was detected both in the and afterglow phases⋅ of GRBs., The high-energy (HE) emission was detected both in the prompt and afterglow phases of GRBs.357 ⊲However. its originos /2009).. still unclear: it could be produced in the internal/external the via leptonic or hadronicthe Povntinemechanisms. Duxor in the Observations of dissipation of (eg. Meszaros &OROO16C.Rees(20112).," However, its origin is still unclear: it could be produced in the internal/external shocks via leptonic or hadronic mechanisms, or in the process of dissipation of the Poynting flux (e.g., \cite{Meszaros1994,Waxman1997,Bahcall2000,Zhang2001,Dermer2004,Fan2008,Panaitescu2008,Zhang2009,Ghisellini2010,Kumar2010,Razzaque2010,Zhang2011, Meszaros2011}) )."358 Detailed information on the high-energy seconds prompt and afterglow emission could. shed. light. on GRB the onset of GRB and its. immediate. interaction with surrounding interstellar meciunm., Detailed information on the high-energy $\gamma$ -ray prompt and afterglow emission could shed light on the onset of GRB and its immediate interaction with surrounding interstellar medium.359 GRB11 observationssyed inH the LE band are. one of∙ the kev science. topics: of miFermi-LAT., GRB observations in the HE band are one of the key science topics of Fermi-LAT.360. Avr.The Fermi. unGamma-ltay 3urst Monitor. (CDM). can initiateDl. autonomous slew of⋅ the spacecraft. tto provide'ovide the best conditionscitions for fora dedicateddedicate CRB observation., The Fermi Gamma-Ray Burst Monitor (GBM) can initiate autonomous slew of the spacecraft to provide the best conditions for a dedicated GRB observation.361 Also. knowledge of time and position of a GRD (either provided by GBM or by other observations) makesakes it feasibl feasibleto (searchsearch for GRBs GRBsin thethe. LAT.LAL photonotc where⋅ they would manifest themselvesas as spaspatially energyHUS and temporally compact clusters of photons (Dandet," Also, knowledge of time and position of a GRB (either provided by GBM or by other observations) makes it feasible to search for GRBs in the LAT photon database where they would manifest themselves as spatially and temporally compact clusters of photons \citep{Band2009}."362 There :has been 20 LAT cletections.. of. GRBs. as of prompt time of⋅↔ writing .(Feb 2011)., There has been 20 LAT detections of GRBs as of the time of writing (Feb 2011).363is shocks of several very bright bursts (CRB process GRB 090510. GRB 090902B. GRB 090926A) Rees more than 100 photons with energies in excess Zhang 100 MeV. allowed one to study the spectral properties of & LIE emission and even their temporal evolution CXbdo Ghisellinial.2009c5.Ackermannet2010.2011).," Observations of several very bright bursts (GRB 080916C, GRB 090510, GRB 090902B, GRB 090926A) containing more than 100 photons with energies in excess of $100$ MeV allowed one to study the spectral properties of the HE emission and even their temporal evolution \citep{Abdo2009a,Abdo2009b,Ackermann2010,Ackermann2011}."364. 1n all these Razzaque the LIE emission demonstrated. a delav (of several & for long bursts. tenth of a second for the short burst 5-rav 090510) with respect to the prompt emission in tho sub-AleV energy range.," In all these cases the HE emission demonstrated a delay (of several seconds for long bursts, tenth of a second for the short burst GRB 090510) with respect to the prompt emission in the sub-MeV energy range."365 The LE emission also lasted much longer., The HE emission also lasted much longer.366" The observations indicated: a significant deviation from the so-called. ""Band. function (Bandetal. 1993).. namely. the presence of a hard power-law component that dominates at highenergies (e.g. (Abdoetal. 2009a)))."," The observations indicated a significant deviation from the so-called 'Band function' \citep{Band1993}, , namely, the presence of a hard power-law component that dominates at highenergies (e.g., \citep{Abdo2009b}) )."367 Observations of the short GRB 000510, Observations of the short GRB 090510368 Observations of the short GRB 000510., Observations of the short GRB 090510369"The configurations of the magnetic fields show that the advection of magnetic fields is very efficient if Pm~| is adopted, unlike the thin disk cases (Lubowetal.1994)..","The configurations of the magnetic fields show that the advection of magnetic fields is very efficient if ${\cal P}_{\rm m}\sim 1$ is adopted, unlike the thin disk cases \citep{1994MNRAS.267..235L}."370 This is because the radial velocity of ADAFs is much higher than that of the thin disk., This is because the radial velocity of ADAFs is much higher than that of the thin disk.371 Such configurations may help launching outflows/jets from ADAFs., Such configurations may help launching outflows/jets from ADAFs.372" For simplicity, we have not considered magnetically driven outflows from ADAFs in this work."," For simplicity, we have not considered magnetically driven outflows from ADAFs in this work."373" The radial velocity of the ADAF will increase if the angular momentum carried away by the outflows is properly taken into account, which will enhance the advection of the magnetic fields in the ADAF."," The radial velocity of the ADAF will increase if the angular momentum carried away by the outflows is properly taken into account, which will enhance the advection of the magnetic fields in the ADAF."374have a dust temperature profile that. drops from 12Ix at the edge of this pre-stellar core down to a temperature of 7Ix at the centre.,"have a dust temperature profile that drops from $\sim 12~{\rm K}$ at the edge of this pre-stellar core down to a temperature of $\sim3757~{\rm K}$ at the centre."376 The combination of the higher gas density and the lower dust. temperature is the likely reason for a substantial freeze-out. of CO in the central regions of LI6891I3. An alternative explanation for these data is that the CO lies along the line of sight and is not actually part of the LIGSOB core., The combination of the higher gas density and the lower dust temperature is the likely reason for a substantial freeze-out of CO in the central regions of L1689B. An alternative explanation for these data is that the CO lies along the line of sight and is not actually part of the L1689B core.377 This can be ruled out since this would require the core to be completely depleted of CO in order that the column density remains roughly constant across the lace of it., This can be ruled out since this would require the core to be completely depleted of CO in order that the column density remains roughly constant across the face of it.378 Jessop&Wared-Phompson(2001) compared. their dedata with earlier. min continuum data of Xndréetal. (1996)., \citet{jessop&wardthompson01} compared their data with earlier mm continuum data of \citet{andre.et.al96}.379. Using a model for LIGSOB with power law profiles for the density. temperature. anc CO abundance. they explored the parameter space that would be consistent with both datasets.," Using a model for L1689B with power law profiles for the density, temperature and CO abundance, they explored the parameter space that would be consistent with both datasets."380 Phey conclude that the CO could be depleted bv up to 95% but they did not have an optical depth measurement for their CO data., They conclude that the CO could be depleted by up to $95\%$ but they did not have an optical depth measurement for their ${\rm C^{18}O}$ data.381 In contrast. our results indicate that the ratio of the column densities towards the centre is only around a [actor of 3 or so.," In contrast, our results indicate that the ratio of the column densities towards the centre is only around a factor of 3 or so."382 Of course. the CO is unlikely to be uniformly. depleted: ancl this factor of 3 represents the column depletion.," Of course, the CO is unlikely to be uniformly depleted and this factor of 3 represents the column depletion."383 The local depletion is likely to be much higher., The local depletion is likely to be much higher.384 This can be investigated. by plotting the column densities that result (rom mocels where the radial density profile. abundance ancl temperature. are all varied.," This can be investigated by plotting the column densities that result from models where the radial density profile, abundance and temperature are all varied."385 To illustrate this in à very simple case. we use a constant temperature model anc vary the abundance such that interior to a certain radius Lycos. the CO is depleted by a [actor 0.95 due to freeze-out.," To illustrate this in a very simple case, we use a constant temperature model and vary the abundance such that interior to a certain radius $R_{\rm freeze}$ the CO is depleted by a factor 0.95 due to freeze-out."386" Instead of using the Bonner-Ebert density cistributions which are inconvenient to work with. we use the parameters of the fit of Evansetal.(2001) in an equivalent. Plummoer-like sphere which has a density profile of the Form (see. Whitworth&Ward-Thompson2001:Whitworth&Bate 2001)) where po,=LO.10""em is the central density and Ay=750AU is an inner radius within which the density is approximately po."," Instead of using the Bonner-Ebert density distributions which are inconvenient to work with, we use the parameters of the fit of \citet{evans.et.al01} in an equivalent Plummer-like sphere which has a density profile of the form (see, \citealt{whitworth&wardthompson01,whitworth&bate02}) ) where $\rho_0=1.0\times 10^6~{\rm cm^{-3}}$ is the central density and $R_0=750~{\rm AU}$ is an inner radius within which the density is approximately $\rho_0$."387 For the purposes here. the difference between the Bonner-Ebert and. Plummer density distributions are of little consequence - both approximate por7 in the envelope and are roughly constant close to the centre.," For the purposes here, the difference between the Bonner-Ebert and Plummer density distributions are of little consequence - both approximate $\rho\sim r^{-2}$ in the envelope and are roughly constant close to the centre."388 Figure 5. shows the column density that results from this simple Plummer sphere model for a range of values Of Fosse with the top curve showing the underlying density distribution., Figure \ref{hole} shows the column density that results from this simple Plummer sphere model for a range of values of $R_{\rm freeze}$ with the top curve showing the underlying density distribution.389 Comparison with Figure 4. shows clearly that values Of su. Of about 40%=5000AU would. be able to reproduce the approximately flat CO abundance with a column depletion factor of about the correct degree.," Comparison with Figure \ref{columns}390 shows clearly that values of $R_{\rm freeze}$ of about $40\arcsec\equiv 5000~{\rm AU}$ would be able to reproduce the approximately flat CO abundance with a column depletion factor of about the correct degree."391 One could of course investigate the parameter space fully and include realistic variations in the freeze-out and temperature but we defer that to a later paper., One could of course investigate the parameter space fully and include realistic variations in the freeze-out and temperature but we defer that to a later paper.392 The close fit of the [lines by a model with only turbulent velocities (1.6. with no infall or outllow) indicates that the CO emitting material is static., The close fit of the lines by a model with only turbulent velocities (i.e. with no infall or outflow) indicates that the CO emitting material is static.393 In contrast. the ICO4=3 δα οj=352 data of CGregersen.&Evans(2000)9 show double peakect line profiles indicating both that the is sellabsorbed and that the gas is undergoing bulk gas motions.," In contrast, the ${\rm394HCO^+}~J=3\rightarrow 2$ and ${\rm H^{13}CO^+}~J=3\rightarrow 2$ data of \citet{gregersen&evans00} show double peaked line profiles indicating both that the $^+$ is self-absorbed and that the gas is undergoing bulk gas motions."395— These emissions are likely to originate from within the denser regions of the core., These emissions are likely to originate from within the denser regions of the core.396 The llatness of the abundance points across the core measured using ((Fisure 4)) indicates that the CO is depleted. from.= the central regions., The flatness of the abundance points across the core measured using (Figure \ref{columns}) ) indicates that the CO is depleted from the central regions.397 Thus one can picture that LI689D. is composed. of a dynamical dense core which is depleted. in CO but in which is present., Thus one can picture that L1689B is composed of a dynamical dense core which is depleted in CO but in which $^+$ is present.398 This is surrounded. by a quiescent surrounding region where both CO and are present., This is surrounded by a quiescent surrounding region where both CO and $^+$ are present.399 Rare isotopes of CO have been observed towards the pre-stellar core LI689D. By using the hyperfine structure of CHOF=2»1 the transition is confirmed to be optically thin and emitted by quicscent gas.," Rare isotopes of CO have been observed towards the pre-stellar core L1689B. By using the hyperfine structure of ${\rm400C^{17}O}~J=2\rightarrow 1$ the transition is confirmed to be optically thin and emitted by quiescent gas."401 “Phis allows an estimate of the total column density as traced by this molecule to be mace., This allows an estimate of the total column density as traced by this molecule to be made.402 A comparison of this value with that inferred. by SCUBA dust. emission measurements reveal the CO to be depleted in the central regions of this object., A comparison of this value with that inferred by SCUBA dust emission measurements reveal the CO to be depleted in the central regions of this object.403 The magnitude and extent of the depletion is estimated by comparison with a simple mocel of a pre-stellar core with an inner depleted region., The magnitude and extent of the depletion is estimated by comparison with a simple model of a pre-stellar core with an inner depleted region.404 We estimate that within 5000 AU of the centre of L1689D. around of the CO has frozen onto grains.," We estimate that within 5000 AU of the centre of L1689B, around of the CO has frozen onto grains."405 The dust temperature at this radius is 2LOW (from figure 5 of Evansetal. 20013)., The dust temperature at this radius is $\simeq 10~{\rm K}$ (from figure 5 of \citealt{evans.et.al01}) ).406 The sublimation temperature of CO is c20[X so potentially freeze-out could. occur in the outer regions of the cloud., The sublimation temperature of CO is $\simeq 20~{\rm K}$ so potentially freeze-out could occur in the outer regions of the cloud.407 In practice of course. the timescale for [reeze-out in the low density outer regions is long and out will occur only in those portions of the cloud with a high local density and low dust temperature.," In practice of course, the timescale for freeze-out in the low density outer regions is long and freeze-out will occur only in those portions of the cloud with a high local density and low dust temperature."408" For the physical parameters of LIGSOB. the rate of freeze-out of CO is given by (Rawlingsctal.1992) πμ... ""πμ where ng and neo are the hydrogen nucleon and. CO densities respectively. noo=28amu is the molecular mass of CO. d. is the ratio of the number density of grains to CO molecules. e is the grain radius. Cis a factor which accounts for electrostatic ellects. Sco=1 is the assumed sticking co-ellicient for CO."," For the physical parameters of L1689B, the rate of freeze-out of CO is given by \citep{rawlings.et.al92}409 10^4 a^2 C where $n_{\rm H}$ and $n_{\rm CO}$ are the hydrogen nucleon and CO densities respectively, $m_{\rm CO}=28~{\rm amu}$ is the molecular mass of CO, $d_{\rm g}$ is the ratio of the number density of grains to CO molecules, $a$ is the grain radius, $C$ is a factor which accounts for electrostatic effects, $S_{\rm CO}=1$ is the assumed sticking co-efficient for CO."410 Using <dae?c2.210-em- and T—LOW vields Using a value of ny=28Wem for LIGSOB (Bacmannetal.2000) we find that in the absence of other CO formation and destruction mechanisms. depletion ofCO is achieved after ~43.400 vears.," Using $<d_{\rm g} a^2>\simeq 2.2\times41110^{-22}~{\rm cm^{-2}}$ and $T\sim 10~{\rm K}$ yields Using a value of $n_{\rm H}=2.8\times10^5$ $^{-3}$ for L1689B \citep{bacmann.et.al00} we find that in the absence of other CO formation and destruction mechanisms, depletion ofCO is achieved after $\sim$ 43,400 years."412 This is approximately half of the nominal free-fall timescale (97.400. vears) for the value of ng quoted above. but we should be aware that the," This is approximately half of the nominal free-fall timescale (97,400 years) for the value of $n_{\rm H}$ quoted above, but we should be aware that the"413Two fundamental parameters of our moclel are the inclination angle 0 of the external clipole field relative to the rotation axis. and 3. the initial ratio of gas to magnetic pressure at the surface.,"Two fundamental parameters of our model are the inclination angle $\theta$ of the external dipole field relative to the rotation axis, and $\beta$, the initial ratio of gas to magnetic pressure at the surface."414 The simulatious start with the interior magnetic field confined in vertical vortices for r«Roy. but with a still unperturbed inclined dipole field for r>1.3Rosy.," The simulations start with the interior magnetic field confined in vertical vortices for $r<R_{\rm QS}$, but with a still unperturbed inclined dipole field for $r > 1.3\,R_{\rm QS}$ ."415 The transition region Roy<r«1.3Ros is filled by a potential field aud is bounced by two current layers (see Fig. 1)).," The transition region $R_{\rm QS}<r<1.3\,R_{\rm QS}$ is filled by a potential field and is bounded by two current layers (see Fig. \ref{fig:init_state}) )."416 The whole coufiguratiou was Chosen by minimizine the total magnetic energy (see Appenclix ??))., The whole configuration was chosen by minimizing the total magnetic energy (see Appendix \ref{sec:FieldConf}) ).417 We solve the following set of non-iceal maguetohydrodyuamiuecal (MHD) equations. using the (seee.g.Dobleretal.," We solve the following set of non-ideal magnetohydrodynamical (MHD) equations, using the \citep[see \eg][]{dob:convstar}."4182006) Here o. @ A. and s are density. velocity. maguetic vector potential auc specific eutropy. respectively.," Here $\varrho$, $\vec{u}$, $\vec{A}$, and $s$ are density, velocity, magnetic vector potential, and specific entropy, respectively."419 Parameters aud functions kept coustant were cg. Pa. 7 and yy corresponding {ο specific heat. gravity potential. kinematic viscosity aud maguetic diffusivity respectively.," Parameters and functions kept constant were $c_p$, $\Phi_{\rm grav}$, $\nu$ and $\eta$ corresponding to specific heat, gravity potential, kinematic viscosity and magnetic diffusivity respectively."420 The remalniug variables e. J. B.S and T represent the sound speed. electric current deusity. iiagnetic flux density. traceless rate-ol-strain teusor. aid temperature. respectively.," The remaining variables $c_s$, $\vec{j}$, $\vec{B}$, $\mathsf{S}$ and $T$ represent the sound speed, electric current density, magnetic flux density, traceless rate-of-strain tensor, and temperature, respectively."421 The leugth scale is in units of the radius of the quark star. Roy. density in uuits of the stars surface density. po. and time is in units of the spin-period. 1/0.," The length scale is in units of the radius of the quark star, $R_{\rm QS}$, density in units of the star's surface density, $\rho_0$, and time is in units of the spin-period, $1/\Omega$."422 This implies that velocities are in units of ZéosQ. while the magnetic field is in units of. /poJtosQ.," This implies that velocities are in units of $R_{\rm QS}\Omega$, while the magnetic field is in units of $\sqrt{\rho_0} R_{\rm QS}\Omega$."423" The strength of the (dipole) magnetic. field⋅ at the surface⋅ can be estimated⋅ for⋅ a given⋅ 7j as. Bg9=;/Saly/3o. where IH, is. the pressure at the surface of the star."," The strength of the (dipole) magnetic field at the surface can be estimated for a given $\beta_0$ as, $B_0^2 = 8\pi P_0/\beta_0$, where $P_0$ is the pressure at the surface of the star."424 Using livedrostatie balance aud tlie perfect gas law tliis becomes. Iu the equation above. pg is the average clensityof the gas close to the surface of the star.," Using hydrostatic balance and the perfect gas law this becomes, In the equation above, $\rho_0$ is the average densityof the gas close to the surface of the star."425 This corona Is supplied by fall-back material following the formation of the QS (Ixeràuuneu. Ouyed.," This corona is supplied by fall-back material following the formation of the QS (Kerännen, Ouyed,"426where the operator L; represeuts the updating of « by including the fiux iu the / direction aud the operator G represents the eravitational acceleration of the fluid.,where the operator $L_i$ represents the updating of $u$ by including the flux in the $i$ direction and the operator $G$ represents the gravitational acceleration of the fluid.427 During the gravitational step. the fux terms in the Euler equatious are ignored.," During the gravitational step, the flux terms in the Euler equations are ignored."428 The density distribution does not change aud only the fluid ποιοτα aud total enerev density are updated., The density distribution does not change and only the fluid momenta and total energy density are updated.429 The stellar density in the cores of elobular aud open clusters is high euough for stellar collisious to take place with significant frequency. (ills&Day1976)., The stellar density in the cores of globular and open clusters is high enough for stellar collisions to take place with significant frequency \citep{hil76}.430. Chiureut observations and simulations sueecst that the iiereer of two maim sequence stars produces a blue strageler (Sillsetal.1997:Saudquist.Bolte.&IIern-quist. 1997).," Current observations and simulations suggest that the merger of two main sequence stars produces a blue straggler \citep{sil97,sbh97}."431. The blue stragelers are out-lwing main sequence stars which lie bevoud the main sequence turnoff in the colour-inagnuitude diagram (CMD) of a star cluster., The blue stragglers are out-lying main sequence stars which lie beyond the main sequence turnoff in the colour-magnitude diagram (CMD) of a star cluster.432 The blue stragelersoo are more mniassive. brighter. aud bluer than the turnoff stars.," The blue stragglers are more massive, brighter, and bluer than the turnoff stars."433 Since more luassive stars evolve faster thin lower mass stars aud are not expected to lie bevoud the tumnoff. this suggests that blue stragelersCoco» nist have formed more receuthly.," Since more massive stars evolve faster than lower mass stars and are not expected to lie beyond the turnoff, this suggests that blue stragglers must have formed more recently."434 Iu principle the merecr of two main sequence stars can produce a voung renmant star provided that sjeuificaut mixing occurs in the process., In principle the merger of two main sequence stars can produce a young remnant star provided that significant mixing occurs in the process.435 The mixing xoduces a higher livdrogen fraction in the core of the remnant than that of the parent stars which have already burut most of the hwdroseen to helium in heir cores., The mixing produces a higher hydrogen fraction in the core of the remnant than that of the parent stars which have already burnt most of the hydrogen to helium in their cores.436 Benz&Ill(1987) used. low resolution SPII nulatious with ~10° particles to simulate he mereius of vy=3/2 polvtropes aud found that hey fully iuixed., \citet{bh87} used low resolution SPH simulations with $\sim10^3$ particles to simulate the merging of $n=3/2$ polytropes and found that they fully mixed.437 However. medimuu resolution SPI simulations with ~10! particles of »=3/2 or 0=3 xolvtropes showed only weals musing (Lombardi.Ra- 1997).," However, medium resolution SPH simulations with $\sim10^4$ particles of $n=3/2$ or $n=3$ polytropes showed only weak mixing \citep{lrs96,sbh97}."438. It is worth noting hat à)=3/2 polvtropes are nore representative of low mass main sequence stars with large convective envelopes while η=3 polvtropes rescluble main sequence stars near the turnoff which rave little mass in thei couvective cuvelopes., It is worth noting that $n=3/2$ polytropes are more representative of low mass main sequence stars with large convective envelopes while $n=3$ polytropes resemble main sequence stars near the turnoff which have little mass in their convective envelopes.439 Tigh resolution SPITI simulations involving ~10° ouwticles have now been applied to simulating stellar collisions (Sillseta.2002)., High resolution SPH simulations involving $\sim10^5-10^6$ particles have now been applied to simulating stellar collisions \citep{sil02}.440. The mereieg stays process ds inostly subsouic and strong shocks are no expected., The merging stars process is mostly subsonic and strong shocks are not expected.441 Tn the absence of shocks. SPII particles will follow flow lues of constant eutropv due to the Laerangian nature of the method.," In the absence of shocks, SPH particles will follow flow lines of constant entropy due to the Lagrangian nature of the method."442 As a result. the particles may experience secLmueutation.," As a result, the particles may experience sedimentation."443 Iu addition. the mixing cau also depend ou the adopted smoothing leneth aud the form of artificial viscosity.," In addition, the mixing can also depend on the adopted smoothing length and the form of artificial viscosity."444"For a SPII finid. the Revuolds μπαυο is of order (À7Ας where AN, is the total uunuber of particles aud IN, is the ummber of particles over which the sinoothing is donc.","For a SPH fluid, the Reynolds number is of order $(N_p/N_s)^{1/3}$, where $N_p$ is the total number of particles and $N_s$ is the number of particles over which the smoothing is done."445" For V,,~10"" aud No-MAN1075. the Revuolds umuuber is. ~.10."," For $N_p\sim10^5$ and $N_s\sim10^2$, the Reynolds number is $\sim10$."446 However. a fluid with a low Revnolds number will tend to experience laminar flow.," However, a fluid with a low Reynolds number will tend to experience laminar flow."447 ence. SPIT may under mix.," Hence, SPH may under mix."448 It is a worthwhile exercise to model the mereine process using Eulerian lydrodvuamical simulations., It is a worthwhile exercise to model the merging process using Eulerian hydrodynamical simulations.449 The differences between Eulerian aud Lagrangian approaches nav lead to very different results on nunxine., The differences between Eulerian and Lagrangian approaches may lead to very different results on mixing.450 As of preseut. no such work has been reported in the literature.," As of present, no such work has been reported in the literature."451 We consider the offtaxis collision of two main sequence stars with Af=OSAL. and R=0.955 FR... which are modeled using 5=—3? polvtropes," We consider the off-axis collision of two main sequence stars with $M=0.8\, M_\sun$ and $R=0.955\,R_\sun$ , which are modeled using $n=3$ polytropes."452 A polvtrope with polytropic iudex 6» has equilibrium, A polytrope with polytropic index $n$ has equilibrium453ils pair production multiplicity is substantially lower than. previous model estimates. for exanple instead of closing the outergap in terms of photon-photon creation process outer eap closed by magnetic pair creation is possible (οἱ.,"its pair production multiplicity is substantially lower than previous model estimates, for example instead of closing the outergap in terms of photon-photon creation process outer gap closed by magnetic pair creation is possible (cf."454 Takata et al., Takata et al.455 2010). which gives less ouleoing pairs.," 2010), which gives less outgoing pairs."456 Llowever. if this is the case. CR contribution can not avoided.," However, if this is the case, CR contribution can not avoided."457 The observed eanmma-ravs in Fermi energy range should be a mixture of CHR. and IC processes., The observed gamma-rays in Fermi energy range should be a mixture of CR and IC processes.458" In fitting data point of view. Ly, is the normalization factor."," In fitting data point of view, $\eta L_{sd}$ is the normalization factor."459" In our model it can be estimated by relating the observed eamma-ray. power L, (to the theoretical IC power. i.e. Napyplet. where Najsp is the total number of MSPs in the globular cluster ancl Lis is the part of spin-down power carriecl away by (he pairs given by Eq."," In our model it can be estimated by relating the observed gamma-ray power $L_{\gamma}$ to the theoretical IC power, i.e. $N_{MSP}L_{e^{\pm}}$, where $N_{MSP}$ is the total number of MSPs in the globular cluster and $L_{e^{\pm}}$ is the part of spin-down power carried away by the pairs given by Eq."460 4., 4.461 Lo we assume that each pulsar has similar spin-down power. e.g. L4;~2. once fs is fixed (hen we can use the above conservation to estimate the total number of MSPs in the globular cluster.," If we assume that each pulsar has similar spin-down power, e.g. $L_{34}\sim 2$, once $f_{e^{\pm}}$ is fixed then we can use the above conservation to estimate the total number of MSPs in the globular cluster."462 Let's, Let's463the statistics of deOliveira-Costaetal.(2001)... aud compute 1) the probability of finding a lower quadrupole moment than the observed one. given the best-fit spectrum. 2) the probability of finding such a strong aliguiment between the quadrupole aud octopole iud 3) the probability of finding such planar iultipoles as secu in the maps.,"the statistics of \citet{de464Oliveira-Costa:2004}, and compute 1) the probability of finding a lower quadrupole moment than the observed one, given the best-fit spectrum, 2) the probability of finding such a strong alignment between the quadrupole and octopole and 3) the probability of finding such planar multipoles as seen in the maps."465 Here we briefly define the various statistics. and refer the interested reader to deOliveira-Costaetal.(2001). for details ou how cach quautity actually is computed.," Here we briefly define the various statistics, and refer the interested reader to \citet{de Oliveira-Costa:2004} for details on how each quantity actually is computed."466 The first statistic is simply the multipole amplitude which is defined in terms of a spherical larmouics ATP.expansion of tle map. The unultipole amplitude is then defined as The next statistic is based on the possibility to define a preferred axis. ny. for cach multipole. namely that axis which maximizes the aneulay moimoeutunm dispersion. The alieuiieut between two modes is then uicasured by takine the dot product of the two preferred directions.," The first statistic is simply the multipole amplitude $\delta T_{l}^2$, which is defined in terms of a spherical harmonics expansion of the map, The multipole amplitude is then defined as The next statistic is based on the possibility to define a preferred axis, $\hat{\mathbf{n}}_l$, for each multipole, namely that axis which maximizes the angular momentum dispersion, The alignment between two modes is then measured by taking the dot product of the two preferred directions."467 The computation of this quantitv is carried out bv computing the spherical harmonic cocfiicicuts iu some coordinate system. aud then rotating these im harmonic space.," The computation of this quantity is carried out by computing the spherical harmonic coefficients in some coordinate system, and then rotating these in harmonic space."468 Since the harmonic space rotation matrices are simple to compute. the complete maximization xocedure becomes relatively inexpensive even for a high-resolution map with several million pixels.," Since the harmonic space rotation matrices are simple to compute, the complete maximization procedure becomes relatively inexpensive even for a high-resolution map with several million pixels."469 The details ou conrputiug these rotation matrices are described by deOliveira-Costaotal.(2001) chooses., The details on computing these rotation matrices are described by \citet{de Oliveira-Costa:2004} .470.. The third quantity we cousider is the degree of dlanarity of a eiven mode., The third quantity we consider is the degree of planarity of a given mode.471 Two different statistics for, Two different statistics for472Numerous astronomical observations suggest that large magnetic fields are associated. with neutron stars (NSs).,Numerous astronomical observations suggest that large magnetic fields are associated with neutron stars (NSs).473 Indeed. evidence for the existence of binary NSs is obtained from binary pulsars. in which one or both NSs are seen to have a large magnetic Ποια.," Indeed, evidence for the existence of binary NSs is obtained from binary pulsars, in which one or both NSs are seen to have a large magnetic field."474 In General Relativity such binary systems cannot be stationary because they emit eravitational waves (Gi\Ws) which extract energy. ancl angular momentum from the binary. inducing it to inspiral and merge.," In General Relativity such binary systems cannot be stationary because they emit gravitational waves (GWs) which extract energy and angular momentum from the binary, inducing it to inspiral and merge."475 During the final stages of the inspiral the CAV emission is expected to be strong enough to be relevant for he detectors now operative at design sensitivities and it »omises to provide important information on the equation of state (EOS) regulating the NS matter (Readοἱal. 2009)., During the final stages of the inspiral the GW emission is expected to be strong enough to be relevant for the detectors now operative at design sensitivities and it promises to provide important information on the equation of state (EOS) regulating the NS matter \citep{Read2009}.476. In aclelition to their importance as sources of GWs. rowever. the merger of binary NSs is likely to. provide important information on the physics of short. eamma-ravy uusts (GRBs).," In addition to their importance as sources of GWs, however, the merger of binary NSs is likely to provide important information on the physics of short gamma-ray bursts (GRBs)."477 The coalescence of the two NSs. in fact. gives rise. either promptly or after some interval. to à system composed of a torus orbiting around a rapidly rotating black ide (BLD) (Baiottietal.2008:Yamamotoοἱ2008)..," The coalescence of the two NSs, in fact, gives rise, either promptly or after some interval, to a system composed of a torus orbiting around a rapidly rotating black hole (BH) \citep{Baiotti08,Yamamoto2008}."478" The complex plasma physics accompanying this event is probably behind the ""engine"" powering GRBs (Piran2004:Aleszaros 2006)."," The complex plasma physics accompanying this event is probably behind the “engine” powering GRBs \citep{Piran:2004ba,Meszaros:2006rc}."479 There is little doubt. therefore. about the importance of assessing the role plaved by magnetic fields in the inspira and merger of binary NSs.," There is little doubt, therefore, about the importance of assessing the role played by magnetic fields in the inspiral and merger of binary NSs."480 Yet. determining this accurately is à remarkably difficult task requiring the solution of the Einstein equations together with those of general-relativistic maenctohycrodvnamies (GRATID).," Yet, determining this accurately is a remarkably difficult task requiring the solution of the Einstein equations together with those of general-relativistic magnetohydrodynamics (GRMHD)."481" So far. only two GIMLBD simulations have been reported. (Ancersonοal.2008:Liuet 2008).. reaching cdillerent conclusions about the importance of very strong magnetic fields (D—1k"" 10€)."," So far, only two GRMHD simulations have been reported \citep{Anderson2008,Etienne08}, reaching different conclusions about the importance of very strong magnetic fields $B\sim 10^{16}-10^{17}\,\G$ )."482 The aim of this Letter is to go. bevom these —qualitative estimates and provide a first quantitative measurement of the influence of magnetic fields on both the inspiral and the merger of magnetized λος., The aim of this Letter is to go beyond these qualitative estimates and provide a first quantitative measurement of the influence of magnetic fields on both the inspiral and the merger of magnetized NSs.483 ὃν considering a large range of magnetic fields. which includes values more realistic than those usec in the works cited above. anc two cdillerent. masses. we find that magnetic Liclds g@encrally grow. the merger. when the turbulent motions triggered. during. the merger. hy the Ixelvin-Helmholtz Gxt) instability. amplify anv initial," By considering a large range of magnetic fields, which includes values more realistic than those used in the works cited above, and two different masses, we find that magnetic fields generally grow the merger, when the turbulent motions triggered during the merger by the Kelvin-Helmholtz (KH) instability amplify any initial"484infrared (ULIR) merecr Arp 220.,infrared (ULIR) merger Arp 220.485 These svstems were chosen because they have been observed in the CO(1-0) molecular line transition with resolutions (full width at half maximo) of 2%2.57 (Yun&Iibbard1999a.Waneetal.1992.Scoville 1997)).," These systems were chosen because they have been observed in the CO(1-0) molecular line transition with resolutions (full width at half maximum) of $2^{\prime\prime}-2.5^{\prime\prime}$ \cite{YH99a,Wang92,Sco97}) )."486 The resulting spatial radial resolution (2060100 pe: Πως 75 kan + +) is sluuilar to the lhydrodvuamical smoothing leneth used in MIT9L (~ 350pe)?.. indicating that the molecular liue observations lave sufficicut resolution to resolve the types of mass concentrations found iu the simulations.," The resulting spatial radial resolution (300–400 pc; $H_o$ = 75 km $^{-1}$ $^{-1}$ ) is similar to the hydrodynamical smoothing length used in MH94 $\sim$ 350, indicating that the molecular line observations have sufficient resolution to resolve the types of mass concentrations found in the simulations."487" The molecular gas surface deusities of cach of these systems are plotted in Figure Ἐν, converted from CO fluxes by adopting a couversion factor of Nyy/loo=3«410%emενας4) (Young&Scoville 1991))."," The molecular gas surface densities of each of these systems are plotted in Figure \ref{fig:radplotA}, converted from CO fluxes by adopting a conversion factor of $N_{H_2} /I_{CO} = 3\times10^{20} \, {\rm488cm^{-2}\,(K \, km \, s^{-1})^{-1}}$ \cite{Young91}) )."489 Detailed studies on each of these systems. which fully discuss their status as late stage merecrs. can be fouud in Schweizer (1996) and ibbard van Corkom (1996) for NCC 3921: Sclaweizer (1982) and Hibboad et al. (," Detailed studies on each of these systems, which fully discuss their status as late stage mergers, can be found in Schweizer (1996) and Hibbard van Gorkom (1996) for NGC 3921; Schweizer (1982) and Hibbard et al. ("4901991) for NCC 7252: and Scoville et al. (,1994) for NGC 7252; and Scoville et al. (4911997) for Arp 220.,1997) for Arp 220.492 For these moderately evolved merger remnants (ages of ~ 0.51 Cr since their tidal tails were launched. ΕΠαναetal.1991.IBbbard&vanCorkom 1996)). the observed eas aud luminosity profiles are used to predict the expected i1unuinositv profile of a 2 Cr old remuaut.," For these moderately evolved merger remnants (ages of $\sim$ 0.5–1 Gyr since their tidal tails were launched, \cite{Hib94,HvG96}) ), the observed gas and luminosity profiles are used to predict the expected luminosity profile of a 2 Gyr old remnant."493 We assume that all of the molecular eas is turued iuto stars at the same radii adopting an exponentially declining starformation ustory.," We assume that all of the molecular gas is turned into stars at the same radii, adopting an exponentially declining starformation history."494 The present bhpunuiuositv profile is allowed to fade due to passive aging effects. and the final Iniinositv profile is the stm of these two populations.," The present luminosity profile is allowed to fade due to passive aging effects, and the final luminosity profile is the sum of these two populations."495 The molecular gas xofiles plotted in Fig., The molecular gas profiles plotted in Fig.496" d. are converted iuto gas lass densities by imultiplviug bv a factor of 1.36 to take iuto consideration the expected coutribution of Ποπα,", \ref{fig:radplotA} are converted into gas mass densities by multiplying by a factor of 1.36 to take into consideration the expected contribution of Helium.497 Optical παλπιοτν profiles have been obtained by Sclaveizer (1982. 1996). Whitinore et al. (," Optical luminosity profiles have been obtained by Schweizer (1982, 1996), Whitmore et al. ("4981993). and IHibbard et al. (,"1993), and Hibbard et al. ("4991991). showing them to be well fitted by an ci! profile over all radii with uo apparent hpuuinositv spikes.,"1994), showing them to be well fitted by an $r^{1/4}$ profile over all radii, with no apparent luminosity spikes."500 The gas surface deusities in AL.pe ?) are converted to optical surface brightuesses(X44;; (9p) by dividing x the stellar mass-to-lieht ratio (M./ L5) expected for a 2 Car old population.," The gas surface densities $\Sigma_{gas}$ in $M_\odot\,{\rm pc}^{-2}$ ) are converted to optical surface brightnesses $\mu_B$ ) by dividing by the stellar mass-to-light ratio $M_*/L_B$ ) expected for a 2 Gyr old population."501 We adopt the stellar imass-o-lieht ratios eiven by de Jong (1995. Table 1 of ch.," We adopt the stellar mass-to-light ratios given by de Jong (1995, Table 1 of ch."502 Lj. which were derived from the population svuthesis uodels of Bruzual Charlot (1993) for an exponentially declining star formation listory. a Salpeter IME. aud Solar netallicity (AL./Lp=0.82ML at 2 Gar).," 4), which were derived from the population synthesis models of Bruzual Charlot (1993) for an exponentially declining star formation history, a Salpeter IMF, and Solar metallicity $M_*/L_B = 0.82 \, M_\odot L_\odot^{-1}$ at 2 Gyr)."503" Noting hat 1 £pe7 corresponds to jp = 27.06 mag > (aclopting Mg.= 15.18). the conversion frou gas surface density to optical surface brightuess is given by ppí(r)=27.06magaresceο2,54log|Xqui)CU Ep)."," Noting that 1 $L_\odot \,{\rm pc}^{-2}$ corresponds to $\mu_B$ = 27.06 mag $^{-2}$ (adopting $M_{B,\odot}=+5.48$ ), the conversion from gas surface density to optical surface brightness is given by $\mu_B(r) =50427.06\,{\rm mag \, arcsec}^{-2} - 2.5\times log[\Sigma_{gas}(r) /505(M_*/L_B) ]$ ."506 The huuinositv profiles of the evolved remnants are estimated. from the observed £-baud. profiles (Scleizer1996.Ihbbardetal. 1991)). allowiug for a fading of FL mag 7 du the Z-baud over the next 2 Cx (Druzual&Charlot1995.Scliwveizer 1996)). aud adding iu the expected contribution of the population formed frou the molecular gas. calculated as above.," The luminosity profiles of the evolved remnants are estimated from the observed $B$ -band profiles \cite{Schwe96a,Hib94}) ), allowing for a fading of +1 mag $^{-2}$ in the $B$ -band over the next 2 Gyr \cite{BC93,Schwe96a}) ), and adding in the expected contribution of the population formed from the molecular gas, calculated as above."507 We emphasize that this should favor the production of a luminous post-merecr »opulation. since it assunes the that none of the molecular eas is lost to stellar winds or SNe aud the adopted IME avors the production of many long-lived low-nass stars.," We emphasize that this should favor the production of a luminous post-merger population, since it assumes the that none of the molecular gas is lost to stellar winds or SNe and the adopted IMF favors the production of many long-lived low-mass stars."508 The results of this exercise are plotted in Figure .., The results of this exercise are plotted in Figure \ref{fig:radplotB}.509 This ot shows that the observed gas deusities in NGC 3921 and NGC 7252. although high. are not high enough to sienificantly affect the present hnuuinositv profiles.," This plot shows that the observed gas densities in NGC 3921 and NGC 7252, although high, are not high enough to significantly affect the present luminosity profiles."510 The motile of NGC 3921 is basically iudistinguishable from an ril profile., The profile of NGC 3921 is basically indistinguishable from an $r^{1/4}$ profile.511 The profile of NGC. 7252 does slow a slight rise at simall radii. but not the clear break predicted ΑΠΟ.," The profile of NGC 7252 does show a slight rise at small radii, but not the clear break predicted by MH94."512 Therefore the resulting luminosity profiles of hese renimeamnts are now and should remain fairly typical of normal elliptical galaxies. and the couclusious of MII91 are not applicable toall mereers of eas-rich salaxies.," Therefore the resulting luminosity profiles of these remnants are now and should remain fairly typical of normal elliptical galaxies, and the conclusions of MH94 are not applicable to mergers of gas-rich galaxies."513 Since both of these systems also obey the EFaber-Jacksou relationship (Lake&Dressler1986)) and NGC 7252 falls upon the fundamental plane defined by normal ellipticals (Ilibbardetal.1991.Tibbard 19953). we couchide that at least some mergers of eas-rich svstenuis can evolve iuto normal elliptical galaxies as far as their optical propertics are concerued.," Since both of these systems also obey the Faber-Jackson relationship \cite{Lake86}) ) and NGC 7252 falls upon the fundamental plane defined by normal ellipticals \cite{Hib94,Hib95}) ), we conclude that at least some mergers of gas-rich systems can evolve into normal elliptical galaxies as far as their optical properties are concerned."514 Since Arp 220 is an extremely dusty object. its optical ininositv profile is poorly suited for a similar analysis.," Since Arp 220 is an extremely dusty object, its optical luminosity profile is poorly suited for a similar analysis."515 Tusteack we use a luninosity profile measured iu the jezr-Anfrared. where the dust obscuration is au order of uaenitude less severe.," Instead, we use a luminosity profile measured in the near-infrared, where the dust obscuration is an order of magnitude less severe."516 Arp 220 was receutlv observed with camera 2 of NICMOS aboard the HIST (Scoville et al., Arp 220 was recently observed with camera 2 of NICMOS aboard the HST (Scoville et al.517 1998). and we use the resulting Aband hwunuinositv xofile. kindly mace available ly N. Scoville.," 1998), and we use the resulting $K-$ band luminosity profile, kindly made available by N. Scoville."518 Since Arp 220 is preseutly undergoing a massive starburst. the facing actor is much less certain thin for the already evolved systems treated above. aud depends scusitively on what action of the current lieht is contributed by receutlv ormed stars.," Since Arp 220 is presently undergoing a massive starburst, the fading factor is much less certain than for the already evolved systems treated above, and depends sensitively on what fraction of the current light is contributed by recently formed stars."519 We adopt a situation biased towards the xoduction of a discrepant Iuninosity profile by assuming hat the entire population was pre-existiug. converting the observed fy baud profile to au evolved B baud profile bx adopting à D...A color of 1. appropriate for a LO Cir old »pulatiou (deJong 1995)).," We adopt a situation biased towards the production of a discrepant luminosity profile by assuming that the entire population was pre-existing, converting the observed $K-$ band profile to an evolved $B-$ band profile by adopting a $B-K$ color of 4, appropriate for a 10 Gyr old population \cite{deJong}) )."520 The contribution due to the »»pulatiou formed from the molecular disk is calculated exactly as before., The contribution due to the population formed from the molecular disk is calculated exactly as before.521 The resulting profile is shown in Figure .., The resulting profile is shown in Figure \ref{fig:radplotB}.522 This figure shows that Arp 220 is predicted to evolve a luminosity profile with a noticeable rise at snall raclii., This figure shows that Arp 220 is predicted to evolve a luminosity profile with a noticeable rise at small radii.523 This is due to the peal in the molecular gas surface density at radii less than 0.5 προ (Fig. 1))., This is due to the peak in the molecular gas surface density at radii less than 0.5 kpc (Fig. \ref{fig:radplotA}) ).524 We conclude that Arp 220 has the potential to evolve a similar feature in its huninosity profile.if indeed all of the current molecular eas is converted iuto stars.," We conclude that Arp 220 has the potential to evolve a similar feature in its luminosity profile, indeed all of the current molecular gas is converted into stars."525 Towever. the expected rise of ~2 mag arcesec? in surface brightness (a factor of —6) is considerably lower than the two orders of maguitude increase predicted by the siuulatious (see Fie.," However, the expected rise of $\sim$ 2 mag $^2$ in surface brightness (a factor of $\sim$ 6) is considerably lower than the two orders of magnitude increase predicted by the simulations (see Fig."526 1 of MIT91)., 1 of MH94).527 From the above exercise. we couclude that neither NCC 3921 nor NGC 7252 are expected to show a siguificaut deviation in them luminosity profiles. aud that the Πιν ΠΠ expected for Arp 220 is considerably lower than the two orders of mmaenitude increase predicted by the simulations of MIT9 1," From the above exercise, we conclude that neither NGC 3921 nor NGC 7252 are expected to show a significant deviation in their luminosity profiles, and that the maximum rise expected for Arp 220 is considerably lower than the two orders of magnitude increase predicted by the simulations of MH94."528.We conchide that the umuerical, We conclude that the numerical529originally proposed by Soltan 1982).,originally proposed by Soltan 1982).530 This high average radiative efficiency. (hough. describes the population of SMDIIs at redshift 2<5. and nothing is known about the radiative efficiency of preealactic quasars in (he early Universe.," This high average radiative efficiency, though, describes the population of SMBHs at redshift $z<5$, and nothing is known about the radiative efficiency of pregalactic quasars in the early Universe."531 The numbers quoted above. in fact. nav suggest that the picture that we have of the low redshift Universe max not apply αἱ earlier (ines.," The numbers quoted above, in fact, may suggest that the picture that we have of the low redshift Universe may not apply at earlier times."532 In (his paper we follow earlier papers in consilering a scenario for the hierarchical assembly of SMDIIS that (races their seeds back to the very [ist generation of stars. in mini-halos above the cosmological Jeans mass collapsing al z20 Irom the high-o peaks of (he primordial density fiel.," In this paper we follow earlier papers in considering a scenario for the hierarchical assembly of SMBHs that traces their seeds back to the very first generation of stars, in mini-halos above the cosmological Jeans mass collapsing at $z\sim 20$ from the $\sigma$ peaks of the primordial density field."533 llowever we introduce (wo new features: (1) we consider more explicitly (he configuration ol the gas from which the accretion occurs: a fat. dense disc of cold gas. likely to form in halos with Tj>10!X and zero metallicity (Oh Iaiman 2002).," However we introduce two new features: (i) we consider more explicitly the configuration of the gas from which the accretion occurs: a fat, dense disc of cold gas, likely to form in halos with $T_{\rm vir}>10^4 {\rm K}$ and zero metallicity (Oh Haiman 2002)."534 And (i) we explore the assumption (hat at early stages. accretion occurs al a “Bondi rate that is higher than the standard rate for €20.1 The evolution of the MDlIIs into the pregalaclic quasars and (heir impact on the re-onization of the universe in (Iiis scenario has been explored recently by Macau 2004.," And (ii) we explore the assumption that at early stages, accretion occurs at a 'Bondi' rate that is higher than the standard rate for $\epsilon\simeq0.1$ The evolution of the MBHs into the pregalactic quasars and their impact on the re-ionization of the universe in this scenario has been explored recently by Madau 2004."535 We here expand ancl deepen the previous studs. focusing in partüeular on the interplay of eas cooling ancl DIL feeding.," We here expand and deepen the previous study, focusing in particular on the interplay of gas cooling and BH feeding."536 In the next section we review the model for assembly of MDIIs in cold dark matter (CDM) cosmogonies., In the next section we review the model for assembly of MBHs in cold dark matter (CDM) cosmogonies.537 We then address the conditions ofthe gas around the MIBIIs in the cores of high redshift mini-halos (833). before and alter the onset of pregalactic quasar activity.," We then address the conditions of the gas around the MBHs in the cores of high redshift mini-halos 3), before and after the onset of pregalactic quasar activity."538 We compute the elobal evolution of the MDBII population aud discuss its implications in 844 and finally summarize our results in 855., We compute the global evolution of the MBH population and discuss its implications in 4 and finally summarize our results in 5.539"55 Unless otherwise stated. all results shown below refer to the currently [avonured. ACDM world model with O4,=0.3. O4=Q.T. h=0.7. ο=0.045. σε=0.93. and n=1."," Unless otherwise stated, all results shown below refer to the currently favoured $\Lambda$ CDM world model with $\Omega_M=0.3$, $\Omega_\Lambda=0.7$, $h=0.7$, $\Omega_b=0.045$, $\sigma_8=0.93$, and $n=1$."540" The main features of a plausible scenario for the hierarchical assembly. growth. and dvnamics of MBIISs in à ΑςΕΟΝ cosmology have been discussed by Volonteri. Haardt. Macau (2003). Volonteri (2004). Macdau (2004). """," The main features of a plausible scenario for the hierarchical assembly, growth, and dynamics of MBHs in a $\Lambda$ CDM cosmology have been discussed by Volonteri, Haardt, Madau (2003), Volonteri (2004), Madau (2004). “"541"Seed"" holes with intermediate masses form as end-product of the very first generation of stars.","Seed"" holes with intermediate masses form as end-product of the very first generation of stars."542 They form in isolation within mini-halos above the cosmological Jeans mass collapsing al zον24 from rare ν-σ peaks of (he primordial density [field (Macau Rees 2001)., They form in isolation within mini-halos above the cosmological Jeans mass collapsing at $z\approx24$ from rare $\nu$ $\sigma$ peaks of the primordial density field (Madau Rees 2001).543 Preealactic seed INIBIIs form within the, Pregalactic seed IMBHs form within the544The limits (hat we place on the radio supernova populations in our target galaxies can be converted (to radio supernova rates. if a few simplilving assumptions are made.,"The limits that we place on the radio supernova populations in our target galaxies can be converted to radio supernova rates, if a few simplifying assumptions are made."545 Conversion of a radio supernova rate to an overall supernova rate is most believable if we can make some estimate of the fraction of supernovae in our galaxies that should produce detectable radio supernovae., Conversion of a radio supernova rate to an overall supernova rate is most believable if we can make some estimate of the fraction of supernovae in our galaxies that should produce detectable radio supernovae.546" The radio emission [rom voung type Il supernovae in ""normal galactic environments (vpically is dominated by (he interaction of the supernova blast wave with (he mass-loss shell from the parent star (Chevalier1932).", The radio emission from young type II supernovae in “normal” galactic environments typically is dominated by the interaction of the supernova blast wave with the mass-loss shell from the parent star \citep{che82}.547. However. in the highest-density regions. such as SSCs in dense molecular clouds. the supernova blast may be dominated by the high-density. interstellar environment (Chevalier&Fransson2001).," However, in the highest-density regions, such as SSCs in dense molecular clouds, the supernova blast may be dominated by the high-density interstellar environment \citep{che01}."548. In such a case. the supernova ejecta will be slowed dramatically. particle acceleration at the shock front will be «quite efficient. ancl radio emission will be both powerlul and long-lived.," In such a case, the supernova ejecta will be slowed dramatically, particle acceleration at the shock front will be quite efficient, and radio emission will be both powerful and long-lived."549 Thus. it appears to be a reasonable inference thatevery (wpe IL supernova in an SSC may produce significant radio emission. Predictions of the supernova rates in our target galaxies mav be derived [rom (heir populations of massive stars and the estimated lifetime of a radio supernova.," Thus, it appears to be a reasonable inference that type II supernova in an SSC may produce significant radio emission, Predictions of the supernova rates in our target galaxies may be derived from their populations of massive stars and the estimated lifetime of a radio supernova."550 The compact starburst regions in our target galaxies must. contain roughly 10 O7-equivalent stars in order to generate the ionizing [hix that energizes the Ha and thermal radio emission., The compact starburst regions in our target galaxies must contain roughly $10^4$ O7-equivalent stars in order to generate the ionizing flux that energizes the $\alpha$ and thermal radio emission.551 In ihe nearby galaxies NGC 253 (Ulvestad&Antonucci1993.1997). and M82 (Ulvestad&Antonueci1993;hronbergetal. 2000).. the radio flux densities of most voung supernova remnants are relatively constant over periods of 815 vears. implying that (μον stay near (heir radio peaks for periods of at least 20.100 vears.," In the nearby galaxies NGC 253 \citep{ulv94,ulv97} and M82 \citep{ulv94,kro00}, the radio flux densities of most young supernova remnants are relatively constant over periods of 8–15 years, implying that they stay near their radio peaks for periods of at least 20–100 years."552 This supports the modeling of (2001).. but is in contrast to Type II radio supernovae in less dense environments. whose flux densities fall off with time / as roughly /.* to 415 (Weileretal.2002).," This supports the modeling of \citet{che01}, but is in contrast to Type II radio supernovae in less dense environments, whose flux densities fall off with time $t$ as roughly $t^{-0.7}$ to $t^{-1.8}$ \citep{wei02}."553. We may use the stellar contents and inferrred radio lifetimes discussed above to estimate (he supernova rates expected in our program galaxies., We may use the stellar contents and inferrred radio lifetimes discussed above to estimate the supernova rates expected in our program galaxies.554 Roughly. (he main-sequence lifetimes of massive voung stus (hal go supernova range [rom ~340 Myr (cf," Roughly, the main-sequence lifetimes of massive young stars that go supernova range from $\sim$ 3–40 Myr (cf."555 Starburst99 models of Leitherer et al., Starburst99 models of Leitherer et al.556 1999). ancl (the supernova rate should peak about 46 Myr alter an instantaneous starburst.," 1999), and the supernova rate should peak about 4–6 Myr after an instantaneous starburst."557" Thus. if we assume Chat the massive stars all Formed in a time of Mr. and that there are ~104 such stars; we should expect a radio supernova rate no higher than ~1 5x107 | at a time about 7x10* vr alter the starburst,"," Thus, if we assume that the massive stars all formed in a time of 1--3 Myr, and that there are $\sim 10^4$ such stars, we should expect a radio supernova rate no higher than $\sim$ $\times 10^{-3}$ $^{-1}$ at a time about $\times 10^6$ yr after the starburst."558 Alternatively. one may take a quasi-instantaneous burst with a total mass of ~LOSM.: the Starburst(99 models (Leithererοἱal.L999) then show a peak supernova rate of slightly less than 10.7 t about G Myr alter the burst.," Alternatively, one may take a quasi-instantaneous burst with a total mass of $\sim 10^6 M_\odot$; the Starburst99 models \citep{lei99} then show a peak supernova rate of slightly less than $10^{-3}$ $^{-1}$ about 6 Myr after the burst."559 Thus. as long as the typical radio supernova stavs above LOS W ! for no more than a lew hundred vears. this is consistent with our lack of detection of anv milliaresecond radio sources in our three target galaxies.," Thus, as long as the typical radio supernova stays above $10^{18}$ W $^{-1}$ for no more than a few hundred years, this is consistent with our lack of detection of any milliarcsecond radio sources in our three target galaxies."560"(PSF) should have a negligible affect on the measured flux at the position of 2MASS 1207 b. However, there does appear to be a slight increase in background at a separations «1"", which is probably the result of 2MASS 1207 A's PSF seeing halo.","(PSF) should have a negligible affect on the measured flux at the position of 2MASS 1207 b. However, there does appear to be a slight increase in background at a separations $<$ 1"", which is probably the result of 2MASS 1207 A's PSF seeing halo."561" We subtract off this small halo contribution by doing a median average of all pixels within a 2.5 pixel radius of the separation of 2MASS 1207 b, and at all position angles except for the aperture that we use for 2MASS 1207 b. We used data from 2010 Mar 31 UT to perform an absolute flux calibration on 2MASS 1207 A using the mid-IR standard HR4888 (?),, which is 16.011 Jy in the Si-2filter]."," We subtract off this small halo contribution by doing a median average of all pixels within a 2.5 pixel radius of the separation of 2MASS 1207 b, and at all position angles except for the aperture that we use for 2MASS 1207 b. We used data from 2010 Mar 31 UT to perform an absolute flux calibration on 2MASS 1207 A using the mid-IR standard HR4888 \citep{1999AJ....117.1864C}, which is 16.011 Jy in the Si-2."562" The best 4-hour period of 2MASS 1207 data (which included 9 of our 15 usable frames) was averaged, and compared to images of HR4888 taken immediately after (HR4888 was found to vary by ~5% over the 3 calibration images)."," The best 4-hour period of 2MASS 1207 data (which included 9 of our 15 usable frames) was averaged, and compared to images of HR4888 taken immediately after (HR4888 was found to vary by $\sim$ over the 3 calibration images)."563" We use a 15 pixel (1.35"") aperture and find that 2MASS 1207 is 0.000343 times as bright as HR4888 in the Si-2 filter, which corresponds to an absolute flux of 5.49 mJy."," We use a 15 pixel (1.35"") aperture and find that 2MASS 1207 is 0.000343 times as bright as HR4888 in the Si-2 filter, which corresponds to an absolute flux of 5.49 mJy."564" We assign an 8 error to this measurement (0.44 mJy), which is conservative given the fidelity of the Si-2 filter and the night's good photometric quality."," We assign an 8 error to this measurement (0.44 mJy), which is conservative given the fidelity of the Si-2 filter and the night's good photometric quality."565" This value is consistent with published values of 5.6 mJy +1 mJy in the Si-2 filter (?),, 5.74 mJy in the slightly bluer Spitzer 8um filter (?),, Spitzer spectroscopy (?),, and roughly consistent with our own measurements using the 2008 Mar 29 UT data (4.3 mJy +0.5 mJy), which suggests that 2MASS 1207 A is probably not wildly variable at these wavelengths, despite the presence of a near edge-on disk."," This value is consistent with published values of 5.6 mJy $\pm$ 1 mJy in the Si-2 filter \citep{2004AA...427..245S}, 5.74 mJy in the slightly bluer Spitzer $\micron$ filter \citep{2006ApJ...639L..79R}, Spitzer spectroscopy \citep{2008ApJ...676L.143M}, and roughly consistent with our own measurements using the 2008 Mar 29 UT data (4.3 mJy $\pm$ 0.5 mJy), which suggests that 2MASS 1207 A is probably not wildly variable at these wavelengths, despite the presence of a near edge-on disk."566" Using 2MASS 1207 A as a PSF, we do a best-fit to determine how much flux is at the near-IR determined position of 2MASS 1207 b (sep=0.773”,PA=125.37°;?)."," Using 2MASS 1207 A as a PSF, we do a best-fit to determine how much flux is at the near-IR determined position of 2MASS 1207 b \citep[sep=0.773"", PA=125.37$^{\circ}$."567" We assume Gaussian error bars (equal on all pixels for the background limited case), and fit within a 5 pixel diameter aperture (see Figure |l) to avoid contamination by 2MASS 1207 A. We do a Monte Carlo evaluation of the background fluctuations by doing a similar aperture measurements at 10000 randomly chosen position angles and separations > 1.5"" from 2MASS 1207 A to avoid residual halo contributions."," We assume Gaussian error bars (equal on all pixels for the background limited case), and fit within a 5 pixel diameter aperture (see Figure \ref{2MASS 1207_image}) ) to avoid contamination by 2MASS 1207 A. We do a Monte Carlo evaluation of the background fluctuations by doing a similar aperture measurements at 10000 randomly chosen position angles and separations $>$ 1.5"" from 2MASS 1207 A to avoid residual halo contributions."568" For each Monte Carlo trial, we subtract the flux of arandomly chosen background aperture from the measured flux at the position of 2MASS 1207 b. This gives us an estimate, for each trial, of what the true flux of 2MASS 1207 b might be."," For each Monte Carlo trial, we subtract the flux of a randomly chosen background aperture from the measured flux at the position of 2MASS 1207 b. This gives us an estimate, for each trial, of what the true flux of 2MASS 1207 b might be."569" We discard all negative flux measurements based on the Bayesian prior that 2MASS 1207 b’s flux must be positive, and use the remaining cases to construct a cumulative distribution function, which calculates the probability that 2MASS 1207 b’s flux is less than a given value (but greater than 0)."," We discard all negative flux measurements based on the Bayesian prior that 2MASS 1207 b's flux must be positive, and use the remaining cases to construct a cumulative distribution function, which calculates the probability that 2MASS 1207 b's flux is less than a given value (but greater than 0)."570" We include an absolute calibration error (see the previous paragraph), although increasing this error up to ~20% has a negligible effect on our results."," We include an absolute calibration error (see the previous paragraph), although increasing this error up to $\sim$ has a negligible effect on our results."571Figure 8 shows the spectral fitting with a power law model for the MCAL data of GRB 080514B. in the energy range 500-5000 keV. using preliminary calibration parameters.,"Figure \ref{080514B_MCAL} shows the spectral fitting with a power law model for the MCAL data of GRB 080514B, in the energy range 500–5000 keV, using preliminary calibration parameters."572 The solid line represents the best-fit model folded with the MCAL response (reduced q=1.1 with [2 degrees of freedom)., The solid line represents the best-fit model folded with the MCAL response (reduced $\chi^2 = 1.1 $ with 12 degrees of freedom).573 The photon index is -2.66133. where quoted errors are at the confidence level.," The photon index is $-2.66^{+0.30}_{-0.25}$, where quoted errors are at the confidence level."574 This value is consistent within errors. with the high energy photon index obtained fitting Konus-Wind and Suzaku-WAM data with a Band model. as reported in ? ?.. 9.. (?).. (??).. 3.2..," This value is consistent within errors with the high energy photon index obtained fitting Konus-Wind and Suzaku-WAM data with a Band model, as reported in \citet{Golenetskii2008GCN7751} \citet{Hanabata2008GCN7752}. \ref{Fig080323}. \citep{Fuschino2008b}. \citep{Fishman1994,Smith2005}, \ref{sensitivity}."575We now consider the behavior of charges in a turbulent magnetic plasma where magnetosonic modes provide (he sites of magnetic mirroring.,	We now consider the behavior of charges in a turbulent magnetic plasma where magnetosonic modes provide the sites of magnetic mirroring.576 This scenario differs [rom first order Fermi acceleration bv shocks in two wavs., This scenario differs from first order Fermi acceleration by shocks in two ways.577 1) Consecutive mirroring events are not coherent. but rather stochastically distributed between heac-on and catch-up.," 1) Consecutive mirroring events are not coherent, but rather stochastically distributed between head-on and catch-up."578 2) The (turbulent. cascade eoverns (he acceleration elliciency: (he svstem picks out a scale where acceleration compeles wilh the cascade., 2) The turbulent cascade governs the acceleration efficiency; the system picks out a scale where acceleration competes with the cascade.579 In the solar corona plasma. (.. the velocity of the magnetic compressions. is the phase speed of the magnetosonic modes. which is roughly the Alfvénn speed for the [ast mode and the sound speed (6ος) for the slow mode.," In the solar corona plasma, $v_c$, the velocity of the magnetic compressions, is the phase speed of the magnetosonic modes, which is roughly the Alfvénn speed for the fast mode and the sound speed $c_s$ ) for the slow mode."580" Typically. thermal electrons in the corona are super-Alfvénnic and non-relativistic. ry«ο, aud ο)0.05; we will solve ihe STFA problem in (his regime."," Typically, thermal electrons in the corona are super-Alfvénnic and non-relativistic, $v_A \ll v_0 \ll c$, and $\beta \backsim 0.05$; we will solve the STFA problem in this regime."581 Recall that the energv exin. from a (vpical reflection is eiven in Eq.(5)) to be c2)., Recall that the energy gain from a typical reflection is given in \ref{Epm}) ) to be $\delta E_\pm = 2m(\pm v_\parallel v_c + v^2_c)$ .582 We deline three parameters: H. the total rate of reflections: R. the rate of head-on reflections: and 22. the rate of catch-up reflections.," We define three parameters: $R$, the total rate of reflections; $R_+$, the rate of head-on reflections; and $R_-$, the rate of catch-up reflections."583 The relation R=Rh+h is automatically satisfied by (his definition as all reflections must be of either the head-on or catel-up (ype., The relation $R = R_+ +R_-$ is automatically satisfied by this definition as all reflections must be of either the head-on or catch-up type.584 This allows us (ο write the approximate acceleration rate as the sum of a coherent term and an incoherent term: where the subscript 5 is used to distinguish our derived. acceleration rate from that, This allows us to write the approximate acceleration rate as the sum of a coherent term and an incoherent term: where the subscript $S$ is used to distinguish our derived acceleration rate from that585 and This gives us all of the ingredients for calculating the steady acceleration [rom (7)) The resulting acceleration rate is where we have added (he additional subsceript 5 to indicate the distinction [from the previously calculated rate.,"Recall that from \ref{erates}) ), From this one easily obtains and This gives us all of the ingredients for calculating the steady acceleration from \ref{sdefine}) ) The resulting acceleration rate is where we have added the additional subscript $b$ to indicate the distinction from the previously calculated rate."586 The steady acceleration rate found in (10)) from the assumption Fo=F =Fis Note that provided 4>45 (his is the largest term in (59))., The steady acceleration rate found in \ref{ratecalcfirst}) ) from the assumption $F_+ = F_- = F$ is Note that provided $A > 4B$ this is the largest term in \ref{sb}) ).587 Indeed. for coronal flare plasma. ο0.1 at electron energy £j aud decreases with increasing energv while Aκ0.1 as well at Ly. but is largely insensitive to eletron energy.," Indeed, for coronal flare plasma, $B \backsim 0.1$ at electron energy $E_0$ and decreases with increasing energy while $A \backsim 0.1$ as well at $E_0$, but is largely insensitive to eletron energy."588 At the onset of the power law regime.," At the onset of the power law regime,"589For planet-planet interactions to produce an eccentric orbit for an isolated planet requires al least one planet to disappear.,For planet-planet interactions to produce an eccentric orbit for an isolated planet requires at least one planet to disappear.590 The missing planet might have collided ancl merged with the remaining planet or effectively. merged due to tidal capture., The missing planet might have collided and merged with the remaining planet or effectively merged due to tidal capture.591 It might also have been ejecled [rom the system or [allen into the central star., It might also have been ejected from the system or fallen into the central star.592 Numerical integrations of svstenis wilh (wo equal mass planets on initially circular orbits by Fordetal.(2001) produce a much ereater Iraction of isolated planets with low eccentricity orbits than is observed., Numerical integrations of systems with two equal mass planets on initially circular orbits by \cite{FHR01} produce a much greater fraction of isolated planets with low eccentricity orbits than is observed.593 These are a consequence of mergers which Fordetal.(2001) assume to occur whenever (he separation between the planets drops below the sum of their radii., These are a consequence of mergers which \cite{FHR01} assume to occur whenever the separation between the planets drops below the sum of their radii.594 Bul the case against planet-planet interactions is even stronger (han the results of Fordetal.(2001) indicate because they jeglect tidal captures., But the case against planet-planet interactions is even stronger than the results of \cite{FHR01} indicate because they neglect tidal captures.595 Taking the tidal capture cross section Lor =1 polvtropes from (1999).. and the relative velocity at infinity as numerically caleulated by (1996).. we estimate a critical impact parameter for (dal capture between (wo ancl three times largerthan (he two radii assumed [or merger bv Rasio&Ford(1996) and 2001).," Taking the tidal capture cross section for $n=1$ polytropes from \cite{KIL99}, and the relative velocity at infinity as numerically calculated by \cite{RAF96}, we estimate a critical impact parameter for tidal capture between two and three times largerthan the two radii assumed for merger by \cite{RAF96} and \cite{FHR01}."596". Next we consider some aspects of the merger process for a svstem whose initial state consists of two planets. each having mass MM, and radius Πρ. moving on coplaner circular orbits with radii ry and rs around a star of mass M, and radius A2,."," Next we consider some aspects of the merger process for a system whose initial state consists of two planets, each having mass $M_p$ and radius $R_p$, moving on coplaner circular orbits with radii $r_1$ and $r_2$ around a star of mass $M_*$ and radius $R_*$."597" We asstume that the final state consists of a single or binary planet with mass 2.M, moving on an orbit with senimajor axis à and eccentricity €. Applving conservation of energv and angular momentum. we relate the final orbit to the initial ones by and llere AF ancl A// are the οποιον and angular momentiun stored internally in either ihe merger product or in the relative orbit of the binary."," We assume that the final state consists of a single or binary planet with mass $2M_p$ moving on an orbit with semimajor axis $a$ and eccentricity $e$ Applying conservation of energy and angular momentum, we relate the final orbit to the initial ones by and Here $\Delta E$ and $\Delta H$ are the energy and angular momentum stored internally in either the merger product or in the relative orbit of the binary."598 Energy. dissipated by impact or through (tidal dissipation is accounted for by an increase of the binding energy., Energy dissipated by impact or through tidal dissipation is accounted for by an increase of the binding energy.599" We estimate AZ and AJ! by noting that when the planets are separated bv less than the Hill radius. ria=CMU,/3M.) ""a. they are effectively a two-body system."," We estimate $\Delta E$ and $\Delta H$ by noting that when the planets are separated by less than the Hill radius, $r_{\rm Hill}\equiv (M_p/3M_*)^{1/3}a$ , they are effectively a two-body system."600evele. and thus do not allow for the fact that the stream will have moved on by the time the material accretes.,"cycle, and thus do not allow for the fact that the stream will have moved on by the time the material accretes."601 Lastly. we rotate the dipole and stream location to mimic the spin and orbital eveles.," Lastly, we rotate the dipole and stream location to mimic the spin and orbital cycles."602 We do not include any emission from the stream itself. further out thanAue. whieh would produce a lower-velocity S-wave on the orbital evele.," We do not include any emission from the stream itself, further out than, which would produce a lower-velocity S-wave on the orbital cycle."603 Note that the above model is very similar to that developed by Ferrario WWickramasinghe (1999). the main difference being that they aimed for model Fourier transforms whereas we aim to compare line profiles.," Note that the above model is very similar to that developed by Ferrario Wickramasinghe (1999), the main difference being that they aimed for model Fourier transforms whereas we aim to compare line profiles."604 Fig., Fig.605 7 shows a sample set of simulated line profiles., 7 shows a sample set of simulated line profiles.606 We acloptecl spin and orbital eveles of 027 s and 3.42 h. a white cbwarf mass of 0.80AL... and à secondary mass of 0.37AL.," We adopted spin and orbital cycles of 927 s and 3.42 h, a white dwarf mass of 0.80, and a secondary mass of 0.37."607.. The inclination was10.. the dipole-olffset. 9. was aand the radius wwas 20fea.," The inclination was, the dipole-offset, $\delta$, was and the radius was 20."608 Phe four panels are for emission at different distances from the white ναι, The four panels are for emission at different distances from the white dwarf.609 A velocity variation at the orbital evele is obvious in all panels of Fig., A velocity variation at the orbital cycle is obvious in all panels of Fig.610 7., 7.611 Note that this is primarily the infall velocity of the stream. as in an XM Ler star. and is not the orbital velocity of the white cwarf. which is much smaller.," Note that this is primarily the infall velocity of the stream, as in an AM Her star, and is not the orbital velocity of the white dwarf, which is much smaller."612 The spin and beat eveles are seen as the much faster variation in the simulated. profiles., The spin and beat cycles are seen as the much faster variation in the simulated profiles.613 The Fourier transforms (Fig., The Fourier transforms (Fig.614 6) revealed a variation of the line profiles over the orbital evele. but when the observed profiles are folded on the orbital evele it is hardly visible (lig.," 6) revealed a variation of the line profiles over the orbital cycle, but when the observed profiles are folded on the orbital cycle it is hardly visible (Fig."615 S)., 8).616 To enhance the variation we have subtracted the phasc-invariant profile (for each velocity we found the phase bin with the lowest. value ancl subtractecl that from. the data)., To enhance the variation we have subtracted the phase-invariant profile (for each velocity we found the phase bin with the lowest value and subtracted that from the data).617 In the portion of the line remaining. only 5 per cent of the original. the slight variation is discernable.," In the portion of the line remaining, only 5 per cent of the original, the slight variation is discernable."618 Since the orbital modulation is so small. it is clear that numust be at a very [ow inclination (as already suggested by he X-ray data)," Since the orbital modulation is so small, it is clear that must be at a very low inclination (as already suggested by the X-ray data)."619 Àut even so. the lack of variation is clearly discrepant with the model. which is caleulated for /=10," But even so, the lack of variation is clearly discrepant with the model, which is calculated for $i = 10$."620" ὃν fitting Gaussians to the phase-resolvect profiles we measure a racial-velocity variation of only 5.8ὃν, which agrees with Buckley ((1995) who report a limit of < 110"," By fitting Gaussians to the phase-resolved profiles we measure a radial-velocity variation of only 5.8, which agrees with Buckley (1995) who report a limit of $<$ 10."621 Uf this is the motion of the white dwarf it implies an inclination of only ((assuming stellar masses of 0.7 and 0.3. M.)., If this is the motion of the white dwarf it implies an inclination of only (assuming stellar masses of 0.7 and 0.3 ).622 Even this could be too high. since. as shown bv the simulations. we expect the orbital motion to be dwarfed. by an infall velocity of anything up to the white-dwarl escape velocity.," Even this could be too high, since, as shown by the simulations, we expect the orbital motion to be dwarfed by an infall velocity of anything up to the white-dwarf escape velocity."623 The observations and simulations can be reconciled only it 1) the system is at an inclination of « (unlikely. since the. probability is only 5. or 2) much of the emission comes from material that has circularised about the white dwarf and clilutes the emission from the infalling stream.," The observations and simulations can be reconciled only if 1) the system is at an inclination of $<$ (unlikely, since the probability is only ), or 2) much of the emission comes from material that has circularised about the white dwarf and dilutes the emission from the infalling stream."624 Thus the variations in, Thus the variations in625the optical path control is achieved with an optical path sensitivity over 0.01μπι. The optical fields emerging from the 8 interferometric arms reach a 8 to one polarization maintaining (PM) coupler to achieve the interferometric mixing.,the optical path control is achieved with an optical path sensitivity over $0.01 \mu m$ The optical fields emerging from the 8 interferometric arms reach a 8 to one polarization maintaining (PM) coupler to achieve the interferometric mixing.626 At the output an InGaAs photodiode detects the interferometric signal that is recorded through a standard 12 bits ADC voltage acquisition system., At the output an InGaAs photodiode detects the interferometric signal that is recorded through a standard 12 bits ADC voltage acquisition system.627" In order to calibrate the imaging properties of our instrument, the first investigation has been to characterize its point spread function."," In order to calibrate the imaging properties of our instrument, the first investigation has been to characterize its point spread function."628" For this purpose, the telescope array has been illuminated by a plane wave using only a single point-like source ( i.e. switching on only 1 fibre of the"," For this purpose, the telescope array has been illuminated by a plane wave using only a single point-like source ( i.e. switching on only 1 fibre of the"629 , 630instabilities with racial oscillations. stationary w=0 wave patterns mark the onset of instabilities (Lemos ct al.,"instabilities with radial oscillations, stationary $\omega=0$ wave patterns mark the onset of instabilities (Lemos et al."631 1991: Sver “Tremaine 1996: Lou 2002: Lou Fan 2002: Shu et al., 1991; Syer Tremaine 1996; Lou 2002; Lou Fan 2002; Shu et al.632 2000: Lou Shen 2003: Shen Lou 2003: Lou Zou 2004) in a composite disc svstem., 2000; Lou Shen 2003; Shen Lou 2003; Lou Zou 2004) in a composite disc system.633 Apparently. there are two unstable regimes. namely. the long wavelength collapse regime and the short wavelength ring fragmentation regine.," Apparently, there are two unstable regimes, namely, the long wavelength collapse regime and the short wavelength ring fragmentation regime."634" ""herefore. the stability criterion for 22. [alls in a range whose width increases with increasing 3."," Therefore, the stability criterion for $D_s^2$ falls in a range whose width increases with increasing $\beta$."635 Both regimes of the collapse instability ancl the ring fragmentation instability are reduced. for larger. values of 2., Both regimes of the collapse instability and the ring fragmentation instability are reduced for larger values of $\beta$.636 As already. noted. for jm0.436. the collapse regime disappears completely.," As already noted, for $\beta>0.436$, the collapse regime disappears completely."637 For sulliciently small values of 3<0.190. the stable range of D. does not exist (see Appendix D for details).," For sufficiently small values of $\beta<-0.130$, the stable range of $D_s^2$ does not exist (see Appendix D for details)."638 As remarked earlier. a composite svstem of two coupled clises is less stable than a single dise system.," As remarked earlier, a composite system of two coupled discs is less stable than a single disc system."639 The introduction of an additional gascous disc with larger 0 and 5g will reduce the overall stable range of D>. while tending to suppress the reginie of collapse for large-scale instabilities (Lou Shen 2003: Lou Zou 2004).," The introduction of an additional gaseous disc with larger $\delta$ and $\eta$ will reduce the overall stable range of $D_s^2$, while tending to suppress the regime of collapse for large-scale instabilities (Lou Shen 2003; Lou Zou 2004)."640 The logarithmic spiral 7i=1 case behaves qualitatively, The logarithmic spiral $m=1$ case behaves qualitatively641ffor GX 339-4 (for MgH=5.8Mo).,for GX 339–4 (for $M_{BH}$ $M_\odot$ ).642" Tomsicketal.(2004) followed XTE J1650-500 down to ~1.4x10?(d/6kpc)? ((for Μηη--δ.δΜ9) withChandra, and also found evidence for spectral softening."," \citet{2004ApJ...601..439T} followed XTE J1650–500 down to $\sim1.4\times 10^{-5} (d/6\,{\rm kpc})^2$ (for $M_{BH}$ $M_\odot$ ) with, and also found evidence for spectral softening."643" Finally, Corbeletal.(2008) show that the power-law spectral index is softer in V404 Cyg in quiescence compared to that in the brighter low—hard state."," Finally, \citet{2008MNRAS.389.1697C} show that the power–law spectral index is softer in V404 Cyg in quiescence compared to that in the brighter low–hard state."644 In Figure 4 we have plotted the power law index from the aand sspectral modelling on dduring the decay towards quiescence and in quiescence., In Figure \ref{plindex} we have plotted the power law index from the and spectral modelling on during the decay towards quiescence and in quiescence.645 The data point at the lowest X-ray flux is determined by combining the quiescent data obtained right after the outburst of early 2008 as well as the quiescence observations used by Corbeletal.(2005).., The data point at the lowest X–ray flux is determined by combining the quiescent data obtained right after the outburst of early 2008 as well as the quiescence observations used by \citet{2005ApJ...632..504C}.646" There is no evidence for a softening of the power-law index towards and in quiescence, although for this source the relatively high neutral hydrogen column density limits the accuracy of the power-law determination."," There is no evidence for a softening of the power–law index towards and in quiescence, although for this source the relatively high neutral hydrogen column density limits the accuracy of the power–law determination."647 A fit of a constant power-law spectral index gives a best-fitting index of 1.704+0.008 with a x?=7 for 9 degrees of freedom., A fit of a constant power–law spectral index gives a best–fitting index of $\pm$ 0.008 with a $\chi^2=7$ for 9 degrees of freedom.648 We have compared the X-ray flux decay rate of wwith that observed during the last part of the outburst decay from the black hole candidate XTE J1908+094 and we find that the fit function describing the decay of iis a good approximation to the last phase of approximately three weeks presented in Jonker et al. (, We have compared the X–ray flux decay rate of with that observed during the last part of the outburst decay from the black hole candidate XTE J1908+094 and we find that the fit function describing the decay of is a good approximation to the last phase of approximately three weeks presented in Jonker et al. (6492004) for XTE J1908+094.,2004) for XTE J1908+094.650" Note however, that due to a reduced sampling that source was only observed twice during those three weeks."," Note however, that due to a reduced sampling that source was only observed twice during those three weeks."651" Nevertheless, it is interesting to see that the sources have a similar decay rate."," Nevertheless, it is interesting to see that the sources have a similar decay rate."652" If other sources also follow the same decay rate in X-rays this could provide constraints on the accretion disc model, such as the accretion disc to ADAF evaporation."," If other sources also follow the same decay rate in X–rays this could provide constraints on the accretion disc model, such as the accretion disc to ADAF evaporation."653 PGJ acknowledges support from a VIDI grant from the Netherlands Organisation for Scientific Research., PGJ acknowledges support from a VIDI grant from the Netherlands Organisation for Scientific Research.654" MAPT acknowledges support from NASA grant GO8-9042A. DS acknowledges an STFC Advanced FellowshipThe National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."," MAPT acknowledges support from NASA grant GO8-9042A. DS acknowledges an STFC Advanced FellowshipThe National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."655"Alaenetic fields are being detected more axd more routinely at the surface of many stars. and are responsible for various plhisical phenomena ikcly to deeply modify our traditional vision of stellar evolution. especially diving thei carly and late stagOs,","Magnetic fields are being detected more and more routinely at the surface of many stars, and are responsible for various physical phenomena likely to deeply modify our traditional vision of stellar evolution, especially during their early and late stages."656" The rescence, of even a relatively weak magnetic field can have an importai onupact on the collapse and yagiuentation of pre*cellar cores (Conuuerconot 2010).. as well as ιαοποιο the rotation rate of stars {seeeg.Alecianetal.2008)."," The presence of even a relatively weak magnetic field can have an important impact on the collapse and fragmentation of prestellar cores \citep{Commercon:2010}, as well as influencing the rotation rate of stars \cite[see e.g.][]{Alecian:2008}."657. Ou he other side of the diaeran its COCback effects may aso plav a kev role iu superuovae aud mechanical enerev deposition in theiu. for iustiuce.," On the other side of the diagram its feedback effects may also play a key role in supernovae and mechanical energy deposition in the, for instance."658 Alaeuetie fields are also an portant actors, Magnetic fields are also an important actor.659"tars, First. we cannot disniss the possibility of a large-scale magetic field being responsible for the quasi-unuiforni rotation1 behaviour iu the bulk of the solar radiation zone. as revealed by pauodes helioseiziuolo:αν(Ef-Darwichetal.2008)."," First, we cannot dismiss the possibility of a large-scale magnetic field being responsible for the quasi-uniform rotation behaviour in the bulk of the solar radiation zone, as revealed by $p$ -modes helioseismology\citep[][]{Eff-Darwich:2008}."660. Secoud. stroug fields (3Mi10 G to 30 κ) are observed via the Zeemau effect i sole fraction of nudrsequence A stars (theApstars.seeAuriereetal2¢ WF). as well as DB stars aud a handful of O sars (Cuunhutetal.2009).," Second, strong fields $300$ G to $30$ kG) are observed via the Zeeman effect in some fraction of main-sequence A stars \citep[the Ap stars, see][]{Auriere:2007}, as well as B stars and a handful of O stars \citep{Grunhut:2009}."661. The binodalitv of rotatioiib observed. among Ap ποια]. A stars sliows the critical effect of magnetic fieldsS on rotation axd therefore also om meridional circulation aud checal transport (seeMathis&Zalu 2005)., The bimodality of rotational observed among Ap normal A stars shows the critical effect of magnetic fields on rotation and therefore also on meridional circulation and chemical transport \citep[see][]{Mathis:2005}.662". Finally. maejetie white dwarfs display fields streneth of 10110? C. and neutron stars host fields in the range 103101 Q1, in both cases detected using SCVCral distinct methods."," Finally, magnetic white dwarfs display fields strength of $10^4 - 10^9$ G, and neutron stars host fields in the range $10^8-10^{15}$ G, in both cases detected using several distinct methods."663" The ασοο-πςαο, ordered nature of these fields (often approxilately dipolar) aud the scaling of their streugtlis as a fuiction «ft their lost properties (according to the fiux coiservation scenario) favour a fossil lapothesis. whose ¢10111 Is rot vet ebucidated."," The large-scale, ordered nature of these fields (often approximately dipolar) and the scaling of their strengths as a function of their host properties (according to the flux conservation scenario) favour a fossil hypothesis, whose origin is not yet elucidated."664 Anoticr fundamental question is the topology of these 1arec-xcae maguetie fields., Another fundamental question is the topology of these large-scale magnetic fields.665 To lave survived since. fie starS formation. a field must be stable ou a dvuausc (Alvwéónn) timescale.," To have survived since the star's formation, a field must be stable on a dynamic (Alfvénn) timescale."666 It was suggested by Prendergast(1956) tat a stellar iiagnetic field in stable axisviunietrie equilibrium iust contain both poloidal (1neridional) aid toroidal (azimuthal) components. since both are mustable oi them own (Tavler1973:Wright1973:Draithw:ute2006:Bonanno&Urpin 2008b).," It was suggested by \cite{Prendergast:1956} that a stellar magnetic field in stable axisymmetric equilibrium must contain both poloidal (meridional) and toroidal (azimuthal) components, since both are unstable on their own \citep{Tayler:1973,Wright:1973,Braithwaite:2006b,Bonanno:2008}."667. This was confine recently by nununerical πιαΊος bw Braithwaite&Spruit(2001):BraitlwaiteNordluud(2006) who showed tat an arbitrary initial feld evolves ou an Alfvéni timescale into a stable configuration: axisviunietric wined poloidaltoroidal folds were found.," This was confirmed recently by numerical simulations by \cite{Braithwaite:2004, Braithwaite:2006a} who showed that an arbitrary initial field evolves on an Alfvénn timescale into a stable configuration; axisymmetric mixed poloidal-toroidal fields were found."668 Tere. we deal with axisvuuuetre. nou force-free magnetic configurations in equiliDrivin inside a conductive fluid.," Here, we deal with axisymmetric, non force-free magnetic configurations in equilibrium inside a conductive fluid."669 We first restrict ourselves to the non-rotatiug case. but results also apply to rotating stars where rotation is uuiforni (Woltjer1959). J which could be thecase if magnetic feld is strong enousli. and where meridional circulation can be ucelected," We first restrict ourselves to the non-rotating case, but results also apply to rotating stars where rotation is uniform \citep{Woltjer:1959b}, , which could be thecase if magnetic field is strong enough, and where meridional circulation can be neglected"670other structures at similar Faraday depth that can not be the Perseus cluster.,other structures at similar Faraday depth that can not be the Perseus cluster.671 The fractional polarization at 351 MHz was estimated by dividing the polarized intensity. integrated over all Faraday depths. by the 408 MHz ? total intensity map converted to 35] MHz using a Galactic synchrotron brightness temperature spectralindex B=-2.8 (???)..," The fractional polarization at 351 MHz was estimated by dividing the polarized intensity, integrated over all Faraday depths, by the 408 MHz \citet{HaslamEtAl1982} total intensity map converted to 351 MHz using a Galactic synchrotron brightness temperature spectralindex $\beta=-2.8$ \citep{ReichReich1988A,ReichReich1988B,PlataniaEtAl1998}."672 Between 10 MHz and 100 MHz. the spectral index is —2.55 (?).. hence the actual spectral index between 408 MHz and 351 MHz is probably closer to —2.7.," Between 10 MHz and 100 MHz, the spectral index is $-2.55$ \citep{Cane1979}, hence the actual spectral index between 408 MHz and 351 MHz is probably closer to $-2.7$."673 The difference with —2.8 1s negligible for the small extrapolation from 408 MHz to 351 MHz., The difference with $-2.8$ is negligible for the small extrapolation from 408 MHz to 351 MHz.:674Gaussian: where $?(1)di 1s the probability of finding a noise value between i and 2+ da. and ji and c are the mean and If the Q and U noise distributions have equal c. zero mean. and uncorrelated.. then the probability of finding a value of |F| between f and f£+df is: The RMS of |F] is equal to the RMS of Q and U. which ts oV2.," where $\mathcal{P}(n)\ \mathrm{d}n$ is the probability of finding a noise value between $n$ and $n+\mathrm{d}n$ , and $\mu$ and $\sigma$ are the mean and If the $Q$ and $U$ noise distributions have equal $\sigma$ , zero mean, and , then the probability of finding a value of $|F|$ between $f$ and $f+\mathrm{d}f$ is: The RMS of $|F|$ is equal to the RMS of $Q$ and $U$, which is $\sigma\sqrt{2}$."675" The mean value of the noise in [F] is The polarized surface brightness of a line of sight. integrated over a range of equidistant Faraday depths όν--- 6. and corrected for the non-zero mean of the noise level. is where B is the area under the restoring beam of the CLEAN divided by Ao=|o,—¢;|."," The mean value of the noise in $|F|$ is The polarized surface brightness of a line of sight, integrated over a range of equidistant Faraday depths $\phi_1\cdots\phi_n$ , and corrected for the non-zero mean of the noise level, is where $B$ is the area under the restoring beam of the RM-CLEAN divided by $\Delta\phi = |\phi_{i+1}- \phi_i|$."676 ? have conducted an absolutely calibrated survey of polarized emission north of declination 230* at 1.41 GHz with the 26 m telescope at the DRAO site 36’ FFWHM., \citet{WollebenEtAl2006} have conducted an absolutely calibrated survey of polarized emission north of declination $-30\degr$ at $1.41$ GHz with the $26$ m telescope at the DRAO site $36\arcmin$ FWHM.677 The integrated 351 MHz polarized intensity overlaid with the polarized intensity contours from ? is shown in Fig. 6.., The integrated 351 MHz polarized intensity overlaid with the polarized intensity contours from \citet{WollebenEtAl2006} is shown in Fig. \ref{brentjens_perseusmosaic_fig:wolleben-overlay}.678 The noise level in the 351 MHz map is approximately 0.5 K. With a spectral index of —2.8. the brightness temperature at 35] MHz should be 50 times higher than at 1.41 GHz.," The noise level in the 351 MHz map is approximately 0.5 K. With a spectral index of $-2.8$, the brightness temperature at 351 MHz should be 50 times higher than at 1.41 GHz."679 This is approximately the case in most of the field. which implies that there is very little depolarization between1.41 GHz and 35] MHz.," This is approximately the case in most of the field, which implies that there is very little depolarization between1.41 GHz and 351 MHz."680 In some places. the polarized intensity is even higher at 351 MHz than one would expect based on the low resolution polarized intensity at 1.41 GHz and a spectral index of —2.8.," In some places, the polarized intensity is even higher at 351 MHz than one would expect based on the low resolution polarized intensity at 1.41 GHz and a spectral index of $-2.8$."681" Examples are the area containing the “front”. “lens”. and ""doughnut"". and the highly polarized region in field A. This is probably caused by beam depolarization in the ? observations due to differences in intrinsic polarization angle at scales well below 30’ that are resolved in the observations presented here."," Examples are the area containing the “front”, “lens”, and “doughnut”, and the highly polarized region in field A. This is probably caused by beam depolarization in the \citet{WollebenEtAl2006}682 observations due to differences in intrinsic polarization angle at scales well below $36\arcmin$ that are resolved in the observations presented here."683 Figure 7 displays the fractional polarization., Figure \ref{brentjens_perseusmosaic_fig:polfraction} displays the fractional polarization.684 It is mostly between and with a maximum of in field A. Although these values are well below the theoretical maximum of70%.. they are relatively high.," It is mostly between and with a maximum of in field A. Although these values are well below the theoretical maximum of, they are relatively high."685 The low fractional polarization between and is an artifact caused by the low resolution (0285) of the ? map. which blends these powerful sources.," The low fractional polarization between and is an artifact caused by the low resolution $0\fdg85$ ) of the \citet{HaslamEtAl1982} map, which blends these powerful sources."686 Because there are no absolutely calibrated polarimetric single dish observations of the fieldMHz.. my maps may lack Q and U features at scales > 90’.," Because there are no absolutely calibrated polarimetric single dish observations of the field, my maps may lack $Q$ and $U$ features at scales $\gtrsim 90\arcmin$ ."687 The are therefore strictly speaking lower limits., The are therefore strictly speaking lower limits.688 The lack ofdepolarization implies that thesynchrotronemitting areas have a Faraday thickness ofless than | rad m7?., The lack ofdepolarization implies that thesynchrotronemitting areas have a Faraday thickness ofless than 1 rad $^{-2}$ .689 This the range of Faraday depths is two orders of magnitude larger., This the range of Faraday depths is two orders of magnitude larger.690 Assuming a line of sight, Assuming a line of sight691encounters between planets but is not telling us much about other properties of motion as for example whether or not the svstem is chaotic.,encounters between planets but is not telling us much about other properties of motion as for example whether or not the system is chaotic.692 Wisdom(1980). basing on the resonance overlapping criterion found a boundary between chaotic and stable motions in the form where m denotes mass ratio of the two biggest bodies in the system., \cite{wisdom80} basing on the resonance overlapping criterion found a boundary between chaotic and stable motions in the form where $m$ denotes mass ratio of the two biggest bodies in the system.693 The coellicient in front of the mass in the original paper was 2. but it has been improved later by numerical simulations (Duncanetal.1989). and. we have adopted here the more recent value 1.5.," The coefficient in front of the mass in the original paper was 2, but it has been improved later by numerical simulations \citep{duncan} and we have adopted here the more recent value 1.5."694 Phe occurrence of the mean motion resonances is therefore very important. for stability of the svstem., The occurrence of the mean motion resonances is therefore very important for stability of the system.695 The width of the resonance regions for the circular restricted. three-body problem has been derived. by Wisdom(19080). (see also Lecarctal. (2001)))., The width of the resonance regions for the circular restricted three-body problem has been derived by \cite{wisdom80} (see also \cite{lecar}) ).696 In Fig., In Fig.697 2 we have plotted the width of 3:2 and 4:3 interior resonances for zero eccentricity given by Lecaretal.(2001). for the Jupiter (vertical dashed lines)., \ref{fig2} we have plotted the width of 3:2 and 4:3 interior resonances for zero eccentricity given by \cite{lecar} for the Jupiter (vertical dashed lines).698 In this Figure the semi-major axis ratio is obtained by dividing the semi-major axis of the Jupiter by that of the small body., In this Figure the semi-major axis ratio is obtained by dividing the semi-major axis of the Jupiter by that of the small body.699 Winter&Murray(1997) have cliscussect the various analytical mocdoels that have been used in the study of interior first order resonances., \cite{winmur} have discussed the various analytical models that have been used in the study of interior first order resonances.700 Analyzing simple pendulum. model they derived. formula. describing masxiniunm deviation of semi-major axis from the nominal value given by exact mean motion commoensurabilitv., Analyzing simple pendulum model they derived formula describing maximum deviation of semi-major axis from the nominal value given by exact mean motion commensurability.701 Their results for uixcl 4:3 resonances are presented also in Fig., Their results for 3:2 and 4:3 resonances are presented also in Fig.702 2 [or comparison (solid. lines)., \ref{fig2} for comparison (solid lines).703 The theoretical prediction illastrated by the pendulum. model. (similar. results were obtained using Hamiltonian approach) that for higher value of the eccentricity of the small body the first order resonance width is larger. has not been verified by. direct. integration (Winter&Murray1997).," The theoretical prediction illustrated by the pendulum model (similar results were obtained using Hamiltonian approach) that for higher value of the eccentricity of the small body the first order resonance width is larger, has not been verified by direct integration \citep{winmur}."704.. The size of the resonant libration region derived by Winter&Alurray(1997). performing numerical integrations of the full equations. of motion is quite similar to that given by Lecarctal.(2001)., The size of the resonant libration region derived by \cite{winmur} performing numerical integrations of the full equations of motion is quite similar to that given by \cite{lecar}.705.. The xot of maximum libration widths has been supplemented with a few examples of the mean motion commensurability known in our Solar System., The plot of maximum libration widths has been supplemented with a few examples of the mean motion commensurability known in our Solar System.706 The cots denote Llilda ane Thule eroups of asteroids., The dots denote Hilda and Thule groups of asteroids.707 This groups are particularly interesting »eause they are locked in the mean motion resonances with Jupiter., This groups are particularly interesting because they are locked in the mean motion resonances with Jupiter.708 Note that asteroids are located exactly within regions predicted by analysis of the restricted. three-bocdy oblem., Note that asteroids are located exactly within regions predicted by analysis of the restricted three-body problem.709 We have performed an analysis of the chaotic regions for our system (the Super-Earth and Jupiter around the vpe star) in the similar way it has been done for asteroid elt., We have performed an analysis of the chaotic regions for our system (the Super-Earth and Jupiter around the Solar-type star) in the similar way it has been done for asteroid belt.710 First. let us evaluate the size of the chaotic region using equation (7)).," First, let us evaluate the size of the chaotic region using equation \ref{wisdom}) )."711 Substituting in this equation for m the mass ratio of Jupiter and Sun we have found that the small body (here the Super-ISarth) should be outside the region of large scale chaos if According to this criterion the Super-Earth would be in the region of large scale chaos when the Jupiter's semi-major axis is smaller than 1.27., Substituting in this equation for $m$ the mass ratio of Jupiter and Sun we have found that the small body (here the Super-Earth) should be outside the region of large scale chaos if According to this criterion the Super-Earth would be in the region of large scale chaos when the Jupiter's semi-major axis is smaller than 1.27.712 Next. we have examined the properties of the planetary orbits in the Lill stable region. which means for the Jupiter semi-major axis having value in the range between 1.24 and 1.45.," Next, we have examined the properties of the planetary orbits in the Hill stable region, which means for the Jupiter semi-major axis having value in the range between 1.24 and 1.45."713 We have found two families of regular orbits separated from each other by the zone of the chaotic behaviour. which is located between 1.277 and 1.289.," We have found two families of regular orbits separated from each other by the zone of the chaotic behaviour, which is located between 1.277 and 1.289."714 Ht should be pointed out here that the fact of the existence of the chaotic zone outside the region indicated by the resonance overlap criterion is not entirely unexpected., It should be pointed out here that the fact of the existence of the chaotic zone outside the region indicated by the resonance overlap criterion is not entirely unexpected.715 I should be kept in mind that criterion given by equation (7)) is quite approximate and. Wisdoni's scaling law is strictly valid only in the asymptotic limit. when the integer p defining first order resonance (p|1):p is much bigger than one (Malhotra1996).," It should be kept in mind that criterion given by equation \ref{wisdom}) ) is quite approximate and Wisdom's scaling law is strictly valid only in the asymptotic limit, when the integer $p$ defining first order resonance $(p+1) : p$ is much bigger than one \citep{malhotra}."716. In our case of the Jupiter mass planet we are concerned. with the first. order resonances with small p., In our case of the Jupiter mass planet we are concerned with the first order resonances with small $p$.717 The families of the stable orbits can be nicely portrayed drawing the structure of the phase space trajectories in a plane on which the polar coordinates are (e. 6) where e is the Super-Earth eccentricity and © is the resonance angle defined as Llere Ase and Ay are the instantaneous mean longitudes of the Super-Earth and Jupiter respectively. and zesz is the longitudes of periastron of the Super-Earth.," The families of the stable orbits can be nicely portrayed drawing the structure of the phase space trajectories in a plane on which the polar coordinates are $e$, $\phi$ ) where $e$ is the Super-Earth eccentricity and $\phi$ is the resonance angle defined as Here $\lambda_{SE}$ and $\lambda_{J}$ are the instantaneous mean longitudes of the Super-Earth and Jupiter respectively, and $\varpi_{SE}$ is the longitudes of periastron of the Super-Earth."718 For given initial values of the eccentricity and resonance angle the motion of planets will follow the closed curve in this plane., For given initial values of the eccentricity and resonance angle the motion of planets will follow the closed curve in this plane.719 In Fig., In Fig.720 3 we have shown three dillerent. phase space trajectories for the three different Jupiter locations at 1.27. 1.28 ancl 1.30.," \ref{fig3} we have shown three different phase space trajectories for the three different Jupiter locations at 1.27, 1.28 and 1.30."721 If the Jupiter is at. 1.27. the Super-Earth has a regular very low cecentricity orbit. (first. panel). if it is placed. at 1.30 the Super-Earth trajectory librates around the stable equilibrium point in the 3:2 resonance zone (third. panel). and in between (when the Jupiter is at. 1.28) the planet experience a chaotic behaviour.," If the Jupiter is at 1.27, the Super-Earth has a regular very low eccentricity orbit (first panel), if it is placed at 1.30 the Super-Earth trajectory librates around the stable equilibrium point in the 3:2 resonance zone (third panel), and in between (when the Jupiter is at 1.28) the planet experience a chaotic behaviour."722 Similar chaotic zone has been found by Llolman&Murray(1996). in the asteroid, Similar chaotic zone has been found by \cite{holmur} in the asteroid723The Large Magellanic Cloud (LAIC) is the first spiral galaxy. in which an elongated feature of stars was identified.,The Large Magellanic Cloud (LMC) is the first spiral galaxy in which an elongated feature of stars was identified.724" In a remarkable drawing of the LAIC as seen with the naked eve made by Sir John Herschel. a bar is clearly shown. which Sir John termed an “axis of light""."," In a remarkable drawing of the LMC as seen with the naked eye made by Sir John Herschel, a bar is clearly shown, which Sir John termed an “axis of light”."725 Also faithfully represented in that drawing is a prominent spiral arm to the north. as well as two “embrvonic§ arms.," Also faithfully represented in that drawing is a prominent spiral arm to the north, as well as two “embryonic” arms."726 The LAIC belongs to the de Vaucouleurs classification bin SD(s)m. and is approximately 50 kpe distant (basedonthedistancemodulus1—M=18.5040.10.followingFreedmanetal. 2001).," The LMC belongs to the de Vaucouleurs classification bin SB(s)m, and is approximately 50 kpc distant \citep[based on the distance modulus $m-M=18.50 \pm 0.10$,727following][]{freedmanetal01}."728. Studies of the LAIC coneur that the inclination to the line-o[-sight of the (warpecl) disk lies in the range 30° 50° while the position angle of the lineof nodes lies between 120° and 150°: an excellent review is provided by vanderMarelet (2008)., Studies of the LMC concur that the inclination to the line-of-sight of the (warped) disk lies in the range $^{\circ}$ $^{\circ}$ while the position angle of the lineof nodes lies between $^{\circ}$ and $^{\circ}$; an excellent review is provided by \citet{vandermareletal08}.729. The study by vanderMarvel&Cioni(2001) using DENIS and 2\LASS surveys vields an inclination of 35° and position angle of 125°. which are the values adopted in this study.," The study by \citet{cioni01}730 using DENIS and 2MASS surveys yields an inclination of $^{\circ}$ and position angle of $^{\circ}$, which are the values adopted in this study."731 The infrared. emission from dust in the diffuse interstellar medium of our Galaxy. has been extensively studied fromIRAS to the Herschel Space Observatory., The infrared emission from dust in the diffuse interstellar medium of our Galaxy has been extensively studied fromIRAS to the Herschel Space Observatory.732 The quantitative Galactic model of Draine&Li(2007) consists of a mixture of amorphous silicate grains and carbonaceous grains. each with a wide size distribution: sizes range [rom molecules containing tens of atoms to large erains greater than | jan. Dust grain temperatures al the surfaces of eiant molecular clouds in our Galaxy are hotter than in their cold interiors 1996).. varving Irom ~ 15 Ix to ~ 7 Ix. respectively.," The quantitative Galactic model of \citet{draine} consists of a mixture of amorphous silicate grains and carbonaceous grains, each with a wide size distribution: sizes range from molecules containing tens of atoms to large grains greater than 1 $\mu$ m. Dust grain temperatures at the surfaces of giant molecular clouds in our Galaxy are hotter than in their cold interiors \citep{greenbergli96}, varying from $\sim$ 15 K to $\sim$ 7 K, respectively."733 Cold dust grains are a crucial component in the ISA of galaxies (e.g.. Block et al.," Cold dust grains are a crucial component in the ISM of galaxies (e.g., Block et al."734 1994)., 1994).735" With regard to the dust temperature - metallicity relations in external svstems. Engelbrachtetal.(2003). [find an anticorrelation between dust temperature and metallicity,"," With regard to the dust temperature - metallicity relations in external systems, \citet{engelbracht08} find an anticorrelation between dust temperature and metallicity."736 Equilibrium dust temperatures are ~ 23 Ix. [or solaa-(wpe metallicities. increasing to 40 Ix at a low metallicity of 12+log(O/1H)~8 (see their Figure 5).," Equilibrium dust temperatures are $\sim$ 23 K for solar-type metallicities, increasing to 40 K at a low metallicity of $12+\log(O/H)\sim8$ (see their Figure 5)."737 First insights [rom a IHerschel study of Messier 33. with its strong metallicity eradient. show that dust temperatures are ~ 25 Ix in the central portion of the stellar disk. declining to  15 Ix in the outer domains (Figure2inBraineetal.2010).," First insights from a Herschel study of Messier 33, with its strong metallicity gradient, show that dust temperatures are $\sim$ 25 K in the central portion of the stellar disk, declining to $\sim$ 15 K in the outer domains \citep[Figure 2 in][]{braine10}."738. llere we study (he power spectrum (PS) of FUR dust emission from the LMC., Here we study the power spectrum (PS) of FIR dust emission from the LMC.739 Power spectra of emission maps are an important but poorly understood diagnostic for interstellar structures and the motions (hat cause them., Power spectra of emission maps are an important but poorly understood diagnostic for interstellar structures and the motions that cause them.740 PS of ILI emission from sections of the Milkv Way are approximately power laws throughout the entire range of observed spatial Irequencies (Crovisier&Dickey1983:Green1993:etal.2001:Khalil2006).," PS of HI emission from sections of the Milky Way are approximately power laws throughout the entire range of observed spatial frequencies \citep{crov83,green93,dickey01,khalil}."741. Similar slopes occur lor PS of Milky Wav CO emission (Stützkietal. 1993).. IRAS and DIRBE 100 jan emission (Gautierοἱal.1992:Schlegel.Finkbeiner.&Davis1998) and III absorption (Deshpande.Dwarakanath.&Goss 2000)..," Similar slopes occur for PS of Milky Way CO emission \citep{stut98}, , IRAS and DIRBE 100 $\mu$ m emission \citep{gautier,schlegel} and HI absorption \citep{des00}. ."742 These power laws resemble theoretical expectations from passive gas motions (Goldman2000) and compressions (Lazarian&Pogosvan2000) , These power laws resemble theoretical expectations from passive gas motions \citep{goldman} and compressions \citep{laz00} 743the present data we see no extra correlation of central black hole mass with specific frequency.,the present data we see no extra correlation of central black hole mass with specific frequency.744 The most sparsely populated area of the correlation is at low mass., The most sparsely populated area of the correlation is at low mass.745" For much less than 10°Mo, the presence of nuclear star clusters (NCs) may also become relevant (Graham&Spitler2009)."," For much less than $10^8 M_{\odot}$, the presence of nuclear star clusters (NCs) may also become relevant \citep{gra09}."746". Both BH and NC are part of the more general class of Central Massive Objects (CMO;seeWehner&Harris2006;Ferrareseetal.2006) and, as Graham&Spitler(2009) show, for some galaxies both are present at once."," Both BH and NC are part of the more general class of Central Massive Objects \citep[CMO; see][]{w3, f2} and, as \citet{gra09} show, for some galaxies both are present at once."747 For intermediate-mass galaxies there is a gradual mass transition zone such that for €2x10’Mo the NC mass tends to become dominant.," For intermediate-mass galaxies there is a gradual mass transition zone such that for $\la 7482 \times 10^7 M_{\odot}$ the NC mass tends to become dominant."749" Inspecting Table 1 of Graham Spitler, we note that for the MW and NGC 7457 particularly,& adding the NC mass to would bring them much closer to the mean E/S relation."," Inspecting Table 1 of Graham Spitler, we note that for the MW and NGC 7457 particularly, adding the NC mass to would bring them much closer to the mean E/S relation."750" For several other galaxies, however (NGC 1023, 1399, 3115, 4697), no important change would result."," For several other galaxies, however (NGC 1023, 1399, 3115, 4697), no important change would result."751 Additional data may help clarify whether or not the NC mass is an essential part of the picture., Additional data may help clarify whether or not the NC mass is an essential part of the picture.752" As said earlier, the basic fact that the GC populations and central black hole mass should be correlated is not surprising, because in rough terms bigger galaxies should have bigger subsystems of all types."," As said earlier, the basic fact that the GC populations and central black hole mass should be correlated is not surprising, because in rough terms bigger galaxies should have bigger subsystems of all types."753" What is more intriguing is that for the E and S galaxies especially, the correlation is so tight and so nearly linear."," What is more intriguing is that for the E and S galaxies especially, the correlation is so tight and so nearly linear."754 BT suggest that mergers may be responsible for the growth of both these subcomponents of the galaxy., BT suggest that mergers may be responsible for the growth of both these subcomponents of the galaxy.755" However, in a more general sense mergers will promote the growth of subcomponents of the galaxy."," However, in a more general sense mergers will promote the growth of subcomponents of the galaxy."756 We speculate that a more relevant link between black holes and globular clusters may be their age., We speculate that a more relevant link between black holes and globular clusters may be their age.757" Both the central BH and the majority of the GCs had their origins at high and nearly similar redshift, during the major stages of hierarchical merging."," Both the central BH and the majority of the GCs had their origins at high and nearly similar redshift, during the major stages of hierarchical merging."758" Most GCs have ages in the range 10—13 Gyr, corresponding to redshifts z~2—72010).."," Most GCs have ages in the range $10 - 13$ Gyr, corresponding to redshifts $z \sim 2 - 7$."759" For large galaxies the seeds of the central black holes are in place before z6 and grow by further gas accretion till z~2 and later (e.g.Alexanderetal.2010;Kellyetal.2010,amongmany others).."," For large galaxies the seeds of the central black holes are in place before $z \sim 6$ and grow by further gas accretion till $z \sim 2$ and later \citep[e.g.][among many others]{a1, h7, vol03, t3, v2, w4, hec04, sil08, mer10, kel10}."760" Both types of structures clearly require formation conditions of extremely high gas density (>10?Mc /pc? within small, parsec-sized volumes), which would have been easiest to generate in large numbers at very early times when gas cloud collisions and hierarchical merging events were frequent and energetic (cf.thediscussionsbyPengetal.2008;Rhode 2005)."," Both types of structures clearly require formation conditions of extremely high gas density $\ga 10^5 M_{\odot}$ $^3$ within small, parsec-sized volumes), which would have been easiest to generate in large numbers at very early times when gas cloud collisions and hierarchical merging events were frequent and energetic \citep[cf. the discussions by][]{p1, r1}."761". In this scenario, and should be closely correlated simply because they are both byproducts of similar extreme conditions at high-density locations during the main period of galaxy formation."," In this scenario, and should be closely correlated simply because they are both byproducts of similar extreme conditions at high-density locations during the main period of galaxy formation."762 A less certain issue would be whether or not there is a direct causal link between them., A less certain issue would be whether or not there is a direct causal link between them.763" For example, the early energy output from the dense, violent star formation conditions in the central regions of a large galaxy that might have accompanied the buildup of the massive BH might also have helped stimulate the formation of dense, massive star clusters throughout the bulge and halo regions (e.g.Harrisal.2005;Kauffmannet 2007)."," For example, the early energy output from the dense, violent star formation conditions in the central regions of a large galaxy that might have accompanied the buildup of the massive BH might also have helped stimulate the formation of dense, massive star clusters throughout the bulge and halo regions \citep[e.g.][]{h9,tan99,kra02,spr05,kau07}."764". The rough equality of total in the two subsystems as shown above may, however, be largely a coincidence."," The rough equality of total in the two subsystems as shown above may, however, be largely a coincidence."765 The reason is that it applies only to the stage of evolution of both subsystems., The reason is that it applies only to the stage of evolution of both subsystems.766" The total mass in the globular cluster system has been continually decreasing since their formation epoch: during the first several ~10’ years of a young massive star cluster's life, it sheds a high fraction of its initial protocluster gas due to SNe and stellar winds, and the highest-mass stars evolve and disappear."," The total mass in the globular cluster system has been continually decreasing since their formation epoch: during the first several $\sim 10^7$ years of a young massive star cluster's life, it sheds a high fraction of its initial protocluster gas due to SNe and stellar winds, and the highest-mass stars evolve and disappear."767 Over the, Over the768There are a variety of astrophysical problems which require both direct. imagery and velocity imagery (ic. images in small intervals of racial velocity) of [aint emission line regions in the close vicinity of clominantly bright. sources.,There are a variety of astrophysical problems which require both direct imagery and velocity imagery (i.e. images in small intervals of radial velocity) of faint emission line regions in the close vicinity of dominantly bright sources.769 For instance. the nebulosities surrounding luminous blue variables (LBVs) are of considerable interest for they are the relics of the most recent eruüptions of these stars.," For instance, the nebulosities surrounding luminous blue variables (LBVs) are of considerable interest for they are the relics of the most recent eruptions of these stars."770 Expanding shells of circumstellar gas have now been found. around. six ealactic LBVs (see Barlow ct al 1994 and Nota et al 1995 for a summary of these observations)., Expanding shells of circumstellar gas have now been found around six galactic LBVs (see Barlow et al 1994 and Nota et al 1995 for a summary of these observations).771 So far. two distinctly clilferent shells have been found (Barlowetal1994). with occultingmask imagerv around P Cvgni (V = 4.8 mag).," So far, two distinctly different shells have been found \cite{ba94} with occulting–mask imagery around P Cygni (V = 4.8 mag)."772 A bright inner shell. zz. 22 aresec diameter. has a racial expansion velocity of 140 iin the line but only 110 iin the exceptionally bright lines.," A bright inner shell, $\approx$ 22 arcsec diameter, has a radial expansion velocity of 140 in the line but only 110 in the exceptionally bright lines."773 A lainter. outer. cemitting shell of zz 1.6 aremin diameter has been shown to » expanding at 185 to give a kinematical age of 2100 vr for a cistance of 1.5 kpe to P Cveni (Meaburnctal1996).," A fainter, outer, emitting shell of $\approx$ 1.6 arcmin diameter has been shown to be expanding at 185 to give a kinematical age of 2100 yr for a distance of 1.8 kpc to P Cygni \cite{me96}."774. The Manchester Echelle spectrometer. MISS (Meaburnal. 1984)... in its imaging mode. and. with an occulting strip in its focal plane. was used for this initial imagery of P Cyent," The Manchester Echelle spectrometer, MES \cite{me84}, in its imaging mode, and with an occulting strip in its focal plane, was used for this initial imagery of P Cygni."775 In this auxiliary mode. MES has a very restrictecl fickofview (1.9 arcmin 1.5r arcmin) on the Isaac Newton 2.5-m telescope consequenthy any nebular ejecta. from P Cveni of larger angular diameter. would remain undetected.," In this auxiliary mode, MES has a very restricted field–of–view (1.9 arcmin $\times$ 1.5 arcmin) on the Isaac Newton 2.5-m telescope consequently any nebular ejecta from P Cygni of larger angular diameter would remain undetected."776 Phe Manchester occulting mask imager (MOXMLI) has now been manufactured: to overcome this restriction., The Manchester occulting mask imager (MOMI) has now been manufactured to overcome this restriction.777 This is a device dedicated to occulting mask imagery and has had its first use on the Nordic Optical telescope (NOT) where iimages of the environs of P Cveni have been obtained over a field area of unprecedented. size., This is a device dedicated to occulting mask imagery and has had its first use on the Nordic Optical telescope (NOT) where images of the environs of P Cygni have been obtained over a field area of unprecedented size.778 The optical lavout of MOMILE at the RitcheyChretien focus of the NOT telescope. is shown in Fie.," The optical layout of MOMI, at the Ritchey–Chretien focus of the NOT telescope, is shown in Fig."779 1., 1.780 The light is collimated for passage through a narrowband interference ilter. centred on a nebular emission line ancl placed in the oipil.," The light is collimated for passage through a narrow–band interference filter, centred on a nebular emission line and placed in the pupil."781" The field is reimagecl on to the ""science CCD.", The field is re–imaged on to the `science' CCD.782 A oessure stepped. optically contacted. FabryPerot etalon can also be included. just before the filter in. the pupil o permit the option of obtaining spatial/radial velocity data “cubes” of emission line regions around dominantlv xieht central sources.," A pressure stepped, optically contacted, Fabry–Perot etalon can also be included just before the filter in the pupil to permit the option of obtaining spatial/radial velocity data `cubes' of emission line regions around dominantly bright central sources."783 Ες option has not vet been used in the present instrument. but was built into. and. proven astrophysicallv. in the forerunner (Aleaburn&White of MOMLI.," This option has not yet been used in the present instrument but was built into, and proven astrophysically, in the forerunner \cite{me82} of MOMI."784 The most critical aspect of the design. is that. the chromium occulting mask. = 4 aresee on the sky. is on the [first surface of the optical chain.," The most critical aspect of the design is that the chromium occulting mask, $\equiv$ 4 arcsec on the sky, is on the first surface of the optical chain."785 With this, With this786In Fig.,In Fig.787" 6 we plot the indices S for this star, for which we obtained a mean Mount Wilson index (9)—1.864 from 1978 to 1995."," \ref{fig.hd22468} we plot the indices $S$ for this star, for which we obtained a mean Mount Wilson index $\langle S \rangle$ =1.864 from 1978 to 1995."788 Cerruti-Solaetal.(1992) reported that HD 22468 had fluxes comparable to those of the brightest solar regions and even approaching those observed in solar flares., \cite{1992A&A...256..185C} reported that HD 22468 had fluxes comparable to those of the brightest solar regions and even approaching those observed in solar flares.789" In fact, several flares were observed with IUE, and can be noticed in Fig. 6.."," In fact, several flares were observed with IUE, and can be noticed in Fig. \ref{fig.hd22468}."790" In particular, reported a flare on 1981, October 3rd. which released a total energy of ~ 103? erg."," In particular, \cite{1989A&A...211..173L} reported a flare on 1981, October $^{\textrm{rd}}$, which released a total energy of $\sim$ $^{32}$ erg."791" Using our calibration, we obtain that S increased in a few hours during that event."," Using our calibration, we obtain that $S$ increased in a few hours during that event."792" Brownetal.(1994) reported another flare on 1994, August 24th."," \cite{1994AAS...185.8519B} reported another flare on 1994, August $^{\textrm{th}}$."793" During that date, the Mount Wilson index presented a variation in six hours."," During that date, the Mount Wilson index presented a variation in six hours."794" In December 1992, Neff(1995) also reported a flare observed by IUE, during which S increased by"," In December 1992, \cite{1995psu..rept.....N} also reported a flare observed by IUE, during which $S$ increased by"795Jo calculate. the power spectrum as a function of redshift) anc cosmological parameters we usec the }iololt.2mannzm: equation: codeeo ‘CMDoeasy.IB2aS. (Doran.2005)i2005).,To calculate the power spectrum as a function of redshift and cosmological parameters we used the Boltzmann equation code CMBeasy \citep{cmbeasy}.796.. UsingNn18 (wowo sitded.sided derivatisderivatives togethergether wiwith convergence tests. we Cian accurately calculatecalcula theh sensitivity.prhivit; onversenee.derivatives.lerivaIOS WIwith respect to cach parameter., Using two sided derivatives together with convergence tests we can accurately calculate the sensitivity derivatives with respect to each parameter.797 Vhese then enter into the Fisher . calculations.: of∙ the parameter estimation.:; taking. into matrixaccount the correlations between parameters.," These then enter into the Fisher matrix calculations of the parameter estimation, taking into account the correlations between parameters."798 The data points are taken to be the power spectrum. evaluated: at the centers of LO (or 11) redshift bins from.=O 2. aal z;=O27|0.1.," The data points are taken to be the power spectrum evaluated at the centers of 10 (or 11) redshift bins from $z=0-2$ , at $z_i=0.2i+0.1$."799 For JDIEALPS. we divide the bin containing +=0.7 into two pieces: z=0.6.0.7] using LRG and 2=0.7.0.8} using EL.," For JDEM-PS, we divide the bin containing $z = 0.7$ into two pieces: $z=[0.6,0.7]$ using LRG and $z=[0.7,0.8]$ using EL."800 The parameter set involves 9 parameters., The parameter set involves 9 parameters.801 Note that when testing the gravitational framework. eexploring bevond-EinsteinYOY-Einsted gravity∙ through€ quantitativeκlative estimationestiall of ~ it EEis crucialrucial to include include:all parametersarameters that could a actin inaz similar manner on the growth and growth rate.," Note that when testing the gravitational framework, exploring beyond-Einstein gravity through quantitative estimation of $\gamma$, it is crucial to include all parameters that could act in a similar manner on the growth and growth rate."802 Pherefore we include a time varving dark energy equation of state wa)—wopesL)e) and massive neutrinos., Therefore we include a time varying dark energy equation of state $w(a)=w_0+w_a(1-a)$ and massive neutrinos.803" The parameter list. and the fiducial value around which theFisher matrix expands. is The values for 256. w, and f are NMADP-5 best. fit parameters (Llinshawetal.2009)."," The parameter list, and the fiducial value around which theFisher matrix expands, is The values for $\Omega_{DE}$, $\omega_b$ and $h$ are WMAP-5 best fit parameters \citep{wmap}."804. Note that the fiducial ~=0.55 is the value. predicted. by €ieneral Relativity for ACDAL (and is quite insensitive to the dark. energy equation of state): the fiducial wy=0.99 is taken to avoid issues of stepping over w=1., Note that the fiducial $\gamma=0.55$ is the value predicted by General Relativity for $\Lambda$ CDM (and is quite insensitive to the dark energy equation of state); the fiducial $w_0=-0.99$ is taken to avoid issues of stepping over $w=-1$.805 Dark energy. perturbations are included in CMDoeasy., Dark energy perturbations are included in CMBeasy.806. We assume there is no spatial curvature., We assume there is no spatial curvature.807" In. the remainder of this section we take the fiducial be,=0.8. and we will investigate the effect. of a different fiducial in the next section."," In the remainder of this section we take the fiducial $b_{EL} = 0.8$, and we will investigate the effect of a different fiducial in the next section."808" Note that the neutrino enerey density [raction is related to the sum of the neutrino masses hy Q,h=Mmy,/94eV.", Note that the neutrino energy density fraction is related to the sum of the neutrino masses by $\Omega_{\nu} h^2 =\sum m_{\nu}/\textrm{94 eV}$.809" Por a reasonable current upper bound Nn,ες0.8 eV (Seljaketal.2009).. this implies ον< 0.006."," For a reasonable current upper bound $\sum m_{\nu} \leq 0.3$ eV \citep{nubound}, this implies $\Omega_{\nu} \leq$ 0.006."810" We take ον=0.002. or Mmy,=0.1 eV as the fiducial."," We take $\Omega_{\nu} = 0.002$, or $\sum m_\nu=0.1$ eV as the fiducial."811 Adding together. the information from the. recdshift slices. independently (note this is not generally a good approximation [or slices thinner than our As=0.2). we obtain the full Fisher matrix.," Adding together the information from the redshift slices independently (note this is not generally a good approximation for slices thinner than our $\Delta z=0.2$ ), we obtain the full Fisher matrix."812 We do not explicitly add any CAB information (except later when adding Stage LLL Fisher matrices. which assume Planck data).," We do not explicitly add any CMB information (except later when adding Stage III Fisher matrices, which assume Planck data)."813 Concentrating on testing the gravitational growth index. we now explore in more detail what allects the constraints on  using information only from the galaxy power spectrum.," Concentrating on testing the gravitational growth index, we now explore in more detail what affects the constraints on $\gamma$ using information only from the galaxy power spectrum."814 The constraints are computed to be =0.048005., The constraints are computed to be = 0.043.815 The importance of including dark cnerey properties. neutrino; masses. and other cosmological; parameters in; the parameter estimation: is: highlighted.: by themuch tighter: constraints: obtained: ifτῇ we neglect theirmn influence. including: only * itself and the galaxy biases.," The importance of including dark energy properties, neutrino masses, and other cosmological parameters in the parameter estimation is highlighted by themuch tighter constraints obtained if we neglect their influence, including only $\gamma$ itself and the galaxy biases."816 In this case we obtain overly optimistic estimates: G(5)panous=0.0006 and alJapeprs=0.0078., In this case we obtain overly optimistic estimates: $\sigma(\gamma)_{\rm BigBOSS} = 0.0096$ and $\sigma(\gamma)_{\rm JDEM-PS} = 0.0078$.817 Thus. taking into account the correlations with other cosmological parameters is essential.," Thus, taking into account the correlations with other cosmological parameters is essential."818 The correlation matrices for the two experiments are shown in Tables 2. and 3: we have replaced the unit diagonal with the uncertainties 7 on cach parameter., The correlation matrices for the two experiments are shown in Tables \ref{cov1} and \ref{cov2}; we have replaced the unit diagonal with the uncertainties $\sigma_i$ on each parameter.819 To obtain an overall view of how tightly correlated a parameter is with the other variables. we employ. the elobal correlation coelficient the largest. correlation of that parameter with. any linear. combinationη. of⋅ all other parameters.," To obtain an overall view of how tightly correlated a parameter is with the other variables, we employ the global correlation coefficient – the largest correlation of that parameter with any linear combination of all other parameters."820" ""his is ⋠given by n= .", This is given by r_i =.821(22)We\ show those vectors in. Table +.., We show those vectors in Table \ref{global}.822 Note. the high degree of correlation. indicating the importance of crosschecks by other data and techniques.," Note the high degree of correlation, indicating the importance of crosschecks by other data and techniques."823 Examining the mareinalizecl parameter. estimations along the diagonals of Tables 20 and 3. we see that as expected. the power spectrum information is especially strong in constraining Qo and h.," Examining the marginalized parameter estimations along the diagonals of Tables \ref{cov1} and \ref{cov2}, , we see that as expected the power spectrum information is especially strong in constraining $\Omega_{DE}$ and $h$."824" One can determine at the 1U'À level the growth index 7 and present of equationstate wy. While we and QO, have uncertainties of order unity."," One can determine at the $\sim10\%$ level the growth index $\gamma$ and present equation of state $w_0$, while $w_a$ and $\Omega_\nu$ have uncertainties of order unity."825" The index and. equation of state parameters estimation is growsimilarh for (he Iwo experunentss 5c.0.55x 0.045. wy=0.99c0.16 and w,=O0cO47 for the ground-based BigBOSS and 5=0.5540.054. «y=0.99+0.14 and dw,=0d0.37 for the space-based JDIZM."," The growth index and equation of state parameters estimation is similar for the two experiments: $\gamma=0.55\pm0.043$ , $w_0=-0.99\pm0.16$ and $w_a=0\pm0.47$ for the ground-based BigBOSS and $\gamma=0.55\pm0.054$, $w_0=-0.99\pm0.14$ and $w_a=0\pm0.37$ for the space-based JDEM."826 We find the usual high anti-correlation between wy and αμ. and astrong correlation between 5 and Cig. uw).," We find the usual high anti-correlation between $w_0$ and $w_a$ , and astrong correlation between $\gamma$ and $(w_0,w_a)$ ."827 Regarding the neutrino mass parameter. neutrino oscillation experiments indicate that neutrinos do have mass (Maltonietal.2008:Ixavser 2008).. but this is not always included. in parameterestimation despite its correlations.," Regarding the neutrino mass parameter, neutrino oscillation experiments indicate that neutrinos do have mass \citep{numass,numass2}, , but this is not always included in parameterestimation despite its correlations."828 We demonstrate the elfect. of neglecting thisingredient. finding that it gives overly optimistic constraints on5 by," We demonstrate the effect of neglecting thisingredient, finding that it gives overly optimistic constraints on$\gamma$ by"829ancl variability of these delays are crucial lor eaining a quantitative understanding of the outflow process.,and variability of these delays are crucial for gaining a quantitative understanding of the outflow process.830 The IUXEXN-600 observations were carried. oul as a part of à monitoring program of microquasars to study their Hare activity across a broad. frequency. range., The RATAN-600 observations were carried out as a part of a monitoring program of microquasars to study their flare activity across a broad frequency range.831 Observations were performed at 3.9. 7.7. 11.2 and 21.7 Gllz: the 311 CGllz data are presented in the top panel of Fig 1.," Observations were performed at 3.9, 7.7, 11.2 and 21.7 GHz; the 3–11 GHz data are presented in the top panel of Fig 1."832 The Hux density calibration was performed using observations of SC286 (1328|30) and PIxX8S1345|12., The flux density calibration was performed using observations of 3C286 (1328+30) and PKS1345+12.833 Although interference sometimes prevented realisation of the maximum sensitivity of the radiometers. daily observations of reference sources indicate that the error in the Dux density measurements for 1915|10 cid not exceed at 2.3. 3.9. and 11.2 Cilz and at 21.7 Cllz.," Although interference sometimes prevented realisation of the maximum sensitivity of the radiometers, daily observations of reference sources indicate that the error in the flux density measurements for 1915+10 did not exceed at 2.3, 3.9, and 11.2 GHz and at 21.7 GHz."834 For further details see Ειςκα. Majorova Dursov (2001).," For further details see Trushkin, Majorova Bursov (2001)."835 We have made use of public data from the Rossi X-ray ‘Timing Experiment (NTIZ) All-Sky. Monitor (ASAL: Levine et al., We have made use of public data from the Rossi X-ray Timing Experiment (XTE) All-Sky Monitor (ASM; Levine et al.836 1996)., 1996).837 These data are available atxte., These data are available at.838mit.edu. The total intensity and ΗΛ X-ray colour (ratio of counts in 512 to 35 keV bands) are presented in Fig 1., The total intensity and HR2 X-ray colour (ratio of counts in 5–12 to 3–5 keV bands) are presented in Fig 1.839 We have observed GRS 1915)105 with the Australia ‘Telescope Compact Array (ATCA) for approximately six yours on 2000 July 26., We have observed GRS 1915+105 with the Australia Telescope Compact Array (ATCA) for approximately six hours on 2000 July 26.840 The override observations were rigecred as a result of the radio [lare observed. by the UNEAN-600 telescope around ALJD 51750., The override observations were triggered as a result of the radio flare observed by the RATAN-600 telescope around MJD 51750.841 At this time the array was ina compact configuration., At this time the array was in a compact configuration.842 As a result. in order to reduce the ellects of field sources (see es. Chaty et al.," As a result, in order to reduce the effects of field sources (see e.g. Chaty et al."843 2001 or radio images of the field) we only used. interferometer xselines z2200m at 4800 MlIIz., 2001 for radio images of the field) we only used interferometer baselines $\geq 2200$ m at 4800 MHz.844 Phe compact configuration ancl relatively poor hour-angle coverage precluced any attempt to confidently calibrate Stokes V. and so we were not able to make a circular polarisation measurement.," The compact configuration and relatively poor hour-angle coverage precluded any attempt to confidently calibrate Stokes V, and so we were not able to make a circular polarisation measurement."845 In Fig 2 we present clual-frequeney total intensity. spectral index and linear polarisation measuremoents of GRS 1915]105 obtained during the ATCA run.," In Fig 2 we present dual-frequency total intensity, spectral index and linear polarisation measurements of GRS 1915+105 obtained during the ATCA run."846 Care has been taken to check the reality of the low-level linear polarisation (ic., Care has been taken to check the reality of the low-level linear polarisation (ie.847 Stokes Q. U) measurements. and we are confident that those presented. in Fig 2 are realistic.," Stokes Q, U) measurements, and we are confident that those presented in Fig 2 are realistic."848 Mapping the data in linear polarisation produces results consistent with averaging the data presented in Fig 2., Mapping the data in linear polarisation produces results consistent with averaging the data presented in Fig 2.849 Note however that for the last few scans the polarisation calibration solutions were not good and those data are not used., Note however that for the last few scans the polarisation calibration solutions were not good and those data are not used.850 The 40-day light curves presented in Fig 1 clearly reveal the existence of radio Ilaring and dramatically varying X-ray Hux and hardness., The 40-day light curves presented in Fig 1 clearly reveal the existence of radio flaring and dramatically varying X-ray flux and hardness.851 It has been previously established (e.g. Foster et al., It has been previously established (e.g. Foster et al.852 1996. Fender et al.," 1996, Fender et al."853 1999) that the prolonged (ic., 1999) that the prolonged (ie.854 more than a few davs) hard states. or ‘plateaux’. are accompanied by [iat or inverted spectrum radio emission aud appear to be always followed by an optically thin radio Iare.," more than a few days) hard states, or `plateaux', are accompanied by flat or inverted spectrum radio emission and appear to be always followed by an optically thin radio flare."855 The inverted-spectrum. emission probably corresponds to a »owerful scl-absorbedl quasi-continuous jet. (Dhawan tal., The inverted-spectrum emission probably corresponds to a powerful self-absorbed quasi-continuous jet (Dhawan et al.856 2000: Fender 20t) and the optically thin post-platcau Hares iwe been directly resolved into relativistically οΠοας components (c.g. Mirabel Itodriguez 1994: Fender tal., 2000; Fender 2001) and the optically thin post-plateau flares have been directly resolved into relativistically outflowing components (e.g. Mirabel Rodriguez 1994; Fender et al.857 1999)., 1999).858 In terms of Fig 1. we consider the plateau phase to he tween MD 51756.51762. and its resultant Dare to be the radio event which peaked around MJD 51762.," In terms of Fig 1, we consider the plateau phase to be between MJD 51756–51762, and its resultant flare to be the radio event which peaked around MJD 51762."859 What is less clear is whether the Dare of ~ ALJD 51750.0 was related to his subsequent plateau or to the shorter X-ray. hard state around ALJD 5174651747., What is less clear is whether the flare of $\sim$ MJD 51750.0 was related to this subsequent plateau or to the shorter X-ray hard state around MJD 51746–51747.860 Lt is the decay of this Hare. and he subsequent. emergence of core-oscillation events. which were studied in detail with our ALCA observations.," It is the decay of this flare, and the subsequent emergence of core-oscillation events, which were studied in detail with our ATCA observations."861 Note hat there appears to be a third. smaller Hare. associated with a drop in the X-ray (lux around. MJD 5176s., Note that there appears to be a third smaller flare associated with a drop in the X-ray flux around MJD 51768.862 The characteristic decay. ancl optically thin spectral index, The characteristic decay and optically thin spectral index863that the population of galaxies luminous in the PLR is subject to strong evolution.,that the population of galaxies luminous in the FIR is subject to strong evolution.864 The redshift clistribution of HAS galaxies (e.g. Saunders et al., The redshift distribution of IRAS galaxies (e.g. Saunders et al.865 1990) is consistent with the density evolution ~(1|2)22? for z<0.25., 1990) is consistent with the density evolution $\sim (1+z)^{6.7\pm 2.3}$ for $z < 0.25$.866 are not to direct information on the FII evolutionSamples at higher deepredshits. enoughbut it is providegenerally accepted that starburst which are the constituent of the FUR galaxies. exhibit a strong ealaxies.increase of the majordensity and/or luminosity to z2.5 with a decay at higher redshifts spatial(e.g. Moorwood 1996. Franceschiniup et al.," Samples are not deep enough to provide direct information on the FIR evolution at higher redshits, but it is generally accepted that starburst galaxies, which are the major constituent of the FIR galaxies, exhibit a strong increase of the spatial density and/or luminosity up to $z \approx 2.5$ with a decay at higher redshifts (e.g. Moorwood 1996, Franceschini et al."867 1997. Jochtold οἱ al.," 1997, Bechtold et al."868 1998)., 1998).869 In agreement with the available data we have used the following parametrization of the evolution of the FUR bluminosities: where we ον2.5. α=1.25.," In agreement with the available data we have used the following parametrization of the evolution of the FIR luminosities: where we put $z_{\star} = 2.5$, $a = 1.25$."870 We have considerc two rates of evolution: 3=put2.5 and 3.0., We have considerd two rates of evolution: $\beta = 2.5$ and $3.0$.871 The XRB and GRB by galaxies to these evolution models is shown in with, The XRB and GRB produced by galaxies subjected to these evolution models is shown in Fig.872 produced.open squares [or subjected7=2.5 and with full for 7=3.0., 1 with open squares for $\beta = 2.5$ and with full squares for $\beta = 3.0$.873 Lt is Fig.conspicuous1. that at above ~5 MeV. the squaresmodel with ο=3.0 produces 20 - 25% of the energies observed.background., It is conspicuous that at energies above $\sim 5$ MeV the model with $\beta = 3.0$ produces 20 - 25 % of the observed background.874 4., \bsk ∋ 4.875 PROSPECTS FOR THE FUTURI ]ndirect methods. which use objects selected at one energy band. to the at another and correlations between predictIuminosities in these backgroundtwo bands are not energycapable to exploit accurate estimates of the background.," PROSPECTS FOR THE FUTURE \ssk ∋ Indirect methods, which use objects selected at one energy band to predict the background at another energy and exploit correlations between luminosities in these two bands are not capable to provide accurate estimates of the background."876 Studies of the soft. X-ray backgroundprovide in the last 20 Clearly show that., Studies of the soft X-ray background in the last 20 years clearly show that.877 However. the calculations demonstrate that vearsthe FIR. galaxies potentially are an importantpresent. constituent of the CIUS.," However, the present calculations demonstrate that the FIR galaxies potentially are an important constituent of the GRB."878the central stars ancl elsewhere (Ganunie1996).,the central stars and elsewhere \citep{Gammie1996}.879. The depth of lavers which are affected bv the MRI turbulence is determined by an ionization equilibrium in which the ionization rate is balanced by the rate of recombination on the grains., The depth of layers which are affected by the MRI turbulence is determined by an ionization equilibrium in which the ionization rate is balanced by the rate of recombination on the grains.880 Interior to ej. the sublimation ol volatile ices eliminates a majority of the recombination sites. enhances the ionization fraction of the disk gas. and increases the depth of the surface laver which is regulated by MARI turbulence.," Interior to $a_{\rm ice}$, the sublimation of volatile ices eliminates a majority of the recombination sites, enhances the ionization fraction of the disk gas, and increases the depth of the surface layer which is regulated by MRI turbulence."881" Across ej. elfective viscosity has a negative gradient. which in quasi steady state. leads (o a positive gradient in X, and 7/7."," Across $a_{\rm ice}$, effective viscosity has a negative gradient, which in quasi steady state, leads to a positive gradient in $\Sigma_g$ and $P_g$."882" For a range of disk mass accretion rates (Mj~10""—10SAL. vr P). the local P; maxinnum at dice provides a natural barrier [or the orbital decay of grains against hydrodynamic drag (INretke&Lin2007)."," For a range of disk mass accretion rates $\dot M_d \sim 10^{-9} - 10^{-8}883M_\odot$ $^{-1}$ ), the local $P_g$ maximum at $a_{\rm ice}$ provides a natural barrier for the orbital decay of grains against hydrodynamic drag \citep{Kretke2007}."884. This structure also reverses (he direction of tvpe I migration due to the effect of horseshoe drag; 2006)., This structure also reverses the direction of type I migration due to the effect of horseshoe drag \citep{Masset2006}.885. Outward migration of embrvos by a few AU may also be induced by an entropy eradient (Paardekooper2010)., Outward migration of embryos by a few AU may also be induced by an entropy gradient \citep{Paarde2010}.886". Thus. similar to grains. orbital evolution. embryos: migration leads to them to congregate near di Where both the gas surface density (X,) ancl pressure (77) attain local maxima in the disk due to (he transition in the efficient turbulent angular momentum transport."," Thus, similar to grains' orbital evolution, embryos' migration leads to them to congregate near $a_{\rm ice}$ where both the gas surface density $\Sigma_g$ ) and pressure $P_g$ ) attain local maxima in the disk due to the transition in the efficient turbulent angular momentum transport."887 Through the accumulation of neighboring embrvos and planetesimals a protoplanet may acquire sufficient mass ~ M) to evolve into a core which efficiently. accretes σας 2003h)., Through the accumulation of neighboring embryos and planetesimals a protoplanet may acquire sufficient mass $\sim M_{\rm crit}$ ) to evolve into a core which efficiently accretes gas \citep{Ida2008b}.888. During the early phase of gas accretion. the Gime scale for a proto-planet to double its mass is longer (han the svnodic periods (relative to its orbit) of planetesimals and embryos within its feeding zone.," During the early phase of gas accretion, the time scale for a proto-planet to double its mass is longer than the synodic periods (relative to its orbit) of planetesimals and embryos within its feeding zone."889 Its (dal perturbation on the nearby embrvos aud (heir interaction with the disk gas induces their orbit to evolve away [rom that of the dominant proto-planet (Zhou&Lin2007:ShiraishiIda2008).," Its tidal perturbation on the nearby embryos and their interaction with the disk gas induces their orbit to evolve away from that of the dominant proto-planet \citep{Zhou2007, ShiraishiIda2008}."890. Therefore slowly (vacliabatically”) growing planets cannot accrete (hose embryos within their expanding feeding zones in contrast to the assumption made in previous models (Pollacketal.1996:Docdson-Robinson2003).," Therefore slowly (“adiabatically”) growing proto-planets cannot accrete those embryos within their expanding feeding zones in contrast to the assumption made in previous models \citep{Pollack1996,891DodsonRobinson2008}."892. This iniüal suppression of giant impacts reduces (he energv released by (he growing proto-planet and promotes the gas accretion rate., This initial suppression of giant impacts reduces the energy released by the growing proto-planet and promotes the gas accretion rate.893" Ii the lanit of neelieibly small thermal feedback and weak proto-planetary torque. proto-planets. uninhibited gas accretion rate (AL,) mav be approximated by the Dondi formula such that the mass doubling Gime scale would be where // is the disk thickness."," In the limit of negligibly small thermal feedback and weak proto-planetary torque, proto-planets' uninhibited gas accretion rate $\dot M_p$ ) may be approximated by the Bondi formula such that the mass doubling time scale would be where $H$ is the disk thickness."894The prolonged period of minimal solar activity. ancl the delaved onset of solar evele 24 have invoked a great deal of interest among solar physicists to identilv precisely when the,The prolonged period of minimal solar activity and the delayed onset of solar cycle 24 have invoked a great deal of interest among solar physicists to identify precisely when the895apparent fulbwidth at lalfanaxiuaun (FWITIMD of each knot is ~OTS while the PSP has FWHAL =07036.,apparent full-width at half-maximum (FWHM) of each knot is $\sim 0\farcs18$ while the PSF has FWHM $= 0\farcs036$.896 Correcting the apparent knot diameters for the width of the PSF eives intrinsic FEWIIMS of ~0717 (~1.51 mrc} for the knots., Correcting the apparent knot diameters for the width of the PSF gives intrinsic FWHMs of $\sim 0\farcs17$ $\sim 1.5$ kpc) for the knots.897 This is consistent with the knots containing clubedded star-forming regious., This is consistent with the knots containing embedded star-forming regions.898 snot Xll is parΙαν covered by the OT and not clearly visible uutil the OT is subtracted., 	Knot 1 is partially covered by the OT and not clearly visible until the OT is subtracted.899 We believe that this knot is a real feature aud not au artifact of tle PSF subtraction since siniar features are not seen when the PSF is used to subtractt stars from the nuages., We believe that this knot is a real feature and not an artifact of the PSF subtraction since similar features are not seen when the PSF is used to subtract stars from the images.900 Fie., Fig.901 2 shows no svsteniatic change in iutensity between the northeas half of the knot. which was not obscured bv the OT. and the soutlsvest half which was OVSCULCE w theOT.," \ref{FIGURE:host} shows no systematic change in intensity between the northeast half of the knot, which was not obscured by the OT, and the southwest half, which was obscured by the."902. This also suggests that the knot is not an artifact of the PSF subtraction., This also suggests that the knot is not an artifact of the PSF subtraction.903 To test the ability of PSF subtraction to reveal structure under the OT we generated a series of artificial stars with the same magnitude as theOT., To test the ability of PSF subtraction to reveal structure under the OT we generated a series of artificial stars with the same magnitude as the.904. These stars were put ou the three knots. aud iu empty parts of the nuage near he OT. then PSFs were fit iud subtracted for cach artificial star.," These stars were put on the three knots, and in empty parts of the image near the OT, then PSFs were fit and subtracted for each artificial star."905 We found that he knots were clearly visible after the artificial stars were removed. although the knots located under artificial stars appeared less centrally concentrated after the artificial stars were subtracted.," We found that the knots were clearly visible after the artificial stars were removed, although the knots located under artificial stars appeared less centrally concentrated after the artificial stars were subtracted."906 Comparing the recovered magnitudes of the isolated artificial stars with those of the artificial stars situated ou the knots suggests that we are able to detect knots with Vy<29 that are located uuder theOT., Comparing the recovered magnitudes of the isolated artificial stars with those of the artificial stars situated on the knots suggests that we are able to detect knots with $V_0 < 29$ that are located under the.907. The OT for is located on the southeast edee of knot #11., 	The OT for is located on the southeast edge of knot 1.908 This knot is the most likely source of the metallic absorption lines seen iu 16 spectra of the OT (Andersen ct citeAC'99:: Kulkarni et cite D99)). although we cau not 11le out the possibility that the absorption lines are chο to a ναν πα. undetected. faint knot located xder theOT.," This knot is the most likely source of the metallic absorption lines seen in the spectra of the OT (Andersen et \\cite{AC99}; Kulkarni et \\cite{KD99}) ), although we can not rule out the possibility that the absorption lines are due to a very small, undetected, faint knot located under the."909. Such absorption systems are often associated with high column deusities of hydrogen. which iu tur1 are associated with star formation.," Such absorption systems are often associated with high column densities of hydrogen, which in turn are associated with star formation."910 We used Eq., We used Eq.911 2 of ATadau et ((1998)) to estimate the star-formation rate (SER) in cach of theknots listed in Table |. assuming a flat (3= 0) spectrum.," 2 of Madau et \cite{MP98}) ) to estimate the star-formation rate (SFR) in each of theknots listed in Table \ref{TABLE:knots}, assuming a flat $\beta = 0$ ) spectrum."912 For 2=1.6 a vest-frame waveleneth of Ay=2800 corresponds to an observed waveleιο of A=7280A. so we computed the observed flax at A=7280 by extrapolating between the V- aud R-baud fluxes.," For $z = 1.6$ a rest-frame wavelength of $\lambda_0 = 2800$ corresponds to an observed wavelength of $\lambda = 7280$, so we computed the observed flux at $\lambda =9137280$ by extrapolating between the $V$ - and $R$ -band fluxes."914 These were computed for cach knot in the same manucr as was done for the host ealaxv (see Sect. 3))., These were computed for each knot in the same manner as was done for the host galaxy (see Sect. \ref{SECTION:magnitudes}) ).915 The estimated SERs for each knot. assuming a Salpeter initial mass function. are listed in Table 2..," The estimated SFRs for each knot, assuming a Salpeter initial mass function, are listed in Table \ref{TABLE:knot_data}."916 For a Scalo initial mass function multiply the SFR by 1.5»., For a Scalo initial mass function multiply the SFR by 1.55.917" The uncertainty iu each SER is ~0.03 At: ο, ", The uncertainty in each SFR is $\sim 0.03$ ${\cal M}_{\sun}$ $^{-1}$.918These SFRs assume that there is no dust. or obscured star formation. i the knots.," These SFRs assume that there is no dust, or obscured star formation, in the knots."919 This is probably a poor assuniptiou if the knots are star-formune regions., This is probably a poor assumption if the knots are star-forming regions.920 Therefore. our derived SERs should be considered a lower limit ou the true SER in each kuot.," Therefore, our derived SFRs should be considered a lower limit on the true SFR in each knot."921 The OT is located at a projected distance of 0765 (=5.6 kpc) from he nucleus of the los ealaxy.," 	The OT is located at a projected distance of $0\farcs65$ $=9225.6$ kpc) from the nucleus of the host galaxy."923 Bloom et ((1998)) calclated that ~50% of the CRBs from NS-NS and BS-NS progenitors in galaxies with a shallow evavitational potevial will occur witli1 5 kpe of the nucleus of the host ealaxy aud ~90% will occur witlin 30 kpc.," Bloom et \cite{BS98}) ) calculated that $\sim 50$ of the GRBs from NS-NS and BS-NS progenitors in galaxies with a shallow gravitational potential will occur within 5 kpc of the nucleus of the host galaxy and $\sim92490$ will occur within 30 kpc."925 This is consisteut with the location ο‘the OT relative to the nucleus of he host galaxy., This is consistent with the location of the OT relative to the nucleus of the host galaxy.926" If the progeuitor is a failed superuova. or a livormova, then the OT should be located within a few Iunudred parsecs of he star-forming reeion (Paczviisski 1998))."," If the progenitor is a failed supernova, or a hypernova, then the OT should be located within a few hundred parsecs of the star-forming region (Paczyńsski \cite{P98}) )."927 The OT is located at a projected distance of 0715 (=1.3 kpc) southeast of the ceutre of the nearest knot ($411). which is larger than the expected separation if the GBD was due to the explosion of a dnassive star. vet consistent with the NS-NS and BS-NS hypotheses.," The OT is located at a projected distance of $0\farcs15$ $= 1.3$ kpc) southeast of the centre of the nearest knot 1), which is larger than the expected separation if the GRB was due to the explosion of a massive star, yet consistent with the NS-NS and BS-NS hypotheses."928 We wish to stress tiat this conclusion depends ou there being uo faint star-forming region directly mucder the OT., We wish to stress that this conclusion depends on there being no faint star-forming region directly under the .929. Further observations willbe needed. after the OT has faded. to determine if there are other sinall star formation regiousOo that are curreutly hidden by the OT..," Further observations willbe needed, after the OT has faded, to determine if there are other small star formation regions that are currently hidden by the ."930"were 12 and 19 hours in 2009 and 2010 respectively, with individual observations lasting approximately 20 minutes.","were 12 and 19 hours in 2009 and 2010 respectively, with individual observations lasting approximately 20 minutes."931" The observing mode was frequency-switching, with a throw of 35.5 km/s within the passband."," The observing mode was frequency-switching, with a throw of 35.5 km/s within the passband."932" The telescope tracked Pluto's sky position but not its velocity, so spectra were shifted in software to the dwarf planet's reference frame."," The telescope tracked Pluto's sky position but not its velocity, so spectra were shifted in software to the dwarf planet's reference frame."933 The spectrometer channel spacing was set to 0.04 km/s and there are systematics of around +0.025 km/s for velocity-drifts of Pluto within each observation that have not been corrected., The spectrometer channel spacing was set to 0.04 km/s and there are systematics of around $\pm$ 0.025 km/s for velocity-drifts of Pluto within each observation that have not been corrected.934" Data are shown on the main-beam antenna temperature brightness scale, with a beam efficiency of 0.75 at 230.538 GHz frequency."," Data are shown on the main-beam antenna temperature brightness scale, with a beam efficiency of 0.75 at 230.538 GHz frequency."935" For each night of data, the shifted spectra were co-added, and then the appropriate positive and negative frequency-switched features were averaged."," For each night of data, the shifted spectra were co-added, and then the appropriate positive and negative frequency-switched features were averaged."936 The region of the net positive line was blanked and replaced with a linear interpolation of surrounding channels., The region of the net positive line was blanked and replaced with a linear interpolation of surrounding channels.937 A running average over 12 channels was then generated from this and subtracted from the pre-blanked spectrum as a baseline., A running average over 12 channels was then generated from this and subtracted from the pre-blanked spectrum as a baseline.938" The 11 nightly spectra made in this way were finally combined, weighted by 1/noise? factors to take into account their different conditions and observing durations."," The 11 nightly spectra made in this way were finally combined, weighted by $^2$ factors to take into account their different conditions and observing durations."939" Rejected data include two earlier nights using a smaller frequency-switch, and two later nights when Pluto had approached within 1.5? of the Galactic Plane."," Rejected data include two earlier nights using a smaller frequency-switch, and two later nights when Pluto had approached within $^{\circ}$ of the Galactic Plane."940" These spectra show ripples and complex baselines, degrading the final line profile."," These spectra show ripples and complex baselines, degrading the final line profile."941 The weight of these rejected data represent only 15 of the total observations made., The weight of these rejected data represent only 15 of the total observations made.942" 'The observations were planned to maximise the height of Pluto from the Galactic Plane, which varies as seen from the Earth, to minimise contamination problems."," The observations were planned to maximise the height of Pluto from the Galactic Plane, which varies as seen from the Earth, to minimise contamination problems."943" Pluto's Galactic latitude in 2009 was +2.2°, and in 2010 was -2.0? to -1.7°, with longitudes of zz+12°."," Pluto's Galactic latitude in 2009 was $^{\circ}$ , and in 2010 was $^{\circ}$ to $^{\circ}$, with longitudes of $\approx +12^{\circ}$."944" Using two periods well spaced around the year (Aug and Apr/May) also caused large changes in Pluto's velocity relative to the telescope, from 4-26 km/s in 2009 to -27 to -19 km/s in 2010."," Using two periods well spaced around the year (Aug and Apr/May) also caused large changes in Pluto's velocity relative to the telescope, from +26 km/s in 2009 to -27 to -19 km/s in 2010."945" Thus any narrow line seen persistently at Pluto's velocity in the co-added data can only be from the planetary atmosphere, as any Galactic background emission is different every night."," Thus any narrow line seen persistently at Pluto's velocity in the co-added data can only be from the planetary atmosphere, as any Galactic background emission is different every night."946" Galactic lines were seen around Pluto's velocity in the spring data, but were much broader than the Pluto line."," Galactic lines were seen around Pluto's velocity in the spring data, but were much broader than the Pluto line."947" Figure 1 shows the co-added Pluto spectrum from 2009/10, in the dwarf planet's rest-frame."," Figure 1 shows the co-added Pluto spectrum from 2009/10, in the dwarf planet's rest-frame."948"A clear line is seen at zero velocity, with 6.50 confidence in the integrated antenna temperature.","A clear line is seen at zero velocity, with $\sigma$ confidence in the integrated antenna temperature."949" There are no other significant narrow features; the next brightest (at -5 km/s), is at only 2.70 over three spectral channels, and was traced to a signal from a pair only of the 11 observing nights."," There are no other significant narrow features; the next brightest (at -5 km/s), is at only $\sigma$ over three spectral channels, and was traced to a signal from a pair only of the 11 observing nights."950" The broad features are residuals from the running-average subtraction process, arising from broad lines of Galactic clouds from different nights."," The broad features are residuals from the running-average subtraction process, arising from broad lines of Galactic clouds from different nights."951" A detail spectrum (Figure 2) shows the very flat baseline achieved near Pluto's velocity, while a test made by blanking a broader region (20 channels instead of 12) showed a change in integrated brightness of only around 10%."," A detail spectrum (Figure 2) shows the very flat baseline achieved near Pluto's velocity, while a test made by blanking a broader region (20 channels instead of 12) showed a change in integrated brightness of only around 10."952". This argues for a robust signal measurement of the narrow line seen, and also implies that broad line wings are faint."," This argues for a robust signal measurement of the narrow line seen, and also implies that broad line wings are faint."953" 'The CO line from Pluto was independently detected in both the 2009 and 2010 epochs, at levels of 3.90 and 5.40 respectively, and with the same integrated brightness within the errors (Figure 2, right panel)."," The CO line from Pluto was independently detected in both the 2009 and 2010 epochs, at levels of $\sigma$ and $\sigma$ respectively, and with the same integrated brightness within the errors (Figure 2, right panel)."954" The antenna temperatures are calibrated to about20 96, based on rms scatter of"," The antenna temperatures are calibrated to about20 , based on rms scatter of"955The light red area shows the 1-c spread around the mean.,The light red area shows the $\sigma$ spread around the mean.956 The mean profile is consistent with zero at all radii; the surface brightness profiles of quiescent galaxies at z2 seem to be well described by Sérrsic profiles., The mean profile is consistent with zero at all radii; the surface brightness profiles of quiescent galaxies at $z\sim2$ seem to be well described by Sérrsic profiles.957" On average the residual correction increases or decreases the total flux of each galaxy in our sample by only a few percent, with an upper limit of7%."," On average the residual correction increases or decreases the total flux of each galaxy in our sample by only a few percent, with an upper limit of."958. The mean contribution of the residual flux to the total flux for all galaxies in our sample is-0., The mean contribution of the residual flux to the total flux for all galaxies in our sample is.959"7%.. Thus, we do not find evidence that indicates that there is missing low surface brightness emission around compact quiescent z~2 galaxies, and we therefore conclude that the small sizes found for these galaxies are correct."," Thus, we do not find evidence that indicates that there is missing low surface brightness emission around compact quiescent $z\sim2$ galaxies, and we therefore conclude that the small sizes found for these galaxies are correct."960" In the previous Section we have shown that the surface brightness profiles of z—2 quiescent galaxies closely follow Sérrsic profiles, and that their sizes are not systematically underestimated due to a lack of sensitivity."," In the previous Section we have shown that the surface brightness profiles of $z\sim2$ quiescent galaxies closely follow Sérrsic profiles, and that their sizes are not systematically underestimated due to a lack of sensitivity."961 We now compare their size distribution and surface to those of low-redshift galaxies., We now compare their size distribution and surface brightness profiles to those of low-redshift galaxies.962" In Figure 6 we brightnessshow the profilesmass-size and magnitude-size relations for the z~2 galaxies and for low-redshift massive elliptical galaxies, taken from(2003)."," In Figure \ref{fig:masssize} we show the mass-size and magnitude-size relations for the $z\sim2$ galaxies and for low-redshift massive elliptical galaxies, taken from."963". The z2 sample has been split into two redshift bins: 1.75<z«2.5 and 1.5<z«1.75 (shown in blue and green, respectively)."," The $z\sim2$ sample has been split into two redshift bins: $1.75 < z < 2.5$ and $1.5 < z <9641.75$ (shown in blue and green, respectively)."965 The low-redshift sample is shown in grey., The low-redshift sample is shown in grey.966 Galaxies at z~2 are significantly smaller than those at ζΞ 0., Galaxies at $z\sim2$ are significantly smaller than those at $z=0$ .967" We fit a power law of the form r,«(1--z)* and find a=—0.94+0.16, which iscomparable to e.g., vanderWeletal.(2008) and vandeSandeetal.(2011),, but slightly steeper than Newmanetal.(2010) and significantly shallower than Buitragoetal.(2008)."," We fit a power law of the form $r_e \propto968(1+z)^\alpha$ and find $\alpha = -0.94\pm0.16$, which iscomparable to e.g., \cite{wel08} and \cite{san11}, but slightly steeper than \cite{new10} and significantly shallower than \cite{bui08}."969". However, the z—2 galaxies span a large range in size; some are while others are as large as z=0 galaxies."," However, the $z\sim2$ galaxies span a large range in size; some are supercompact, while others are as large as $z=0$ galaxies."970" Following supercompact,Shenetal.(2003),, we quantify this range using Which is defined as the 1-σ spread in log;r, around the Clog,,r,,median mass-size relation, which we fix to the z=0 slope."," Following \cite{she03}, we quantify this range using $\sigma_{\log_{10} r_e}$, which is defined as the $\sigma$ spread in $\log_{10} r_e$ around the median mass-size relation, which we fix to the $z=0$ slope."971" Note that we define the scatter in log;o basis, not the natural logarithm as used by Shenetal.(2003)."," Note that we define the scatter in $_{10}$ basis, not the natural logarithm as used by \cite{she03}."972". Itis equal to 0.24+ for our entire sample, while Shenetal.(2003) find values around jog,r,=0.16 for early-type galaxies at z=0.1 in the same massrange."," It is equal to $0.24\pm0.06$ for our entire sample, while \cite{she03} find values around $\sigma_{\log_{10} r_e} = 0.16$ for early-type galaxies at $z=0.1$ in the same massrange."973 The values for the two high-redshift subsamples are 0.21+0.11 at 1.5<z«1.75 and 0.19+0.07 at 1.75«z< 2.5.," The values for the two high-redshift subsamples are $0.21\pm0.11$ at $1.5 < z < 1.75$ and $0.19\pm0.07$ at $1.75 < z <9742.5$ ."975" These values are upper limits, since they"," These values are upper limits, since they"976qe evolved star VY CMa and is estimated to be at the level of ——although strongly correlated between the three bands: see 7. for details).,the evolved star VY CMa and is estimated to be at the level of (although strongly correlated between the three bands; see \citealt{Truch09} for details).977 While the maps represent the optimal weighting of the data across all spatial scales. the largest scales are less constrained due to various systematic effects. particularly because of the lack of cross-linking in the scans of this particular field.," While the maps represent the optimal weighting of the data across all spatial scales, the largest scales are less constrained due to various systematic effects, particularly because of the lack of cross-linking in the scans of this particular field."978 This can produce residual large-scale fluctuating patterns across the map., This can produce residual large-scale fluctuating patterns across the map.979 To suppress these spurious signals. all maps have been filtered to remove frequencies contributing to the large-scale noise without affecting the sources. corresponding to scales in excess of about 10 aremin (approximately the size of the detector array projected on the sky).," To suppress these spurious signals, all maps have been filtered to remove frequencies contributing to the large-scale noise without affecting the sources, corresponding to scales in excess of about 10 arcmin (approximately the size of the detector array projected on the sky)."980 This procedure. already used for the BLAST GOODS-S (BGS) maps (e.g. 2)» also explicitly sets the mean of each map to zero.," This procedure, already used for the BLAST GOODS-S (BGS) maps (e.g. \citealt{Devlin09}) ) also explicitly sets the mean of each map to zero."981We obtain Lo noise values of 27.3. 21.3 and 15.6 mJy at 250. 350. andmicron.. respectively. in the central 0.8 deg? of the observed field.,"We obtain $1\sigma$ noise values of 27.3, 21.3 and 15.6 mJy at 250, 350, and, respectively, in the central 0.8 $\rmn{deg}^2$ of the observed field."982 This is an area comparable to the BGS-Deep map., This is an area comparable to the BGS-Deep map.983 An outer. shallower region of about 0.3 deg? total area was also observed. albeit with higher noise: we calculate Lo map r.m.s.," An outer, shallower region of about 0.3 $\rmn{deg}^2$ total area was also observed, albeit with higher noise: we calculate $1\sigma$ map r.m.s."984 of 63.8. 58.6. and 39.0 mJy for this region.," of 63.8, 58.6, and 39.0 mJy for this region."985 Due to the high noise of the wider area. we will only include it in the stacking analyses where it can be appropriately weighted: source extraction will only be performed in the deeper region. to ensure robust identifications.," Due to the high noise of the wider area, we will only include it in the stacking analyses where it can be appropriately weighted; source extraction will only be performed in the deeper region, to ensure robust identifications."986 Analysis of the signal-to-noise maps shows that the observations of A3112 are not strongly dominated by confusion., Analysis of the signal-to-noise maps shows that the observations of A3112 are not strongly dominated by confusion.987 The distribution of signal-to-noise is well described by a Gaussian with a wing towards high positive values (due to bright sources in the maps). whose width only slightly exceeds pure instrumental noise.," The distribution of signal-to-noise is well described by a Gaussian with a wing towards high positive values (due to bright sources in the maps), whose width only slightly exceeds pure instrumental noise."988 We evaluate the contribution of confusion to total measured noise. measured as Mosul/7p. a82446...26%... and at 250. 350. andmicron.. respectively.," We evaluate the contribution of confusion to total measured noise, measured as $\sigma_{\rmn{conf}}/\sigma_{\rmn{map}}$, as, and at 250, 350, and, respectively."989 Individual BLAST sources are extracted from each map using a source-tinding algorithm which identifies peaks in a smoothed map produced by convolving the flux density map. weighted by he inverse of the variance. with the point spread function (PSF) of BLAST.," Individual BLAST sources are extracted from each map using a source-finding algorithm which identifies peaks in a smoothed map produced by convolving the flux density map, weighted by the inverse of the variance, with the point spread function (PSF) of BLAST."990 Peaks with a signal-to-noise ratio (SNR) of at least 3 were selected as sources. their flux then calculated as the value of the beam-convolved flux density map at the position of the peak.," Peaks with a signal-to-noise ratio (SNR) of at least 3 were selected as sources, their flux then calculated as the value of the beam-convolved flux density map at the position of the peak."991 The xositional uncertainty is calculated as in 2.. where we assume the slope of the number counts from ? and a minimum uncertainty of 5 uresec is imposed. equal to the intrinsic pointing uncertainty of he instrument.," The positional uncertainty is calculated as in \citet{Ivison07}, where we assume the slope of the number counts from \citet{Patanchon09} and a minimum uncertainty of 5 arcsec is imposed, equal to the intrinsic pointing uncertainty of the instrument."992 The positional uncertainty lies below 8. 9. and 13 aresee for 5o (tor more) sources at 250. 350. andmieron.. respectively.," The positional uncertainty lies below 8, 9, and 13 arcsec for $\sigma$ (or more) sources at 250, 350, and, respectively."993 Sources at different wavelengths are associated on the basis of their peak position. as in »2z: the source catalogues are provided in appendix.," Sources at different wavelengths are associated on the basis of their peak position, as in \citet{Devlin09}; the source catalogues are provided in appendix."994 We identify 86. 74 and 46 sources at 250. 350. andmicron.. respectively. with SNR values in excess of 3 and up to 23.," We identify 86, 74 and 46 sources at 250, 350, and, respectively, with SNR values in excess of 3 and up to 23."995 We use the results of ο to predict the number of sources expected for a blank field with the same area and depth of the A3112 observations., We use the results of \citet{Patanchon09} to predict the number of sources expected for a blank field with the same area and depth of the A3112 observations.996" We find expected numbers of 37.15, 22.τν. and 20.15 sources at 250. 350. andmicron.. respectively."," We find expected numbers of $37^{+18}_{-14}$, $22^{+30}_{-16}$, and $20^{+41}_{-17}$ sources at 250, 350, and, respectively."997 Despite the relatively large errors on the predicted number counts. the number of sources detected at exceeds the expected number by a factor 2.," Despite the relatively large errors on the predicted number counts, the number of sources detected at exceeds the expected number by a factor 2."998 The excess decreases at longer wavelengths. to become completely consistent with the expected counts atmicron.," The excess decreases at longer wavelengths, to become completely consistent with the expected counts at."999. This suggests that the excess detected is mainly due to the presence of the cluster. whose population should be detected preferentially atmicron.," This suggests that the excess detected is mainly due to the presence of the cluster, whose population should be detected preferentially at."1000. Robust association of optical counterparts to BLAST sources is usually possible only for very robust detections and relies on a number of assumptions and on the use of ancillary data (ef., Robust association of optical counterparts to BLAST sources is usually possible only for very robust detections and relies on a number of assumptions and on the use of ancillary data (cf.1001 Section 3.2))., Section \ref{indsources}) ).1002 Moreover. the map noise will prevent the detection of individual faint sources.," Moreover, the map noise will prevent the detection of individual faint sources."1003 It is nevertheless possible to obtain robust statistical information about the average BLAST flux density of a sample of counterparts by the BLAST maps on a provided catalogue., It is nevertheless possible to obtain robust statistical information about the average BLAST flux density of a sample of counterparts by the BLAST maps on a provided catalogue.1004 This technique has been described in great detail by ?.. and was successfully used to determine the intensity of the far-IR background (FIRB) using BLAST data in the BGS field (2:: 2:: 25).," This technique has been described in great detail by \citet{Marsden09}, and was successfully used to determine the intensity of the far-IR background (FIRB) using BLAST data in the BGS field \citealt{Devlin09}; \citealt{Marsden09}; \citealt{Pascale09}) )."1005 Stacking will be applied in Sections 3.1.. 3.4.. and 3.5..," Stacking will be applied in Sections \ref{cmembs}, \ref{blaststackrad}, and \ref{blaststackmag}."1006 BLAST data were combined with optical spectroscopy to identify cluster members based on their dynamical state. and UV to near-IR photometry used to characterize the cluster members on the basis of their photometrie properties.," BLAST data were combined with optical spectroscopy to identify cluster members based on their dynamical state, and UV to near-IR photometry used to characterize the cluster members on the basis of their photometric properties."1007 This allows us to investigate the FIR star-formation activity of cluster galaxies together with their unobscured UV star-formation and classification from multi-band photometry., This allows us to investigate the FIR star-formation activity of cluster galaxies together with their unobscured UV star-formation and classification from multi-band photometry.1008 We collected spectroscopic redshifts in the field of A3112 from dedicated observations of the AAOmega Spectrometer on the Anglo-Australian Telescope (AAT), We collected spectroscopic redshifts in the field of A3112 from dedicated observations of the AAOmega Spectrometer on the Anglo-Australian Telescope (AAT).1009 AAOmega (2). is the new fiber-fed spectrograph for the 2dF robot fibre positioner., AAOmega \citep{Sharp06} is the new fiber-fed spectrograph for the 2dF robot fibre positioner.1010 It has 392 fibers covering a total area of 2 deg: a dichroic allow continuous coverage of the spectral region from 3700-8800 with a resolution AfoA=1300., It has 392 fibers covering a total area of 2 $^2$; a dichroic allow continuous coverage of the spectral region from 3700–8800 with a resolution $\lambda/\delta\lambda = 1300$.1011 Data were collected on November 24. 2009 (Proposal ID A103. PE: FGB) for a total of 5.1 hours.," Data were collected on November 24, 2009 (Proposal ID A103, PI: FGB) for a total of 5.1 hours."1012 Target selection was based on available photometry (see below) to select potential cluster members from optical and near-IR colours., Target selection was based on available photometry (see below) to select potential cluster members from optical and near-IR colours.1013 Two instrumental setups allowed us to collect 683 spectra., Two instrumental setups allowed us to collect 683 spectra.1014 Data reduction was performed using the standard AAOmega pipeline. 2dfdr. available on the AAOmegawebsite!.," Data reduction was performed using the standard AAOmega pipeline, 2dfdr, available on the AAOmega."1015. Redshifts were extracted using standard IRAF tasks and then checked individually., Redshifts were extracted using standard IRAF tasks and then checked individually.1016 We obtained redshifts for 578 out of 683 spectra (a success rate of )). with 550 reliable non-stellar spectra.," We obtained redshifts for 578 out of 683 spectra (a success rate of ), with 550 reliable non-stellar spectra."1017 Additional spectroscopic redshifts in the tield were collected from the NASA/TPAC Extragalactic Database (NED)., Additional spectroscopic redshifts in the field were collected from the NASA/IPAC Extragalactic Database (NED).1018 Most of these are from the Las Campanas Redshift Survey (LCRS. ?)) and from the Two Degree Field survey (2dF. 2)).," Most of these are from the Las Campanas Redshift Survey (LCRS, \citealt{Shectman96}) ) and from the Two Degree Field survey (2dF, \citealt{Colless01}) )."1019 A total of 188 redshifts were obtained for this field. with a redshift coverage from 0—0.22. from which we identitied 90 galaxies not covered by the AAOmega observations for a total of 640 redshifts within 0.8 degrees from the cluster centre.," A total of 188 redshifts were obtained for this field, with a redshift coverage from 0–0.22, from which we identified 90 galaxies not covered by the AAOmega observations for a total of 640 redshifts within 0.8 degrees from the cluster centre."1020 Optical photometry in the Harris D. and // passbands is available from. dedicated. observations of ΑΦΗΣ carried out at the Las Campanas Observatory Swope Telescope within the Las Campanas/AAT Rich Cluster Survey (LARCS: ?: 2))., Optical photometry in the Harris $B$ and $R$ passbands is available from dedicated observations of A3112 carried out at the Las Campanas Observatory Swope Telescope within the Las Campanas/AAT Rich Cluster Survey (LARCS: \citealt{Pimbblet01}; \citealt{Pimbblet02}) ).1021 These two filters cover the spectral region from 3500 to AA.. which includes a number of important spectral features for galaxies at low redshift.," These two filters cover the spectral region from 3500 to , which includes a number of important spectral features for galaxies at low redshift."1022 In particular. the two filters bracket the break for galaxies at the redshift of A312. thus providing a," In particular, the two filters bracket the break for galaxies at the redshift of A3112, thus providing a"1023In 1921. P. Ewalel suggested a method to compute the forces in a periodic particle distribution (he aimed. explicitely to compute the potential in. atomic lattices in solids).,"In 1921, P. Ewald suggested a method to compute the forces in a periodic particle distribution (he aimed explicitely to compute the potential in atomic lattices in solids)."1024 For a more recent discussion see Llernquist. Bouchet Suto (1991): for applications to large scale structure simulations see Ixatz. Weinberg Herndquist (1996). or Davé.. Dubinski IHernquist (1997).," For a more recent discussion see Hernquist, Bouchet Suto (1991); for applications to large scale structure simulations see Katz, Weinberg Hernquist (1996), or Davé,, Dubinski Hernquist (1997)."1025 Poisson's equation for à svstem of IN. particles which ave infinitely replicated in all directions with period £ is (n being an integer vector) and can be solved in general using the appropriate Cireen's function 6: which is for the gravitational potential G(r)=l/r. or Gtk)=AAT in Fourier space.," Poisson's equation for a system of $N$ particles which are infinitely replicated in all directions with period $L$ is ${\bf n}$ being an integer vector) and can be solved in general using the appropriate Green's function $\cal G$: which is for the gravitational potential ${\cal G}({\bf r}) = 1/r$, or $\hat{\cal G}({\bf k}) = 4\pi / {k^2}$ in Fourier space."1026 The sum in Eqn., The sum in Eqn.1027 3 converges very slowly. which strongly limits its numerical applicability.," 3 converges very slowly, which strongly limits its numerical applicability."1028 However. Ewald (1921) realised. that convergence can be improved considerably by splitting the Cireen's function into a short range Gs and a long range part ο and solving the first in real space and the latter one in Fourier space: llere @ is a scalingὃν factor in units of inverse lengthe and erf(r) is the error function with erle(r) its complement. With the Cireen's function split into two parts. Gr. the potential reads as ó(r)=óst(r)|ór(r): Finally. the force (sav exerted onto particle 7) is with lo achieve goo accuracy with reasonable computational effort. twpical values are à=2/L. lrn£L|<3.6L and k=2xsh/L with h being an integer vectorwith 10 (see e.g. Llernquist et al.," However, Ewald (1921) realised, that convergence can be improved considerably by splitting the Green's function into a short range ${\cal G}_S$ and a long range part ${\cal G}_L$ and solving the first in real space and the latter one in Fourier space: Here $\alpha$ is a scaling factor in units of inverse length and ${\rm erf}(x)$ is the error function with ${\rm1029erfc}(x)$ its complement, With the Green's function split into two parts, ${\cal G} = {\cal G}_S1030+ {\cal G}_L$ , the potential reads as $\phi({\bf r}) =\phi_S({\bf r}) + \phi_L({\bf r})$: Finally, the force (say exerted onto particle$i$ ) is with To achieve good accuracy with reasonable computational effort, typical values are $\alpha = 2/L$, $|{\bf r} - {\bf n}L| < 3.6 L$ and ${\bf k} \equiv 2 \pi1031{\bf h} / L$ with ${\bf h}$ being an integer vectorwith $|{\bf h}|^21032< 10$ (see e.g. Hernquist et al."1033 1991: note. the second component in their Eqn.," 1991; note, the second component in their Eqn."1034 2.14b is missing a factor [rn]: it is identical to our Eqn. 8))," 2.14b is missing a factor $ | {\bf r} -1035{\bf n}L|$; it is identical to our Eqn. \ref{force}) )."1036 For à particle pair with separation r. f(vr) includes the contributions from all. L-periodic pairs with identical separation.," For a particle pair with separation ${\bf1037r}$, ${\bf f}({\bf r})$ includes the contributions from all $L$ -periodic	 	 pairs with identical separation."1038 To obtain the pure periodic Εμ). one has to subtract the direct. interaction of the central pair. which leads to To obtain the periodic correction for the pofeatiad of the particle pair. Ομ). one proceeds similar ancl again subtracts the isolated solution from. the Ewald solution (Eqn. 6)).," To obtain the pure periodic ${\bf f}_{\rm cor}({\bf r})$, one has to subtract the direct interaction of the central pair, which leads to To obtain the periodic correction for the of the particle pair, $\phi_{\rm cor}({\bf r})$, one proceeds similar and again subtracts the isolated solution from the Ewald solution (Eqn. \ref{potential}) ),"1039" We compute the correction terms fio, and Ooo» for particle pairs on a Cartesian grid. covering our whole simulation box. placing particle 1 in the central node aud particle 2 on dillerent ericl points. and obtain so a table of pairwise force values."," We compute the correction terms ${\bf f}_{\rm cor}$ and $\phi_{\rm1040cor}$ for particle pairs on a Cartesian grid covering our whole simulation box, placing particle 1 in the central node and particle 2 on different grid points, and obtain so a table of pairwise force values."1041 Particle-mesh methods assign particle properties to. mesh points. solve the interaction equations on the grid and interpolate the solution back onto the particles (see c.g. Lockney Eastwood 1955).," Particle-mesh methods assign particle properties to mesh points, solve the interaction equations on the grid and interpolate the solution back onto the particles (see e.g. Hockney Eastwood 1988)."1042 Tvpicallv. the density at each erid point is determined. from. the particle positions ancl masses using CIC (veloud in cell?) or TSC (triangular shaped: cloud”) schemes. whose assignment functions are triangles or euacdratic splines. respectively (again Llockney Eastwood 1988. chap.," Typically, the density at each grid point is determined from the particle positions and masses using CIC (“cloud in cell”) or TSC (“triangular shaped cloud”) schemes, whose assignment functions are triangles or quadratic splines, respectively (again Hockney Eastwood 1988, chap."1043 5)., 5).1044 For the gravitational. N-body problem one solves Poisson's equation., For the gravitational N-body problem one solves Poisson's equation.1045 Usually this is done in Fourier space. since the dilferential operation there acts as a simple multiplication.," Usually this is done in Fourier space, since the differential operation there acts as a simple multiplication."1046 Fherefore. the density. distribution on the erid is transformed. into A-space using FLL and convolved with the appropriate Green's function to solve for the potential.," Therefore, the density distribution on the grid is transformed into $k$ -space using FFT and convolved with the appropriate Green's function to solve for the potential."1047 To obtain the forces. we convolve the Fourier transform of the density with the Green's function for theforce’: separately for the vw. jj ancl z-componoent.," To obtain the forces, we convolve the Fourier transform of the density with the Green's function for the; separately for the $x$ , $y$ and $z$ -component."1048 Finally. inverseFLT returns potential and force at cach grid. point.," Finally, inverseFFT returns potential and force at each grid point."1049Poincaré waves (solid lines) are much higher (han (he rotational Ireequency. Qo. ancl thev increase with increasing ».,"Poincaré waves (solid lines) are much higher than the rotational frequency, $\Omega_0$, and they increase with increasing $n$."1050 For example. the frequency of the 5=3 harmonics is ~50Ομ. which corresponds to oscillations with period  13 hours.," For example, the frequency of the $n=3$ harmonics is $\sim 50 \, \Omega_0$, which corresponds to oscillations with period $\sim$ 13 hours."1051 On the other haud. the frequencies of magnetic Rosshy waves are in general much lower than the rotational Ireequency.," On the other hand, the frequencies of magnetic Rossby waves are in general much lower than the rotational frequency."1052 The absolute value of the frequency. of fast magnetic Rossby waves (dotted line inthe zoon) decreases will increasing n». such as happens with the WD Rossby wave.," The absolute value of the frequency of fast magnetic Rossby waves (dotted line inthe zoom) decreases with increasing $n$, such as happens with the HD Rossby wave."1053 Fie., Fig.1054 2 shows the dependence of the s=1.902 and s=1.5»3 harmonies on € and ay (i.e. Che magnetic field strength).," 2 shows the dependence of the $s=1, n=2$ and $s=1, n=3$ harmonics on $\epsilon$ and $\alpha_0$ (i.e. the magnetic field strength)."1055 The frequency (in absolute value) of fast. (dotted) and slow (dashed) magnetic Rossby. modes significantly increases when the magnetic field is increased., The frequency (in absolute value) of fast (dotted) and slow (dashed) magnetic Rossby modes significantly increases when the magnetic field is increased.1056 We also see (that the lrequencey. difference between the vn=2.3 harmonics of [ast magnetic Rossby waves is almost independent of αμ.," We also see that the frequency difference between the $n=2,3$ harmonics of fast magnetic Rossby waves is almost independent of $\alpha_0$."1057 On the contrary. (hie η=2.3 harmonics of slow magnetic Rossby waves diverge when the magnetic field is increased.," On the contrary, the $n=2,3$ harmonics of slow magnetic Rossby waves diverge when the magnetic field is increased."1058 It (urns out that the frequencies of magnetic Rossby waves have almost no dependence on e. therefore they are not shown in the top panel.," It turns out that the frequencies of magnetic Rossby waves have almost no dependence on $\epsilon$, therefore they are not shown in the top panel."1059" On the other hand. the frequency of magnetic Poincaré waves depends on € (lower panel) aud significantly decreases wilh increasing, e. llo"," On the other hand, the frequency of magnetic Poincaré waves depends on $\epsilon$ (lower panel) and significantly decreases with increasing $\epsilon$."1060wever. these frequencies have almost no dependence on ay. therefore magnetic Poincaré waves are only slightly affected by the magnetic field.," However, these frequencies have almost no dependence on $\alpha_0$, therefore magnetic Poincaré waves are only slightly affected by the magnetic field."1061 In the weakly stable overshoot region (or sav upper tachocline) reduced gravity is much smaller. being 0.055 envs 7 (Schecteretal.2001).," In the weakly stable overshoot region (or say upper tachocline) reduced gravity is much smaller, being 0.05–5 $\cdot$ $^{-2}$ \citep{sch01}."1062. Then. for the thickness //5—5 * 105 em we gel e= 27. 2.7 - 100.," Then, for the thickness $H_0$ =5 $\cdot$ $^8$ cm we get $\epsilon=$ 27 – 2.7 $\cdot$ $^{3}$ ."1063 Therefore. ¢>1 is a good approximation in the overshoot region.," Therefore, $\epsilon \gg10641$ is a good approximation in the overshoot region."1065 In (his case. we follow the calculation of Longuet-Ilieeins(1968) (this paper considers," In this case, we follow the calculation of \citet{lon68} (this paper considers"1066us that RL sources are a distinct ACN population that shows fundamental differences in BLR structure and kinematics.,us that RL sources are a distinct AGN population that shows fundamental differences in BLR structure and kinematics.1067 El indicates that the weak radio euissiou frou: RO sources is unrelated to the RL phenomenon., E1 indicates that the weak radio emission from RQ sources is unrelated to the RL phenomenon.1068 El also shows evidence for a parameter space separation between steep aud flat spectruui RE sources., E1 also shows evidence for a parameter space separation between steep and flat spectrum RL sources.1069 This shows up in Table 2 as a possible RL correlation between FWHAL aandReyyp., This shows up in Table 2 as a possible RL correlation between FWHM and.1070. Steep spectrmm RL sources represent the opposite extrem from NLSy1 while flat spectra sources are more similar to RQ sources., Steep spectrum RL sources represent the opposite extremum from NLSy1 while flat spectrum sources are more similar to RQ sources.1071 The teudeney for steep spectimm RL sources to show the broadest Dahner profiles aud the weakest eenission is confirmed in two large (with some source overlap between themselves aud our RL sample) survevs (Brotherton 1996: Corbin 1997)., The tendency for steep spectrum RL sources to show the broadest Balmer profiles and the weakest emission is confirmed in two large (with some source overlap between themselves and our RL sample) surveys (Brotherton 1996; Corbin 1997).1072 Siebert et al (1998) provide evidence that this separation may also be preseut iu the N-ray spectral index., Siebert et al (1998) provide evidence that this separation may also be present in the X-ray spectral index.1073 NLSvil ave a ΠΟ aud population A οκποπ with the narrowest Dahner line profiles. stroug eenission and a stroug soft X-ray excess;," NLSy1 are a RQ and population A extremum with the narrowest Balmer line profiles, strong emission and a strong soft X-ray excess."1074 NLSy1 show a clear coutimuity ancl correlation with broader line Sevtert 1 galaxies in all measured parameters., NLSy1 show a clear continuity and correlation with broader line Seyfert 1 galaxies in all measured parameters.1075 This challenges the idea that they represent a unique or disjoint AGN population., This challenges the idea that they represent a unique or disjoint AGN population.1076 BAL quasars (1 RQ BAL in the BOY? sample) occupy au EL domain that is similar to the NLSvl., BAL quasars (4 RQ BAL in the BG92 sample) occupy an E1 domain that is similar to the NLSy1.1077 Other (BAL) studies (e.g. Boroson Myers 1992) have also suggestedOO that BAL quasars show Daliier line FWHALZ 3000 fand moderate to strong qpueasures., Other (BAL) studies (e.g. Boroson Myers 1992) have also suggested that BAL quasars show Balmer line $ \simlt$ 3000 and moderate to strong measures.1078 Much less clear is whether RO population B sources represeut a disjoint AGN population or show a sinooth continuation of the population A correlations., Much less clear is whether RQ population B sources represent a disjoint AGN population or show a smooth continuation of the population A correlations.1079 Figure 1 ob for example shows a breakdown of the FWIDMC(IL/) correlation at the nominal boundary between population A aud D. Whatever the relation )etween the two RQ populations. their line profiles show striking differences.," Figure \ref{fig:e1} b for example shows a breakdown of the – ) correlation at the nominal boundary between population A and B. Whatever the relation between the two RQ populations, their line profiles show striking differences."1080 Figure 2. shows a comparison of the aand lue profiles for prototwpe NLSvi source I Zw 1 aud NGC 5518 which is a typical broad line Sevfert ealaxy., Figure \ref{fig:profiles} shows a comparison of the and line profiles for prototype NLSy1 source I Zw 1 and NGC 5548 which is a typical broad line Seyfert galaxy.1081 While differeuces in the Baler profiles are striking. the most mipressive difference is the apparcut sineaticdecoupling of the aud yprofiles in I Zw 1.," While differences in the Balmer profiles are striking, the most impressive difference is the apparent kinematic of the and profiles in I Zw 1."1082 Our previous work (MS96) sugeested that properly NLR corrected jin RQ sources (see Suleutie Marziani 1999) is always blueshifted relative to aand the ACN rest frame., Our previous work (MS96) suggested that properly NLR corrected in RQ sources (see Sulentic Marziani 1999) is always blueshifted relative to and the AGN rest frame.1083 Fie., Fig.1084 2. shows the line shift vs. pe) diagram., \ref{fig:civ} shows the line shift vs. ) diagram.1085 We have normalized the sshift by W(CIVALS19)) iu this plot because we see a complementary trend in El for WECIVALS19)) to be stnallest in the NLSvi population., We have normalized the shift by ) in this plot because we see a complementary trend in E1 for ) to be smallest in the NLSy1 population.1086 Figure 3. coufirms that csseutially all RQ sources show a lueshitt while RL sources show equal red and blueshifts with züuplitudes generally less than x10?suus.., Figure \ref{fig:civ} confirms that essentially all RQ sources show a blueshift while RL sources show equal red and blueshifts with amplitudes generally less than $\pm$ $^3$.1087 Correlations of sshift with acad ((see Table 2) indicate that it is likely to be au miportaut El correlate., Correlations of shift with and (see Table 2) indicate that it is likely to be an important E1 correlate.1088 We will explore it further iu succeeding papers., We will explore it further in succeeding papers.1089 El correlates AGN spectroscopic data in a way that removes much of the apparent “randomness” of line properties., E1 correlates AGN spectroscopic data in a way that removes much of the apparent “randomness” of line properties.1090 It also redefines input paraiucters for photoionization and kincuatical models., It also redefines input parameters for photoionization and kinematical models.1091 We have ouly, We have only1092core.,.1093.. However. even before that stage. the turbulent dviiuuo may have auplified the magnetic field enough that it could reduce or evel suppress fraeieutation (Machidaetal.2008:etal. 2011).. thus potentially infiuencing the initial 1iass function of stars in both primordial aud coutemporary star formation.," However, even before that stage, the turbulent dynamo may have amplified the magnetic field enough that it could reduce or even suppress fragmentation \citep{MachidaEtAl2008,HennebelleTeyssier2008,BuerzleEtAl2010,PetersEtAl2011}, thus potentially influencing the initial mass function of stars in both primordial and contemporary star formation."1094 We thauk Sebastien Fromane for helpful discussions ou the spectral analysis of uou-periodic datasets., We thank Sebastien Fromang for helpful discussions on the spectral analysis of non-periodic datasets.1095 CLF. has received. funding from the European/ Research Council under the Europeau Conuuuuitv* Seveuth Framework. Programme (FP7/2007-2013 Grant Aereenment uo., C.F. has received funding from the European Research Council under the European Community's Seventh Framework Programme (FP7/2007-2013 Grant Agreement no.1096 217060) for the research presented in this work., 247060) for the research presented in this work.1097 C.F.. R.D.. aud R.S.Kk. acknowledge subsidies from the Baden-Wiurttembere-Stiftune (eraut P-LS-SPII/15) and from the Germani Duudesiuinisteum fir Bildung uud Forschung via the ASTRONET project STAR FORMAT (eraut 05AO00VILA).," C.F., R.B., and R.S.K. acknowledge subsidies from the Baden-Würrttemberg-Stiftung (grant P-LS-SPII/18) and from the German Bundesministerium fürr Bildung und Forschung via the ASTRONET project STAR FORMAT (grant 05A09VHA)."1098" S.S. thanks the German Science Foundation (DFC) for financial support via the priority proerun 1177 ""Witnesses of Cosmic Distory: Formation aud Evolution of Black Toles. Cialaxies aud their Επιποιοι (erant KL 1358/10)."," S.S. thanks the German Science Foundation (DFG) for financial support via the priority program 1177 “Witnesses of Cosmic History: Formation and Evolution of Black Holes, Galaxies and their Environment” (grant KL 1358/10)."1099 D.S. thanks for finding from the European Conumuitvs Seveutl Framework Progranuue (FP?/2007-2013) uuder eraut aerecnent No 229517., D.S. thanks for funding from the European Community's Seventh Framework Programme (FP7/2007-2013) under grant agreement No 229517.1100 R.D. is funded by the Enuuyv-Noether evant (DEG)Rechenzentum BA 3706., R.B. is funded by the Emmy-Noether grant (DFG) BA 3706.1101Supercomputing time at the Leibniz (projects pr32lo aud hl221) aud the Forschuugszeutru Jüllieh (projects hhd20 aud hhdll) are eratefully ackuowledeed., Supercomputing time at the Leibniz Rechenzentrum (projects pr32lo and h1221) and the Forschungszentrum Jüllich (projects hhd20 and hhd14) are gratefully acknowledged.1102 The software used in this work was du part developed bv the DOLE-supported ASC Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago., The software used in this work was in part developed by the DOE-supported ASC / Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.1103 Figure 3. was produced with the open-source visualization software VISIT.., Figure \ref{fig:snapshots} was produced with the open-source visualization software .1104Our prescription to account for stellar evolution in the simulations implies a substantial change of the multiphase Collective model” by 1100.,Our prescription to account for stellar evolution in the simulations implies a substantial change of the multiphase “effective model” by SH03.1105 In. this model. gas particles are assumed to have a cold neutral and a hot. ionized phase in pressure equilibrium. the former providing the reservoir for star formation.," In this model, gas particles are assumed to have a cold neutral and a hot ionized phase in pressure equilibrium, the former providing the reservoir for star formation."1106 We have mocified the criterion. dependent on local density ancl temperature. to establish the relative amount of such two phases. so as to account for the eradual SN energy. release. which mocifies the temperature of the hot phase and the evaporation of the cold one.," We have modified the criterion, dependent on local density and temperature, to establish the relative amount of such two phases, so as to account for the gradual SN energy release, which modifies the temperature of the hot phase and the evaporation of the cold one."1107 This means Το include new energy terms in Eq.(10) of SLHI03. which describes the evolution of the internal energy of the hotphase component. while making the cooling function dependent on local metallicity.," This means to include new energy terms in Eq.(10) of SH03, which describes the evolution of the internal energy of the hot–phase component, while making the cooling function dependent on local metallicity."1108" “Phe metaldependence of cooling. that we introduce using the tables from Sutherland DopitaDopitz (1993).9903). also enters ini determiningqGerpminine the, onset, of⋅ he⋅ thermal instability (eq.22] of SLLO3) and. the value of the density threshold for star formation (eq.23] of 81103)."," The metal–dependence of cooling, that we introduce using the tables from Sutherland Dopita (1993), also enters in determining the onset of the thermal instability (eq.[22] of SH03) and the value of the density threshold for star formation (eq.[23] of SH03)."1109 511059 also provided a phenomenological description for ealactic winds. which are triggered by SN energy release and whose strength is regulated by two parameters.," SH03 also provided a phenomenological description for galactic winds, which are triggered by SN energy release and whose strength is regulated by two parameters."1110 A first one gives the rate of gas ejection by winds. according to the relation. Aly=pM. where A. is the star formation rate.," A first one gives the rate of gas ejection by winds, according to the relation, $\dot M_W=\eta \dot1111M_*$, where $\dot M_*$ is the star formation rate."1112 The second one. determines the fraction of SN energy. that »owers the winds. ERITDESovMe. where esx ds the energy feedback provided by the SN under LRA assumption or each. AZ. of stars formed.," The second one, determines the fraction of SN energy that powers the winds, ${1\over11132}\dot M_W v_W^2= \chi \epsilon_{SN}\dot M_*$, where $\epsilon_{SN}$ is the energy feedback provided by the SN under IRA assumption for each $M_\odot$ of stars formed."1114 In our implementation of the winds. we also account for the energy. contributed from all he SN treated. without the LRA. namely SNIL in the 20M. mass range and SNIa.," In our implementation of the winds, we also account for the energy contributed from all the SN treated without the IRA, namely SNII in the $M_\odot$ mass range and SNIa."1115 lt is worth noticing that the star formation and. SN eedback scheme. originally introduced. by SIIO3. has been already demonstrated to provide the correct. fraction. of murvons locked in stars. the correct cosmic star [ormation listory (Springel Llernquist 2003h). the correct amount of neutral hydrogen in. high. columndensity. absorbing. systems at high redshift (Nagamine ct al.," It is worth noticing that the star formation and SN feedback scheme, originally introduced by SH03, has been already demonstrated to provide the correct fraction of baryons locked in stars, the correct cosmic star formation history (Springel Hernquist 2003b), the correct amount of neutral hydrogen in high column–density absorbing systems at high redshift (Nagamine et al."1116 2003). and to reproduce the esie Xray scaling properties of galaxy clusters (Borgani et al.," 2003), and to reproduce the basic X–ray scaling properties of galaxy clusters (Borgani et al."1117 2003)., 2003).1118 As such. it represents a good starting point for a simulation study of the ICM chemical enrichment.," As such, it represents a good starting point for a simulation study of the ICM chemical enrichment."1119that the intrisic non-Iensed flux of the sources is just. below 0.95 Js. a small gravitational amplification of the tux by a [actor of10-15% would. push these ΑΝ above the selection threshold of the survey.,"that the intrisic non-lensed flux of the sources is just below 0.95 Jy, a small gravitational amplification of the flux by a factor of$\approx$ would push these AGN above the selection threshold of the survey."1120 The ability of clusters of ealaxies to gravitationally lens extended sources is clearly demonstrated by the luminous ares which are background galaxies (Le. extended sources) distorted. ancl amplified. by the cluster potential (c.g. Smail et al., The ability of clusters of galaxies to gravitationally lens extended sources is clearly demonstrated by the luminous arcs which are background galaxies (i.e. extended sources) distorted and amplified by the cluster potential (e.g. Smail et al.1121 1995)., 1995).1122 Although in our case the geometry. is different. ie. the distance of the lens is not much smaller than that of the source. we are only considering small magnification factors which could involve even only one radio lobe. whose size is smaller or similar to that of galaxies in the optical.," Although in our case the geometry is different, i.e. the distance of the lens is not much smaller than that of the source, we are only considering small magnification factors which could involve even only one radio lobe, whose size is smaller or similar to that of galaxies in the optical."1123 To summarize. the present cata do not. allow us to explore in more details the validity of the possible scenarios.," To summarize, the present data do not allow us to explore in more details the validity of the possible scenarios."1124 However. according to the above discussion. we can conclude that. based on their colours and A-band: magnitudes. the excess red galaxies around the AGN are likely to be at 1κο<2. hence favoring the scenario where racdio-loucl AGN at 1.5<2 live in rich environments. or the case where the red galaxies are at 2<στον.," However, according to the above discussion, we can conclude that, based on their colours and $K$ -band magnitudes, the excess red galaxies around the AGN are likely to be at $1<z<2$, hence favoring the scenario where radio-loud AGN at $1.5<z<2$ live in rich environments, or the case where the red galaxies are at $z<z_{AGN}$."1125 The former case is relevant to understand the genesis of the AGN phenomenon. the relationships between radio-quiet ancl racto-Loucd ACN and for the use of radio-loud AGN as tracers of cluster and massive structures of galaxies at high-z.," The former case is relevant to understand the genesis of the AGN phenomenon, the relationships between radio-quiet and radio-loud AGN and for the use of radio-loud AGN as tracers of cluster and massive structures of galaxies at $z$."1126 Follow-up optical and near-LR spectroscopy ancl submillimetre photometry of the selected. EROs is under wav in order to unveil their nature and their role in the framework of galaxy. formation and evolution., Follow-up optical and near-IR spectroscopy and submillimetre photometry of the selected EROs is under way in order to unveil their nature and their role in the framework of galaxy formation and evolution.1127 We are grateful to J. Cohen and P. Hall for providing the colours and the magnitudes of their galaxy. samples. to D. Thompson for uscful information on the CADIS photometric system and on the Lfdy colors of the CADIS EltOs. to C. Lidman for the information on the IRAC2b photonietric calibration. to D. Silva for his assistance in the use of the SUSI data taken in service observing. and to €. Bruzual ad 8. Charlot for providing their spectral synthesis mocdols.," We are grateful to J. Cohen and P. Hall for providing the colours and the magnitudes of their galaxy samples, to D. Thompson for useful information on the CADIS photometric system and on the $H-K$ colors of the CADIS EROs, to C. Lidman for the information on the IRAC2b photometric calibration, to D. Silva for his assistance in the use of the SUSI data taken in service observing, and to G. Bruzual and S. Charlot for providing their spectral synthesis models."1128 We acknowledge the anonymous. referee. for the useful comments., We acknowledge the anonymous referee for the useful comments.1129 This research has mace use of the NASA/LPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Acronautics ancl Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."1130 L.I. acknowledges the support of a research erant [rom during the development of this project., L.P. acknowledges the support of a research grant from during the development of this project.1131have placed limits on HI columns of <10!9 ((e.g.. Burns. White. and Havnes 1981: Valentijn and Giovanelli 1982: MeNamara. Bregman. and O'Connell 1990: Jaffe 1991. 1992: Dwarakanath 11994: O'Dea. Gallimore. Daum 1995: O'Dea. Pavne lxocevski 1998).,"have placed limits on HI columns of $\lae 10^{19}$ (e.g., Burns, White, and Haynes 1981; Valentijn and Giovanelli 1982; McNamara, Bregman, and O'Connell 1990; Jaffe 1991, 1992; Dwarakanath 1994; O'Dea, Gallimore, Baum 1995; O'Dea, Payne Kocevski 1998)."1132 HI absorption has been detected in only a few clusters ((A4126/3C84 - Crane 11932: Jaffe 1990; Sijbring 1993: A7S8O0/Ivdra A - Tavlor 1996; A2597/PINS2322-123 - ODea. Baum Gallimore 1994: Tavlor 11999) with implied masses of atomic hydrogen in the range ~10*—105 ML..," HI absorption has been detected in only a few clusters (A426/3C84 - Crane 1982; Jaffe 1990; Sijbring 1993; A780/Hydra A - Taylor 1996; A2597/PKS2322-123 - O'Dea, Baum Gallimore 1994; Taylor 1999) with implied masses of atomic hydrogen in the range $\sim 10^7 - 10^8$ $_\odot$."1133 The strong dominance of molecular relative to atomic gas in these clusters is in contrast (ο the situation in normal galaxies where roughly ~15% of the cold gas is thought to be in molecular form (e.g.. Doselli. Lequeux ancl Gavazzi 2002): however. the bulk of the molecular gas in galaxies may be undetected (e.g.. Allen 11997).," The strong dominance of molecular relative to atomic gas in these clusters is in contrast to the situation in normal galaxies where roughly $\sim 15\%$ of the cold gas is thought to be in molecular form (e.g., Boselli, Lequeux and Gavazzi 2002); however, the bulk of the molecular gas in galaxies may be undetected (e.g., Allen 1997)."1134 The 21 cm limits on the HI column are subject to two caveats: 1) the limits depend linearly on the electron excitation temperature. which may be higher (han previously thought.," The 21 cm limits on the HI column are subject to two caveats: 1) the limits depend linearly on the electron excitation temperature, which may be higher than previously thought."1135 2) the absorbing gas may have a verv small velocity dispersion. producing a very narrow saturated absorption line. ancl (hus a lower absorption equivalent width than estimated [or the optically thin case.," 2) the absorbing gas may have a very small velocity dispersion, producing a very narrow saturated absorption line, and thus a lower absorption equivalent width than estimated for the optically thin case."1136 The UV region provides a powerlul but as vet not fully exploited probe of the ICM., The UV region provides a powerful but as yet not fully exploited probe of the ICM.1137 The aabsorplion cross-section is >[0* times larger than that of the 21 em line (e.g.. Baheall Ekers 1969: Laor 1997).," The absorption cross-section is $\gae 10^7$ times larger than that of the 21 cm line (e.g., Bahcall Ekers 1969; Laor 1997)."1138" In addition. it is not subject to the two caveats which affect the 21 em line. as 1) the absorption is practically independent of the gas temperature (the extremely short lifetime of the excited η=2 level ensures that effectivelv all ΠΠ is in the ground level): 2) sienificant absorption is expected even if the absorber has no velocity dispersion (the absorption will occur in the Lorentzian. or ""damping wings)."," In addition, it is not subject to the two caveats which affect the 21 cm line, as 1) the absorption is practically independent of the gas temperature (the extremely short lifetime of the excited $n=2$ level ensures that effectively all HI is in the ground level); 2) significant absorption is expected even if the absorber has no velocity dispersion (the absorption will occur in the Lorentzian, or “damping” wings)."1139 Thus. lis superior to the 21 em line as a probe of the cold gas content of the ICM.," Thus, is superior to the 21 cm line as a probe of the cold gas content of the ICM."1140 In fact. Noekemoer (1998) presented IIST/FOS UV spectra of the quasar in the center of AI030 and placed imits on column densities in the range 10!—1055 lor a wide range of molecular. atomic. aud ionized species Chat max be associated with the ICM.," In fact, Koekemoer (1998) presented HST/FOS UV spectra of the quasar in the center of A1030 and placed limits on column densities in the range $10^{11} - 10^{13}$ for a wide range of molecular, atomic, and ionized species that may be associated with the ICM."1141 Miller. Dregman. Wnezek (2001) obtained limits of Lot—LOM on column densities lor several absorption lines (Fell. Mell) through lines of sight towards the outer parts of six clusters.," Miller, Bregman, Knezek (2001) obtained limits of $10^{12} - 10^{13}$ on column densities for several absorption lines (FeII, MgII) through lines of sight towards the outer parts of six clusters."1142 Laor (1997) has used a low resolution (~250 !) HST/FOS spectrum of the line in (he center of the Perseus cluster (published bv Johustone Fabian 1995) to set an, Laor (1997) has used a low resolution $\sim 250$ ) /FOS spectrum of the line in the center of the Perseus cluster (published by Johnstone Fabian 1995) to set an1143Conventional Vega. magnitudes. are usec throughout this work.,Conventional Vega magnitudes are used throughout this work.1144" Except where explicitly. stated to the contrary a cosmology is adopted throughout such that H5270 Ixm t+ + and O3,20.3. Q420.7."," Except where explicitly stated to the contrary a cosmology is adopted throughout such that $\rm_0$ =70 Km $^{-1}$ $^{-1}$ , and $\Omega_{\rm{M}}$ =0.3, $\Omega_{\rm\Lambda}$ =0.7."1145 The imaging data were obtained with the 3.6m Canacda-France-Llawaii (CELL) Telescope and the CELDI21I. camera., The imaging data were obtained with the 3.6m Canada-France-Hawaii (CFH) Telescope and the CFH12K camera.1146 CELIL2ES is a close packed mosaic of 12 back side illuminated AIP Lincoln Laboratories CCDs giving a total field of view in a single observation of 28arcmin and a pixel size of 0.206arcsec., CFH12K is a close packed mosaic of 12 back side illuminated MIT Lincoln Laboratories CCDs giving a total field of view in a single observation of $\times$ 28arcmin and a pixel size of 0.206arcsec.1147 The cosmetic quality of the mosaic is good with only one chip significantly alleeted by bad columns., The cosmetic quality of the mosaic is good with only one chip significantly affected by bad columns.1148 Coupled with the excellent. seeing routinely obtained at CELUT. the fine pixel scale allows excellent. morphological classification oL objects leading to a reduction of the number of marginally extended non quasar sources in the candidate list.," Coupled with the excellent seeing routinely obtained at CFHT, the fine pixel scale allows excellent morphological classification of objects leading to a reduction of the number of marginally extended non quasar sources in the candidate list."1149 Observations were obtained in the 17 and Z bancs during the six nights 12-17. September 1999. interleaved with observations for alternate projects., Observations were obtained in the $VI$ and $Z$ bands during the six nights 12-17 September 1999 interleaved with observations for alternate projects.1150 Details are given in table 1.., Details are given in table \ref{obsdata}.1151 Figure 1. shows the pointing lavout for the Canada-France Deep Field (CEDE: MeCracken 2001)) 2215.00 field from which the observations presented. here ave drawn., Figure \ref{cfdf field} shows the pointing layout for the Canada-France Deep Field (CFDF: McCracken 2001)) 2215+00 field from which the observations presented here are drawn.1152 Anderson (2001) subsequently identified a high luminosity z=4.99 quasar within the 2215|00 field using data from SDSS., Anderson (2001) subsequently identified a high luminosity $z=4.99$ quasar within the 2215+00 field using data from SDSS.1153 The recovery of this objects as a canciadte quasar is discussed in section 4.. ," The recovery of this objects as a candiadte quasar is discussed in section \ref{pSDSSJ2216+0013, a1154quasar at z=4.99}. ."1155The data form. part of an extended: survey project of ~6.2dee= oobserved in the Wi and Z bands., The data form part of an extended survey project of $\sim$ observed in the $VI$ and $Z$ bands.1156 J band imaging from the public NOAQO field at 04:30(J2000) will be used in conjunction with V and Z band imaging obtained with CFEILI2lx. Canclicate selection and spectroscopic followup observations have been carried out on a rregion of the survey centred on the Canada-Lranee Deep field region at | 00:30(.2000) for which V7 and Z data have been obtained., $I$ band imaging from the public NOAO field at $-$ 04:30(J2000) will be used in conjunction with $V$ and $Z$ band imaging obtained with CFH12K. Candidate selection and spectroscopic followup observations have been carried out on a region of the survey centred on the Canada-France Deep field region at $+$ 00:30(J2000) for which $VI$ and $Z$ data have been obtained.1157 Figure 2. shows the [iter transmission curves [or the filters used in the survey., Figure \ref{cfht viz filters} shows the filter transmission curves for the filters used in the survey.1158 With accurate star/galaxy separation. possible here due to the excellent seeing during the observations and the well matched: pixel scale of the CrUI2K camera. the dominate. contaminant population within candidate lists for most high. redshift quasar colour selection surveys is low mass stars.," With accurate star/galaxy separation, possible here due to the excellent seeing during the observations and the well matched pixel scale of the CFH12K camera, the dominate contaminant population within candidate lists for most high redshift quasar colour selection surveys is low mass stars."1159 The extended red. tail of the 4 filter is therefore not ideal for this work since the overlap between the £ and Z filters reduces the power of (πο Z colour index to identify low mass stars., The extended red tail of the $I$ filter is therefore not ideal for this work since the overlap between the $I$ and $Z$ filters reduces the power of the $I-Z$ colour index to identify low mass stars.1160 However. the VZZ filter set. and the associated colour-colour diagram (Figure 4)). does provide sullicent. cliscrimination between quasars at recdshift z24.8 and low mass stars.," However, the $VIZ$ filter set, and the associated colour-colour diagram (Figure \ref{cfht viz}) ), does provide sufficent discrimination between quasars at redshift $z>4.8$ and low mass stars."1161 The data were processed using and the package., The data were processed using and the package.1162 Twilight fat field frames are required to facilitate the removal of interference. [ringes from. the / and Z data., Twilight flat field frames are required to facilitate the removal of interference fringes from the $I$ and $Z$ data.1163 Fringe removal is carried out. by iteratively scaling anc subtracting a master fringe frame to minimise residual background: variations., Fringe removal is carried out by iteratively scaling and subtracting a master fringe frame to minimise residual background variations.1164 “Phe process is automated hrough the use of software created as part of the Isaac rewton Telescope. Wide. Field Camera (INPWEC) data reduction pipeline (Irwin and Lewis (2001))., The process is automated through the use of software created as part of the Isaac Newton Telescope Wide Field Camera (INTWFC) data reduction pipeline (Irwin and Lewis (2001)).1165 The master ringe frame is compiled on a nightly. basis from Lat fielded sky observations., The master fringe frame is compiled on a nightly basis from flat fielded sky observations.1166 Fringing in the Z band is present at the level before processing., Fringing in the $Z$ band is present at the level before processing.1167 Residual [ringing remains at he level prior to image alignment and stacking., Residual fringing remains at the level prior to image alignment and stacking.1168 Figure 3 gives an example of the results of the data reduction orocedure., Figure 3 gives an example of the results of the data reduction procedure.1169 Object catalogue eoneration. morphological classification and the computation of απ astrometric solution was performed using elements. of the INPAWEC pipeline (Irwin and Lewis 2001)) and with the aid of the image analysist," Object catalogue generation, morphological classification and the computation of an astrometric solution was performed using elements of the INTWFC pipeline (Irwin and Lewis 2001)) and with the aid of the image analysis."1170ool? Aperture photometry is derived. based on a median secing radius aperture (Q.Saresec) anc is used for V7 ancl Z images., Aperture photometry is derived based on a median seeing radius aperture (0.8arcsec) and is used for $VI$ and $Z$ images.1171 An aperture correction is derived on a chip-by-chip bases from a curve of erowth analysis of bright. non saturated stars in the field.," An aperture correction is derived on a chip-by-chip bases from a curve of growth analysis of bright, non saturated stars in the field."1172 The World Coordinate System (WC€S) is used to merge the object catalogues. across the V4 and Z filters. with the Z catalogue used as the reference., The World Coordinate System (WCS) is used to merge the object catalogues across the $VI$ and $Z$ filters with the $Z$ catalogue used as the reference.1173 Flux calibrations isachieved. with observations of standard starfields [rom the Landolt (1992) catalogue (primarily LO5 ancl LILO)., Flux calibrations isachieved with observations of standard starfields from the Landolt (1992) catalogue (primarily L95 and L110).1174 Observations were recorded. at several locations across the CCD mosaic., Observations were recorded at several locations across the CCD mosaic.1175 Atos3 the optical colour indices of quasars are dominated, At $z>3$ the optical colour indices of quasars are dominated1176dynamical friction from galaxy wakes (e.g.Schipper1974:Lea&DeYoung1976:Miller1986:Justetal. 1990): or some combination of these (e.g.Brüggen2003:Dennis&Chandran2005:Conroy&Ostriker 2008).,"dynamical friction from galaxy wakes \citep[e.g.][]{schipper74,ly76,miller86,jdkbm90}; or some combination of these \citep[e.g.][]{brueggen03,dc05,co08}."1177. Much of recent work has focused on either thermal conduction. AGN heating. or both.," Much of recent work has focused on either thermal conduction, AGN heating, or both."1178 Thermal conduction alone cannot be the solution to the cooling flow problem across the full range of masses due to its steep temperature dependence (Voigt&Fabian2004:Kaastraetal.2004:Popeet 2006).," Thermal conduction alone cannot be the solution to the cooling flow problem across the full range of masses due to its steep temperature dependence \citep{vf04,kaastra04,ppkf06}."1179. Even in hot systems where conduction is potent. fine tuning of the suppression factor f (the fraction by which Spitzer conductivity is reduced due to. e.g.. tangled magnetic tield lines) is required (Bregman&David1988).," Even in hot systems where conduction is potent, fine tuning of the suppression factor $f$ (the fraction by which Spitzer conductivity is reduced due to, e.g., tangled magnetic field lines) is required \citep{bd88}."1180". Moreover. if he thermal conduction has the same temperature dependence as he Spitzer conductivity (i.e. if f is a constant) for a given ICM atmosphere. the resulting equilibria are thermally unstable (e.g. Bregman&David1988:Soker2003:KimNarayan2003b:: see ""urther discussion in Sections ??. and ??))."," Moreover, if the thermal conduction has the same temperature dependence as the Spitzer conductivity (i.e. if $f$ is a constant) for a given ICM atmosphere, the resulting equilibria are thermally unstable (e.g. \citealt{bd88,soker03,kn03b}; see further discussion in Sections \ref{sec:balance} and \ref{sec:conduction}) )."1181 More recently. however. there has been renewed interest in the yossibility that thermal conduction may provide sufficient heating o stably counteract the effects of radiative cooling.," More recently, however, there has been renewed interest in the possibility that thermal conduction may provide sufficient heating to stably counteract the effects of radiative cooling."1182 This has gone hand-in-hand with a dramatic increase in our understanding of ‘dilute’ (i.e. only weakly collisional) plasmas. due to the appreciation that even very weak magnetic fields introduce an anisotropy into heat fluxes.," This has gone hand-in-hand with a dramatic increase in our understanding of `dilute' (i.e. only weakly collisional) plasmas, due to the appreciation that even very weak magnetic fields introduce an anisotropy into heat fluxes."1183 One important consequence is that the criterion for convective instability ehanges to one of temperature. rather than entropy. increasing downwards (Balbus2000.2001).," One important consequence is that the criterion for convective instability changes to one of temperature, rather than entropy, increasing downwards \citep{balbus00,balbus01}."1184 Quataert(2008). generalised Balbus's (2000) analysis and found that a heat-flux buoyaney-driven instability (ABI) occurs. for upwardly-increasing temperature protiles as well. so long as the magnetic field is not entirely horizontal torthogonal to gravity).," \citet{quataert08} generalised Balbus's (2000) analysis and found that a heat-flux buoyancy-driven instability (HBI) occurs for upwardly-increasing temperature profiles as well, so long as the magnetic field is not entirely horizontal (orthogonal to gravity)."1185 Balbus&Reynolds(2008) conjectured that the nonlinear HBI is self-regulating and drives a reverse convective thermal flux. both of which may mediate the stabilisation of cooling cores.," \citet{br08} conjectured that the nonlinear HBI is self-regulating and drives a reverse convective thermal flux, both of which may mediate the stabilisation of cooling cores."1186 Numerica simulations of the HBI have been performed by (2008).. Bogdanoviéetal.(2009) and Parrish.Quataert&Sharma(2009) with applications to the ICM.," Numerical simulations of the HBI have been performed by \citet{pq08}, \citet{brbp09} and \citet{pqs09} with applications to the ICM."1187 It was discovered that the HBI acts to rapidly reorient field lines to insulate the core. undermining the role of thermal conduction and causing a cooling catastrophe to occur.," It was discovered that the HBI acts to rapidly reorient field lines to insulate the core, undermining the role of thermal conduction and causing a cooling catastrophe to occur."1188 Subsequent work has shown that a moderate amount of turbulent driving may help regulate the HBI and allow therma channels to remain open. potentially stabilising the core agains collapse (Sharmaetal.2009:Parrish.Quataert&Sharma2010:Ruszkowski&Oh 2010).," Subsequent work has shown that a moderate amount of turbulent driving may help regulate the HBI and allow thermal channels to remain open, potentially stabilising the core against collapse \citep{scqp09,pqs10,ro10}."1189 There are several reasons to believe that AGNs play an important role in regulating cooling., There are several reasons to believe that AGNs play an important role in regulating cooling.1190 In many clusters. the AGN energy output inferred from radio-emitting plasma outflows and cavities is similar to the X-ray cooling rate of thecentral gas (Fabianetal.2000:Kanov.Sarazin&Hicks2006:McNamara2007:Formanetal. 2007).," In many clusters, the AGN energy output inferred from radio-emitting plasma outflows and cavities is similar to the X-ray cooling rate of thecentral gas \citep{fabian00,ksh06,mn07,forman07}."1191.. Moreover. 2TO% of cool-core clusters harbour radio sources at their centres. while 25% of clusters without cool cores are radio loud (Burns19901. providing strong circumstantial evidence for a connection between the processes that fuel the radio emission (such as AGNs) and the X-ray emission from the cooling gas.," Moreover, $\gtrsim 70\%$ of cool-core clusters harbour radio sources at their centres, while $\lesssim 25\%$ of clusters without cool cores are radio loud \citep{burns90}, providing strong circumstantial evidence for a connection between the processes that fuel the radio emission (such as AGNs) and the X-ray emission from the cooling gas."1192 Models of self-regulated heating from AGN have been constructed (e.g.Ciotti&OstrikerBrüggen&Scannapieco2009) in which AGN activity is triggered by cooling-induced gas accretion toward cluster centres. increasing AGN heating and halting the collapse.," Models of self-regulated heating from AGN have been constructed \citep[e.g.][]{co01,rb02,bm03,kb03,hb04,go08,bs09} in which AGN activity is triggered by cooling-induced gas accretion toward cluster centres, increasing AGN heating and halting the collapse."1193 Episodic outflows from AGN are thought not only to quench cooling and condensation in clusters. but also to limit the maximum luminosity of galaxies and regulate the growth of black holes at their centres (Binney2005).," Episodic outflows from AGN are thought not only to quench cooling and condensation in clusters, but also to limit the maximum luminosity of galaxies and regulate the growth of black holes at their centres \citep{binney05}."1194. While AGN activity is fundamentally linked with the observed presence of radio bubbles and/or X-ray cavities (e.g.Birzanetal.2004:Dunn&Fabian2006).. it is currently unclear how the AGN energy is actually thermalised (e.g. see the introduction of Voit&Donahue2005 for a review of possibilities).," While AGN activity is fundamentally linked with the observed presence of radio bubbles and/or X-ray cavities \citep[e.g.][]{brmwn04,df06}, it is currently unclear how the AGN energy is actually thermalised (e.g. see the introduction of \citealt{vd05} for a review of possibilities)."1195 This question can only be answered once knowledge of the effective viscosity of the ICM is acquired., This question can only be answered once knowledge of the effective viscosity of the ICM is acquired.1196 The ICM hosts subsonic turbulence and magnetic fields with energy density comparable to that of the motions., The ICM hosts subsonic turbulence and magnetic fields with energy density comparable to that of the motions.1197 Both of these should affect the viscosity of the ICM (seereviewbySchekochihin&Cowley2006)., Both of these should affect the viscosity of the ICM \citep[see review by][]{sc06}.1198. In particular. the presence of a magnetic tield alters the form of the viscosity when the ratio of the ion cyclotron and collision frequencies is much greater than unity (Braginskil1965)... a condition amply satisfied in galaxy clusters.," In particular, the presence of a magnetic field alters the form of the viscosity when the ratio of the ion cyclotron and collision frequencies is much greater than unity \citep{braginskii65}, a condition amply satisfied in galaxy clusters."1199 As a result. the transport properties of the ICM become strongly dependent on both the zeometry and strength of the magnetic field. as Well as on microscale plasma instabilities that are likely to occur ubiquitously in the ICM (e.g.firehoseandmirror:Schekochihinal.2010).," As a result, the transport properties of the ICM become strongly dependent on both the geometry and strength of the magnetic field, as well as on microscale plasma instabilities that are likely to occur ubiquitously in the ICM \citep[e.g. firehose and mirror;][]{sckhs05,lyutikov07,sckrh08,scrr10,rsrc10}."1200 In this paper. we investigate the effect these plasma effects might have on the large-scale transport. properties of the ICM.," In this paper, we investigate the effect these plasma effects might have on the large-scale transport properties of the ICM."

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