ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 This should be possible for some of the closer clusters. such as Virgo.," This should be possible for some of the closer clusters, such as Virgo."3 Note that the solutions for X and 7 given in this paper do not depend upon the normalisations of the gas and total eravitational masses., Note that the solutions for $\Sigma$ and $\tau$ given in this paper do not depend upon the normalisations of the gas and total gravitational masses.4 Lf desired. these can be determined by fixing the overall temperature and. luminosity.," If desired, these can be determined by fixing the overall temperature and luminosity."5 The analysis of Gunn Thomas (1996) shows that the gas density will be slightly. lower and the total mass density slightly higher than in the equivalent single-phase analysis., The analysis of Gunn Thomas (1996) shows that the gas density will be slightly lower and the total mass density slightly higher than in the equivalent single-phase analysis.6 Cooling flows in individual galaxies are much better resolved. than those in clusters and. so look like promising candidates for the kind. of modelling discussed here., Cooling flows in individual galaxies are much better resolved than those in clusters and so look like promising candidates for the kind of modelling discussed here.7 However. the lower temperatures leads to complications such as à varving slope. a. for the cooling function. ancl larger corrections for absorption and emission Lying outside the pass-band of the detector.," However, the lower temperatures leads to complications such as a varying slope, $\alpha$, for the cooling function, and larger corrections for absorption and emission lying outside the pass-band of the detector."8 Phere may also be sources of mass and energy injection into the flow., There may also be sources of mass and energy injection into the flow.9 This paper was prepared. using the facilities of the SEPARLININ minor node at Sussex., This paper was prepared using the facilities of the STARLINK minor node at Sussex.10 Ht was written while PAT was holding a Nullielcl Foundation Science. Research Fellowship., It was written while PAT was holding a Nuffield Foundation Science Research Fellowship.112811133 We assume an emulsion. of density phases which comove with the flow., We assume an emulsion of density phases which comove with the flow.12 The clistribution is described by the volume fraction. f(p.y./). such that fdp is the fractional volume," The distribution is described by the volume fraction, $f(\rho,\rvec,t)$, such that $f\,\dd\rho$ is the fractional volume"13the optical. it is unlikely to be sufficient to suppress the optica broad lines entirely unless the absorber has an anomalously high dust-to-gas ratio (c.f.,"the optical, it is unlikely to be sufficient to suppress the optical broad lines entirely unless the absorber has an anomalously high dust-to-gas ratio (c.f."14 MCG-6-30-15: 25)., MCG-6-30-15; \cite{reynolds97}) ).15 The X-ray spectrum also shows line emission between 6 and 7 keV. The centroid energy of this line is not well constrained. so we assume that the emission is from neutral iron Κά at a res energy of 6.4 keV. though it may originate from ionised iron.," The X-ray spectrum also shows line emission between 6 and 7 keV. The centroid energy of this line is not well constrained, so we assume that the emission is from neutral iron $\alpha$ at a rest energy of 6.4 keV, though it may originate from ionised iron."16 À gaussian fit to this feature with the centroid energy fixed at 6.4 keV gives an equivalent width (EW) of 160 eV. On the face of it. this relatively modest iron Ka EW argues against a hidden Compton thick nucleus (Bianchi et al.," A gaussian fit to this feature with the centroid energy fixed at 6.4 keV gives an equivalent width (EW) of 160 eV. On the face of it, this relatively modest iron $\alpha$ EW argues against a hidden Compton thick nucleus (Bianchi et al."17 2008)., 2008).18 We found. however. that it was possible to fit a second. heavily absorbed componen from new Monte-Carlo models of X-rays in heavily obscured AGN. incorporating line emission (Brightman et al..," We found, however, that it was possible to fit a second, heavily absorbed component from new Monte-Carlo models of X-rays in heavily obscured AGN, incorporating line emission (Brightman et al.,"19 in. preparation). similar to those presented by e.g. ? and ?..," in preparation), similar to those presented by e.g. \cite{ghisellini94} and \cite{krolik94}."20 For simplicity given the limitations of the data we fit the extreme case of 4a coverage (ie. a spherical distribution of matter). and we constrained the column density of the heavily absorbed component to Ny=2.10ο. ?.," For simplicity given the limitations of the data we fit the extreme case of $\pi$ coverage (i.e. a spherical distribution of matter), and we constrained the column density of the heavily absorbed component to $N_{\rm{H}}=2 \times 10^{24}|^{+65}_{-0.8}$ $^{-2}$."21 Although it is not formally required in the fit. this demonstrates a clear scenario in which the multiwavelength data ean be reconciled. as such a column is easily sufficient to suppress the direct nuclear emission in the optical.," Although it is not formally required in the fit, this demonstrates a clear scenario in which the multiwavelength data can be reconciled, as such a column is easily sufficient to suppress the direct nuclear emission in the optical."22 By correcting for absorption in his heavily obscured component. we can make another prediction of the intrinsic luminosity of the nucleus and calculate the corresponding bolometric luminosity.," By correcting for absorption in this heavily obscured component, we can make another prediction of the intrinsic luminosity of the nucleus and calculate the corresponding bolometric luminosity."23 This now agrees well with estimations from {ο and £Ly2jau. TheNe," This now agrees well with estimations from $L_{\rm [O\,III]}$ and $L_{\rm{12{\umu}m}}$."24wton observation also shows no clear signs of variability., The observation also shows no clear signs of variability.25 The spectropolarimetric survey of ?. reveals that this AGN hosts a HBLR. in full agreement with our conclusion of a heavily buried nuclear contribution.," The spectropolarimetric survey of \cite{tran03} reveals that this AGN hosts a HBLR, in full agreement with our conclusion of a heavily buried nuclear contribution."26 The same electrons which scatter the broad optical lines into the line of sight can then also scatter nuclear X- accounting for the apparently unobscured nature of the soft spectrum., The same electrons which scatter the broad optical lines into the line of sight can then also scatter nuclear X-rays accounting for the apparently unobscured nature of the soft spectrum.27 Our analysis of the spectrum of NGC 3147 shows no absorption above the Galactic column. however it features a small (130 eV) iron Καὶ line at 6.4. keV. These results are in full agreement with those of Bianchi et al. (," Our analysis of the spectrum of NGC 3147 shows no absorption above the Galactic column, however it features a small (130 eV) iron $\alpha$ line at 6.4 keV. These results are in full agreement with those of Bianchi et al. ("282008).,2008).29 ὁ used the small EW of the iron Ka line to argue against the Compton thick nature of this source and came to the conclusion that NGC 3147 has an intrinsic absence of a BER., \cite{bianchi08} used the small EW of the iron $\alpha$ line to argue against the Compton thick nature of this source and came to the conclusion that NGC 3147 has an intrinsic absence of a BLR.30 However. as we have shown in IRASFOI475-0740. this conclusion is not necessarily robust.," However, as we have shown in IRASF01475-0740, this conclusion is not necessarily robust."31 Fitting the Monte Carlo models to the spectrum once again shows that a very heavily obscured component is consistent with the spectrum., Fitting the Monte Carlo models to the spectrum once again shows that a very heavily obscured component is consistent with the spectrum.32" Here the column density is constrained to be Vy=91077|L7 em7, with the main constraint coming from the iron Ka line."," Here the column density is constrained to be $N_{\rm{H}}=9 \times 10^{23}|^{+419}_{-1.0}$ $^{-2}$, with the main constraint coming from the iron $\alpha$ line."33" We can again use this second component to predict the intrinsic luminosity of this source and calculate the corresponding bolometric luminosity. which also agrees well with estimations from Loup and Zs;u,."," We can again use this second component to predict the intrinsic luminosity of this source and calculate the corresponding bolometric luminosity, which also agrees well with estimations from $L_{\rm [O\,III]}$ and $L_{\rm{12{\umu}m}}$."34 Other evidence suggests. however. that this may not be the correct interpretation in this ease.," Other evidence suggests, however, that this may not be the correct interpretation in this case."35 The 2ks observation of NGC 3147 shows that the measured flux of the nucleus is ~2 times that measured by indicating that it is variable in nature and hence it must be being observed directly. rather than in scattered light.," The 2ks observation of NGC 3147 shows that the measured flux of the nucleus is $\sim2$ times that measured by indicating that it is variable in nature and hence it must be being observed directly, rather than in scattered light."36 Furthermore. the imaging reveals no evidence for extra-nuclear X-ray sources that could be producing the unabsorbed X-ray protile or the variability. and spectral extraction from the nuclear region in the image does not reveal a significantly harder X-ray spectrum then the larger beam (Table 4).," Furthermore, the imaging reveals no evidence for extra-nuclear X-ray sources that could be producing the unabsorbed X-ray profile or the variability, and spectral extraction from the nuclear region in the image does not reveal a significantly harder X-ray spectrum then the larger beam (Table 4)."37 The spectrum of NGC 3486 is well fitted by a power- absorbed by the Galactic column only and so appears to be another unabsorbed Seyfert 2 galaxy., The spectrum of NGC 3486 is well fitted by a power-law absorbed by the Galactic column only and so appears to be another unabsorbed Seyfert 2 galaxy.38 However. an excess at hard energies points to a different scenario.," However, an excess at hard energies points to a different scenario."39 The spectrum is also well fitted by the addition of a Compton reflection model(tpzzrmon. ?)) with an underlying thermal component (ravrionz). so this could in fact be a Compton thick Seyfert 2 which looks unabsorbed.," The spectrum is also well fitted by the addition of a Compton reflection model, \cite{nandra07}) ) with an underlying thermal component ), so this could in fact be a Compton thick Seyfert 2 which looks unabsorbed."40 We also attempt to add a strongly absorbed transmission component, We also attempt to add a strongly absorbed transmission component41particle distribution rather than three. which is what is seen iu the simulations of (Liou&Zook1999).,"particle distribution rather than three, which is what is seen in the simulations of \citep{liou99}."42. We have found that the morphology of dusty material trapped in exterior resonances with a planet cau be strongly dependent upon the eccentricity aud orbital phase of the planet., We have found that the morphology of dusty material trapped in exterior resonances with a planet can be strongly dependent upon the eccentricity and orbital phase of the planet.43" The morphology of the e Eridani dust ring ποσα» to be reproduced by dust particles captured iuto the 5:3 and 3:2 exterior resonances with a moderate eccentricity. 60,~0.2. plauet near periastrou."," The morphology of the $\epsilon$ Eridani dust ring seems to be reproduced by dust particles captured into the 5:3 and 3:2 exterior resonances with a moderate eccentricity, $e_p \sim 0.3$, planet near periastron."44 When the planet was more massive (jg5«10 !) or amore lughly eccentric (greater than 0.1). the resonances closest to the planet did not capture and hold particles as efücieuntlv.," When the planet was more massive $\mu > 5 \times 10^{-4}$ ) or more highly eccentric (greater than 0.4), the resonances closest to the planet did not capture and hold particles as efficiently."45 When the planet mass was below a Saturn mass. the 1:3 resonance also contributed to the particle distribution resulting iu donunant svimimetric four peaks.," When the planet mass was below a Saturn mass, the 4:3 resonance also contributed to the particle distribution resulting in dominant symmetric four peaks."46 The asvuunetries observed in the ε Eridaui iniage were not observed in the dust distribution so we consider planets less massive than Saturn uulikelv to explain the particle distribution., The asymmetries observed in the $\epsilon$ Eridani image were not observed in the dust distribution so we consider planets less massive than Saturn unlikely to explain the particle distribution.47 When the planet ecceutricity was below 0.15. the azuuuthal deusitv variations in the dust distribution were too low to account for the morphology im the e Exidani disk.," When the planet eccentricity was below 0.15, the azimuthal density variations in the dust distribution were too low to account for the morphology in the $\epsilon$ Eridani disk."48 Our model differs from that proposed by Ozernoyctal.(2000). in a number of wavs., Our model differs from that proposed by \citet{ozernoy} in a number of ways.49 The period of the orbital planet in our model is 280 vears which should cause the pattern to revolve by about the star by ~L.3%yr|. faster than the ~O.7°vr+ estimated by Ozcruoyctal. (2000).. where the model planet senmunajor axis is  GOAT rather than at  LOAT.," The period of the orbital planet in our model is 280 years which should cause the pattern to revolve by about the star by $\sim 1.3^\circ {\rm yr}^{-1}$, faster than the $\sim 0.7^\circ {\rm yr}^{-1}$ estimated by \citet{ozernoy}, , where the model planet semi-major axis is $\sim 60$ AU rather than at $\sim 40$ AU."50 Our model plauet is located to the north of the star. rather than to the west of the stax.," Our model planet is located to the north of the star, rather than to the west of the star."51 Our model planet mass is similar to theirs but at a moderate eccentricity., Our model planet mass is similar to theirs but at a moderate eccentricity.52 Our inodol dust conceutratious are a result of segregation in the phase of a resouaut auele. rather than caused by a large libration amplitude.," Our model dust concentrations are a result of segregation in the phase of a resonant angle, rather than caused by a large libration amplitude."53 Furthermore. because of the eccentricity of the planet. our nocdel predicts that the morphology of the dusty ring will vary. as well as revolve as the planet orbits about the star.," Furthermore, because of the eccentricity of the planet, our model predicts that the morphology of the dusty ring will vary, as well as revolve as the planet orbits about the star."54 The initial conditions of our simulations cause niuiv of he particles to begin trapped im the resonances aud at arly low eccentricity., The initial conditions of our simulations cause many of the particles to begin trapped in the resonances and at fairly low eccentricity.55 It is possible that low eccentricity dlanetesimals exist iu the svstem aud that they are the source of the cust particles that we sec., It is possible that low eccentricity planetesimals exist in the system and that they are the source of the dust particles that we see.56 Alternatively he planctesimals in the svstem are further out and the dust particles become trapped in the resonances as they spiral in toward the planet., Alternatively the planetesimals in the system are further out and the dust particles become trapped in the resonances as they spiral in toward the planet.57 Further simulations would )0 required to differcutiate between these possibilities., Further simulations would be required to differentiate between these possibilities.58 Iu either case. resonaut capture iuto these resonances is ess likely and less prolonged when the planet mass or eccentricity is high.," In either case, resonant capture into these resonances is less likely and less prolonged when the planet mass or eccentricity is high."59 If the planet mass is too low then it cannot be responsible for cleariug a gap or central region im he dust distribution., If the planet mass is too low then it cannot be responsible for clearing a gap or central region in the dust distribution.60 Additional and more massive plauets would be required to do this., Additional and more massive planets would be required to do this.61 It is possible that high ecceutricity planets are common in the outskirts of extra-solar svstems., It is possible that high eccentricity planets are common in the outskirts of extra-solar systems.62 If κο then the resulting dust distributions would not only revolve (Ozernoyetal.2000).. but will also be depeudent upon the orbital phase of the planets.," If so then the resulting dust distributions would not only revolve \citep{ozernoy}, but will also be dependent upon the orbital phase of the planets."63 This is an exciting prospect because there would be variations im the dust morphology on observable timescales., This is an exciting prospect because there would be variations in the dust morphology on observable timescales.64 This work would not have been carried out without helpful discussions with StepheuThorudike. Joel Creenu aud Dau Watson.," This work would not have been carried out without helpful discussions with StephenThorndike, Joel Green and Dan Watson."65"papers report the features to be ""transient"" (i.e. detected in only a subset of the data).",papers report the features to be “transient” (i.e. detected in only a subset of the data).66 The large number of moderately significant detections (see sect. 2) , The large number of moderately significant detections (see sect. \ref{sect:analysis}) )67might be considered as evidence to support the reality of the features., might be considered as evidence to support the reality of the features.68" However. when considering results presented for individual datasets drawn from a much larger population of available data one should be aware of the distorting etfects of ""publication bias.” also Known as the “tile-drawer etfect"" — the tendency for positive results to be published and negative results to go unreported (""filed away)."," However, when considering results presented for individual datasets drawn from a much larger population of available data one should be aware of the distorting effects of “publication bias,” also known as the “file-drawer effect” – the tendency for positive results to be published and negative results to go unreported (“filed away”)."69 See 9 ?.. 9? and ? ‘or genera discussion of publication bias. and also 9? and ?. for a more recen discussion of the importance of publication bias in the contex of medical trials.," See \citet{sterling59}, , \citet{rosenthal79} and \citet{begg88} for general discussion of publication bias, and also \citet{stern97} and \citet{naylor97} for a more recent discussion of the importance of publication bias in the context of medical trials."70" One way to test for the presence of publication bias is through a ""funnel plot originally proposed to aid analyses of medical trials (2).. which compares the size of a to its estimate of the strength of the effect."," One way to test for the presence of publication bias is through a “funnel plot,” originally proposed to aid meta-analyses of medical trials \citep{egger97}, which compares the size of a to its estimate of the strength of the effect."71 All estimates of the strength of an effect should be symmetrically scattered around the true value. with smaller trials providing less precise estimates and so larger scatter.," All estimates of the strength of an effect should be symmetrically scattered around the true value, with smaller trials providing less precise estimates and so larger scatter."72 In the absence of publication bias this will result in asymmetric funnel-shaped plot because the estimates of the effect strength are independent of the sample sizes. but the scatter is larger or smaller samples.," In the absence of publication bias this will result in a symmetric funnel-shaped plot because the estimates of the effect strength are independent of the sample sizes, but the scatter is larger for smaller samples."73 If publication bias is present. experiments or observations are less likely to be published if the estimate of the effect strength is low tor of low significance). i.e. the bias is against oublishing non-detections. leading to an asymmetric funnel plot in which the strength of the effect is correlated with the sample size.," If publication bias is present, experiments or observations are less likely to be published if the estimate of the effect strength is low (or of low significance), i.e. the bias is against publishing non-detections, leading to an asymmetric funnel plot in which the strength of the effect is correlated with the sample size."74 The funnel plot produced from biased literature is the same as from he equivalent unbiased literature but with less “interesting” results (i.e. less significant or less strong results. which lie on one side of he funnel) systematically removed.," The funnel plot produced from biased literature is the same as from the equivalent unbiased literature but with less “interesting” results (i.e. less significant or less strong results, which lie on one side of the funnel) systematically removed."75 This paper describes a simple meta-analysis of the published detections of highly shifted. narrow X-ray lines in active galaxies using a funnel plot-like analysis.," This paper describes a simple meta-analysis of the published detections of highly shifted, narrow X-ray lines in active galaxies using a funnel plot-like analysis."76 The conventional funnel plot would not be appropriate in the present context because there is expected to be considerable intrinsic heterogeneity in the strength and properties of the shifted lines. which means there is no single. true value for the strength of the “effect.”," The conventional funnel plot would not be appropriate in the present context because there is expected to be considerable intrinsic heterogeneity in the strength and properties of the shifted lines, which means there is no single, true value for the strength of the “effect.”"77 But the principle of the funnel plot should still hold: the estimated strengths of the lines should be independent of the quality of the data used to find them (in this context quality means essentially the signal-to-noise of the data. which is of course closely related to the size of the photon sample that constitutes the spectrum).," But the principle of the funnel plot should still hold: the estimated strengths of the lines should be independent of the quality of the data used to find them (in this context quality means essentially the signal-to-noise of the data, which is of course closely related to the size of the photon sample that constitutes the spectrum)."78 The strength of a real line should be independent of the exposure time and detector sensitivity used to measure it., The strength of a real line should be independent of the exposure time and detector sensitivity used to measure it.79 The starting point for the meta-analysis was a search for published claims of highly shifted. narrow. emission or absorption lines in the X-ray spectra of Seyfert galaxies and quasars.," The starting point for the meta-analysis was a search for published claims of highly shifted, narrow, emission or absorption lines in the X-ray spectra of Seyfert galaxies and quasars."80 For the purposes of the present study. these features are defined as intrinsically emission or absorption features found in 1.5—20 keV X-ray spectra of Seyfert galaxies or quasars. that have been identitied with prominent transitions in the X-ray band (e.g. Ka lines of Fe. Ca. Ar. S. Si. Mg) leading to inflow/outflow velocities0.," For the purposes of the present study, these features are defined as intrinsically emission or absorption features found in $1.5-20$ keV X-ray spectra of Seyfert galaxies or quasars, that have been identified with prominent transitions in the X-ray band (e.g. $\alpha$ lines of Fe, Ca, Ar, S, Si, Mg) leading to inflow/outflow velocities."8105. In the particular case of iron. emission lines were accepted if outside the range 6.1—7.3 keV (corresponding to 6.4—6.9 keV Ka lines from Fe at r/o=£0.05). and absorption lines if outside the range 6.4—7.3 keV (corresponding to 6.7—6.9 keV resonances in H and He-like Fe xxv-xxvi.," In the particular case of iron, emission lines were accepted if outside the range $6.1-7.3$ keV (corresponding to $6.4-6.9$ keV $\alpha$ lines from Fe at $v/c = \pm 0.05$ ), and absorption lines if outside the range $6.4-7.3$ keV (corresponding to $6.7-6.9$ keV resonances in H and He-like Fe )."82 The slightly different ranges for absorption and emission correspond to the ditferent species expected to dominate in each ease., The slightly different ranges for absorption and emission correspond to the different species expected to dominate in each case.83 This is criterion provides a simple and uniform. albeit arbitrary. way to distinguish between the highly shifted. isolated. narrow line features that are the focus of this paper and more mildly shifted structures likely to be more directly linked to a broad emission tor absorption) complex centred around 6.4—6.9 keV. The main result and conclusion of the paper would not be significantly changed if the r/c threshold was increased (e.g. to v/c2O.1).," This is criterion provides a simple and uniform, albeit arbitrary, way to distinguish between the highly shifted, isolated, narrow line features that are the focus of this paper and more mildly shifted structures likely to be more directly linked to a broad emission (or absorption) complex centred around $6.4-6.9$ keV. The main result and conclusion of the paper would not be significantly changed if the $v/c$ threshold was increased (e.g. to $v/c \ge 0.1$ )."84" An initial list of papers was constructed from all refereed articles listed in the NASA Astrophysics Data System (ADS°)) sublished between 1995 and 2007 Cinclusive}. selecting papers with abstract text that matched the Boolean exoression ""narrowand X-ray lineand (redshiftedor blueshiftedi"""," An initial list of papers was constructed from all refereed articles listed in the NASA Astrophysics Data System ) published between 1995 and 2007 (inclusive), selecting papers with abstract text that matched the Boolean expression “narrow X-ray line (redshifted blueshifted).”"85 The resulting 135 »upers Were then examined individually to select only those that reported new detections of the type of feaures under investigation., The resulting $135$ papers were then examined individually to select only those that reported new detections of the type of features under investigation.86 This provided a list of 12 such papers., This provided a list of $12$ such papers.87" By following their ""paper rails"" (citations to/from the articles) it was possible to add a further I+ papers. yielding a total of 26 papers oresenting detections 38shifted. narrow lines."," By following their “paper trails” (citations to/from the articles) it was possible to add a further $14$ papers, yielding a total of $26$ papers presenting detections $38$shifted, narrow lines."88 Table | lists all the line features found hrough this literature search., Table \ref{table} lists all the line features found through this literature search.89 The X-ray absorption systems repor'ed in the gravitationally ensed Broad Absorption Line (BAL) quasars (2222). were treated separately.," The X-ray absorption systems reported in the gravitationally lensed Broad Absorption Line (BAL) quasars \citep[][]{chartas02, chartas03, chartas07b, chartas07a} were treated separately."90 In every published BAL case at least one componen of the absorption system was reported as resolved and broad. anc so did not match the criterion above.," In every published BAL case at least one component of the absorption system was reported as resolved and broad, and so did not match the criterion above."91 Also. gravitationally lensec BAL quasars arguably represent a rather special sample of objects within which there known high velocity absorption systems. and so there are good reasons to treat these objects as distinct from the sample of Seyferts and non-BAL quasars.," Also, gravitationally lensed BAL quasars arguably represent a rather special sample of objects within which there known high velocity absorption systems, and so there are good reasons to treat these objects as distinct from the sample of Seyferts and non-BAL quasars."92 For completeness these are included in Table [.. but are not considered in the discussion that follows.," For completeness these are included in Table \ref{table}, but are not considered in the discussion that follows."93 The strength of the feature reported by ? was given inboth equivalent width (EW) and photon flux terms bu the relative uncertainties stated for each are different: 13 and SS per cent. respectively.," The strength of the feature reported by \citet{gallo05} was given inboth equivalent width $EW$ ) and photon flux terms but the relative uncertainties stated for each are different: $13$ and $55$ per cent, respectively."94 Given the modest etfect of this feature on the fi statistic (Ay?= 8.1). the larger of these two uncertainties woulc appear to be the more plausible. and this value is used in Table I.. although it should be noted that the conclusion of the present paper does not depend on this one value.," Given the modest effect of this feature on the fit statistic $\Delta \chi^2 = 8.1$ ), the larger of these two uncertainties would appear to be the more plausible, and this value is used in Table \ref{table}, although it should be noted that the conclusion of the present paper does not depend on this one value."95 In the cases of Mrk 766 (2). and NGC 3516 (2).. the line strengths were given only in flux terms.," In the cases of Mrk 766 \citep{turner04} and NGC 3516 \citep{turner02}, the line strengths were given only in flux terms."96 In order to provide a better comparison with the other lines these were converted in EW terms using the flux density of the continuum a the location of the lines. found by fitting the relevant data’.," In order to provide a better comparison with the other lines these were converted in $EW$ terms using the flux density of the continuum at the location of the lines, found by fitting the relevant ."97.. The 5.9 keV absorption line in NGC 3516 has nopublished EW (2).. although spectral fitting of the publically available data yielded an," The $5.9$ keV absorption line in NGC 3516 has nopublished $EW$ \citep{nandra99}, , although spectral fitting of the publically available data yielded an"98Disruption of a three-body exavitational svstem is an enigmatic dvnanmical process. statistics of which is mostly unexplored. at long timescales especially.,"Disruption of a three-body gravitational system is an enigmatic dynamical process, statistics of which is mostly unexplored, at long timescales especially."99" Valtonen(1988). supposed that the lifetime distribution for a three-body. system is an exponenUally decaving ΠΙΠΟΙΟΙ. in analogy with ""radioactive decav."," \cite{V88} supposed that the lifetime distribution for a three-body system is an exponentially decaying function, in analogy with “radioactive decay”."100 Recent statistical numerical studies by MikkolaTanikawa(2007) of (his process in the equal-anass problem revealed new important data on the statistics of the disruption (mes and raised new questions on the nature of this process., Recent statistical numerical studies by \cite{MT07} of this process in the equal-mass problem revealed new important data on the statistics of the disruption times and raised new questions on the nature of this process.101 Mikkola&Tanikawa(2007) have explored a statistics of the disruption times 74 in the, \cite{MT07} have explored a statistics of the disruption times $T_\mathrm{d}$ in the102"to be unabsorbed, but in fact are hiding a deeply buried AGN and that the unabsorbed profile is not necessarily nuclear.","to be unabsorbed, but in fact are hiding a deeply buried AGN and that the unabsorbed profile is not necessarily nuclear."103" The observed equivalent width of the iron Ka line is strongly dependent on the relative contributions of the various components and, in addition, the geometry and physical parameters of the obscuring material."," The observed equivalent width of the iron $\alpha$ line is strongly dependent on the relative contributions of the various components and, in addition, the geometry and physical parameters of the obscuring material."104" In this paper we present an analysis of six apparently unabsorbed Seyfert 2 galaxies, aiming to probe the nature of these objects."," In this paper we present an analysis of six apparently unabsorbed Seyfert 2 galaxies, aiming to probe the nature of these objects."105 The parent sample for our study is the extendedJRAS 12 um sample of ?.., The parent sample for our study is the extended 12 $\umu$ m sample of \cite{rush93}.106 From this we select objects defined as Seyfert 2s by the NASA/IPAC Extragalactic Database (NED)., From this we select objects defined as Seyfert 2s by the NASA/IPAC Extragalactic Database (NED).107 We perform our own post-hoc analysis of the optical spectra below., We perform our own post-hoc analysis of the optical spectra below.108" We require there to be good quality X-ray data ASCA) with a column density indicating that they are unabsorbed in X-rays (Ng.<10°? cm~?), from the literature or our own analysis."," We require there to be good quality X-ray data ) with a column density indicating that they are unabsorbed in X-rays $N_{\rm{H}} < 10^{22}$ $^{-2}$ ), from the literature or our own analysis."109" The NED optical classifications may not necessarily be robust, so we searched the literature to deselect sources with incorrect or ambiguous optical classification."," The NED optical classifications may not necessarily be robust, so we searched the literature to deselect sources with incorrect or ambiguous optical classification."110" The remaining six objects, IRASF 01475-0740, NGC 3147, NGC 3486, NGC 3660, NGC 3976, NGC 4501 form our sample: they are all unambiguously classified as AGN using line ratio diagnostics (Line ratios and references given in Table 1 and BPT diagram presented in Fig. 1))."," The remaining six objects, IRASF 01475-0740, NGC 3147, NGC 3486, NGC 3660, NGC 3976, NGC 4501 form our sample: they are all unambiguously classified as AGN using line ratio diagnostics (Line ratios and references given in Table \ref{obsdat} and BPT diagram presented in Fig. \ref{bptdiag}) )."111" We compiled optical spectroscopy, spectropolarimetry and line ratio data for our objects from the literature and present these data with references in Table 1 and Fig. 1,,"," We compiled optical spectroscopy, spectropolarimetry and line ratio data for our objects from the literature and present these data with references in Table \ref{obsdat} and Fig. \ref{bptdiag},"112 which plots the ratio ΠΠ] A5007/Hf versus the ratio ΠΠ] A6584/Ha. We plot these on top of the catalogue of low redshift SDSS galaxies of ?.., which plots the ratio ] $\lambda5007$ $\beta$ versus the ratio ] $\lambda6584$ $\alpha$ We plot these on top of the catalogue of low redshift SDSS galaxies of \cite{kauffmann03}.113" The demarcation line is that defined by ? separating AGN and starbursts, hence showing that their AGN classification is not in doubt."," The demarcation line is that defined by \cite{kewley01} separating AGN and starbursts, hence showing that their AGN classification is not in doubt."114" Additionally, we plot the line ratios of NGC 6810, a Seyfert 2 also unabsorbed in X-rays, but shown to have a dubious Seyfert 2 classification due to broader-than-normal optical lines produced by a super-wind (?).."," Additionally, we plot the line ratios of NGC 6810, a Seyfert 2 also unabsorbed in X-rays, but shown to have a dubious Seyfert 2 classification due to broader-than-normal optical lines produced by a super-wind \citep{strickland07}."115 The authors used an XMM—Newton observation to show that the X-ray emission from this galaxy was probably due to X-ray binaries., The authors used an $XMM-Newton$ observation to show that the X-ray emission from this galaxy was probably due to X-ray binaries.116" We carried out the X-ray analysis on observations of IRASF 01470-0740, NGC 3147, NGC 3486, NGC 3976 and NGC 4501; observations of NGC 3147 and NGC 4501 and an observation of NGC 3660 (X-ray observational information listed in table 2))."," We carried out the X-ray analysis on observations of IRASF 01470-0740, NGC 3147, NGC 3486, NGC 3976 and NGC 4501; observations of NGC 3147 and NGC 4501 and an observation of NGC 3660 (X-ray observational information listed in table \ref{xobsdat}) )."117" For data we use tasks to perform spectral extractions on the EPIC-pn data, for theChandra ACIS-S data we use tasks and use data products from the database."," For data we use tasks to perform spectral extractions on the EPIC-pn data, for the ACIS-S data we use tasks and use data products from the database."118" All spectra were grouped with a minimum of 20 counts per bin, with the exception of theChandra observation of NGC 4501, where the spectrum was grouped using a minimum of 7 counts."," All spectra were grouped with a minimum of 20 counts per bin, with the exception of the observation of NGC 4501, where the spectrum was grouped using a minimum of 7 counts."119 Spectral fitting was carried out usingv11., Spectral fitting was carried out using.120"3. We estimate the bolometric luminosities, Lpoi, of the six AGN from their 12 jum and [Ο1Π]] fluxes using bolometric corrections, K12\m calculated from template SEDs of ? and «jorij] published by ?.."," We estimate the bolometric luminosities, $L_{\rm{Bol}}$ , of the six AGN from their 12 ${\umu}$ m and ] fluxes using bolometric corrections, $\kappa_{\rm 12\umu m}$ calculated from template SEDs of \cite{mrr08} and $\kappa_{\rm [O\,III]}$ published by \cite{heckman04}."121" We also estimate bolometric luminosities from the unabsorbed 2-10 keV (HX) luminosity using the mean bolometric correction, KHx, of ?.."," We also estimate bolometric luminosities from the unabsorbed 2-10 keV (HX) luminosity using the mean bolometric correction, $\kappa_{\rm HX}$, of \cite{vasudevan07}."122 Table 3 presents the observed luminosities with these estimates and Fig., Table \ref{lumins} presents the observed luminosities with these estimates and Fig.123 2 plots the estimated bolometric luminosities from the rr]] luminosity against the estimated bolometric luminosities from the 2-10 keV luminosity., \ref{lumfig} plots the estimated bolometric luminosities from the ] luminosity against the estimated bolometric luminosities from the 2-10 keV luminosity.124" From this analysis, all six objects appear to be significantly under-luminous in the 2-10 keV X-ray band by factors of 10-100."," From this analysis, all six objects appear to be significantly under-luminous in the 2-10 keV X-ray band by factors of 10-100."125" For a typical Seyfert 2 this is easily understood, as we expect the 2-10 keV X-rays to be suppressed by absorption, but of course for our sample it seems the evidence for that absorption in terms of the spectral shape is absent."," For a typical Seyfert 2 this is easily understood, as we expect the 2-10 keV X-rays to be suppressed by absorption, but of course for our sample it seems the evidence for that absorption in terms of the spectral shape is absent."126" As discussed below, despite the lack of any obvious soft X- absorption in these objects, some of the X-ray spectra present evidence for hidden hard components."," As discussed below, despite the lack of any obvious soft X-ray absorption in these objects, some of the X-ray spectra present evidence for hidden hard components."127 This allows an additional estimate of the intrinsic AGN bolometric luminosity show as the red arrows in Fig. 2.., This allows an additional estimate of the intrinsic AGN bolometric luminosity show as the red arrows in Fig. \ref{lumfig}.128 These are discussed in more detail below., These are discussed in more detail below.129 The spectrum of IRASF 01475-0740 iswell fitted by a power-law absorbed by a column of Ny=4.1x10?! cm.While this will produce some reddening and extinction in, The spectrum of IRASF 01475-0740 iswell fitted by a power-law absorbed by a column of $N_{\rm{H}} = 4.1 \times 10^{21}$ $^{-2}$ .While this will produce some reddening and extinction in130 , 131"μιαν lan Yau 707242 TOs 10HE,κLxLol? 7Tüjun [jin-selected. 70g ~~1 TOjau 70721 10& TÜpiaa. ddeg.JD | ",$\mu$ $\mu$ $\mu$ $\mu$ $\mu$ $10^{11}\lsun \la L_{ir} \la 10^{12}\lsun$ $\mu$ $\mu$ $\mu$ $z\sim1$ $\mu$ $\mu$ $10^{\prime}\times 10^{\prime}$ $\mu$ $^2$ $^{-1}$ 132 11.05 μα., 1.05$\mu$ G.133are statistical only: additional systematic error arises from uncertainty in the effective area of the LAT (about below | GeV. at 1-10 GeV. and above 10 GeV) and the structure of the diffuse background.,"are statistical only; additional systematic error arises from uncertainty in the effective area of the LAT (about below 1 GeV, at 1–10 GeV, and above 10 GeV) and the structure of the diffuse background."134 Systematic uncertainties in the upper limits stem primarily from uncertainty in the background model and are comparable in magnitude to those associated with the assumed beaming factor discussed below., Systematic uncertainties in the upper limits stem primarily from uncertainty in the background model and are comparable in magnitude to those associated with the assumed beaming factor discussed below.135 Our study gives six candidate sub-luminous pulsars (marked in bold)., 	Our study gives six candidate sub-luminous pulsars (marked in ).136 Three in the uniform sample (plus PSR J1932+1059 = B1929-10) have parallax distance measurements., Three in the uniform sample (plus PSR J1932+1059 = B1929+10) have parallax distance measurements.137 These are particularly important as the parallax constraints control a major factor in the luminosity uncertainty. allowing us to probe the effects of beaming geometry and gap emissivity.," These are particularly important as the parallax constraints control a major factor in the luminosity uncertainty, allowing us to probe the effects of beaming geometry and gap emissivity."138 For the others we must rely at present on the DM distance estimates., For the others we must rely at present on the DM distance estimates.139 These pulsars are displayed in Table | and Figure |., These pulsars are displayed in Table 1 and Figure 1.140 Figure | also shows several other nearby non-recycled LAT-detected pulsars. highlighting the separation of our sub-luminous set from this sample.," Figure 1 also shows several other nearby non-recycled LAT-detected pulsars, highlighting the separation of our sub-luminous set from this sample."141 For this figure we have assumed fo=1 for all pulsars., For this figure we have assumed $f_\Omega=1$ for all pulsars.142 The plotted luminosity errors are dominated by the distance uncertainties. but do include the statistical flux errors.," The plotted luminosity errors are dominated by the distance uncertainties, but do include the statistical flux errors."143 Of course. systematic errors and non-unity fo may add additional uncertainty for individual pulsars.," Of course, systematic errors and non-unity $f_\Omega$ may add additional uncertainty for individual pulsars."144 For simple dipole models (e.g. the OG model) the pulse profile and the expected radiation on the Earth line-of-sight are determined by the magnetic inclination angle à and the viewing angle ¢., 	For simple dipole models (e.g. the OG model) the pulse profile and the expected radiation on the Earth line-of-sight are determined by the magnetic inclination angle $\alpha$ and the viewing angle $\zeta$.145 If these angles are known. we can predict 7-ray pulse profiles and fluxes for specific models and correct observations to the true L..," If these angles are known, we can predict $\gamma$ -ray pulse profiles and fluxes for specific models and correct observations to the true $L_\gamma$."146 Unfortunately these are poorly known in many cases., Unfortunately these are poorly known in many cases.147 The sub-luminous candidates treated here are known radio pulsars. so the magnetic impact angle }=¢ «Is believed to be small.," The sub-luminous candidates treated here are known radio pulsars, so the magnetic impact angle $\beta=\zeta-\alpha$ is believed to be small."148 In the context of the rotating vector model (RadhakrishnanandCooke1969) radio polarization data can constrain the viewing angles., In the context of the rotating vector model \citep{rc69} radio polarization data can constrain the viewing angles.149" In most cases. the small range of phase illuminated by the radio pulse allows only an estimate of the magnetic impact angle where the maximum rate of the polarization position angle (PA) sweep V(o) occurs at o4,4,4 Near the closest approach to the magnetic axis."," In most cases, the small range of phase illuminated by the radio pulse allows only an estimate of the magnetic impact angle where the maximum rate of the polarization position angle (PA) sweep $\Psi(\phi)$ occurs at $\phi_{\rm d\psi, max}$, near the closest approach to the magnetic axis."150 Here the sign of the sweep is meaningful. determining whether the line of sight is closer to or farther from the positive rotation axis than the observed magnetic pole (at inclination à).," Here the sign of the sweep is meaningful, determining whether the line of sight is closer to or farther from the positive rotation axis than the observed magnetic pole (at inclination $\alpha$ )."151 Occasionally. when the radio pulse is very broad or when the pulse profile presents an inter-pulse. the radio polarization can make meaningful estimates of both o and c. from fits to the full polarization sweep where the polarization. has the absolute position. angle Wy at oy.," Occasionally, when the radio pulse is very broad or when the pulse profile presents an inter-pulse, the radio polarization can make meaningful estimates of both $\alpha$ and $\zeta$, from fits to the full polarization sweep where the polarization has the absolute position angle $\Psi_0$ at $\phi_0$."152 Keithetal.(2010) have. recently. presented several examples of constraining fits of Eq. (, \citet{keith10} have recently presented several examples of constraining fits of Eq. (1534) to high quality polarization data.,4) to high quality polarization data.154 As described by Everett&Weis-berg (2001).. while nearly all authors fit to Eqs. (," As described by \citet{ew01}, while nearly all authors fit to Eqs. ("155"3) and (4). given the standard astronomical convention of position angle measurement (increasing N through E) these equations are inconsistent with pulsar angles increasing from the positive spin axis (the ""RVM convention problem’).","3) and (4), given the standard astronomical convention of position angle measurement (increasing N through E) these equations are inconsistent with pulsar angles increasing from the positive spin axis (the `RVM convention problem')."156 To be consistent. one must actually use αρυι=7αμα and συ]=nva.," To be consistent, one must actually use $\alpha_{EW01} = \pi-\alpha_{RVM}$ and $\beta_{EW01} = -\beta_{RVM}$ ."157 Usually this correction is only a formality. but as fits to the x-ray emission improve. including details of sweep-back and magnetospheric currents. the signs can be important.," Usually this correction is only a formality, but as fits to the $\gamma$ -ray emission improve, including details of sweep-back and magnetospheric currents, the signs can be important."158 Thus in the figures and discussion to follow. we convert all *RVM’-fit angles to the consistent Everett&Weisberg(2001) convention; we encourage future workers to do the same.," Thus in the figures and discussion to follow, we convert all `RVM'-fit angles to the consistent \citet{ew01}159 convention; we encourage future workers to do the same."160 Other phenomenological constraints may be extracted from the radio data., 	Other phenomenological constraints may be extracted from the radio data.161" For example. radio emission is. generally believed to be produced within the ""open zone' above the polar cap."," For example, radio emission is generally believed to be produced within the `open zone' above the polar cap."162 For a static aligned dipole the half opening angle covered by this radio beam is radians for modest emission altitudes γε=2rh/Pe., For a static aligned dipole the half opening angle covered by this radio beam is radians for modest emission altitudes $h_{LC} = 2\pi h/Pc$.163 Η the observed radio pulse fills this cone we can write ες in terms of the pulse width ITz2p if the radio emission does not fill the open zone this provides a lower limit for the emission height., If the observed radio pulse fills this cone we can write $h_{LC}$ in terms of the pulse width $W \approx 2\rho$ if the radio emission does not fill the open zone this provides a lower limit for the emission height.164 It has also been shown that. due to a combination of field line sweep back and aberration. the phase of the center of the radio pulse o; should lead the phase of the max PA sweep rate by (Blaskiewiez.Cordes&Wasserman1991:Dyks2008. eg.," It has also been shown that, due to a combination of field line sweep back and aberration, the phase of the center of the radio pulse $\phi_I$ should lead the phase of the max PA sweep rate by \citep[eg.]{bcw91,d08}. ."165).. Observationally we identify o; with the mid-point of the pulse at of its peak and op is identified with oy in an RVM fit., Observationally we identify $\phi_I$ with the mid-point of the pulse at of its peak and $\phi_{\rm PA}$ is identified with $\phi_0$ in an RVM fit.166 The true phase of minimum magnetic angle is between Op and Ορος: ," The true phase of minimum magnetic angle is between $\phi_I$ and $\phi_{\rm d\psi, max}$."167These expressions assume simple static dipoles and low altitudes., These expressions assume simple static dipoles and low altitudes.168 We have checked against detailed numerical simulations of swept-back dipole magnetospheres and find that the actual pulse intensity center and phase of maximum PÀ sweep are both sensitive to. details of the magnetic field structure. especially conditions at the light cylinder that define the edge of the open zone (Craigetal. 2011).," We have checked against detailed numerical simulations of swept-back dipole magnetospheres and find that the actual pulse intensity center and phase of maximum PA sweep are both sensitive to details of the magnetic field structure, especially conditions at the light cylinder that define the edge of the open zone \citep{crj11}."169". These differences are modest at re:<0.05,", These differences are modest at $h_{LC} < 0.05$.170 For objects indicating higher altitude radio emission detailed comparison with the numerical results can be important., For objects indicating higher altitude radio emission detailed comparison with the numerical results can be important.171 In practice. radio pulse profiles may represent “patehy’ illumination of the radio zone (LyneandManchester1988).. even for these young pulsars.," 	In practice, radio pulse profiles may represent `patchy' illumination of the radio zone \citep{lm88}, even for these young pulsars."172 This complicates our estimates of JV and o;., This complicates our estimates of $W$ and $\phi_I$.173 In Table Il. we list both Wyy. the full width of the radio pulse at of the peak intensity and Vy. an estimate of the pulse width at of the peak.," In Table 1, we list both $W_{10}$, the full width of the radio pulse at of the peak intensity and $W_1$, an estimate of the pulse width at of the peak."174 These measurements were made on archival GGHz profiles (see Table references)., These measurements were made on archival GHz profiles (see Table references).175" The VY, estimate is necessarily approximate. especially for the lower S/N pulse profiles."," The $W_1$ estimate is necessarily approximate, especially for the lower S/N pulse profiles."176 At such low flux levels. extended pulsed emission. may be generated by interstellar scattering tails. weak emission components unassociated with the main dipole cap or even non-linearities in the measurement system.," At such low flux levels, extended pulsed emission may be generated by interstellar scattering tails, weak emission components unassociated with the main dipole cap or even non-linearities in the measurement system."177 Nevertheless. for at least a few of these pulsars. this broader width captures weak components of the pulse coming from the principal emission zone.," Nevertheless, for at least a few of these pulsars, this broader width captures weak components of the pulse coming from the principal emission zone."178 Further. in some cases. the assumption of pure dipole geometry and even the identification of the radio beam with the open zone are suspect.," Further, in some cases, the assumption of pure dipole geometry and even the identification of the radio beam with the open zone are suspect."179 However. despite all of these caveats. these radiomeasurements do provide some phenomenological constraints on the range of allowable « and C. even when values for individual pulsars are suspect.," However, despite all of these caveats, these radiomeasurements do provide some phenomenological constraints on the range of allowable $\alpha$ and $\zeta$ , even when values for individual pulsars are suspect."180"A new URD for (νο OD2 using the distance modulus of 10,50 is shown in Figure 12.",A new HRD for Cyg OB2 using the distance modulus of 10.80 is shown in Figure 12.181 Here it can be seen how the location of the dwarf stars are well fit to a 2 million vear old isochrone. as eiven by LeJeune Schaerer (2001). with a spread in age of between 1 aud 3 παοι years old.," Here it can be seen how the location of the dwarf stars are well fit to a 2 million year old isochrone, as given by LeJeune Schaerer (2001), with a spread in age of between 1 and 3 million years old."182 The luminosity spread secu in the B15 and B2 dwarts looks to be greater and in the direction of beiug over-Iuniuous than is secu for the ο aud carlicr B dwarfs., The luminosity spread seen in the B1.5 and B2 dwarfs looks to be greater and in the direction of being over-luminous than is seen for the O and earlier B dwarfs.183 There may be increased contamination from foreground stars at these dwarf star masses which are slightly lower than their hotter dwarts., There may be increased contamination from foreground stars at these dwarf star masses which are slightly lower than their hotter dwarfs.184 As first stated in AIT9L. a reasonably well-defined main sequence is seen in the IIRD of Cre OD2.," As first stated in MT91, a reasonably well-defined main sequence is seen in the HRD of Cyg OB2."185" However, contamination is present in the form of several evolved giant and superelant stars (Fie."," However, contamination is present in the form of several evolved giant and supergiant stars (Fig."186 12)., 12).187" This indicates contamination frou, a non-coeval population within the Cve OB2 sample region.", This indicates contamination from a non-coeval population within the Cyg OB2 sample region.188 Despite this contamination. we can concentrate ou the df stars to estimate the approximate age of the primary constitueuts of the cluster.," Despite this contamination, we can concentrate on the dwarf stars to estimate the approximate age of the primary constituents of the cluster."189 The presence the O5 and O5.5 dwarf stars indicates an age not ereater than 1 to 1.5 ταμοι vears old., The presence the O5 and O5.5 dwarf stars indicates an age not greater than 1 to 1.5 million years old.190 However. the O5 dwarf (A387) comes from the CPR2002 study and has vet to be confirmed as being part of the main (νο OB2 cluster.," However, the O5 dwarf (A37) comes from the CPR2002 study and has yet to be confirmed as being part of the main Cyg OB2 cluster."191 Fiatlermore. the luminosity of the 05.5 V (|NET91] 516). puts it far from the main sequence (or possibly at a nearer distance). making it also suspect as a reliable age measure for the bulk of the cluster.," Furthermore, the luminosity of the O5.5 V ([MT91] 516), puts it far from the main sequence (or possibly at a nearer distance), making it also suspect as a reliable age measure for the bulk of the cluster."192 A well behaved population of dwarf stars is found starting at around OG and OF and by O7.5 and Os there exists a particularly uuuerous sample of dwart stars., A well behaved population of dwarf stars is found starting at around O6 and O7 and by O7.5 and O8 there exists a particularly numerous sample of dwarf stars.193 These mid-O dwarfs put a strong upper limit on the age of the cluster of about 3 willion vears., These mid-O dwarfs put a strong upper limit on the age of the cluster of about 3 million years.194 An age of 2 million vears. with a spread of perlaps one million vears. represents a reasonable age for the bulk of the Cre OB2 cluster based ou its most massive dwarf stars.," An age of 2 million years, with a spread of perhaps one million years, represents a reasonable age for the bulk of the Cyg OB2 cluster based on its most massive dwarf stars."195 hDudeed. a vouug age of not more than a few million vears is needed to explain several very high bhuuinositv blue supereiant stars (such as (νο OB2 #77. the O3 If) and the possible WolfRavet star menibers. the positions of which are also shown in 112.," Indeed, a young age of not more than a few million years is needed to explain several very high luminosity blue supergiant stars (such as Cyg OB2 7, the O3 If) and the possible Wolf-Rayet star members, the positions of which are also shown in 12."196 The Wolt-Bavet stars where placed on the WRD in 112 based on observations of the stars apparent maguitude and line of sieht extinction as measured by Schutz Vacca (1991) for WRILS: Dougherty ot (2000) for WRIIG and Massey et ((2001) and references therein for WR Ltt., The Wolf-Rayet stars where placed on the HRD in 12 based on observations of the stars apparent magnitude and line of sight extinction as measured by Schmutz Vacca (1991) for WR145; Dougherty et (2000) for WR146 and Massey et (2001) and references therein for WR 144.197 The temperature of the three WR stars were sinaply asstuned to be 50.000 IK. The temperature. apparent magnitude aud line of sight extinction of MAVC3LOA was estimated by Iofinaun et ((2002).," The temperature of the three WR stars were simply assumed to be 50,000 K. The temperature, apparent magnitude and line of sight extinction of MWC349A was estimated by Hofmann et (2002)."198 Citing proximity areuments. Massey. et ((2001) dismissed the Van der IIucht et (01981) claim that all που WR stars where members of (νο OB2.," Citing proximity arguments, Massey et (2001) dismissed the Van der Hucht et (1981) claim that all three WR stars where members of Cyg OB2."199 Ouly WR LLL is near the optically distinct cluster as first studied by MT91., Only WR 144 is near the optically distinct cluster as first studied by MT91.200 However. if a new exteuded radius for the Cre OD2 cluster is to be considered. then WR 115. WR L16 and possibly NWC 319À should be re-cousidered as possible iienibers.," However, if a new extended radius for the Cyg OB2 cluster is to be considered, then WR 145, WR 146 and possibly MWC 349A should be re-considered as possible members."201 Let's cousider the possibility of non-coeval ποιους appearing in the IRD shown in 112., Let's consider the possibility of non-coeval members appearing in the HRD shown in 12.202" Over the temperature rauge Log T,;; = L30 to L37. four carly D. stars lie well above the main sequence. A39 (B2 V). 612 (BI IID. 575 (B1.5V). and 793 (DBI.5IIT?)"," Over the temperature range Log $_{eff}$ = 4.30 to 4.37, four early B stars lie well above the main sequence, A39 (B2 V), 642 (B1 III), 575 (B1.5V), and 793 (B1.5III?)"203 in 112., in 12.204 Two show near mfrared excesses (575. 793). consistent with them being Be stars.," Two show near infrared excesses (575, 793), consistent with them being Be stars."205 A39 is likely foreground., A39 is likely foreground.206" Au additional 11 giant and supereiaut stars. Iving mostly in the tempcrature rauge Log T.7, = L10 to 153. are undeniably older than the few milliiou vears we've assigned to the dwarf stars."," An additional 11 giant and supergiant stars, lying mostly in the temperature range Log $_{eff}$ = 4.40 to 4.53, are undeniably older than the few million years we've assigned to the dwarf stars."207 Of these eleven older. evolved stars. over half of them originated," Of these eleven older, evolved stars, over half of them originated"208disk truncatiou as a universal phenomenon aud do not. in general explain how and why truucations nüeht occur.,"disk truncation as a universal phenomenon and do not, in general, explain how and why truncations might occur."209 If we assume that a significant fraction of prescut-day cluster SOs originally had triucatious. then we cau look for something which could erase truucations.," If we assume that a significant fraction of present-day cluster S0s originally had truncations, then we can look for something which could erase truncations."210 One possibility is “larassimeut™ (e.g.Mooreetal.1996 1999).. where a galaxys motion through the (evolving) cluster potential leads to repeated tidal shocks.," One possibility is “harassment” \citep[e.g.,][]{moore96,moore99}, where a galaxy's motion through the (evolving) cluster potential leads to repeated tidal shocks."211 The “high-surtace-brightuess” model galaxy in Mooreetal.(1999.theirFie.?) had its imitially sinele-exponcutial xofile transformed iuto a wild antitruncation by this xocess: this suggests that a Type IT profile could “flatten out” iuto a Type I profile (butseeCuedin2003).," The “high-surface-brightness” model galaxy in \citet[][their Fig.~7]{moore99} had its initially single-exponential profile transformed into a mild antitruncation by this process; this suggests that a Type II profile could “flatten out” into a Type I profile \citep[but212see][]{gnedin03}."213. More detailed simulations involving disks with initially runcated profiles are needed to see if this is a viable uechanisui. (, More detailed simulations involving disks with initially truncated profiles are needed to see if this is a viable mechanism. (214Of course. this cannot be the wav to orn Tvpe I profiles. since they are also fouud iu the field.),"Of course, this cannot be the way to form Type I profiles, since they are also found in the field.)"215" Alternatively, wo can couskder cluster-based uechanisius which ο prevent Type II profiles roni foriung in the first place."," Alternatively, we can consider cluster-based mechanisms which might prevent Type II profiles from forming in the first place."216 If as suggested by Erwinetal.(2008) and Exwiuctal.(2012).. truncations iu SOs and carly-type spirals are predonmünantlv related o the same sort of barOuter Lindblad Resonance (OLR) interactions which produce outer rings. then oue possible scenario nüght be the following.," If, as suggested by \citet{erwin08} and \citet{erwin12}, truncations in S0s and early-type spirals are predominantly related to the same sort of bar–Outer Lindblad Resonance (OLR) interactions which produce outer rings, then one possible scenario might be the following."217 The long dynamical times i outer disks sugeest — and simulations agree — that large-scale chauges there such as outer-ring formation require several Cir (seereferencesinButa&Combes1996)., The long dynamical times in outer disks suggest – and simulations agree – that large-scale changes there such as outer-ring formation require several Gyr \citep[see references in][]{buta96}.218. Since bi-OLR interactions are strengthened by the presence of sieuficant eas in the outer disk (gas. beiug dvuamically cooler than the stellar disk. more readily absorbs augular momentum from the bar). the of eas should weaken OLR effects.," Since bar-OLR interactions are strengthened by the presence of signficant gas in the outer disk (gas, being dynamically cooler than the stellar disk, more readily absorbs angular momentum from the bar), the of gas should weaken OLR effects."219 One possible signature of this might be a teudency of eas-deficicnt spirals to less likely to have Type II profiles aud outer If SOs iu the field were able to retain eas in thei outer disk for longer periods. they would be more likely to show the effects of har-OLR interactions and develop Twpe II profiles.," One possible signature of this might be a tendency of gas-deficient spirals to be less likely to have Type II profiles and outer If S0s in the field were able to retain gas in their outer disk for longer periods, they would be more likely to show the effects of bar-OLR interactions and develop Type II profiles."220 S0. ealaxies in Virgo. on the other hand. could have lost their eas xwtieululv iu the outer disk — eurlier on due to. e.g.. a conibiation of ran-pressure strippius ancl straugulation.," S0 galaxies in Virgo, on the other hand, could have lost their gas – particularly in the outer disk – earlier on due to, e.g., a combination of ram-pressure stripping and strangulation."221 We thank Dave Wiliiau for helpful couuneuts. Micliacl Pollen for the initial version of the catalog-parsing codo. and the referee for sugecsting the possible relevance of depletion.," We thank Dave Wilman for helpful comments, Michael Pohlen for the initial version of the catalog-parsing code, and the referee for suggesting the possible relevance of depletion."222 P.E. was supported by DEG Priority Program 1177., P.E. was supported by DFG Priority Program 1177.223wavelengths.,wavelengths.224 The HI images in conjunction with the optical data will be used to investigate a variety of scientific questions including he star formation feedback on the neutral ISM. threshold for star ormation. baryonic TF relation and dark matter distribution in ow mass galaxies.," The HI images in conjunction with the optical data will be used to investigate a variety of scientific questions including the star formation feedback on the neutral ISM, threshold for star formation, baryonic TF relation and dark matter distribution in low mass galaxies."225 The optical properties of the FIGGS sample. GMRT observations and the main science drivers for the survey are described.," The optical properties of the FIGGS sample, GMRT observations and the main science drivers for the survey are described."226 The GMRT integrated HI column density maps and he HI spectra for the sample galaxies are presented., The GMRT integrated HI column density maps and the HI spectra for the sample galaxies are presented.227 The global HI properties of the FIGGS sample. derived from the GMRT observations. and their comparison with the optical properties of he sample galaxies are also presented.," The global HI properties of the FIGGS sample, derived from the GMRT observations, and their comparison with the optical properties of the sample galaxies are also presented."228 A detailed comparison of he gas distribution. Kinematics and star formation in the sample galaxies will be presented in the companion papers.," A detailed comparison of the gas distribution, kinematics and star formation in the sample galaxies will be presented in the companion papers."229 The observations presented in this paper were made with the Giant Metrewave Radio Telescope (GMRT)., The observations presented in this paper were made with the Giant Metrewave Radio Telescope (GMRT).230 The GMRT is operated by the National Center for Radio Astrophysics of the Tata Institute of Fundamental Research., The GMRT is operated by the National Center for Radio Astrophysics of the Tata Institute of Fundamental Research.231 Partial support for this work was provided by ILTP grant B-3.13., Partial support for this work was provided by ILTP grant B-3.13.232We exclude the PC2 images [rom our analvsis.,We exclude the PC2 images from our analysis.233 This guarantees a homogeneous V. and { photometry dataset. aud represents only a small loss in imaged area.," This guarantees a homogeneous $V$ and $I$ photometry dataset, and represents only a small loss in imaged area."234" The 3 WEs are 800x800 pixel arravs with a pixel size of 0"".1.", The 3 WFs are $\times$ 800 pixel arrays with a pixel size of $^{\prime\prime}$ .1.235 The resulting total field-of-view is approximately 32 Q'.," The resulting total field-of-view is approximately 32 $\sq^{\,\prime}$."236 The ZST/WF instrumental F555W and F8I4JW. magnitudes were derived with DAOPIIOT IL. and calibrated to Landolts V. and £ with an aceuracy of 0.020.03 mag following the usual procedures (Alcock et al.," The /WF instrumental F555W and F814W magnitudes were derived with DAOPHOT II, and calibrated to Landolt's $V$ and $I$ with an accuracy of 0.02–0.03 mag following the usual procedures (Alcock et al."237 2001)., 2001).238 Nearly all of the LAIC clump giants were delected with very. high signal-to-noise., Nearly all of the LMC clump giants were detected with very high signal-to-noise.239 In order to check our V. and 7 zero points. we cross-correlated. a subset of our data with OGLE II ground-based (Udalski et al.," In order to check our $V$ and $I$ zero points, we cross-correlated a subset of our data with OGLE II ground-based (Udalski et al."240 2000)., 2000).241 A comparison reveals zero-point differences of AW=0.11 and A=0.02 mag in the sense that OGLE II is brighter., A comparison reveals zero-point differences of $\Delta V = 0.11$ and $\Delta I = 0.02$ mag in the sense that OGLE II is brighter.242 Aleoek et al. (, Alcock et al. (2431999) compared MACIIO eround-based photometry wilh an early reduction of these same standardized LST/WE data (a dilferent subset). and found that NLACIIO was brighter bv AW=0.06 and AZ=0.00 mag.,"1999) compared MACHO ground-based photometry with an early reduction of these same standardized /WF data (a different subset), and found that MACHO was brighter by $\Delta V = 0.06$ and $\Delta R = 0.00$ mag."244 We speculate that the aperture corrections applied to the OGLE II and MACIIO V. data may be affected at the level bv crowding errors., We speculate that the aperture corrections applied to the OGLE II and MACHO ground-based $V$ data may be affected at the level by crowding errors.245 For example. it is possible that the W-band “sky” around bright stars used to determine the aperture corrections was on average underestimated. because nearby neighbor stars were oversubtiracted. (Alcock οἱ al.," For example, it is possible that the $V$ -band “sky” around bright stars used to determine the aperture corrections was on average underestimated because nearby neighbor stars were oversubtracted (Alcock et al."246 1999: Udalski et al., 1999; Udalski et al.247 2000)., 2000).248 We tentativelv adopt a formal systematic calibration error of 0.02 mag in V. and 7 based on the stated accuracy of /WFE calibrations., We tentatively adopt a formal systematic calibration error of 0.02 mag in $V$ and $I$ based on the stated accuracy of /WF calibrations.249 We detected approximately 36.000 stars in A and 109.000 stars in V. and /.," We detected approximately 36,000 stars in $K$ and 109,000 stars in $V$ and $I$."250 Of these approximately 27.000 have A/< 20. and 7.000 have V.«&21.," Of these approximately 27,000 have $K < 20$ , and 7,000 have $V < 21$."251 Cross-correlating these latter source lists vields 4745 stars with A. /. and V. mags.," Cross-correlating these latter source lists yields 4745 stars with $K$, $I$, and $V$ mags."252 Figure 1 shows the resulting A.(V —A) colormagnitude diagram (CMD) in (he region around the red clamp., Figure 1 shows the resulting $K$ $(V-K)$ color-magnitude diagram (CMD) in the region around the red clump.253 Only 2353 stars appear within the limits of Fig., Only 2353 stars appear within the limits of Fig.254 1. and we restrict our subsequent analyses to these stars.," 1, and we restrict our subsequent analyses to these stars."255 As discussed bv Alves (2000: and refs., As discussed by Alves (2000; and refs.256 therein). not all of the stars appearing in Fig.," therein), not all of the stars appearing in Fig."257 1 are bona fide red clump eiants., 1 are bona fide red clump giants.258 The first-ascent. red. giant branch (RGB) is identified as the roughly vertical sequence running [rom (V.—A)ου 22al Nο183.5 to (V—N)e 3.1 at ANe14.5., The first-ascent red giant branch (RGB) is identified as the roughly vertical sequence running from $(V-K) \sim$ 2.2 at $K \sim 18.5$ to $(V-K) \sim$ 3.1 at $K \sim 14.5$.259 It has a width of about 0.2 mag in (V.—A) color. and no distnet component branches are evident.," It has a width of about 0.2 mag in $(V-K)$ color, and no distinct component branches are evident."260 Some of these stars are second-ascent asvmptotie giant branch (AGB) stars., Some of these stars are second-ascent asymptotic giant branch (AGB) stars.261 The distinction between an RGB and AGB star is not necessarily clear ina CMD of mixed-age field populations like Chis one., The distinction between an RGB and AGB star is not necessarily clear in a CMD of mixed-age field populations like this one.262 The sought-after horizontal branch ved clamp is the most prominent feature in Fie., The sought-after horizontal branch red clump is the most prominent feature in Fig.263 l., 1.264 It appears at A.17 and lies mostly blueward of the giant branch. iie. (V—A)S2.4.," It appears at $K \sim 17$ and lies mostly blueward of the giant branch, i.e. $(V-K) \simlt 2.4$."265 The detailed structure of (he chunp seen here is well understood in the context of stellar evolution theory (Girardi Salaris 2001)., The detailed structure of the clump seen here is well understood in the context of stellar evolution theory (Girardi Salaris 2001).266 The overdensity of stars with A~IT and colorsthat associate them with the giant branch (i.e. V—NK2.4 to 2.5), The overdensity of stars with $K \sim 17$ and colorsthat associate them with the giant branch (i.e. $V-K \sim 2.4$ to 2.5)267A method of redshift estimation. based on the AW: diagram of the BCRR. οςο. 6C* and TORS radio galaxies has been developed.,"A method of redshift estimation, based on the $K-z$ diagram of the 3CRR, 6CE, 6C* and 7CRS radio galaxies has been developed."268 Redshift: probability. density functions vs* sources which are identified with a near-infrared counterpart. ie. for 66 of the 68 members of the sample.," Redshift probability density functions are derived for all of the 6C** sources which are identified with a near-infrared counterpart, i.e. for 66 of the 68 members of the sample."269 Comparison of the resulting redshift estimates with the subset of spectroscopic redshifts shows that our method. is [airlv. robust whenever emission-Iine and/or non-stellar contributions t« the A-magnitudes can be neglected., Comparison of the resulting redshift estimates with the subset of spectroscopic redshifts shows that our method is fairly robust whenever emission-line and/or non-stellar contributions to the $K$ -magnitudes can be neglected.270 Phe estimated: redshift) distribution has a median redshift of zsc1.6.," The estimated redshift distribution has a median redshift of $z_{\rm med} \simeq 1.6271$."272 However. we find that the quasars have their redshifts significantly uncder-estimated by our method.," However, we find that the quasars have their redshifts significantly under-estimated by our method."273 This is explained by the fact that the method is based on the Az relation. which is only valid. for racio ealaxies.," This is explained by the fact that the method is based on the $K-z$ relation, which is only valid for radio galaxies."274 Removing the quasars from the distribution results in à median estimated redshift of z=1.7., Removing the quasars from the distribution results in a median estimated redshift of $z \approx 1.7$.275 This is similar to that of the 6C* sample (tuner7 19) and is significantly higher than that of unfiltered. complete surveys at the same Hux censity limit.," This is similar to that of the 6C* sample $z_{\rm med} \approx 1.9$ ) and is significantly higher than that of unfiltered, complete surveys at the same flux density limit."276 We conclude that the filtering criteria were ellective in biasing the 6C sample to objects at high-redshift., We conclude that the filtering criteria were effective in biasing the 6C** sample to objects at high-redshift.277 The τούς clistribution of the most luminous sources sample is compared. with the predictions. of the stecp-spectrum RLF model of Jarvis et al. , The redshift distribution of the most luminous sources in the 6C** sample is compared with the predictions of the steep-spectrum RLF model of Jarvis et al. (27820010].,2001c).279 We find that the 6C°* data is consistent with a constant co-moving space density at ooο)2.5. and mocerate declines bv factors of ~ 4 can be excluded at the ~23o level.," We find that the 6C** data is consistent with a constant co-moving space density at $z280\,\,\gtsim\,\, 2.5$, and moderate declines by factors of $\sim$ 4 can be excluded at the $\sim 2 - 3\sigma$ level."281 Although Jarvis et al. (, Although Jarvis et al. (2822001€) excluded: these declines at the ~deo level. the additional data from 6€* brovide an independent measure.,"2001c) excluded these declines at the $\sim 4\sigma$ level, the additional data from 6C** provide an independent measure."283 Thus. the two independent studies are in quantitative agreement with the result that any decline at high recdshift is shallow.," Thus, the two independent studies are in quantitative agreement with the result that any decline at high redshift is shallow."284 We note that our. result is. based. on a redshift distribution which is uncertain for the following reasons: (i) a significant fraction of the sample is not. identified spectroscopically: Gi) the method of redshift estimation relies on the limited. statisties of the 3CRR. GCE. TORS and 6C radio galaxies at z73 (the scatter in the Ac diagram at these high redshifts is still poorly defined): and (ii) the redshift’ estimates of quasars are systematically uncder-estimated by our method.," We note that our result is based on a redshift distribution which is uncertain for the following reasons: (i) a significant fraction of the sample is not identified spectroscopically; (ii) the method of redshift estimation relies on the limited statistics of the 3CRR, 6CE, 7CRS and 6C* radio galaxies at $z > 3$ (the scatter in the $K-z$ diagram at these high redshifts is still poorly defined); and (iii) the redshift estimates of quasars are systematically under-estimated by our method."285 Most. of these are likely to lead to an under-estimate of the true median. redshift of the complete sample., Most of these are likely to lead to an under-estimate of the true median redshift of the complete sample.286 Fhus. although spectroscopically incomplete. with the 6C** sample we have additional strong constraints on the high-redshift space density. with a sample that increases the number of powerful steep-spectrum. sources. from complete samples at z22. bv a [actor of  2.," Thus, although spectroscopically incomplete, with the 6C** sample we have additional strong constraints on the high-redshift space density, with a sample that increases the number of powerful steep-spectrum sources, from complete samples at $z > 2$, by a factor of $\sim$ 2."287 The work presented here could be significantly improved bv obtaining spectroscopic. redshifts for a larger fraction of sources in the 6C** sample., The work presented here could be significantly improved by obtaining spectroscopic redshifts for a larger fraction of sources in the 6C** sample.288 This would be particularly important for the faintest sources (Az I9mmag). since these are the most probable z2 candidates.," This would be particularly important for the faintest sources $K \,\,\gtsim\,\, 19$ mag), since these are the most probable $z > 2$ candidates."289 With spectra of these sources we should be able to obtain a tighter constrain on the co-moving space density at z>>2., With spectra of these sources we should be able to obtain a tighter constrain on the co-moving space density at $z > 2$.290 We thank Isobel Hook and. Ross MeLure for very useful comments., We thank Isobel Hook and Ross McLure for very useful comments.291 MJC€ acknowledges the support. from. the Portuguese FundactikzmarkmainBoclyEnd6950°mainDBodsStart6051a00 para a Cienneia e a Tecnologia. and the receipt of a NOVA Marie Curie Early Stage ‘Training Fellowship.," MJC acknowledges the support from the Portuguese Funda\c{c}\\tikzmark{mainBodyEnd6950}\~\tikzmark{mainBodyStart6951}{a}oo para a Ciênncia e a Tecnologia, and the receipt of a NOVA Marie Curie Early Stage Training Fellowship."292 She also eratefully acknowledges the generous hospitality of the. Institute for Computational Cosmology and the Extragalactic Cosmology. Research Croup. at Durham. University.," She also gratefully acknowledges the generous hospitality of the Institute for Computational Cosmology and the Extragalactic Cosmology Research Group, at Durham University."293 ΙΝΔΙΟ acknowledges the Roval Society for a University Research Fellowship., KMB acknowledges the Royal Society for a University Research Fellowship.294than the quiescent bin bv =Sa-level. where the quiescent bin is the müniumn time bin m cach observation.,"than the quiescent bin by $\ge 5\sigma$ -level, where the quiescent bin is the minimum time bin in each observation."295 Under this criterion. we detect six. three. and two flares from DoAr 21. ROXs 21 and ROXs 31. respectively.," Under this criterion, we detect six, three, and two flares from DoAr 21, ROXs 21 and ROXs 31, respectively."296" These are labelled as ""E ""El ἝἜληον αμα ""CES with arrows in Figure 1d and 2.. while the quiesceut pliases are indicated by ""QU."," These are labelled as “F”, “F1”, “F2”, and “F3” with arrows in Figure \ref{fig:lc_c1} and \ref{fig:lc_a1-3}, while the quiescent phases are indicated by “Q”."297 An extremely large flare with the maxima flux of ~LOO times higher than the quiescent level is detected from ROXs 31 (Figure 2ec) aud is separately treated in §33.[., An extremely large flare with the maximum flux of $\sim$ 100 times higher than the quiescent level is detected from ROXs 31 (Figure \ref{fig:lc_a1-3}c c) and is separately treated in 3.4.298 The backerounc-subtracted ACTS spectra of cach source using all the data of obs Cl are shown in Figure 3.., The background-subtracted ACIS spectra of each source using all the data of obs C1 are shown in Figure \ref{fig:spec}.299 Several enüssiou lines from highly ionized elements such as O. Ne. Na. Mg. Si. S. Ar. Ca. and Fe iare secu. hence the spectra are characteristic of a thin thermal plasima.," Several emission lines from highly ionized elements such as O, Ne, Na, Mg, Si, S, Ar, Ca, and Fe are seen, hence the spectra are characteristic of a thin thermal plasma."300 We therefore fit the spectra with a thin thermal plasina model allowing the abundance of each clement with prominent lines (indicated by arrows in Figure 3)) to be free., We therefore fit the spectra with a thin thermal plasma model allowing the abundance of each element with prominent lines (indicated by arrows in Figure \ref{fig:spec}) ) to be free.301 For the study of time variation of the spectra. we separately fit the flare and quiescent spectra as shown in Figure 1..," For the study of time variation of the spectra, we separately fit the flare and quiescent spectra as shown in Figure \ref{fig:lc_c1}."302 A sinele-temiperature (1-T) thin-thermal plasiua model (Alewe 1985) is not acceptable for all the spectra except DoAr 21 and the flare plhiase of ROXs 31., A single-temperature (1-T) thin-thermal plasma model (Mewe 1985) is not acceptable for all the spectra except DoAr 21 and the flare phase of ROXs 31.303 We therefore fit. to-teniperature (2-T) plasma models for ROXs 21 and RONs 31 and find acceptable fits., We therefore fit two-temperature (2-T) plasma models for ROXs 21 and ROXs 31 and find acceptable fits.304 Table 2. shows the best-fit paramcters., Table \ref{tab:spec} shows the best-fit parameters.305 All the clemental abundances are below the solar photospheric values (Anders&CGrevesse1989) except for Ne in DoAr 21 and Na in ROXs 21., All the elemental abundances are below the solar photospheric values \citep{Anders1989} except for Ne in DoAr 21 and Na in ROXs 21.306 The best-fit abundances are displaved in Fieure L with the order of the FIP., The best-fit abundances are displayed in Figure \ref{fig:abund_doar21} with the order of the FIP.307 We investigate the δα overabundance secu in RONs 21 in further detail. because Na is a relatively rare clement compared with the major IN nuclei aud the Na Lyuiuro line lies near the line of the more abundant clement Ne.," We investigate the Na overabundance seen in ROXs 21 in further detail, because Na is a relatively rare element compared with the major 4N nuclei and the Na $\alpha$ line lies near the $\beta$ line of the more abundant element Ne."308" Figure baa shows the best-fit result when the Na mudauce is correctively varied with the ""other"" clucuts.", Figure \ref{fig:spec_roxs21_1keV}a a shows the best-fit result when the Na abundance is correctively varied with the “other” elements.309 À sienificant data excess is found at 2ο.» keV. uear the Lyiman-.) line of Ne (1.211 keV).," A significant data excess is found at $\approx$ 1.2 keV, near the $\beta$ line of Ne (1.211 keV)."310 The intensity ratio of Lyman.) to -o lines of Ne should increase with increasing plasma temperature. but stay almost constant at zz 7 for plasma temperatures hieher than 1 keV (Alewe 1985).," The intensity ratio of $\beta$ to $\alpha$ lines of Ne should increase with increasing plasma temperature, but stay almost constant at $\approx$ 7 for plasma temperatures higher than 1 keV (Mewe 1985)."311 Therefore 1/7 of the Lyian-a flux should be au upper-Iiuüt to the Lymanu-.) line ux., Therefore 1/7 of the $\alpha$ flux should be an upper-limit to the $\beta$ line flux.312 Since the flux of the Ne Lyian-a line ds failv strong. the line flix is well coustraincd as seen in Fig Sa.," Since the flux of the Ne $\alpha$ line is fairly strong, the $\beta$ line flux is well constrained as seen in Fig 5a."313 Thus the excess flux. should originate from another clement: the Lyiman-a line of Na (1.236 keV) is the best candidato., Thus the excess flux should originate from another element; the $\alpha$ line of Na (1.236 keV) is the best candidate.314 Allowing the Na abundance to be a free parameter iuproved the fit as demonstrated in Figure Shh. The reduction of 47 (4.0.f.) from 111 (97) to 108 (96) is significant with ~98 confidence level for the F-test (Bevineton&Robinson1992)., Allowing the Na abundance to be a free parameter improved the fit as demonstrated in Figure \ref{fig:spec_roxs21_1keV}b b. The reduction of $\chi^2$ $d.o.f.$ ) from 114 (97) to 108 (96) is significant with $\sim$ 98 confidence level for the $F$ -test \citep{Bevington1992}.315. Using the CUS spectra of DoAr 21. we investigate the lousg-teriü (ντ) behavior of the spectral piriuueters.," Using the GIS spectra of DoAr 21, we investigate the long-term $\sim$ yr) behavior of the spectral parameters."316 Since the euerev resolution does not allow us to separate the emission lines from tle kev elements. we correctively varied the abundiauces.," Since the energy resolution does not allow us to separate the emission lines from the key elements, we correctively varied the abundances."317" The best-fit huninosity shows long-term variability in the range (28)<10 ere i, ", The best-fit luminosity shows long-term variability in the range $ (2-8)\times10^{31}$ erg $^{-1}$ .318The best-fit temperatures aud abundances also vary frou 2.5 to { keV and 0.15 to 0.5 solar. respectively: both are possibly correlated to the N-rav. Iuninosity a4. are shown in Figure Gaa aud 6bb. The Ny value is consistent with beiug coustaut. although weak variability within a factor zz 2 cannot boe rejected.," The best-fit temperatures and abundances also vary from 2.5 to 4 keV and 0.15 to 0.5 solar, respectively; both are possibly correlated to the X-ray luminosity as are shown in Figure \ref{lx_kt_abund}a a and \ref{lx_kt_abund}b b. The $N_{\rm H}$ value is consistent with being constant, although weak variability within a factor $\approx$ 2 cannot be rejected."319 We note that although the 1-T model iu the medium resolution N-rav. spectra may lead to an error in abundance by a factor ~2 (Reale.Peres.&Orlando 2001).. such an effect is significantly reduced with 2-T or iuulti-T. models.," We note that although the 1-T model in the medium resolution X-ray spectra may lead to an error in abundance by a factor $\sim$ 2 \citep {Reale2001}, such an effect is significantly reduced with 2-T or multi-T models."320 Since we used. 2-T inodels for 2 out of the 3 sources. the artificial abundance uncertaity should be better constrained.," Since we used 2-T models for 2 out of the 3 sources, the artificial abundance uncertainty should be better constrained."321 For the study of the spectral evolution diving the flare. we make aud fit time-sliced GIS spectra of the giant flare from ROXs 31 in obs A2 (Figure 2cc).," For the study of the spectral evolution during the flare, we make and fit time-sliced GIS spectra of the giant flare from ROXs 31 in obs A2 (Figure \ref{fig:lc_a1-3}c c)."322 The time intervals are shown in Figure 7 (left)., The time intervals are shown in Figure \ref{fig:flare_roxs31} (left).323 The spectra show stroug Iko ΕΣ ος like Fe at z6.7 keV. hence we let the Fe abundance vary.," The spectra show strong $\alpha$ emission of He-like Fe at $\approx$ 6.7 keV, hence we let the Fe abundance vary."324 The flare spectra are well fitted with a 1-T model., The flare spectra are well fitted with a 1-T model.325 The best-fit parameters for cach time interval ive shown in Figure 7 (right)., The best-fit parameters for each time interval are shown in Figure \ref{fig:flare_roxs31} (right).326 As shown in the figure. this fare has an unusual time profile.," As shown in the figure, this flare has an unusual time profile."327 The ποτ curve increases slowly over ~15 ks (see also Skinucr 2000)., The light curve increases slowly over $\sim$ 15 ks (see also Skinner 2000).328" The temperaturereaches its peak value before the flux παπα, and stavs nearly constant durius the decay pliase."," The temperaturereaches its peak value before the flux maximum, and stays nearly constant during the decay phase."329" The absorption (Ny,) and abuucdances davamatically change from quicscent to flare: Nyy in the flare is", The absorption $N_{\rm H}$ ) and abundances dramatically change from quiescent to flare; $N_{\rm H}$ in the flare is330in the p Oph cloud. are found to possess features of a similar nature.,"in the $\rho$ Oph cloud, are found to possess features of a similar nature."331 The two isolated HAEBEs (which might also be called post-HAEBEs since they no longer embedded in nebulosity) HD 31618 aud HD 163296. were shown to have spectra sinilar to that of Hale-Bopp aud Levy 1999).," The two isolated HAEBEs (which might also be called post-HAEBEs since they no longer embedded in nebulosity) HD 31648 and HD 163296, were shown to have spectra similar to that of Hale-Bopp and Levy \citep{sitko99}."332. However. the stellar cisς dust in these systems emitted relatively 1iore strongly near 9.1 jan than the comets.," However, the stellar disk dust in these systems emitted relatively more strongly near 9.4 $\mu$ m than the comets."333 — should be rememberect t1at in the case of a comet. we are dealiug witl material with a sinele well«efined distance fro the sui. while this is no the case for a disk.," It should be remembered that in the case of a comet, we are dealing with material with a single well-defined distance from the sun, while this is not the case for a disk."334 Neve‘theless. have shown that over ar auge of a [actor of three in heliocentric clistaice. the spectral shape of Hae-Bopp was. witl the exceytion of one subfeaure. Invariant.," Nevertheless, \citet{wooden00} have shown that over a range of a factor of three in heliocentric distance, the spectral shape of Hale-Bopp was, with the exception of one subfeature, invariant."335 Furtbermore. a comparison of the ISO s»ecira of the HAEBE star HD 100516 ancl Hale-Bopp indicate that. wlile the overall spectral everey distributious a'e signilicautly different in the way expectec (nore cool dust iu the stellar disk). the shapes of tie narrow [eatures were still quite similar (lalfaitetal.1998)..," Furthermore, a comparison of the ISO spectra of the HAEBE star HD 100546 and Hale-Bopp indicate that, while the overall spectral energy distributions are significantly different in the way expected (more cool dust in the stellar disk), the shapes of the narrow features were still quite similar \citep{malfait98}. ."336 However. he shape of tle 10 p ur band is seusitive to tle abundance aud fkoiu of pyroxenes. as evidenced in the variety «of specral shapes of pyroxene-rich interplanetary ¢ust particles (IDPs) tlat have )een analvzed in the Iaboratory (SaulordandWalker1985:Bradleyetal.1992).," However, the shape of the 10 $\mu$ m band is sensitive to the abundance and form of pyroxenes, as evidenced in the variety of spectral shapes of pyroxene-rich interplanetary dust particles (IDPs) that have been analyzed in the laboratory \citep{sw85,brad92}."337. It is ineresting hat the one subfeattre that observed to vary with heliocentric distauce in Hale-Bopp |las veer attributec| to crystalline pyroxeues (Woodenetal.2000)., It is interesting that the one subfeature that observed to vary with heliocentric distance in Hale-Bopp has been attributed to crystalline pyroxenes \citep{wooden00}.338. This suggests that perlaps he nalu reason [o “the excess 9.Ε jun emission of the disks compared to comets is due to the prese oL dust at temperatures much higher than those generally observed in solar system ComeS. wh allows the crysalliue pyroxenes to be enhanced. but on a grauder scale than was seeu in Hale-Bop," This suggests that perhaps the main reason for the excess 9.4 $\mu$ m emission of the disks compared to comets is due to the presence of dust at temperatures much higher than those generally observed in solar system comets, which allows the crystalline pyroxenes to be enhanced, but on a grander scale than was seen in Hale-Bopp."339 Iu Figure 7 we show the 10 jiu baud of TW Hya. normalized to the continuum by dividiug by a gray body that matches the observed flux at 8 and 13 pain. Also shown are the data for HD 31618. HD 163296. comet Hale-Bopp. aud comet Levy from Sitkoetal.(1999).," In Figure 7 we show the 10 $\mu$ m band of TW Hya, normalized to the continuum by dividing by a gray body that matches the observed flux at 8 and 13 $\mu$ m. Also shown are the data for HD 31648, HD 163296, comet Hale-Bopp, and comet Levy from \citet{sitko99}."340. Although the data for TW Ηνα has a modest SNB. it is apparent that TW Hya does uot have the same shape as the other objects.," Although the data for TW Hya has a modest SNR, it is apparent that TW Hya does not have the same shape as the other objects."341 Iun particular. the well-clefi1ος erystalline olivine feature seen at 11.2 imi in both the comets and HD 381618 aud HD 163296 seems to considerably less-defined in TW Hya.," In particular, the well-defined crystalline olivine feature seen at 11.2 $\mu$ m in both the comets and HD 31648 and HD 163296 seems to considerably less-defined in TW Hya."342 At best. a modes Change of spectral slope is presen at that waveleugthi," At best, a modest change of spectral slope is present at that wavelength."343 In. Figure 8 we compare the spectrtm of TW Ηνα with three HAEBESs (Sitsoetal.2000) whose emission features are represelalive ¢X the wide variety that are seen in tLene PMS objects (excludiug objects dominated by orgaue enSSIOL eatures)., In Figure 8 we compare the spectrum of TW Hya with three HAEBEs \citep{sitko00} whose emission features are representative of the wide variety that are seen in these PMS objects (excluding objects dominated by organic emission features).344 That of RY Ta exhibits a silicate feature whitch is stroug aud relatively featureless. except for a single maxim nlear 9.7 jun. I1 HD 163296 (again) the oliviue feature Is )resell. but lot as prominent as in cojets such as Hae-Bopp aud Levy.," That of RY Tau exhibits a silicate feature which is strong and relatively featureless, except for a single maximum near 9.7 $\mu$ m. In HD 163296 (again) the olivine feature is present, but not as prominent as in comets such as Hale-Bopp and Levy."345 Finally. tlie spectrum o CHD:SD1ST las a weak silicate baud. bul nore coiet-Ise (flat-topped) in shape.," Finally, the spectrum of HD 35187 has a weak silicate band, but more comet-like (flat-topped) in shape."346 Ii Figure 9 we asO COLipare he spectrum of TW Hya [9] that of RW Ar. a PMS Ce star with ci'cumstellar CO einissiOn lu he first vibrational overtoe band (SitkoadHanson2000).," In Figure 9 we also compare the spectrum of TW Hya to that of RW Aur, a PMS G5e star with circumstellar CO emission in the first vibrational overtone band \citep{sh00}."347. Because the streugth of the silicate xul with respectto the uderlying continuun in RW Aur is only one-cuarter that of TW Hya. for," Because the strength of the silicate band with respectto the underlying continuum in RW Aur is only one-quarter that of TW Hya, for"348understood.,understood.349 The X-ray spectrum of iin the Low State can be described. by a hard underlving power law plus a rellection Component (Done et al., The X-ray spectrum of in the Low State can be described by a hard underlying power law plus a reflection component (Done et al.350 1992) and a weak soft. excess (Baluceiisska ancl Llasinger 1991) which has been identified with blackbody emission. [rom he accretion disk (Dalucciüsska-Church et al., 1992) and a weak soft excess (Bałuccińsska and Hasinger 1991) which has been identified with blackbody emission from the accretion disk (Bałuccińsska-Church et al.351 1995)., 1995).352 In he Soft State. the spectrum. is dominated. by à. strong hermal component. produced. by enhanced. emission. [rom he accretion disk (Dotani et al.," In the Soft State, the spectrum is dominated by a strong thermal component produced by enhanced emission from the accretion disk (Dotani et al."353 1996: Cul οἱ αἱ., 1996; Cui et al.354 1997)., 1997).355 X-ray dips are also observed in N-1. which usually last several minutes but have been up to S hrs in length.," X-ray dips are also observed in X-1, which usually last several minutes but have been up to 8 hrs in length."356 During the clips. there is a spectral hardening and the Ix absorption edge of iron may be seen showing that they are due to photoelectrie absorption (e.g. Ixitamoto ct al.," During the dips, there is a spectral hardening and the K absorption edge of iron may be seen showing that they are due to photoelectric absorption (e.g. Kitamoto et al."357 1984)., 1984).358 Spectral fitting of the X-ray spectra of hhas shown that the column density can increase [rom the non-dip value of ~6.0107 LL atom = to —11077 Il atom em 7 in dipping (Baluccitisska-Church οἱ al.," Spectral fitting of the X-ray spectra of has shown that the column density can increase from the non-dip value of $\rm {\sim 6.0\times 10^{21}}$ H atom $^{-2}$ to $\rm {\sim 1\times 10^{23}}$ H atom cm $^{-2}$, in dipping (Bałuccińsska-Church et al."359 1997: Ixitamoto et al., 1997; Kitamoto et al.360 1984)., 1984).361 Spectral. fitting of. dip data. for exemple fromL is consistent with neutral absorber although. ciscriminating between cold. and. warn absorber was not possible (Baluccitisska-Chureh et al.," Spectral fitting of dip data, for example from is consistent with neutral absorber although discriminating between cold and warm absorber was not possible (Bałuccińsska-Church et al."362 1997)., 1997).363 ]xitamoto et al. (, Kitamoto et al. (3641984). from analysis of a high quality spectrum of a dip. found an Fe edge implying an ionization state < be V. Lt was realised at an carly stage that dipping tends to occur at about phase zero of the 5.6 d orbital evele. Le. near to superior conjunction of the black hole. but dipping was also seen. for example. at ó ~0.71 (Remillard Canizares 1984).,"1984) from analysis of a high quality spectrum of a dip, found an Fe edge implying an ionization state $<$ Fe V. It was realised at an early stage that dipping tends to occur at about phase zero of the 5.6 d orbital cycle, i.e. near to superior conjunction of the black hole, but dipping was also seen, for example, at $\phi \sim$ 0.71 (Remillard Canizares 1984)."365 Possible causes of dipping that have been suggested (see Remillard Canizares 1984) ave: 1) absorption asociated with Roche lobe overflow. matter having however. to be far out of the orbital plane: 2) absorption taking place in the stream flowing from the companion towards the compact object inferred from ο 11 À 4686 measurements (Bolton 1975: Treves ct al.," Possible causes of dipping that have been suggested (see Remillard Canizares 1984) are: 1) absorption asociated with Roche lobe overflow, matter having however, to be far out of the orbital plane; 2) absorption taking place in the stream flowing from the companion towards the compact object inferred from He II $\lambda$ 4686 measurements (Bolton 1975; Treves et al."366 1980): 3) absorption in the wind of the companion 4) absorption in obs in the wind of the companion (Ixitamoto et al., 1980); 3) absorption in the wind of the companion 4) absorption in blobs in the wind of the companion (Kitamoto et al.367 1984)., 1984).368 Ixitamoto et al., Kitamoto et al.369" used the duration of short dips to estimate he size of an absorbing cloud as  10"" em. Le. à relatively small region. and hence associated the absorber with clouds or “blobbiness” in the stellar wind of the companion."," used the duration of short dips to estimate the size of an absorbing cloud as $\sim$ $^9$ cm, i.e. a relatively small region, and hence associated the absorber with clouds or “blobbiness” in the stellar wind of the companion."370 The esent position is that the physical state and origin of the absorber is not at all well understood., The present position is that the physical state and origin of the absorber is not at all well understood.371 This contrasts with N-rav. dipping in low mass X-ray rinaries in which it is generally accepted that dipping is due o absorption in the bulge in the outer accretion dise where he aceretion low from the companion impacts (White Swank 1982)., This contrasts with X-ray dipping in low mass X-ray binaries in which it is generally accepted that dipping is due to absorption in the bulge in the outer accretion disc where the accretion flow from the companion impacts (White Swank 1982).372 Spectral evolution in dipping could not »: explained in terms of absorption of a single emission component. since in particular sources. the spectrum may xcome harder. remain energy. independent or even become softer during dipping.," Spectral evolution in dipping could not be explained in terms of absorption of a single emission component, since in particular sources, the spectrum may become harder, remain energy independent or even become softer during dipping."373 However. this behaviour has been explained. by assuming (wo emission regions: point-like Xxackbody. emission. [rom the surface of the neutron star plus extended: Comptonized emission. from the accretion disk corona (Church Batueciisska-Chureh 1995: Church et al.," However, this behaviour has been explained by assuming two emission regions: point-like blackbody emission from the surface of the neutron star plus extended Comptonized emission from the accretion disk corona (Church Bałuccińsska-Church 1995; Church et al."374 1997. 1998a. 1998b).," 1997, 1998a, 1998b)."375 This model is able to explain the varied and complex spectral evolution in dipping in different sources., This model is able to explain the varied and complex spectral evolution in dipping in different sources.376 The presence of X-ray eclipses in. 0748-676 shows that deepest. dipping occurs at orbital phase ~0.9 consistent with the position of impact on the disk of an accretion [low trailing sideways from the inner Lagrangian point in the binary frame., The presence of X-ray eclipses in 0748-676 shows that deepest dipping occurs at orbital phase $\sim$ 0.9 consistent with the position of impact on the disk of an accretion flow trailing sideways from the inner Lagrangian point in the binary frame.377 Thus. dipping is much better unclerstoocl in LAINBs than inX-I.," Thus, dipping is much better understood in LMXBs than in."378. SXDs in general have strong stellar winds ancl exhibit strong orbital-related decreases in. X-ray intensity due to absorption in the wind., SXBs in general have strong stellar winds and exhibit strong orbital-related decreases in X-ray intensity due to absorption in the wind.379 In an/ZXOSAT observation of the archetypal eclipsing SND. 1700-371. it was. found that large. smooth increases in column density took place between orbital phases 0.8 and 1.2 (Llaberl White IWwallman 1989: hereafter. HINIxS9).," In an observation of the archetypal eclipsing SXB, 1700-371, it was found that large, smooth increases in column density took place between orbital phases 0.8 and 1.2 (Haberl, White Kallman 1989; hereafter HWK89)."380. This observation was useful since a full orbital evele of 3.41 day was covered., This observation was useful since a full orbital cycle of 3.41 day was covered.381 In 1700-371. the binary separation is the smallest known. with the compact object orbiting the primary star at 1.4 stellar radii (van Genderen 1977). leading to the dramatic changes in absorption with orbital phase.," In 1700-371, the binary separation is the smallest known, with the compact object orbiting the primary star at 1.4 stellar radii (van Genderen 1977), leading to the dramatic changes in absorption with orbital phase."382 The increases in column density could be well modelled by absorption in a stellar wind obeving a CALS velocity law (Castor. Abbott ία 1975).," The increases in column density could be well modelled by absorption in a stellar wind obeying a CAK velocity law (Castor, Abbott Klein 1975)."383 Xn additional sharp increase in column density at phase ~0.6 could. be well modelled. as a gas. stream originating on the companion. possibly on a tidal bulge (LIWISS9).," An additional sharp increase in column density at phase $\sim$ 0.6 could be well modelled as a gas stream originating on the companion, possibly on a tidal bulge (HWK89)."384 In the companion of N-1. 226868. the radial velocity curve of the Le LL A 4686 emission line is shifted by about 120° with respect to absorption in the companion. also indicating the presence of an accretion stream (llutchines et al.," In the companion of X-1, 226868, the radial velocity curve of the He II $\lambda$ 4686 emission line is shifted by about $\degmark$ with respect to absorption in the companion, also indicating the presence of an accretion stream (Hutchings et al."385 1973)., 1973).386 In the case ofl. a strong decrease in X-ray intensity with orbital phase has not been seen: however. Ixitamoto et al. (," In the case of, a strong decrease in X-ray intensity with orbital phase has not been seen; however, Kitamoto et al. ("3871990) showed evidence for a dependence of column density of the quiescent (non-dip) spectra on orbital phase. but with column density increasing [rom ~6105 LH atom 2 o only 2.10E 11 atom 7. àt least one order of magnitude less than in 1700-371.,"1990) showed evidence for a dependence of column density of the quiescent (non-dip) spectra on orbital phase, but with column density increasing from $\rm {\sim 6\times 10^{21}}$ H atom $^{-2}$ to only $\rm {2\times 10^{22}}$ H atom $^{-2}$, at least one order of magnitude less than in 1700-371."388 Apart [rom this. the only absorption events seen are the X-ray dips.," Apart from this, the only absorption events seen are the X-ray dips."389 There has been no systematic study of the phase of dipping. and in this work we present. such a systematic studs. and compare dipping in wwith absorption ellects in other SXDs.," There has been no systematic study of the phase of dipping, and in this work we present such a systematic study, and compare dipping in with absorption effects in other SXBs."390 lt has not previously been possible to make a survey of the distribution of dips with orbital phase because of uncertainties in the available ephemeris., It has not previously been possible to make a survey of the distribution of dips with orbital phase because of uncertainties in the available ephemeris.391 Phe ephemoeris previously available of Gics Bolton (1982) gave an orbital period of 5.59974 +0.00008 cays. and this precision implies an uncertainty in phase of 40.02 evcle at the present time.," The ephemeris previously available of Gies Bolton (1982) gave an orbital period of 5.59974 $\pm $ 0.00008 days, and this precision implies an uncertainty in phase of $\pm$ 0.02 cycle at the present time."392 Llowever. Ninkov. Walker Yang (1987) presented evidence [or a period increase with time. and phases caleulated. using," However, Ninkov, Walker Yang (1987) presented evidence for a period increase with time, and phases calculated using"393Vjin Fig.9as resultsfrom (his analysis. For comparison. we showthe CTEQ3 inpu,parton distribution functions at $Q_0=1.6$ GeV of a global QCD analysis for various hard394Vjin Fig.9as resultsfrom (his analysis. For comparison. we showthe CTEQ3 input,parton distribution functions at $Q_0=1.6$ GeV of a global QCD analysis for various hard395differences just by assuming a fixed luminosity Linactive ): These values are also listed in Table Ι..,differences just by assuming a fixed luminosity $L_{\mathrm{active}}\simeq L_{\mathrm{inactive}}$ ): These values are also listed in Table \ref{tab_means}.396 Since. differences in. radius are computed through temperature ratios and we are curying out a differential study. the results are almost independent of the 7r scale adopted. which is still controversial for M-type stars.," Since differences in radius are computed through temperature ratios and we are carrying out a differential study, the results are almost independent of the $T_{\rm eff}$ scale adopted, which is still controversial for M-type stars."397 A cross-check with the Zr calibrations of Schmidt-Kaler (1982) and de Jager Nieuwenhuijzen (1987) expectedly yielded the same results., A cross-check with the $T_{\rm eff}$ calibrations of Schmidt-Kaler (1982) and de Jager Nieuwenhuijzen (1987) expectedly yielded the same results.398 The trend of lower temperatures and larger radii for active stars has high statistical significance since Ssimilar differentials are obtained in all luminosity bins., The trend of lower temperatures and larger radii for active stars has high statistical significance since similar differentials are obtained in all luminosity bins.399 In Fig., In Fig.400 3 we show the calculated radius differences translated into spectral type bins and together with the radius differences obtained from a sample of eclipsing binaries with data of best quality (taken from Lóppez-Morales et al., \ref{fig_cat2} we show the calculated radius differences translated into spectral type bins and together with the radius differences obtained from a sample of eclipsing binaries with data of best quality (taken from Lóppez-Morales et al.401 2006 and Lóppez-Morales 2007)., 2006 and Lóppez-Morales 2007).402 Note that the two radius differentials have slightly different meanings., Note that the two radius differentials have slightly different meanings.403 While eclipsing binary values come from the direct comparison of radius measurements with the predictions of theoretical models of Baratfe et al. (, While eclipsing binary values come from the direct comparison of radius measurements with the predictions of theoretical models of Baraffe et al. (4041998) — which do not include the effects of stellar activity. — the values for single stars are computed from the difference between active and inactive samples.,"1998) – which do not include the effects of stellar activity, – the values for single stars are computed from the difference between active and inactive samples."405 As can be seen. the differentials from these two completely independent approaches are in very good agreement.," As can be seen, the differentials from these two completely independent approaches are in very good agreement."406" To rule out a possible effect of the My, binning procedure itself on the differences reported. we used an alternative approach using a polynomial fit."," To rule out a possible effect of the $M_{\rm bol}$ binning procedure itself on the differences reported, we used an alternative approach using a polynomial fit."407 We first calculated a best-fitting third order polynomial of the form Ty=fCM4po)) for the inactive star sample., We first calculated a best-fitting third order polynomial of the form $T_{\rm eff}=f(M_{\rm bol})$ for the inactive star sample.408 Then. for each individual active star we calculated the temperature difference (and radius difference) between the observed value and the one predicted by the polynomial fit.," Then, for each individual active star we calculated the temperature difference (and radius difference) between the observed value and the one predicted by the polynomial fit."409 A graphical representation of the results ts provided in Fig. 4..," A graphical representation of the results is provided in Fig. \ref{fig_cat4},"410 where the individual temperatures and radius differences are shown together with à running average of 10 points and averages computed for each My bin., where the individual temperatures and radius differences are shown together with a running average of 10 points and averages computed for each $M_{\rm bol}$ bin.411 The actual values are given in Table 2.. ana are very similar to those in Table 1.. which were calculated from My; bins.," The actual values are given in Table \ref{tab_polynomial}, and are very similar to those in Table \ref{tab_means}, which were calculated from $M_{\rm bol}$ bins."412 An interesting feature of Fig., An interesting feature of Fig.413 4. i8 the relatively large scatter at Moo«8.5 mag. which occurs near the value where models predict the change between fully convective stars and those with a radiative core.," \ref{fig_cat4} is the relatively large scatter at $M_{\rm bol}<8.5$ mag, which occurs near the value where models predict the change between fully convective stars and those with a radiative core."414 A further useful check of the results comes from restricting the analysis to only stars that have been explicitly. classified as single through high-resolution spectroscopy (Gizis et al., A further useful check of the results comes from restricting the analysis to only stars that have been explicitly classified as single through high-resolution spectroscopy (Gizis et al.415 2002)., 2002).416" Although the statistics are less significant (with 127 inactive and 19 active stars). the mean temperature and radius differences for each My, bin are within one sigma of those in Table 1.."," Although the statistics are less significant (with 127 inactive and 19 active stars), the mean temperature and radius differences for each $M_{\rm bol}$ bin are within one sigma of those in Table \ref{tab_means}."417 In Table | we also include the mean differences in the (V—K) colour index. which are a direct consequence of the different average 7. values between active and inactive stars.," In Table \ref{tab_means} we also include the mean differences in the $(V-K)$ colour index, which are a direct consequence of the different average $T_{\rm eff}$ values between active and inactive stars."418" The relatively large scatter of these means (especially that of the first Mj, bin) may be caused by the variability in the V. band of the stars in the active sample.", The relatively large scatter of these means (especially that of the first $M_{\rm bol}$ bin) may be caused by the variability in the $V$ band of the stars in the active sample.419 This stems from the existence of surface spots with various cycles that could affect single-epoch V-band measurements but not K-band measurements that are more immune to spot-induced variability., This stems from the existence of surface spots with various cycles that could affect single-epoch $V$ -band measurements but not $K$ -band measurements that are more immune to spot-induced variability.420 Note that no obvious or only marginal colour differences for active stars have been reported before (see. e.g.. Hawley et al.," Note that no obvious or only marginal colour differences for active stars have been reported before (see, e.g., Hawley et al."421 1996: Bochanski et al., 1996; Bochanski et al.422 2007)., 2007).423 This would seem to stand in contradiction with our results. but it ts not.," This would seem to stand in contradiction with our results, but it is not."424 It is important to emphasise that such colour comparisons are, It is important to emphasise that such colour comparisons are425front itself rather than being an overdensity within the H region.,front itself rather than being an overdensity within the H region.426 The extended photoionisation front and photon-dominated region seen edge-on in MI7 was modelled by ὁ as being supported by magnetic pressure (observations by ? suggest a LOS magnetic field D>100 7G).," The extended photoionisation front and photon-dominated region seen edge-on in M17 was modelled by \citet{PelBalBroEA07} as being supported by magnetic pressure (observations by \citealt{BroTro01}427 suggest a LOS magnetic field $B>100\,\mu$ G)."428 Unfortunately much of this H region is heavily obscured at optical wavelengths and so straightforward inspection of Lla images is not particularly revealing., Unfortunately much of this H region is heavily obscured at optical wavelengths and so straightforward inspection of $\mathrm{H}\alpha$ images is not particularly revealing.429 In the extensive IIa survey of the Carina nebula by ?.. only the elephant trunk in “Pos.," In the extensive $\mathrm{H}\alpha$ survey of the Carina nebula by \citet*{SmiBalWal10}, only the elephant trunk in `Pos."430 23 of fig., 23' of fig.431 | shows evidence of a bright ridge in front of the trunk. suggesting that such features are not common.," 1 shows evidence of a bright ridge in front of the trunk, suggesting that such features are not common."432 A clearer example is seen in NGC 6357 (2). where a prominent bright ridge in front of an elephant trunk is the brightest Lla structure in the field of view., A clearer example is seen in NGC 6357 \citep{BohTapRotEA04} where a prominent bright ridge in front of an elephant trunk is the brightest $\mathrm{H}\alpha$ structure in the field of view.433 A similar structure is found in front of a massive pillar in NGC 3603 (2).., A similar structure is found in front of a massive pillar in NGC 3603 \citep{BraGreChuEA00}.434 In both of these observations. however. the adjacent massive star cluster is expected to drive out-flowing gas and so it is unclear if the bright ridges are ram-pressure or magnetically contined.," In both of these observations, however, the adjacent massive star cluster is expected to drive out-flowing gas and so it is unclear if the bright ridges are ram-pressure or magnetically confined."435 The larger simulation domains used here (compared to MLIO) allow us to follow the evolution of the dense gas for significantly longer. up to 700 Kyr in some simulations.," The larger simulation domains used here (compared to ML10) allow us to follow the evolution of the dense gas for significantly longer, up to $700\,$ kyr in some simulations."436 We find that in the absence of dense gas further from the star. a pillar-like structure will be flattened into a cometary globule with a dense head and low density tail.," We find that in the absence of dense gas further from the star, a pillar-like structure will be flattened into a cometary globule with a dense head and low density tail."437 To study this in more detail. clump configuration 2 was set up with an extra clump a further parsee from the radiation. source.," To study this in more detail, clump configuration 2 was set up with an extra clump a further parsec from the radiation source."438 It was found that the pillar-like stage of the evolution lasted 600 Kyr in this model. until all of the dense gas reached the position of the furthest clump.," It was found that the pillar-like stage of the evolution lasted $\sim600\,$ kyr in this model, until all of the dense gas reached the position of the furthest clump."439 Subsequently a cometary morphology developed again., Subsequently a cometary morphology developed again.440 This experiment shows that while the pillar's lifetime can be extended somewhat. it cannot survive indefinitely unless there is a long dense filament pointing away from the radiation source.," This experiment shows that while the pillar's lifetime can be extended somewhat, it cannot survive indefinitely unless there is a long dense filament pointing away from the radiation source."441 This supports the general picture of elephant trunks shown in fig., This supports the general picture of elephant trunks shown in fig.442" IS of ο as structures which ‘orm and disperse over about 10""yr."," 15 of \citet{SmiPovWhiEA10} as structures which form and disperse over about $10^6\,$ yr."443 With our non-gravitating simulations we cannot model the star formation with also occurs. but the simulations by ? and ? show that the compression induced by RDI produces gravitationally unstable fragments which would ikely form stars.," With our non-gravitating simulations we cannot model the star formation with also occurs, but the simulations by \citet{GriBurNaaEA10} and \citet{BisWhiWueEA10} show that the compression induced by RDI produces gravitationally unstable fragments which would likely form stars."444 While it may seem unlikely that the three pillars in M16 should all have originated from elongated overdensities »ointing back towards the ionising stars. recent observations from he Herschel observatory (e.g.2) have shown that molecular clouds appear to have a distinctly filamentary structure.," While it may seem unlikely that the three pillars in M16 should all have originated from elongated overdensities pointing back towards the ionising stars, recent observations from the Herschel observatory \citep[e.g.][]{MolSwiBalEA10} have shown that molecular clouds appear to have a distinctly filamentary structure."445 Simulations of THD turbulence generated by colliding flows (2). show similar structures. and simulations of small-scale non-ideal isothermal THD turbulence (2?) also show significantly more linear structure compared to purely hydrodynamical models.," Simulations of MHD turbulence generated by colliding flows \citep{BanVazHenEA09}446 show similar structures, and simulations of small-scale non-ideal isothermal MHD turbulence \citep{DowOSul09} also show significantly more linear structure compared to purely hydrodynamical models."447 We have performed a series of R-MHD simulations. of the johoetoionisation of dense clumps of gas and their evolution from villar-like to cometary globule-like structures., We have performed a series of R-MHD simulations of the photoionisation of dense clumps of gas and their evolution from pillar-like to cometary globule-like structures.448 Our results for je emissivity of ionised gas agree very well with those of ?.. Palhowing that a dense. ionised. bar-shaped region standing off from le ionisation front is a generic feature of strongly magnetised johetoionisation in a clumpy medium for a perpendicular field orientation.," Our results for the emissivity of ionised gas agree very well with those of \citet{HenArtDeCEA09}, showing that a dense, ionised, bar-shaped region standing off from the ionisation front is a generic feature of strongly magnetised photoionisation in a clumpy medium for a perpendicular field orientation."449 This ridge can be as bright as. or even brighter than. 1e photoionisation front when observed in recombination radiation e.g. Ho) and its presence or absence can be used as a diagnostic of de strength of any large scale magnetic field which may be present.," This ridge can be as bright as, or even brighter than, the photoionisation front when observed in recombination radiation (e.g. $\mathrm{H}\alpha$ ) and its presence or absence can be used as a diagnostic of the strength of any large scale magnetic field which may be present."450 Bright ridges or ribbons are observed in some H regions (e.g.222). although they are not common.," Bright ridges or ribbons are observed in some H regions \citep[e.g.][]{BraGreChuEA00, BohTapRotEA04, SmiBalWal10}, although they are not common."451 An overdense ridge could also be produced by ram-pressure confinement. and more detailed modelling is required to find observational signatures which could distinguish these different continement mechanisms.," An overdense ridge could also be produced by ram-pressure confinement, and more detailed modelling is required to find observational signatures which could distinguish these different confinement mechanisms."452 Comparing to observations of M16 (?). there is no such ribbon or ridge. suggesting the ambient field measured by ?— is not dominant ia the ionised gas.," Comparing to observations of M16 \citep{HesScoSanEA96} there is no such ribbon or ridge, suggesting the ambient field measured by \citet{SugWatTamEA07} is not dominant in the ionised gas."453 This conclusion is strengthened when we consider the magnetic field orientation observed in MI6 (?.fig.9)., This conclusion is strengthened when we consider the magnetic field orientation observed in M16 \citep[][fig.~9]{SugWatTamEA07}.454 The results presented here show that both RDI and acceleration of clumps by the rocket effect tend to align the magnetic field in dense neutral gas with the radiation propagation direction., The results presented here show that both RDI and acceleration of clumps by the rocket effect tend to align the magnetic field in dense neutral gas with the radiation propagation direction.455 In our models a field configuration. similar to the observed one is clearly seen when the initial field strength is IB)=20μα: the simulation with |D|50Ο is consisten with observations. and the simulation with |D|160j/G is no consistent.," In our models a field configuration similar to the observed one is clearly seen when the initial field strength is $\vert\mathbf{B}\vert\simeq20\,\mu$ G; the simulation with $\vert\mathbf{B}\vert\simeq50\,\mu$ G is consistent with observations, and the simulation with $\vert\mathbf{B}\vert\simeq160\,\mu$ G is not consistent."456 Our simulations thus suggest an ambient field strength of |B]τι50//G around the M16 pillars.," Our simulations thus suggest an ambient field strength of $\vert\mathbf{B}\vert\lesssim 50\,\mu$ G around the M16 pillars."457 The morphology of the structures which develop due to RDI and the rocket effect is also affected by a strong magnetic field. partly due to shielding by the dense ionised ridge and partly by the effect of the field within the pillar or globule.," The morphology of the structures which develop due to RDI and the rocket effect is also affected by a strong magnetic field, partly due to shielding by the dense ionised ridge and partly by the effect of the field within the pillar or globule."458 Inspection of La images of elephant trunks and globules in the literature (e.g.??) suggests that the features seen in the strong field simulations are not common. although the uniform initial field configurations considered here are certainly somewhat artificial.," Inspection of $\mathrm{H}\alpha$ images of elephant trunks and globules in the literature \citep[e.g.][]{HesScoSanEA96, SmiBalWal10} suggests that the features seen in the strong field simulations are not common, although the uniform initial field configurations considered here are certainly somewhat artificial."459 Additionally many H regions have significantly higher gas pressure than that modelled in our simulations. in which case the magnetic field must also be correspondingly stronger to dominate the dynamics.," Additionally many H regions have significantly higher gas pressure than that modelled in our simulations, in which case the magnetic field must also be correspondingly stronger to dominate the dynamics."460 Comparing these results with our earlier simulations in MLIO. the larger simulation domains used here show that the pillar-like structures which form will ultimately evolve to cometary structures in the absence of dense gas further from the star.," Comparing these results with our earlier simulations in ML10, the larger simulation domains used here show that the pillar-like structures which form will ultimately evolve to cometary structures in the absence of dense gas further from the star."461 The lifetimes of pillars in our models are /500 Kvyr. although this depends significantly on the initial mass and concentration (and presumably velocity. cf. 23) ," The lifetimes of pillars in our models are $t \lesssim 500\,$ kyr, although this depends significantly on the initial mass and concentration (and presumably velocity, cf. \citealt{GriBurNaaEA10}) )"462of the dense gas clumps., of the dense gas clumps.463 Finally we emphasise. in agreement with previous authors (222).. that a strong magnetic field has a very significant influence on the dynamics of the photoionisation process. and many of these effects should be easily observable.," Finally we emphasise, in agreement with previous authors \citep{Wil07,KruStoGar07,HenArtDeCEA09}, that a strong magnetic field has a very significant influence on the dynamics of the photoionisation process, and many of these effects should be easily observable."464 Given the difficulty of measuring the full 3D magnetic tield in the ISM. comparison to detailed numerical simulations such as these offers an indirect means to constrain the tield strength and orientation in and around H regions.," Given the difficulty of measuring the full 3D magnetic field in the ISM, comparison to detailed numerical simulations such as these offers an indirect means to constrain the field strength and orientation in and around H regions."465 JM's work has been part funded by the Irish Research Council for Science. Engineering and Technology: also by a grant from the Dublin Institute for Advanced Studies. and by Science Foundation Ireland.," JM's work has been part funded by the Irish Research Council for Science, Engineering and Technology; also by a grant from the Dublin Institute for Advanced Studies, and by Science Foundation Ireland."466 AJL's work was funded by a Schróddinger Fellowship from the Dublin Institute for Advanced Studies., AJL's work was funded by a Schröddinger Fellowship from the Dublin Institute for Advanced Studies.467 JM acknowledges support from an Argelander Fellowship during the writing of this paper., JM acknowledges support from an Argelander Fellowship during the writing of this paper.468 Figures were generated using the visualisation tool., Figures were generated using the visualisation tool.469 The authors wish to acknowledge the SFI/HEA Irish Centre for High-End Computing (CHEC) for the provision of computational facilities and support., The authors wish to acknowledge the SFI/HEA Irish Centre for High-End Computing (ICHEC) for the provision of computational facilities and support.470 We thank the referee for useful suggestions and for pointing out an error in an earlier draft., We thank the referee for useful suggestions and for pointing out an error in an earlier draft.471 Results from. the following MHD test) problems ean. be found at together with the HD test problems., Results from the following MHD test problems can be found at together with the HD test problems.472 These adiabatic, These adiabatic473 Wicringa ct ((1993) were the first to note structure on arcuinute scales iu the linearly polarized coupoucut of the galactic radio background at 325 MIIz. observed with the The small-scale structure in the maps of polarized intensity P. (with polarized brüehtuess temperatures Tia of up to 10 I) does NOT have a counterpart in total intensity. or Stokes J. down to very low lanits.," Wieringa et (1993) were the first to note structure on arcminute scales in the linearly polarized component of the galactic radio background at 325 MHz, observed with the The small-scale structure in the maps of polarized intensity $P$, (with polarized brightness temperatures $T_{\rm b,pol}$ of up to 10 K) does NOT have a counterpart in total intensity, or Stokes $I$, down to very low limits."474 Because the total Stokes 7 of the galactic radio background has au estimated Fypor of the order of 30 50 I& at 325 AIIIz. which must be very smooth aud therefore filtered out completely in the WSRT measurements. theapparent polarization percentage of the sinallscale features can become very nuch larger than.," Because the total Stokes $I$ of the galactic radio background has an estimated $T_{\rm b,pol}$ of the order of 30 – 50 K at 325 MHz, which must be very smooth and therefore filtered out completely in the WSRT measurements, the polarization percentage of the small-scale features can become very much larger than."475. The absence of corresponding simall-scale structure iu Stokes I led Wicrinea et (Gbid.), The absence of corresponding small-scale structure in Stokes $I$ led Wieringa et (ibid.)476 to propose that the stnall-scale structure in polarizec intensity P is due to Faraday rotation modulation., to propose that the small-scale structure in polarized intensity $P$ is due to Faraday rotation modulation.477 Iu this picture. svuchrotron radiation generated in the Galactic halo reaches τις hrough a magneto-ionic screen. viz.," In this picture, synchrotron radiation generated in the Galactic halo reaches us through a magneto-ionic screen, viz."478 the wari relatively jearby ISAL, the warm relatively nearby ISM.479 Structure in the electron deusitv and/or naenetic field in the ISM causes spatial variations in the Rotation Measure (RM) of the screen., Structure in the electron density and/or magnetic field in the ISM causes spatial variations in the Rotation Measure (RM) of the screen.480 Hence. the augle of lnear volarization of the svuchrotronu cnussion frou the wo is rotated by different amounts along different lines of sieht.," Hence, the angle of linear polarization of the synchrotron emission from the halo is rotated by different amounts along different lines of sight."481 Even if the polarized emission in the halo were otally smooth. in iutensitv as well as augle. the screen would produce structure in Stokes Q aud C.," Even if the polarized emission in the halo were totally smooth, in intensity as well as angle, the screen would produce structure in Stokes $Q$ and $U$."482 Sumalbscale structure in the polarized galactic radio vackeround Las recently been observed also at other yequencies, Small-scale structure in the polarized galactic radio background has recently been observed also at other frequencies.483 At 1120 MITZ. Gray et (1998. 1999) used he DRAO svuthesis telescope to study the phenomenon at 1 resolution.," At 1420 MHz, Gray et (1998, 1999) used the DRAO synthesis telescope to study the phenomenon at $^{\prime}$ resolution."484 Uvaniker et (1999) used the Effelshere clescope at 1.1 CHIz. to map the polarized emission at Y resolution over about 11008.," Uyaniker et (1999) used the Effelsberg telescope at 1.4 GHz, to map the polarized emission at $^{\prime}$ resolution over about $^\Box$."485" Duncan et ((1998) discuss radio polarization data at 1.1. 2.1 and Ls CIIz with the Parkes radio telescope au the VLA. at 55, 10 and 15 resolution."," Duncan et (1998) discuss radio polarization data at 1.4, 2.4 and 4.8 GHz with the Parkes radio telescope and the VLA, at $^{\prime}$, $^{\prime}$ and $^{\prime}$ resolution."486 All these observations support the interpretation iu terms of modulation of emission originating at larger distances. bv a relatively nearby Faraday screcu.," All these observations support the interpretation in terms of modulation of emission originating at larger distances, by a relatively nearby Faraday screen."487 The distributions of polarized intensity auc angle aay therefore be used o study the structure of the Faraday screen., The distributions of polarized intensity and angle may therefore be used to study the structure of the Faraday screen.488 Iu particular. polarization observations," In particular, polarization observations"489"are intrinsically weak, they have a very similar shape for both distributions, each with a characteristic ‘turnover’ scale.","are intrinsically weak, they have a very similar shape for both distributions, each with a characteristic `turnover' scale."490" Matching this scale in the two correlation functions corresponds to w,~0.04+0.01kpckms for the six haloes, which we adopt as a fiducial value."," Matching this scale in the two correlation functions corresponds to $w_{v}\sim0.04\pm0.01\,\wvunits$ for the six haloes, which we adopt as a fiducial value."491" We !caution that although the scales on which we match the one-dimensional correlation functions are somewhat smaller than the smoothing scales we adopt to create the random distributions, this does not guarantee that our choice of w, is unaffected by our choice of smoothing."," We caution that although the scales on which we match the one-dimensional correlation functions are somewhat smaller than the smoothing scales we adopt to create the random distributions, this does not guarantee that our choice of $w_{v}$ is unaffected by our choice of smoothing."492" Clearly, there are other ways of fixing w,."," Clearly, there are other ways of fixing $w_{v}$."493" In practice, however, our conclusions are not highly sensitive to the value of this parameter."," In practice, however, our conclusions are not highly sensitive to the value of this parameter."494 Values of the order of ων~ 0.01--1.0kpckm!s result in very similar £(A) correlation functions.," Values of the order of $w_{v}\sim0.01$ $1.0\,\wvunits$ result in very similar $\xi({\Delta})$ correlation functions."495 Values lower than 0.01kpckm recover very little signal.," Values lower than $0.01\,\wvunits$ recover very little signal."496" Values above lkpckm15 treat !1skms""! velocity differences as equivalent to >1kpc separations in space, and so make the cumulative correlation function very noisy on small scales for only a marginal increase in the overall signal. ("," Values above $1\,\wvunits$ treat $1\,\kms$ velocity differences as equivalent to $>1\,\kpc$ separations in space, and so make the cumulative correlation function very noisy on small scales for only a marginal increase in the overall signal. ("497"This noise, in turn, increases the scatter between signals measured by different observers.)","This noise, in turn, increases the scatter between signals measured by different observers.)"498" We find that our choice of w,~0.04kpcΊκπι1s is a reasonable compromise."," We find that our choice of $w_{v}\sim0.04\,\wvunits$ is a reasonable compromise."499" Our method for choosing w, can be compared with that of Starkenburg (2009), who determine the equivalent of Wy in their metric to be the ratio of the Spaghetti survey limits in radial distance and velocity (0.26kpckm! s)."," Our method for choosing $w_{v}$ can be compared with that of Starkenburg (2009), who determine the equivalent of $w_{v}$ in their metric to be the ratio of the Spaghetti survey limits in radial distance and velocity $0.26\,\wvunits$ )."500 Either value is acceptable to illustrate our approach and compare to simulations., Either value is acceptable to illustrate our approach and compare to simulations.501" We therefore adopt w,~0.04kpcΚαιτς,"," We therefore adopt $w_{v}\sim0.04\,\wvunits$."502 shows computed for 2392 BHB stars in the Xue (2008) sample3;; grey points)., shows computed for 2392 BHB stars in the Xue (2008) sample; grey points).503 Stars at small separations in the metric of (A< 4kpc) show significant clustering.," Stars at small separations in the metric of $\Delta<4\,\mathrm{kpc}$ ) show significant clustering."504" The amplitude of the signal increases if we restrict the sample to larger galactocentric distances, r>20 kpc (black points)"," The amplitude of the signal increases if we restrict the sample to larger galactocentric distances, $r>20$ kpc (black points)."505 At larger distances substructure is expected to be dynamically young and to have undergone less phase mixing., At larger distances substructure is expected to be dynamically young and to have undergone less phase mixing.506 Our finding of stronger clustering for more distant halo stars is in qualitative agreement with the results of Xueal. (2011)., Our finding of stronger clustering for more distant halo stars is in qualitative agreement with the results of Xue (2011).507" Although we appear to recover a significant clustering signal in3,, we have only one SDSS survey."," Although we appear to recover a significant clustering signal in, we have only one SDSS survey."508 The observed signal may be an artifact of the particular structures covered by the SDSS footprint., The observed signal may be an artifact of the particular structures covered by the SDSS footprint.509" Other parts of the halo may be smoother or more structured, or may appear so when viewed from different points around the Solar circle."," Other parts of the halo may be smoother or more structured, or may appear so when viewed from different points around the Solar circle."510 We will address this issue of sample variance in the following section using our mock catalogues., We will address this issue of sample variance in the following section using our mock catalogues.511 We show two further permutations of the Xue (2008) sample in3., We show two further permutations of the Xue (2008) sample in.512". The first of these (red open circles) includes stars close to the Galactic plane, |Z|<4 kpc."," The first of these (red open circles) includes stars close to the Galactic plane, $|Z|<4\,\mathrm{kpc}$ ."513 These were excluded from the main sample of Xue (2008) to excise the thick disc., These were excluded from the main sample of Xue (2008) to excise the thick disc.514" Although only ~150 stars are excluded by the cut on |Z], they make a substantial difference to the correlation function, suppressing the clustering signal on scales below A«8kpc."," Although only $\sim150$ stars are excluded by the cut on $|Z|$, they make a substantial difference to the correlation function, suppressing the clustering signal on scales below $\Delta<\sim8\,\mathrm{kpc}$."515 In the SDSS data the majority of low-|Z| stars are at small heliocentric radii., In the SDSS data the majority of $|Z|$ stars are at small heliocentric radii.516" These stars constitute a foreground ‘screen’ with a relatively smooth distribution, which may dilute the signal of correlated stars."," These stars constitute a foreground `screen' with a relatively smooth distribution, which may dilute the signal of correlated stars."517 The final sample shown in (blue open squares) includes all stars from the main sample (grey points) and a further nine BHB stars identified as globular cluster members by Xue (2009)., The final sample shown in (blue open squares) includes all stars from the main sample (grey points) and a further nine BHB stars identified as globular cluster members by Xue (2009).518" Two of these are from one cluster, and seven from another."," Two of these are from one cluster, and seven from another."519 Including these stars marginally increases the clustering signal in the smallest-separation bin., Including these stars marginally increases the clustering signal in the smallest-separation bin.520" This shows that the technique is sensitive to the clustering of starson these scales, which correspond"," This shows that the technique is sensitive to the clustering of starson these scales, which correspond"521A relativistic jet is a clear taxonomical characteristic of extragalactic sources detected 1n. y-rays.,A relativistic jet is a clear taxonomical characteristic of extragalactic sources detected in $\gamma$ -rays.522 Sources with jets pointing close to our line of sight are called blazars and are the brightest and most dominant population of active galactic nuclet (AGN) in the y-ray sky (e.g.22)..," Sources with jets pointing close to our line of sight are called blazars and are the brightest and most dominant population of active galactic nuclei (AGN) in the $\gamma$ -ray sky \citep[e.g.][]{Fichtel_1994,1fgl}."523 The radiation mechanism of the y-ray emission in blazars is widely believed to be inverse Compton scattering of ambient photons. either from inside the jet (synchrotron-self-Compton or SSC; e.g. Bloom Marscher 1986). or from a source external to the jet (external Compton scattering or EC). where the source of seed photons could be the aceretion disk (e.g.?).. the broad-line region (e.g. 2).. or perhaps the dusty torus ?)..," The radiation mechanism of the $\gamma$ -ray emission in blazars is widely believed to be inverse Compton scattering of ambient photons, either from inside the jet (synchrotron-self-Compton or SSC; e.g. Bloom Marscher 1986), or from a source external to the jet (external Compton scattering or EC), where the source of seed photons could be the accretion disk \citep[e.g.][]{dermer_1993}, the broad-line region \citep[e.g.][]{sikora_1994}, , or perhaps the dusty torus \citep[e.g.][]{blazejowski_2000}."524 We refer to ? for a review of theoretical models for blazar emission., We refer to \citet{boettcher_review_2010} for a review of theoretical models for blazar emission.525 Despite all the theoretical modeling efforts. the precise location in general of the y-ray emission within sources is still controversial. which in turn makes the origin of the seed photons for inverse-Compton scattering unclear.," Despite all the theoretical modeling efforts, the precise location in general of the $\gamma$ -ray emission within sources is still controversial, which in turn makes the origin of the seed photons for inverse-Compton scattering unclear."526 The proposed models can be roughly divided into two categories: those with y-rays originating relatively close to the black hole and the accretion disk. inside the broad-line region (BLR). and those with y-rays originating in the radio jet. at distances of several parsecs and well beyond the BLR.," The proposed models can be roughly divided into two categories: those with $\gamma$ -rays originating relatively close to the black hole and the accretion disk, inside the broad-line region (BLR), and those with $\gamma$ -rays originating in the radio jet, at distances of several parsecs and well beyond the BLR."527 Studies based on data from the instrument onboard the Compton Gamma Ray Observatory triggered an open debate about the location of the emission site in blazars., Studies based on data from the instrument onboard the Compton Gamma Ray Observatory triggered an open debate about the location of the emission site in blazars.528 The most popular opinion was that y-rays are produced within the BLR region via EC (e.g. Sikora et al., The most popular opinion was that $\gamma$ -rays are produced within the BLR region via EC (e.g. Sikora et al.529 1994)., 1994).530 However. other studies found that high levels of emission occurred after the ejections of superluminal jet components (?).. and that y-ray detected sources tend to have ongoing high frequency radio flares (2??)..," However, other studies found that high levels of emission occurred after the ejections of superluminal jet components \citep{jorstad_2001}, , and that $\gamma$ -ray detected sources tend to have ongoing high frequency radio flares \citep{valtaoja_1995,valtaoja_1996,anne_2003}."531 These results led the authors to conclude that strong y-ray emission in blazars was produced in growing shocks in the relativistic jets at parsec-scale distances from the black hole., These results led the authors to conclude that strong $\gamma$ -ray emission in blazars was produced in growing shocks in the relativistic jets at parsec-scale distances from the black hole.532 Well beyond the central BLR. the only source of seed photons appears to be the jet itself. implying that SSC is the main radiation mechanism for the strongest flares in blazars.," Well beyond the central BLR, the only source of seed photons appears to be the jet itself, implying that SSC is the main radiation mechanism for the strongest flares in blazars."533 For a historical perspective of the results obtained during the era. we refer to ? The dramatically improved y-ray data from the Large Area Telescope (LAT) onboard the hhas opened up the possibility of testing results obtained from the era regarding the origin of y-rays.," For a historical perspective of the results obtained during the era, we refer to \citet{aller_2010}534 The dramatically improved $\gamma$ -ray data from the Large Area Telescope (LAT) onboard the has opened up the possibility of testing results obtained from the era regarding the origin of $\gamma$ -rays."535" Several studies. based on the first year of LAT operations. have shown that: (1) the y-ray and the averaged radio flux densities are significantly correlated (22222222?)., and (11) blazars detected at y-rays are more likely to have larger Doppler factors (2???) and larger apparent opening angles (?) than those not detected by LAT."," Several studies, based on the first year of LAT operations, have shown that: (i) the $\gamma$ -ray and the averaged radio flux densities are significantly correlated \citep{kovalev_2009,giroletti_2010,ghirlanda_2010,mahony_2010,nieppola_2010,angelakis_2010,richards_2010,linford_2011,arshakian_2011}, and (ii) blazars detected at $\gamma$ -rays are more likely to have larger Doppler factors \citep{lister_2009,savolainen_2010,tornikoski_2010} and larger apparent opening angles \citep{pushkarev_2009} than those not detected by LAT."536 This observational evidence strongly suggests that radio and y-ray emission have a co-spatial origin., This observational evidence strongly suggests that radio and $\gamma$ -ray emission have a co-spatial origin.537 To locate and identify the region. where the bulk of y- emission is produced. and to provide details about its connection to the radio jet. an analysis of simultaneous radio and y-ray light curves Is necessary. ?.. ," To locate and identify the region where the bulk of $\gamma$ -ray emission is produced, and to provide details about its connection to the radio jet, an analysis of simultaneous radio and $\gamma$ -ray light curves is necessary. \citet{pushkarev_2010}, ,"538using data from the MOJAVE survey (?) and the monthly binned y-ray light curves provided by the 11-month LAT catalogue (?).. reported that radio flux-density variations lag significantly behind those in y-rays. with delays ranging from one to eight months (1n the observers frame).," using data from the MOJAVE survey \citep{lister_mojave} and the monthly binned $\gamma$ -ray light curves provided by the 11-month LAT catalogue \citep{1fgl}, reported that radio flux-density variations lag significantly behind those in $\gamma$ -rays, with delays ranging from one to eight months (in the observer's frame)."539 This suggests that there is a correlation between the parsec-scale radio emission and the strength of the y-rays., This suggests that there is a correlation between the parsec-scale radio emission and the strength of the $\gamma$ -rays.540 However. the method employed by the authors did not allow them to clearly characterize the sequence and the structure of individual flares.," However, the method employed by the authors did not allow them to clearly characterize the sequence and the structure of individual flares."541 Furthermore. we highlight two caveats about the interpretation of radio/gamma correlation analyses. which have often been interpreted too simplistically.," Furthermore, we highlight two caveats about the interpretation of radio/gamma correlation analyses, which have often been interpreted too simplistically."542 First. it 15 well-known that there is usually a considerable delay between mm and em radio flares.," First, it is well-known that there is usually a considerable delay between mm and cm radio flares."543 Thus. although em-flares would tend to peak after the y-ray flares. mm-flares would show shorter delays or possibly even peak before the y-rays.," Thus, although cm-flares would tend to peak after the $\gamma$ -ray flares, mm-flares would show shorter delays or possibly even peak before the $\gamma$ -rays."544 The very important second caveat is that a correlation analysis tends to measure the distance between the peaks. especially if the flares have different timescales (as the radio and the y-ray flares tend to have).," The very important second caveat is that a correlation analysis tends to measure the distance between the peaks, especially if the flares have different timescales (as the radio and the $\gamma$ -ray flares tend to have)."545 However. a radio flare starts to grow a considerable time before it peaks.," However, a radio flare starts to grow a considerable time before it peaks."546 The of a millimeter flare coincides with the ejection of a new VLBI component from the radiocore (e.g.. Savolainen et al.," The of a millimeter flare coincides with the ejection of a new VLBI component from the radiocore (e.g., Savolainen et al."547 2002)., 2002).548 This is the epoch thatmust be compared with the y-ray flaring. not theepoch," This is the epoch thatmust be compared with the $\gamma$ -ray flaring, not theepoch"549discussion in Chapter 15 of Numerical Recipes ((?))).,discussion in Chapter 15 of Numerical Recipes \citep{numrec1992}) ).550" We performed these tests, and the results can be summarized as follows: after correction for the degree of freedom, the adjustments by 4th order polynomials are slightly better (the chi-square obtained with straight lines are larger by a factor which ranges from 2 to only a few percent)."," We performed these tests, and the results can be summarized as follows: after correction for the degree of freedom, the adjustments by 4th order polynomials are slightly better (the chi-square obtained with straight lines are larger by a factor which ranges from 2 to only a few percent)."551" In other words the straight lines are not the best choice, but they cannot be excluded on statistical grounds."," In other words the straight lines are not the best choice, but they cannot be excluded on statistical grounds."552" Testing different orders of polynomials from 2 to 6, in all cases the 4th order is slightly better than all the others except for NGC1042 for which the 3rd order polynomial gives a slightly smaller (396)) normalized chi-square."," Testing different orders of polynomials from 2 to 6, in all cases the 4th order is slightly better than all the others except for NGC1042 for which the 3rd order polynomial gives a slightly smaller ) normalized chi-square."553 This last result still suggests that an inflection point is present., This last result still suggests that an inflection point is present.554" ?) argued that the bends in the slopes of radial abundance gradients in the disks of spiral galaxies is a consequence of a systematic error involving the excitation parameter P, which would produce realisticO abundance in high-excitation regions, but would overestimate this abundance in low excitation regions."," \cite{Pilyugin03} argued that the bends in the slopes of radial abundance gradients in the disks of spiral galaxies is a consequence of a systematic error involving the excitation parameter P, which would produce realistic abundance in high-excitation regions, but would overestimate this abundance in low excitation regions."555" For a measured value of R23, there are two possible solutions for theO abundance."," For a measured value of R23, there are two possible solutions for the abundance."556" For instance, according to Pilyugin, the positive slope of the observedO abundance by ?) in the external regions of the galaxy NGC1365 would be changed into a negative slope, and the break would disappear, if one used the lower branch of the graph ofO abundance versus R23 for these regions."," For instance, according to Pilyugin, the positive slope of the observed abundance by \cite{RoyWalsh97} in the external regions of the galaxy NGC1365 would be changed into a negative slope, and the break would disappear, if one used the lower branch of the graph of abundance versus R23 for these regions."557" The question, then, is when one should move from one branch to the other, to use correctly the method."," The question, then, is when one should move from one branch to the other, to use correctly the method."558" According to Pilyugin's calibrations (?),, the upper branch corresponds to 12+log(O/H)>8.15."," According to Pilyugin's calibrations \citep{Pilyugin03}, the upper branch corresponds to $12+log(O/H) > 8.15$."559" As can be seen in Figure 6, the slope breaks are situated above this abundance threshold for two of the galaxies studied here (NGC1042 and NGC6907), so they cannot be explained by a change of theO abundance versus R23 law."," As can be seen in Figure \ref{fig6}, the slope breaks are situated above this abundance threshold for two of the galaxies studied here (NGC1042 and NGC6907), so they cannot be explained by a change of the abundance versus R23 law."560 One argument given by Pilyugin to justify the existence of bends above the previous limit is that the less the value of the excitation parameterP the more the oxygen abundance obtained with the R23-method is overestimated (?)., One argument given by Pilyugin to justify the existence of bends above the previous limit is that the less the value of the excitation parameter the more the oxygen abundance obtained with the R23-method is overestimated \citep{Pilyugin03}.561" However, even with an increasing behavior of the excitation parameter at radii larger than the radius identified as the corotation radius for NGC1042, the overall decreasing behavior of the radial metallicity distribution is not flattened by the almost linear radial increase of the excitation parameterP."," However, even with an increasing behavior of the excitation parameter at radii larger than the radius identified as the corotation radius for NGC1042, the overall decreasing behavior of the radial metallicity distribution is not flattened by the almost linear radial increase of the excitation parameter."562 It only presents a small step in the region where the corotation seems to occur., It only presents a small step in the region where the corotation seems to occur.563" Another important point to be considered is that, in contrast to the R23 or the p-method the [O 111]/[N 1] method is single- and for the galaxies observed here we verified minima and inflexions independently of the R23-based methods."," Another important point to be considered is that, in contrast to the R23 or the the [O ]/[N ] method is single-valued and for the galaxies observed here we verified minima and inflexions independently of the R23-based methods."564" One might worry that [O rmu]/[N τῇ could depend upon a second parameter near co-rotation, and the behavior of the excitation parameter might even suggest a correlation with the corotation radius, but such an hypothesis would have to be investigated in a large sample of galaxies, and for the moment, it is not supported by any theoretical prediction."," One might worry that [O ]/[N ] could depend upon a second parameter near co-rotation, and the behavior of the excitation parameter might even suggest a correlation with the corotation radius, but such an hypothesis would have to be investigated in a large sample of galaxies, and for the moment, it is not supported by any theoretical prediction."565 This is not the first time that breaks in the radial distribution of metallicity are associated to the corotation radius., This is not the first time that breaks in the radial distribution of metallicity are associated to the corotation radius.566" The break observed by Roy Walsh in NGC1365 is real, as it has been observed using differentcalibrations"," The break observed by Roy Walsh in NGC1365 is real, as it has been observed using differentcalibrations"567that it undergoes a supernova. the other Qvhich will become the runaway) has already gone through some or all of its O star phase. but is unlikely to have become a WR star yet.,"that it undergoes a supernova, the other (which will become the runaway) has already gone through some or all of its O star phase, but is unlikely to have become a WR star yet."568 The statistics of observed WR binaries back this up: of 20 Galactic WR binaries with measured masses (van der Hucht 2001). at least [4 have as companions O stars which are massive enough to subsequently go through a WR phase. but only one (WR20a. which. as a system of an 83M. and an 82M. star (Bonanos et al.," The statistics of observed WR binaries back this up: of 20 Galactic WR binaries with measured masses (van der Hucht 2001), at least 14 have as companions O stars which are massive enough to subsequently go through a WR phase, but only one (WR20a, which, as a system of an $83 \, \msun$ and an $82 \, \msun$ star (Bonanos et al."569 2004). is highly unusual) contains two WR stars.," 2004), is highly unusual) contains two WR stars."570 One potential way of getting around this is to invoke a population with very uneven mass ratios. such that the less massive star still has most of its O star lifetime left after the SN of its companion.," One potential way of getting around this is to invoke a population with very uneven mass ratios, such that the less massive star still has most of its O star lifetime left after the SN of its companion."571 However. in order to be an O star on the main sequence at solar metallicity a mass of at least 17M is required.," However, in order to be an O star on the main sequence at solar metallicity a mass of at least $17 \msun$ is required."572 This then requires the other star to be extremely ..massive. and so is unlikely to be a common situation.," This then requires the other star to be extremely massive, and so is unlikely to be a common situation."573 Furthermore. the masses of the components in massive binaries are correated (Garmany. Conti Massey 1980). with very few unevolved systems having mss ratios q=Μ.Μ below 0.3.," Furthermore, the masses of the components in massive binaries are correlated (Garmany, Conti Massey 1980), with very few unevolved systems having mass ratios $q = M_{2}/M_{1}$ below 0.3."574 However. the above assumes that all supernovae are equal — that is. that the distribution of parameters which produce WR evolution in the and the distribution of parameters which produce an asymmetrie SN explosion which unbinds the system are completely unrelated.," However, the above assumes that all supernovae are equal -- that is, that the distribution of parameters which produce WR evolution in the and the distribution of parameters which produce an asymmetric SN explosion which unbinds the system are completely unrelated."575 In reality this is unlikely to be the case., In reality this is unlikely to be the case.576 The likelihood of a system producing a runaway WR star depends on a number of factors., The likelihood of a system producing a runaway WR star depends on a number of factors.577 First. the star whieh survives bevond the SN of its companion must be. or become. massive enough to undergo a WR phase.," First, the star which survives beyond the SN of its companion must be, or become, massive enough to undergo a WR phase."578 In the case that mass transfer occurs. the mass limit for a star to go through a WR phase is lowered somewhat (Dray Tout 2005) due to accretion of He-enhanced matter and subsequent thermohaline mixing.," In the case that mass transfer occurs, the mass limit for a star to go through a WR phase is lowered somewhat (Dray Tout 2005) due to accretion of He-enhanced matter and subsequent thermohaline mixing."579 Mass ransfer also promotes the likelihood of the secondary undergoing a WR phase by raising its mass. although the amount of mass transfer which can occur before the secondary is spun up to rotational break-up velocity is a matter of debate (Packet 1981. Dewi 2005).," Mass transfer also promotes the likelihood of the secondary undergoing a WR phase by raising its mass, although the amount of mass transfer which can occur before the secondary is spun up to rotational break-up velocity is a matter of debate (Packet 1981, Dewi 2005)."580 Second. this star must then be given a large enough velocity by the SN of the other that it is observable as a runaway.," Second, this star must then be given a large enough velocity by the SN of the other that it is observable as a runaway."581 Here we use 30kms| as the velocity threshold above which a star is classified as runaway., Here we use $30 {\rm kms}^{-1}$ as the velocity threshold above which a star is classified as runaway.582 Recent work by Pfahl et al. (, Recent work by Pfahl et al. (5832001) suggests hat the velocity distribution of observed X-ray binary systems is indicative of a bimodal distribution of SN kicks. with smaller Kicks originating from systems which underwent mass transfer earlier in heir evolution.,"2001) suggests that the velocity distribution of observed X-ray binary systems is indicative of a bimodal distribution of SN kicks, with smaller kicks originating from systems which underwent mass transfer earlier in their evolution."584 Early mass transfer (i.e. short initial period) has also been linked with more nearly conservative mass transfer (Langer 2005)., Early mass transfer (i.e. short initial period) has also been linked with more nearly conservative mass transfer (Langer 2005).585 In addition. simulations of WR production in these systems (Dray Tout 2005) suggest that early mass transfer increases the ikelihood of the secondary undergoing a WR phase.," In addition, simulations of WR production in these systems (Dray Tout 2005) suggest that early mass transfer increases the likelihood of the secondary undergoing a WR phase."586 A further notable effect is that. since masses of stars in massive binaries are correlated. primaries of high mass tend to have relatively massive secondaries.," A further notable effect is that, since masses of stars in massive binaries are correlated, primaries of high mass tend to have relatively massive secondaries."587 Therefore. in the population as a whole. as the primary mass increases. so does the likelihood of the secondiry being massive enough to undergo a WR phase.," Therefore, in the population as a whole, as the primary mass increases, so does the likelihood of the secondary being massive enough to undergo a WR phase."588 Above a primary mass of 40M. or so. nearly all solar metallicity systems which interact and avoid merging will have secondaries which can go through a WR phase.," Above a primary mass of $40 \, \msun$ or so, nearly all solar metallicity systems which interact and avoid merging will have secondaries which can go through a WR phase."589 Therefore the initial primary mass distribution of systems which can make WR runaways is strongly Figureου ο, Therefore the initial primary mass distribution of systems which can make WR runaways is strongly skewed towards high initial primary mass.590 πρ πα ο Ht)Mnails aleto thesthireBho!ο αλ which tthenrenplia ing «Kthesmadsstous diftardesehy hausbalw ROR anas ds onsale hdirilaestis last ο μα Me. Massie ο ΟΙ hioke rhe aise Heleerton applede 900 dàllotitl Heuvel et al.," However, $40 \, \msun$ is also the threshold above which the remaining cores of stars after losing mass in RLOF and winds at solar metallicity are above $8 \, \msun$ – i.e. massive enough to collapse to a black hole rather than a neutron star (Fryer 1999, van den Heuvel et al."591 2000)., 2000).592 The level of wind mass loss for WR stars is a matter of debate. particularly. when rotation is considered (e.g. Maeder Meynet 2000).," The level of wind mass loss for WR stars is a matter of debate, particularly when rotation is considered (e.g. Maeder Meynet 2000)."593 However it is notable that. in all five of the Galactic WC+O systems with measured masses. the WC star (which is a stripped core nearing the very end of its lifetime) is more massive than 8M...," However it is notable that, in all five of the Galactic WC+O systems with measured masses, the WC star (which is a stripped core nearing the very end of its lifetime) is more massive than $8 \, \msun$."594 These systems also have massive O star companions (23 — 34M.. van der Hucht 2001) which will quite probably still be O stars at the SNe of their companions and later undergo a WR phase. making them prime examples of potential contributors to both the (BSS scenario) O and WR runaway populations.," These systems also have massive O star companions $23$ – $34\, \msun$, van der Hucht 2001) which will quite probably still be O stars at the SNe of their companions and later undergo a WR phase, making them prime examples of potential contributors to both the (BSS scenario) O and WR runaway populations."595 If a neutron star (NS) is formed in the SN explosion. the rest of the mass of the pre-explosion star is lost.," If a neutron star (NS) is formed in the SN explosion, the rest of the mass of the pre-explosion star is lost."596 However. when a black hole (BH) is formed the amount of matter which falls back is not limited by the maximum NS mass.," However, when a black hole (BH) is formed the amount of matter which falls back is not limited by the maximum NS mass."597 There is a peak in the system SN mass loss for pre-SN primary core masses somewhere around the boundary between NS and BH formation (Fig., There is a peak in the system SN mass loss for pre-SN primary core masses somewhere around the boundary between NS and BH formation (Fig.598 1)., 1).599 The most massive stars form BHs at the end of their lifetimes., The most massive stars form BHs at the end of their lifetimes.600 These are also the stars which lose most or all of jeir envelopes in winds or RLOF before they explode., These are also the stars which lose most or all of their envelopes in winds or RLOF before they explode.601 This again limits the mass loss in the explosion. since the matter external © the core has already been lost.," This again limits the mass loss in the explosion, since the matter external to the core has already been lost."602 Consequently. BH-forming exslosions involve smaller mass loss. relative to the remaining sysem mass. than do most NS-forming explosions.," Consequently, BH-forming explosions involve smaller mass loss, relative to the remaining system mass, than do most NS-forming explosions."603 If Kicks are basically related to anisotropic mass loss in the SN explosion. as opposed to anisotropic neutrino emission. then a smaller amount of mass lost relative to the system mass should also correspond to a smaller average kick velocity.," If kicks are basically related to anisotropic mass loss in the SN explosion, as opposed to anisotropic neutrino emission, then a smaller amount of mass lost relative to the system mass should also correspond to a smaller average kick velocity."604 Even if the Kick velocity is unchanged. the amount of mass lost affects the final velocities imparted to the stars as the source of the mass loss is not at the centre of mass of the system (see e.g. Tauris Takens 1998).," Even if the kick velocity is unchanged, the amount of mass lost affects the final velocities imparted to the stars as the source of the mass loss is not at the centre of mass of the system (see e.g. Tauris Takens 1998)."605 If the BH-forming explosion mechanism is basically similar to that forming a NS the imparted velocity to the companion to a BH-forming star. therefore. is still on average lower.," If the BH-forming explosion mechanism is basically similar to that forming a NS the imparted velocity to the companion to a BH-forming star, therefore, is still on average lower."606 Is this effect. combined with the greater likelihood of BH-forming stars to have companions which will become WR stars. enough. to," Is this effect, combined with the greater likelihood of BH-forming stars to have companions which will become WR stars, enough to"607independent of time as described above. then under the same change of variables. the creation rate becomes independent of halo mass.,"independent of time as described above, then under the same change of variables, the creation rate becomes independent of halo mass."608 The resulting pdf f(Inz|AZ) is now only valid if we are examining the distribution of halos at fixed mass., The resulting pdf $f(\ln\nu|M)$ is now only valid if we are examining the distribution of halos at fixed mass.609 Following the notation adopted above. this is given by: Because the mass of cach halo is assumed to monotonically increase with time. within anv interval of mass and time. an infinite number of “creation events’ occur.," Following the notation adopted above, this is given by: Because the mass of each halo is assumed to monotonically increase with time, within any interval of mass and time, an infinite number of `creation events' occur."610 This means that he joint. probability of the existence of a halo in6o/fi mass and time cannot be properly. normalised., This means that the joint probability of the existence of a halo in mass and time cannot be properly normalised.611 Equation 16. gives the link between two pdfs. the mass 'unction and the ereation rate using a mathematical trick to cope with an un-nornialised prior in time.," Equation \ref{eq:bayes2} gives the link between two pdfs, the mass function and the creation rate using a mathematical trick to cope with an un-normalised prior in time."612 Phe numerator of his equation is the joint distribution of halos in mass and ime. FC.DdAdl=fCALIOMf(0)di.," The numerator of this equation is the joint distribution of halos in mass and time, $f(M,t)\,dM\,dt=f(M|t)\,dM\,f(t)\,dt$."613 The denominator is not a function of time: it. only normalises the resulting ormula so ΤΠ} integrates to unity.," The denominator is not a function of time: it only normalises the resulting formula so $f(t|M)\,dt$ integrates to unity."614" Following this argument.e Bgiven a mass function. multiplving EAby do,(E)/dl creates a function withbo/h the correct mass and. time »haviour."," Following this argument, given a mass function, multiplying by $d\delta_c(t)/dt$ creates a function with the correct mass and time behaviour."615 This joint distribution function (not a pdf) has he same mass dependence as (ΑΠΕ and the same time dependence as f(]M)., This joint distribution function (not a pdf) has the same mass dependence as $f(M|t)$ and the same time dependence as $f(t|M)$.616 As an example we consider the fitting function. of Sheth Yormen (1999) to the multiplicity function determined. from the results of N-bock simulations for different cosmological parameters: where ν΄=«ty anda pare parameters., As an example we consider the fitting function of Sheth Tormen \shortcite{sheth} to the multiplicity function determined from the results of N-body simulations for different cosmological parameters: where $\nu'=a^{1/2}\nu$ and $a$ $p$ are parameters.617 Note that Sheth Tormen cisplavecl this formula. using a cillerent notation to that adopted here. although parameters e and p are the same in both cases.," Note that Sheth Tormen displayed this formula using a different notation to that adopted here, although parameters $a$ and $p$ are the same in both cases."618 ; is determined by requiring hat the integral of f(Invf) over all Inv gives unity., $A$ is determined by requiring that the integral of $f(\ln\nu|t)$ over all $\ln\nu$ gives unity.619 Sheth ‘Tormen found best fit parameters e=0.707 and p=0.3 for heir simulations and group finding algorithm., Sheth Tormen found best fit parameters $a=0.707$ and $p=0.3$ for their simulations and group finding algorithm.620" The standard ""S multiplicity function has @=1I. p=0 and vl=1/2."," The standard PS multiplicity function has $a=1$, $p=0$ and $A=1/2$."621 Unless stated otherwise. by standard. PS theory. we refer o the adoption of this multiplicity function combined with op-hat filtering (to caleulate στ).," Unless stated otherwise, by standard PS theory, we refer to the adoption of this multiplicity function combined with top-hat filtering (to calculate $\sigma_M^2$ )."622" In order to convert this function to provide a model of both the time and mass of iilo creation events. all we need to do is to multiply. by (dà,fdi."," In order to convert this function to provide a model of both the time and mass of halo creation events, all we need to do is to multiply by $d\delta_c/dt$."623" For standard PS theory. writing v explicitly in terms of m3, and 3, we find that the joint distribution of the existence ofa halo in mass and time reduces to: Although not. normalised. such a formula integrated Over anv [wo areas of the mass-time plane will provide the Correct. relativὉ number densities."," For standard PS theory, writing $\nu$ explicitly in terms of $\sigma_M^2$ and $\delta_c$ we find that the joint distribution of the existence of a halo in mass and time reduces to: Although not normalised, such a formula integrated over any two areas of the mass-time plane will provide the correct relative number densities."624 Note tha this is the same formula as obtained by simply multipIving the mass function with the creation time distribution a fixed mass., Note that this is the same formula as obtained by simply multiplying the mass function with the creation time distribution at fixed mass.625 Fhis would be inconsistent within a Bavesian [framework and. would produce a joint. density function whic1 lacks the correct mass and time behaviour: the form of cach conditional pef is altered. by the other., This would be inconsistent within a Bayesian framework and would produce a joint density function which lacks the correct mass and time behaviour: the form of each conditional pdf is altered by the other.626 Care should therefore be taken when using the creation rate in models which also include the mass function., Care should therefore be taken when using the creation rate in models which also include the mass function.627 So lar. we have only been concerned. with the epoch at which a halo is created.," So far, we have only been concerned with the epoch at which a halo is created."628 However. there is an important distinction between major mergers and the slow aceretion of mass when applving the results in models of certain cosmological phenomena.," However, there is an important distinction between major mergers and the slow accretion of mass when applying the results in models of certain cosmological phenomena."629 For instance. only violent. merger events are thought to be important for starbursts ane quasar activation.," For instance, only violent merger events are thought to be important for starbursts and quasar activation."630 In. paper L we showed that for standard. PS heorv with a sharp fA-space filter. if mass jumps in a xwticeular trajectory correspond to merger events. then the distribution of mergers is the same as that of the build-up of matter from all types of creation event.," In paper I, we showed that for standard PS theory with a sharp $k$ -space filter, if mass jumps in a particular trajectory correspond to merger events, then the distribution of mergers is the same as that of the build-up of matter from all types of creation event."631 This is because he trajectories are Brownian random walks which have the special property that their form is independent of the initial »oint., This is because the trajectories are Brownian random walks which have the special property that their form is independent of the initial point.632 Given only the mass function and the assumptions outlined above it is not possible to determine how each clump increases in mass. onlv the distribution of times at which it reaches a certain mass.," Given only the mass function and the assumptions outlined above it is not possible to determine how each clump increases in mass, only the distribution of times at which it reaches a certain mass."633 More information about the build-up. of individual clumps is required before. the distribution. of major mergers can be determined., More information about the build-up of individual clumps is required before the distribution of major mergers can be determined.634 Such information is available in PS theory and. follows from the argument that each trajectory gives the history of the halo masses in which a particular small mass element resides., Such information is available in PS theory and follows from the argument that each trajectory gives the history of the halo masses in which a particular small mass element resides.635 A direct approach to modelling structure formation is to simulate the evolution of the mass density of the Universe using a clistribution of softened particles., A direct approach to modelling structure formation is to simulate the evolution of the mass density of the Universe using a distribution of softened particles.636 We have run three such simulations using the Hydra. N-body. hvedrocdynanmies code (Couchman.Phomas&Pearce1995) with 128* dark matter particles to model the build-up of halos for three cillerent cosmological models. described in Table 1..," We have run three such simulations using the Hydra N-body, hydrodynamics code \cite{couchman} with $128^3$ dark matter particles to model the build-up of halos for three different cosmological models, described in Table \ref{tab:cosmo}."637 1n order to determine the at whieh halos are created within these simulations. we output particle positions at a laree number of times.," In order to determine the at which halos are created within these simulations, we output particle positions at a large number of times."638 For the ΕςΟΔΙΕ simulation. we output particle positions at 362. cillerent epochs. separated by approximately equal intervals in time.," For the $\Gamma$ CDM simulation, we output particle positions at 362 different epochs, separated by approximately equal intervals in time."639 For, For640TOY equation. the central density is closer to Spy (Witten. 1984). and a arger increase of volume can be accommocdated.,"TOV equation, the central density is closer to $5\rho_0$ (Witten, 1984), and a larger increase of volume can be accommodated."641 Typically. a few high density (pz po) clumps break olf the star or condense out of the accretion stream.," Typically, a few high density $\rho>\rho_0$ ) clumps break off the star or condense out of the accretion stream."642 Most of these are quickly acereted by he black hole., Most of these are quickly accreted by the black hole.643 At the conclusion of tidal stripping a starlet of ~LO24. remains relatively lar from the black hole. and on occasion is tically injected into a highly elliptic orbitsuch is the case in run C. where he Q.O55AL. starlet visible in Figs.," At the conclusion of tidal stripping a starlet of $\sim 10^{-2}M_\odot$ remains relatively far from the black hole, and on occasion is tidally injected into a highly elliptic orbit—such is the case in run C, where the $0.055M_\odot$ starlet visible in Figs."644 2 and 3 is in a bound orbit and still moving away [rom the black hole with orbital speed of 5.4.101 Ems at he end of the simulation (see Fig.," 2 and 3 is in a bound orbit and still moving away from the black hole with orbital speed of $5.4\times10^4\,$ km/s at the end of the simulation (see Fig."645 4 for the center-of-mass velocity vector)., 4 for the center-of-mass velocity vector).646 The same tidal interaction has also substantially spun up the starlet to a rotational period of 1.3 ms (at the end of simulation C)., The same tidal interaction has also substantially spun up the starlet to a rotational period of 1.3 ms (at the end of simulation C).647 In the six runs presented. here. the amount of mass cjected remains unresolved.," In the six runs presented here, the amount of mass ejected remains unresolved."648 At most. à few individual SPELL particles (of mass ωςΑΔ) cach. see Table 1) are in unbound trajectories.the number of SPIEL particles ejected varies from none (zero) to three for the runs of Table. 1.," At most, a few individual SPH particles (of mass $M_{\rm SS}/N$ each, see Table 1) are in unbound trajectories—the number of SPH particles ejected varies from none (zero) to three for the runs of Table 1."649 For instance. at the end of run C only one SPL particle (to the left of ther km tick-mark in Fie.," For instance, at the end of run C only one SPH particle (to the left of the $x=-200\,$ km tick-mark in Fig."650 2. at jy&—390 km) is on a clearly outbouncl trajectory. its terminal velocity (at infinitv) will be about 45000 kms. ancl its velocity is so high that it will not only leave the erstwhile binary. but the Galaxy and the Virgo cluster as well.," 2, at $y\approx -390\,$ km) is on a clearly outbound trajectory, its terminal velocity (at infinity) will be about $45 000\,$ km/s, and its velocity is so high that it will not only leave the erstwhile binary, but the Galaxy and the Virgo cluster as well."651 The future of one more particle is unclecided. it may or may not be bound to the starlet. whose fate is sealed.," The future of one more particle is undecided, it may or may not be bound to the starlet, whose fate is sealed."652 The starlet is doomed to undergo a close encounter with the black hole., The starlet is doomed to undergo a close encounter with the black hole.653 , 65448 d was not found in any other subset.,48 d was not found in any other subset.655 A 63.8-d modulation period fits the 1949-1959 data reasonably well., A 63.8-d modulation period fits the 1949–1959 data reasonably well.656 The residual spectrum of this data set is shown in the bottom panel of Fig. 6.., The residual spectrum of this data set is shown in the bottom panel of Fig. \ref{v14sp}.657 The CCD data do not contradict either the 50 or the 64-d modulation period., The CCD data do not contradict either the 50 or the 64-d modulation period.658" The star is far from the centre, but a close companion (IV75;Arp1962) at 6” distance with brightness similar to V18 at minimum light may contaminate the photometric data of low-resolution observations."," The star is far from the centre, but a close companion \citep[IV 75;][]{arp} at 6"" distance with brightness similar to V18 at minimum light may contaminate the photometric data of low-resolution observations."659" Oo41 noticed considerable changes in the shape and amplitude of the light curve, which were attributed to probable Blazhko effect in KK71."," Oo41 noticed considerable changes in the shape and amplitude of the light curve, which were attributed to probable Blazhko effect in KK71."660 The pulsation-period variation of the star is very complex as can be seen in Fig. 7.., The pulsation-period variation of the star is very complex as can be seen in Fig. \ref{v18oc}.661 The total range of the period variation is larger than 0.0001 d and abrupt period changes occur on time-scales of several thousands of days., The total range of the period variation is larger than 0.0001 d and abrupt period changes occur on time-scales of several thousands of days.662 The seasonal light curves of all the photographic and B band CCD data are plotted in Fig. 8.., The seasonal light curves of all the photographic and $B$ band CCD data are plotted in Fig. \ref{v18.1}.663 The observations from different sources are denoted by different symbols., The observations from different sources are denoted by different symbols.664" The pulsation light curve of V18 varies significantly from one season to the next, while the scatter of the light curves in the individual panels does not exceed what is expected from the inhomogeneous, photographic data."," The pulsation light curve of V18 varies significantly from one season to the next, while the scatter of the light curves in the individual panels does not exceed what is expected from the inhomogeneous, photographic data."665" The V band CCD light curves of Stormetal. (1991),, Cohen&Matthews(1992) and K00 (not plotted in Fig. 8))"," The $V$ band CCD light curves of \cite{s91}, , \cite{cm92} and K00 (not plotted in Fig. \ref{v18.1}) )"666 also show significant variations., also show significant variations.667" 'The amplitude of the pulsation was smaller than 0.5- in 1977 and 1979, while no light variability was detected in 1943 and 1984 at all."," The amplitude of the pulsation was smaller than 0.5-mag in 1977 and 1979, while no light variability was detected in 1943 and 1984 at all."668" There are three stars close to V18: IV74, IV75 and IV76 with 14.60, 15.82 and 17.19 B magnitude, respectively (Arp1962)."," There are three stars close to V18: IV74, IV75 and IV76 with 14.60, 15.82 and 17.19 $B$ magnitude, respectively \citep{arp}."669" The mean magnitude of V18 measured in the 1984 Las Campanas observations is, however, 15.35-mag."," The mean magnitude of V18 measured in the 1984 Las Campanas observations is, however, 15.35-mag."670" This brightness value is close to the mean brightness of V18 in any other season, but it is significantly different from the magnitudes of the neighbouring constant stars."," This brightness value is close to the mean brightness of V18 in any other season, but it is significantly different from the magnitudes of the neighbouring constant stars."671" Therefore, we exclude the possibility that wrong identification of V18 can account for the constancy of its light curve in these years."," Therefore, we exclude the possibility that wrong identification of V18 can account for the constancy of its light curve in these years."672" As the seasonal light curves of V18 do not show large inhomogeneity the modulation cycle has to be relatively long, about hundreds of days."," As the seasonal light curves of V18 do not show large inhomogeneity, the modulation cycle has to be relatively long, about hundreds of days."673" However, we failed to determine any modulation period when analysing the light curves of different segments of the data."," However, we failed to determine any modulation period when analysing the light curves of different segments of the data."674 The phase modulation cannot be separated from the irregular period changes (note the different periods of each column in Fig. 8))., The phase modulation cannot be separated from the irregular period changes (note the different periods of each column in Fig. \ref{v18.1}) ).675" The residual spectrum of any data subset, which is long enoughfor analysis, is dominated by signals arising"," The residual spectrum of any data subset, which is long enoughfor analysis, is dominated by signals arising"676"= VW615 Cas is a Be binary system which exhibits radio outbursts at regular 26.5r, day intervals which is believed to be the orbital period (οσον Tavlor 1978. Tavlor Caegory 1982).","$\equiv$ V615 Cas is a Be binary system which exhibits radio outbursts at regular 26.5 day intervals which is believed to be the orbital period (Gregory Taylor 1978, Taylor Gregory 1982)."677 This source has created particular excitement because of the possible association with the strong. 100 MeV. eununia-ay source 2€60 135|OL (Diguani IIeruisen 1983).," This source has created particular excitement because of the possible association with the strong, 100 MeV gamma-ray source 2CG 135+01 (Bignami Hermsen 1983)."678 In fact. wwas discovered do A-ravs while searching for the counterpart of ον 135|01 (Digna 11981).," In fact, was discovered in X-rays while searching for the counterpart of 2CG 135+01 (Bignami 1981)."679 Confuxiou was initially introduced by the preseuce of the quasar QSO 02LL|622 which is less than 1255 away rom aand which contributes to the (bleuded) enmuüssion as ucasured by OSSE and COMPTEL., Confusion was initially introduced by the presence of the quasar QSO 0241+622 which is less than 5 away from and which contributes to the (blended) emission as measured by OSSE and COMPTEL.680 In the past vears. ROSAT. ASCA and RXTE observatious have resolved the cluission from the Be star aud the quasar.," In the past years, ROSAT, ASCA and RXTE observations have resolved the emission from the Be star and the quasar."681" Moreover. the atest EQRET analysis has reduced the error circle of 20°C 135|01 = 2EG Jo216119 = BEC JO21L|6108 toLL"". lus excluding he quasar as a counterpart candidate as well as making the association with ((which is aapart) quite uncertain (Ixniffen 11997. Πανπα 11999)!"," Moreover, the latest EGRET analysis has reduced the error circle of 2CG 135+01 $\equiv$ 2EG J0241+6119 $\equiv$ 3EG J0241+6103 to, thus excluding the quasar as a counterpart candidate as well as making the association with (which is apart) quite uncertain (Kniffen 1997, Hartman 1999)!"682 Nevertheless. the N-rayv cussion niechauisia of hhas inspired both. theoreticians as well as observers.," Nevertheless, the X-ray emission mechanism of has inspired both, theoreticians as well as observers."683 While there is general agreement that the radio outbursts are produced by svuchrotrou radiation. it has been argued that it is unlikely that the N-ray euissiou stems from the sale electron population. but rather could originate by inverse Compton chussion (CTavlor 11996. ILurison 22000).," While there is general agreement that the radio outbursts are produced by synchrotron radiation, it has been argued that it is unlikely that the X-ray emission stems from the same electron population, but rather could originate by inverse Compton emission (Taylor 1996, Harrison 2000)."684 This picture is supported by extensive inulti-wavelength observations which are reported iu Strick (1998) ancl Warrison (2000)., This picture is supported by extensive multi-wavelength observations which are reported in Strickman (1998) and Harrison (2000).685 Very recently. Apparay (2001) has imodelled this inverse Compton cussion in more detail.," Very recently, Apparao (2001) has modelled this inverse Compton emission in more detail."686 Based on the assuniptiou that the electron population is rather steady (tha is. lenoriug the ~0.5 phase shüft between N-ray aud. radio peak). he derived the properties of this electrou population (such as spectral index. maguctic feld and dift velocity) from the observed radio peak flux. aud from that calculates the LO200 keV spectruu.," Based on the assumption that the electron population is rather steady (that is, ignoring the $\sim$ 0.5 phase shift between X-ray and radio peak), he derived the properties of this electron population (such as spectral index, magnetic field and drift velocity) from the observed radio peak flux, and from that calculates the 10–200 keV spectrum."687" While Apparao (2001) fuds eeneral agreement of his model with the data. he also predicts a clear spectral break caused by the peak of the optical photon field (from the Be star) around €10 eV aud boosted by the Loreutz factor For the values he used. Apparao (2001) derives 5,~ 13. and thus a break at ~52e20 keV. The inverse Compton photon spectrum is therefore given by (see thick line is Fie. 3))"," While Apparao (2001) finds general agreement of his model with the data, he also predicts a clear spectral break caused by the peak of the optical photon field (from the Be star) around $\epsilon \sim 10$ eV and boosted by the Lorentz factor For the values he used, Apparao (2001) derives $\gamma_c \sim 43$ , and thus a break at $\sim \gamma_c^2 \epsilon \sim 20$ keV. The inverse Compton photon spectrum is therefore given by (see thick line is Fig. \ref{broadsp}) )"688 and Iu this short note we have investigated pointedoss? N-rav Timing Explorer (RNTE: Swank 1998) data to test this clear prediction for the spectral shape of the X-ray cuuission., and In this short note we have investigated pointed X-ray Timing Explorer (RXTE; Swank 1998) data to test this clear prediction for the spectral shape of the X-ray emission.689 Note that while IHarrisou ((2000) eive a complete description ofthese same RATE data with respect to photon flux aud temporal variability. no clear statement is made about the spectral slope.," Note that while Harrison (2000) give a complete description of these same RXTE data with respect to photon flux and temporal variability, no clear statement is made about the spectral slope."690 We have analyzed 11 poiuted RNTE obscrvatious of ttaken between March 1 aud. March 30. 1996 (see Tab.," We have analyzed 11 pointed RXTE observations of taken between March 1 and March 30, 1996 (see Tab."691 1 for a log of all observations)., \ref{fitres} for a log of all observations).692 All data have been retrieved from the TEASARC archive., All data have been retrieved from the HEASARC archive.693" Similar to IEuisou ((2000) we have used. ""Standard 27 mode spectra of the Proportional Counter Array (PCA: Jahoda 11996) which have full ΟΠΟΙΟΥ: resolution aud l6-sec", Similar to Harrison (2000) we have used “Standard 2” mode spectra of the Proportional Counter Array (PCA; Jahoda 1996) which have full energy resolution and 16-sec694though more similar between runs. they are still larger in ltun A (Figure 4).,"though more similar between runs, they are still larger in Run A (Figure 4)."695 For consecutive data points. in Run A positive Dux variations are 56 of the events and the fux increment is. |744 on average. while negative changes are 44% ofthe events and have an average magnitude of 27%.," For consecutive data points, in Run A positive flux variations are 56 of the events and the flux increment is $+74$ on average, while negative changes are 44 of the events and have an average magnitude of $-27$."696. Similarly. for Run B. positive changes (52% of the events) have an average magnitude of |42%.. while the average Lux decrement (48 of the events) is. 18," Similarly, for Run B, positive changes (52 of the events) have an average magnitude of $+42$, while the average flux decrement (48 of the events) is $-18$."697 We present a comparison with the ionized-gas surveys of UC and LIC regions by Wood&Churchwell(1989) and Wurtzetal.(1904)protostars., We present a comparison with the ionized-gas surveys of UC and HC regions by \cite{WC89} and \cite{Kurtz94}.698".. Figure 5 shows normalized bistogranis of 2-em luminosity (Ss 4,7) obtained for both runs using the time steps previously shown in figures 1 to 4 and the co-added: observed samples from Wood&Churehwell(1989). and Kurtzetal.(1904). taking the SI sources for which they report a 2-cm [lux and a distance."," Figure 5 shows normalized histograms of 2-cm luminosity $S_{\rm 2cm}d^2$ ) obtained for both runs using the time steps previously shown in figures 1 to 4 and the co-added observed samples from \cite{WC89} and \cite{Kurtz94}, taking the 81 sources for which they report a 2-cm flux and a distance."699 Simulation and observations roughly agree. but neither Run A nor DB can reproduce the high-Iuminosity end of the observed. distribution: 20 of the observed. sources have Sound50 Jy kpc. while only 1% of the regions in the simulation steps of Run A and in no step in Run D have luminosities above this threshold.," Simulation and observations roughly agree, but neither Run A nor B can reproduce the high-luminosity end of the observed distribution: 20 of the observed sources have $S_{\rm 2cm}d^2>50$ Jy $^2$, while only 1 of the regions in the simulation steps of Run A and in no step in Run B have luminosities above this threshold."700 These bright UC regions likely correspond to stages in which accretion o the ionizing protostar(s) is completely shut olf. therefore hey do not correspond to the analyzed regions fron he simulation in which the protostars are still accreting.," These bright UC regions likely correspond to stages in which accretion to the ionizing protostar(s) is completely shut off, therefore they do not correspond to the analyzed regions from the simulation in which the protostars are still accreting."701 In the range 20«Shad<50 Jy kpc. Run A matches νοτος than Run B the observed luminosities.," In the range $20<S_{\rm 2cm}d^2<50$ Jy $^2$, Run A matches better than Run B the observed luminosities."702" This may be cause Run B does not. produce any star with a mass AL,c30M..", This may be because Run B does not produce any star with a mass $M_\star>30~\Msun$.703 Simulations identical to Run Bo produce ügher-mass stars in the presence of magnetic fields (Petersetal.2011). and may also produce more massive stars in he purely raciation-hydrodynamical case with hiehcr-mass initial clumps., Simulations identical to Run B produce higher-mass stars in the presence of magnetic fields \citep{Peters11} and may also produce more massive stars in the purely radiation-hydrodynamical case with higher-mass initial clumps.704 We plan to perform in the future a study on he robustness of our results for different initial conditions., We plan to perform in the future a study on the robustness of our results for different initial conditions.705 For the smallest luminosities. Run B matches better than tun A the observed 55 of regions that have Saud?<5 Jv kpc. but over-estimates (35 %)) the observed fraction (14% )) of regions with 5<Sood10.," For the smallest luminosities, Run B matches better than Run A the observed 55 of regions that have $S_{\rm 2cm}d^2<5$ Jy $^2$, but over-estimates (35 ) the observed fraction (14 ) of regions with $5<S_{\rm 2cm}d^2<10$."706 Petersοἱal.(2010b) presented statistics of the morphologies of the regions in Run A and 1tun D and found that Run D agrees better with observed surveys., \cite{Peters10b} presented statistics of the morphologies of the regions in Run A and Run B and found that Run B agrees better with observed surveys.707 Although Run A can be interpreted as a moce of isolated massive star formation. the treatment of fragmentation is more realistic in Run B (see Paper D).," Although Run A can be interpreted as a mode of isolated massive star formation, the treatment of fragmentation is more realistic in Run B (see Paper I)."708 Moreover. most high-mass stars form in clusters (Zinnecker&Yorke 2007)..," Moreover, most high-mass stars form in clusters \citep{ZY07}. ."709When the potential is correct. we find to a σοοι approximation that AZx[1sso]. while with an incorrec potential the curve of ZZ can be fitted near its minimum with NEoxIsísy)*|k2 with &=0.,"When the potential is correct, we find to a good approximation that $\Delta710E\propto|1-s/s_0|$, while with an incorrect potential the curve of $\Delta711E$ can be fitted near its minimum with $\Delta E\propto\sqrt{(1-s/s_0)^2+k^2}$ with $k\ne0$."712 Thus in a practical case the quantity & can be used as a measure of how far the potential emploved. dilfers from its true value., Thus in a practical case the quantity $k$ can be used as a measure of how far the potential employed differs from its true value.713 lig., Fig.714 3 explores the effect of varving both the parameters θα and λα of the potential used in the reconstruction of this orbit., \ref{erms2} explores the effect of varying both the parameters $b/a$ and $GM/a$ of the potential used in the reconstruction of this orbit.715 Although an error in one parameter can to some extent be olfset by an error in the other parameter. there is a sharp minimum in AZ when the parameters take their true values.," Although an error in one parameter can to some extent be offset by an error in the other parameter, there is a sharp minimum in $\Delta E$ when the parameters take their true values."716 Thus from this one orbit one can determine both the Ilattening and the mass of the potential with precision., Thus from this one orbit one can determine both the flattening and the mass of the potential with precision.717 Fig., Fig.718 4 illustrates how high this precision is bv showing isodensity contours of the Mivamoto-Nagai mocels that have b/a=0.2 (top) and 0.22., \ref{Miyamfig} illustrates how high this precision is by showing isodensity contours of the Miyamoto-Nagai models that have $b/a=0.2$ (top) and $0.22$.719 Although these models are extremely similar. analvsis of the orbit segment easily. distinguishes them.," Although these models are extremely similar, analysis of the orbit segment easily distinguishes them."720 Changing the length of the orbit segment analysed by a [actor 3 either side of the length shown in Fig., Changing the length of the orbit segment analysed by a factor 3 either side of the length shown in Fig.721 1. does not change the results significantly., \ref{onsky1} does not change the results significantly.722 In particular. the algorithm identifies incorrect. potentials even with quite short. orbit segments.," In particular, the algorithm identifies incorrect potentials even with quite short orbit segments."723 These tests demonstrate that with high-quality. data the algorithm has no cilliculty recovering the truce orbit if the potential is the correct one. and. provides. excellent discrimination against incorrect potentials.," These tests demonstrate that with high-quality data the algorithm has no difficulty recovering the true orbit if the potential is the correct one, and provides excellent discrimination against incorrect potentials."724 Dilferent. streams will be sensitive to cifferent aspects of the Galactic potential streams near the plane will. be sensitive to the structure of the disc. while the Magellanic Steam must be insensitive to the dise but sensitive to the total Galactic mass.," Different streams will be sensitive to different aspects of the Galactic potential – streams near the plane will be sensitive to the structure of the disc, while the Magellanic Steam must be insensitive to the disc but sensitive to the total Galactic mass."725 The Galactic potential would be mos strongly constrained. by a plot like that shown in Fig., The Galactic potential would be most strongly constrained by a plot like that shown in Fig.726 3 in which the quantity contoured is the sum of AL for al available streams., \ref{erms2} in which the quantity contoured is the sum of $\Delta E$ for all available streams.727 Obscuration by cust near the plane sometimes breaks a stream into two or more segments with an unobservec gap between them., Obscuration by dust near the plane sometimes breaks a stream into two or more segments with an unobserved gap between them.728 In principle cach segment. could be analysed. independently. but this would not be the optima procedure.," In principle each segment could be analysed independently, but this would not be the optimal procedure."729 First. all segments should be used to determine the path of the underlving orbit on the sky. including through obscured regions.," First, all segments should be used to determine the path of the underlying orbit on the sky, including through obscured regions."730 Then the algorithm should. be used. to reconstruct cach segment. and the rms energy variation along all reconstructed segments. (treated. as a whole) minimized within the space spanned by the initial distances assumed for cach section.," Then the algorithm should be used to reconstruct each segment, and the rms energy variation along all reconstructed segments (treated as a whole) minimized within the space spanned by the initial distances assumed for each section."731 In practice the input cata will be less exact than in the tests above because tidal streams have finite widths on the sky and do not exactly trace a single orbit., In practice the input data will be less exact than in the tests above because tidal streams have finite widths on the sky and do not exactly trace a single orbit.732 En favourable cases they delineate an orbit with great. precision: for example the tidal stream of a globular cluster such as that of Pal 5 is precisely bounded by the eluster's orbit2003)., In favourable cases they delineate an orbit with great precision: for example the tidal stream of a globular cluster such as that of Pal 5 is precisely bounded by the cluster's orbit.733. The tidal streams of larger bodies. such as the Ser dwarf galaxy. are broad. ancl their relation to the orbit of the progenitor is less evident.," The tidal streams of larger bodies, such as the Sgr dwarf galaxy, are broad, and their relation to the orbit of the progenitor is less evident."734 Correspondinelv. in the case of a globular cluster. the uncertainty in the radial velocity along the stream will be limited to measurement errors in the velocities of stars. while in the case of a dwarf galaxy one has to contend also with the velocity dispersion of the progenitor. which both widens the range of velocities encountered at any point on the stream. and. leads to a systematic olfset. between the mean velocity in the stream and the velocity of the underlying orbit.," Correspondingly, in the case of a globular cluster, the uncertainty in the radial velocity along the stream will be limited to measurement errors in the velocities of stars, while in the case of a dwarf galaxy one has to contend also with the velocity dispersion of the progenitor, which both widens the range of velocities encountered at any point on the stream and leads to a systematic offset between the mean velocity in the stream and the velocity of the underlying orbit."735 When applying the algorithm to a stream such as tha of Pal 5. one would not use the positions and. velocities of individual stars directly. but would make a judgment as to the relation of the orbit to the stellar stream. and estimate the racial velocity on the orbit by averaging the velocities of nearby stars and then making a correction for the expectec olfset. between the stream: and orbital velocities.," When applying the algorithm to a stream such as that of Pal 5, one would not use the positions and velocities of individual stars directly but would make a judgment as to the relation of the orbit to the stellar stream, and estimate the radial velocity on the orbit by averaging the velocities of nearby stars and then making a correction for the expected offset between the stream and orbital velocities."736 When al this had been done. one would contrive that the trajectory of the orbit on the sky and in radial velocity was smooth.," When all this had been done, one would contrive that the trajectory of the orbit on the sky and in radial velocity was smooth."737 Therefore. in tests it is not appropriate to ad independen," Therefore, in tests it is not appropriate to ad independent"738adopted in the Paris uecting (1985: see Table 7 iu the Πασν manual. Ferland 1996).,"adopted in the Paris meeting (1985; see Table 7 in the Hazy manual, Ferland 1996)."739 We used LAIC average abuudances of PNs quoted in Staughelui.etal.(2000) and SAIC average abundances from MeCall(1998) excep for thieCIT ratio which comes from Leisy&Dennetelel(1996).. where we hacl selected from their Table EMi oulv the low-error data.," We used LMC average abundances of PNs quoted in \citet{sta00}740 and SMC average abundances from \citet{srmc} except for the C/H ratio which comes from \citet{ld96}, where we had selected from their Table 3 only the low-error data."741 The stellar ionizing spectra is asstuned to be a blackbody with temperatures aiid luminosities from the ILburuneo evolutionary tracks for the appropriate ealaxian population by Vassiliacis&Wood(199[)., The stellar ionizing spectrum is assumed to be a blackbody with temperatures and luminosities from the H-burning evolutionary tracks for the appropriate galaxian population by \citet{vw}.742. Our iuodel ↻⋜∐⋅⋜↧⋯↸∖↑↸∖↥⋅↴∖↴⋜⋯∖↴∖↴⋯⊔⋯⋜∐⋅↕∑↸∖≼↧↕∐, Our model parameters are summarized in Table 5.743⊺⋜∏⋝↕↸∖⋅⊐∙ The predictious of frou Cloudy models are shown iu Fig↥⋅↸∖∐↕≯∪↥⋅↑↕∐∖≋⋀∖∐⊲∙∫⇀⋀∖∐⊲∙⋜⋯≼↧ Galactic PNs.," The predictions of from Cloudy models are shown in Figure 11 for the SMC, LMC, and Galactic PNs."744 The outcomes of tje Cloudy mode]s are rather iuseusitive to the adopted stellar properties aud he assuued density (c.e.. Fig.," The outcomes of the Cloudy models are rather insensitive to the adopted stellar properties and the assumed density (e.g., Fig."745 12)., 12).746 But frev are very sensitive to the adopted abundance ratios. especially he oxveen abundance.," But they are very sensitive to the adopted abundance ratios, especially the oxygen abundance."747 The moel values of are shown as a function of stellar teniperature over he ranee of temperatures encountered by a st:αἱ of wmass = 2AL... as it evolves from the ACD tip to its uaxinmn post-ACGD temperature.," The model values of are shown as a function of stellar temperature over the range of temperatures encountered by a star of mass = 2, as it evolves from the AGB tip to its maximum post-AGB temperature."748 At stelar teuperatures m excess of ~1tPW the ratio is iu rough agreement with the LOSEut observations (sce also Carnett&Dinerscin(1989. 1988))).," At stellar temperatures in excess of $\sim10^5$ K the ratio is in rough agreement with the present observations (see also \citet{gda,gdb}) )."749 It is worth iotiug tiat the Galactic models presented here are not to be conrpared directly with the known observed distribution. witherut accounine for selection effects that lauor Galactic PN statistics.," It is worth noting that the Galactic models presented here are not to be compared directly with the known observed distribution, without accounting for selection effects that hamper Galactic PN statistics."750 We should carity that the resuts of Figure 1l are valid for the input abundances., We should clarify that the results of Figure 11 are valid for the input abundances.751 If. for example. we chose to use as input the average oxveen albnucdaucο for SAIC ONs from Leisy&Deunucfeld(1996) iusteacd of that of Stasiuska.Richer.&MeCall(1998).. the over enaission line would be as high as 7.2.," If, for example, we chose to use as input the average oxygen abundance for SMC ONs from \citet{ld96} instead of that of \citet{srmc}, the over emission line would be as high as 7.2."752 The cooliug processes that cletermine T.(O!|) in the SMC. EMC aud Galactic PNs are noteworthy.," The cooling processes that determine $_e$ $\rm^{++})$ in the SMC, LMC and Galactic PNs are noteworthy."753" In the Galaxy the primary cooants of PNs with hot σστα] stars are the optical forbidden lines of ABOUT and other lines of O! alc OF,", In the Galaxy the primary coolants of PNs with hot central stars are the optical forbidden lines of $\lambda$ 5007 and other lines of $\rm^+$ and $\rm^{++}$.754 However. in cuvirommenuts in which Ο/Π is as low as in the SMC. the primary coolants may. become 1Itraviolet intercombination lines of C! and C!.," However, in environments in which O/H is as low as in the SMC, the primary coolants may become ultraviolet intercombination lines of $\rm^+$ and $\rm^{++}$."755 It will be interesti18o to coufiru these predictions with fuure UV observations., It will be interesting to confirm these predictions with future UV observations.756 Dokxe reaching conclusious alu further speculation. let us explore tje weight of our assunmptious.," Before reaching conclusions and further speculation, let us explore the weight of our assumptions."757 ο is that fjese 1nodoels have const:uit. hydrogen density. which is arbitrarily assuned equal to LOOyom," One is that these models have constant hydrogen density, which is arbitrarily assumed equal to 1000 $^{-3}$."758" 5m running Cloudy for a coistaut dcusity model. the outer radius is determined bs ""the Strocmigen spiere."," By running Cloudy for a constant density model, the outer radius is determined by the Stroemgen sphere."759 This makes ourn models larger as stellay temperature increases. reaching a maxili ali then decliI5 with tje A5 inteusiVv ratio.," This makes our models larger as stellar temperature increases, reaching a maximum and then decline with the $\lambda$ intensity ratio."760 Bu in order to compare the radii of our models with our dat: vawe shoud luae a nuoclel for cach jebula w1 he correct radius (.0.. assunidug a density profile that reproxduces theit PN. for exauple. as deteined iu hydrodynamic calculations).," But in order to compare the radii of our models with our data we should make a model for each nebula with the correct radius (i.e., assuming a density profile that reproduces that PN, for example, as determined in hydrodynamic calculations)."761 This will be dojo du cetail in a futuret papel., This will be done in detail in a future paper.762 For now. let us examine what a (iference iu the average bydrogen deusitv wil make in the 5007 flux ratio.," For now, let us examine what a difference in the average hydrogen density will make in the 5007 flux ratio."763 Iu Figure 12 we pot the ratio against the oxygen abundance for the carly aud tιο hottest models in tιο SAIC! (triangles). LMC (squares). aud the Galaxy (circles).," In Figure 12 we plot the ratio against the oxygen abundance for the early and the hottest models in the SMC (triangles), LMC (squares), and the Galaxy (circles)."764 The early models correspond models 1. 6. and 11 of Table 5. xwi hlog Ny=3 ?.," The early models correspond to models 1, 6, and 11 of Table 5, with log $_{\rm H}$ =3 $^{-3}$."765 The hot models have been calculated for log Nyp=2.5. 3. and 3.5 cmi7.," The hot models have been calculated for log $_{\rm H}$ =2.5, 3, and 3.5 $^{-3}$."766 We see that varving he density does not affect very uch the studied mteusitv ratio. especialv for the low oxvecn models.," We see that varying the density does not affect very much the studied intensity ratio, especially for the low oxygen models."767" T1ο (iferent hot SAIC models in Figure 12 are all within 6 ))—1. ancl their outer radii are 0.11. 0,25. aud 0.55 parsees for log Nj 23.5. 3. and 2.5 sim? respectively."," The different hot SMC models in Figure 12 are all within $\delta$ )=1, and their outer radii are 0.11, 0.25, and 0.55 parsecs for log $_{\rm H}$ =3.5, 3, and 2.5 $^{-3}$ respectively."768 Al models described iu tus xjer have filling factor (6) equal to unity., All models described in this paper have filling factor $\epsilon$ ) equal to unity.769 Chaneine the value of the filling factor will also change the otter radius of the nebular models., Changing the value of the filling factor will also change the outer radius of the nebular models.770 For example. the hot Calactic model (with log Nj 22) with e-0.5 has an οuter radius about 1.3 times larecr than the model with e=1.," For example, the hot Galactic model (with log $_{\rm H}$ =3) with $\epsilon$ =0.5 has an outer radius about 1.3 times larger than the model with $\epsilon$ =1."771 We fud that the A5007/TE3 intensity ratio with the differeut radius and filling factor is 95% the ratio in the unity filling factor model. thus the filling factor assumption is not very important the discussion of the 5007 line intensity with respect to )," We find that the $\lambda$ intensity ratio with the different radius and filling factor is $\%$ the ratio in the unity filling factor model, thus the filling factor assumption is not very important the discussion of the 5007 line intensity with respect to ."772Vallisueri for sugeestious in the carly development of the work.,Vallisneri for suggestions in the early development of the work.773 We thank Chad Calley. Mansi Wasliwal. Ίνα. Mandel anc Eran Ofek for careful reading of the manuscript. aud Yanbei Chen. Curt Cutler. Weu-Fai Fone. Chris Hirata. Tom Prince. Bernard Schutz. hip Thorne. Linging Wen aud Stan Whitcomb for helpful ciscussious.," We thank Chad Galley, Mansi Kasliwal, Ilya Mandel and Eran Ofek for careful reading of the manuscript, and Yanbei Chen, Curt Cutler, Wen-Fai Fong, Chris Hirata, Tom Prince, Bernard Schutz, Kip Thorne, Linqing Wen and Stan Whitcomb for helpful discussions."774 Simulations were performed using the Suuuvvale cluster at CITA. which is fuuded bv NSERC aud CIAR.," Simulations were performed using the Sunnyvale cluster at CITA, which is funded by NSERC and CIAR."775 SAIN’s research was carried out at the Jet Propulsion Laboratory. California Lhustitute of Technology under a coutract with National Aeronautics and Space Acininistration.," SMN's research was carried out at the Jet Propulsion Laboratory, California Institute of Technology under a contract with National Aeronautics and Space Administration."776"In Figure 10, the nuclear flows for the reactions that bridge from A«12 (the pp-chain region) to A>12 CNO region) are shown as a function of the temperature(the before (Tg> 3) and after (Το 3) the onset of the rp-process.","In Figure 10, the nuclear flows for the reactions that bridge from $A777< 12$ (the pp-chain region) to $A \ge 12$ (the CNO region) are shown as a function of the temperature before $T_9 > 3$ ) and after $T_9 <7783$ ) the onset of the $\nu$ p-process."779 The nuclear flows at Tg=2.5 for the relevant N-Z region are also shown in Figure 11., The nuclear flows at $T_9 = 2.5$ for the relevant $N$ $Z$ region are also shown in Figure 11.780" Here, the nuclear flow is defined as the difference between the time-derivatives of abundances for the forward and inverse reactions of a given channel."," Here, the nuclear flow is defined as the difference between the time-derivatives of abundances for the forward and inverse reactions of a given channel."781" It is clear that, at a high temperature > the triple-a process (withtherateofCaughlan(Το&3),Fowler plays a dominant role for the breakout from the pp-chain region."," It is clear that, at a high temperature $T_9 > 3$ ), the $\alpha$ process \citep[with the rate of][]{Caug1988} plays a dominant role for the breakout from the pp-chain region."782" We find, however, a couple of 2-body reaction sequences and ""Be(o, p)!°B(a, p)'?C compete with thetriple-a ""Be(a,*5)!C(o,p)!4Nprocess during the vp-process phase?."," We find, however, a couple of 2-body reaction sequences $^7$ $(\alpha,783\gamma)^{11}$ $(\alpha, p)^{14}$ N and $^7$ $(\alpha,784p)^{10}$ $(\alpha, p)^{13}$ C compete with the$\alpha$ process during the $\nu$ p-process ."785. 'Table 2 lists the reaction rates and decay timescales for the relevant isotopes at 79=2.5 and 2.0., Table 2 lists the reaction rates and decay timescales for the relevant isotopes at $T_9 = 2.5$ and 2.0.786" It is clear that ""Be(o,*)!!C, 4 orders of magnitude slower than ! C(o, p)'*N, governs the former sequence."," It is clear that $^7$ $(\alpha,787\gamma)^{11}$ C, 4 orders of magnitude slower than $^{11}$ $(\alpha,788p)^{14}$ N, governs the former sequence."789" For the latter, p)!?C, although a factor of 10 smaller than ""Be(a, p)!°B, 1°B(a,mainly controls the reaction flow, which takes away nuclear abundances from !°B formed by the endothermic reaction."," For the latter, $^{10}$ $(\alpha, p)^{13}$ C, although a factor of 10 smaller than $^7$ $(\alpha, p)^{10}$ B, mainly controls the reaction flow, which takes away nuclear abundances from $^{10}$ B formed by the endothermic reaction."790" Figure 10 shows that ""Be(o,4)!!C and !°B(a, p)'?C exhibit similar roles to triple-a in the temperature range relevant to the vp-process."," Figure 10 shows that $^7$ $(\alpha, \gamma)^{11}$ C and $^{10}$ $(\alpha, p)^{13}$ C exhibit similar roles to $\alpha$ in the temperature range relevant to the $\nu$ p-process."791" Therefore, we select these three reactions for the sensitivity tests."," Therefore, we select these three reactions for the sensitivity tests."792" Note that the unstable isotope !!C produced is followed by H1 C(o, (seeHayakawaetal.2010,forarecentex-perimentalp)!4Nevaluationofthisrate) before decaying back to B. All the data of these three reactions, from Wagoner B(o,p)9Ο) and Caughlan&Fowler(1988,in the REACLIB compilation, are based on remainder)experimental information of single resonance states."," Note that the unstable isotope $^{11}$ C produced is followed by $^{11}$ $(\alpha,793p)^{14}$ N \citep[see][for a recent experimental evaluation of this794rate]{Haya2010} before decaying back to $^{11}$ B. All the data of these three reactions, from \citet[][for795$^{10}$ $(\alpha, p)^{13}$ and \citet[][for the796remainder]{Caug1988} in the REACLIB compilation, are based on experimental information of single resonance states."797 Contribution from (possible) resonances at higher excitation energies could thus sizably change these rates., Contribution from (possible) resonances at higher excitation energies could thus sizably change these rates.798The analytic fitting formula lor this case A(T;ερ} has the same form as equations (4.1). (L2). (4.3).,"The analytic fitting formula for this case $K(\Gamma_{ij}, \epsilon_{r})$ has the same form as equations (4.1), (4.2), (4.3)."799 We have presented a calculation of the enhancement of the resonant thermonuclear reaction rates lor extremely dense stellar plasmas., We have presented a calculation of the enhancement of the resonant thermonuclear reaction rates for extremely dense stellar plasmas.800 The calculation has been carried out. by adopting (he screening potential derived from (he Monte Carlo computations of (he classical one-component plasma., The calculation has been carried out by adopting the screening potential derived from the Monte Carlo computations of the classical one-component plasma.801 We have summarized our numerical restults by an accurate analylic fitting formula to facilitate applications., We have summarized our numerical results by an accurate analytic fitting formula to facilitate applications.802 The present results will be useful if the PC + PC fusion reaction contains narrow resonances in (he astrophvsical enerev range., The present results will be useful if the $^{12}$ C + $^{12}$ C fusion reaction contains narrow resonances in the astrophysical energy range.803 We wish to thank Ix. Langanke5 for making5 the preprint available (o us prior to its publication., We wish to thank K. Langanke for making the preprint available to us prior to its publication.804 We also (hank Y. Ovanaei for allowing us to use the least-squares fitting program SALS., We also thank Y. Oyanagi for allowing us to use the least-squares fitting program SALS.805 We are most 5grateful to our referee for many valuable comments on the original5 manuscript which helped us tremendously in revising the manuscript., We are most grateful to our referee for many valuable comments on the original manuscript which helped us tremendously in revising the manuscript.806 This work is financially supported in part by Grants-in-Aid of the Japanese Ministry of Education. Culture. Sports. Science. and Technology under contracts 13640245. 13740129.," This work is financially supported in part by Grants-in-Aid of the Japanese Ministry of Education, Culture, Sports, Science, and Technology under contracts 13640245, 13740129."807To extract the Ia line profile. we fit the continuum in the wavelength range of 5400-75004 with a power-law.,"To extract the $\alpha$ line profile, we fit the continuum in the wavelength range of $-$ with a power-law."808 Apparent emission lines are masked curing the fit., Apparent emission lines are masked during the fit.809" This vields a spectral index of o—1.03 Cf,xv °).", This yields a spectral index of $\alpha=1.03$ $f_\nu\propto \nu^{-\alpha}$ ).810 Extrapolating (ie power-law to short wavelengths predicts a continuum hieher than the observed one. suggesting the presence of reddening or intrinsic sleepenine of spectrum in the UV band.," Extrapolating the power-law to short wavelengths predicts a continuum higher than the observed one, suggesting the presence of reddening or intrinsic steepening of spectrum in the UV band."811 Figure 2 shows the line prolile of Ha (left panel)., Figure 2 shows the line profile of $\alpha$ (left panel).812 The profile is apparently skewed to the red. and the blue peak is higher than the red one.," The profile is apparently skewed to the red, and the blue peak is higher than the red one."813 These features are signatures of lines originating [rom a relativistic disk., These features are signatures of lines originating from a relativistic disk.814 A circular disk model is considered first., A circular disk model is considered first.815 Defore making a detailed fit. we ean estimate approximately the inner. outer radii and (he inclination bv using the (wo peak positions. the maxinmnmn extension of the red ancl blue wings. following the expression of Doppler [actor for circular disk model given by ClI39.," Before making a detailed fit, we can estimate approximately the inner, outer radii and the inclination by using the two peak positions, the maximum extension of the red and blue wings, following the expression of Doppler factor for circular disk model given by CH89."816" The two peaks al 6265 and in the source rest frame suggest an outer radius of around 1107,. and an inclination angle /&38"" (7=0"" for a face-on disk)."," The two peaks at 6265 and in the source rest frame suggest an outer radius of around $r_g$ , and an inclination angle $i\sim38$ ( $i=0$ for a face-on disk)."817" The maximum extensions of blue and red wings αἱ and suggest an inner radius of around 90r,.", The maximum extensions of blue and red wings at and suggest an inner radius of around $r_g$.818 Thus most of the line flux is produced in a narrow annulus of radii around 100 ry., Thus most of the line flux is produced in a narrow annulus of radii around 100 $r_g$.819" A detailed fit to the Ha profile with a circular. relativistic Keplerian disk model (CII89) with a power-law distribution of emissivity (F(r)xr 7) vields parameters: rj,=98ry. Fou=227 ry.=418. =q28Tkm2 qo=4.6."," A detailed fit to the $\alpha$ profile with a circular, relativistic Keplerian disk model (CH89) with a power-law distribution of emissivity $F(r)\propto r^{-\beta}$ ) yields parameters: $r_{in}=98\;r_g$, $r_{out}=227\;r_g$, $i=41$, $\sigma=1287$, $q=4.6$."820 Narrow lines ([OI]. [NI]. [SII] and narrow Ila) were masked during the fit.," Narrow lines ([OI], [NII], [SII] and narrow $\alpha$ ) were masked during the fit."821 However. this fit does not reproduce either the right position of the red peak or the relative height of the two peaks as displaved in Fig 2. which is a strong indication for the deviation [rom axi-svinmetry of the disk. in emissivity or/and in kinematics.," However, this fit does not reproduce either the right position of the red peak or the relative height of the two peaks as displayed in Fig 2, which is a strong indication for the deviation from axi-symmetry of the disk, in emissivity or/and in kinematics."822 Several (vpes of asymmetry. such as an elliptical disk. presence of hot spols or a spiral aym may reproduce the line profile.," Several types of asymmetry, such as an elliptical disk, presence of hot spots or a spiral arm may reproduce the line profile."823 In order to compare parameters for (his object with those presented in EIIO3 and $03. elliptical disk models (Eracleous et al.," In order to compare parameters for this object with those presented in EH03 and S03, elliptical disk models (Eracleous et al."824 1995) are (hen (ried., 1995) are then tried.825 The profile can be reasonably well fitted using (his model., The profile can be reasonably well fitted using this model.826" The derived disk parameters are: rj,=62ry. ry,=Siry. d=39° 0=1335 boq24. 62-025 and Oy=I8""."," The derived disk parameters are: $r_{in}=62$$r_g$, $r_{out}=87$$r_g$, $i=35$, $\sigma=1335$, $q=2.7$, $e=0.25$ and $\phi_0=18$."827. An additional svinmetrie broad component of FWIIM ~ 3300 kins | is required to reproduce the whole Ha line., An additional symmetric broad component of FWHM $\sim$ 3300 km $^{-1}$ is required to reproduce the whole $\alpha$ line.828 The equivalent width (EW) of this component is214., The equivalent width (EW) of this component is.829. Because the profile of broad [Lo is similar to that of La and blends with II5 line. we simply rescale the best fitted IIa. model by a factor of 0.45 in the flux density (to match the 1L? line profile.," Because the profile of broad $\beta$ is similar to that of $\alpha$ and blends with $\gamma$ line, we simply rescale the best fitted $\alpha$ model by a factor of 0.45 in the flux density to match the $\beta$ line profile."830 This vields Ha /IL2~2.9 for the double-peaked component. which is typical of Sevfert galaxies and quasars as well as double-peaked line emitters.," This yields $\alpha$ $\beta\simeq2.9$ for the double-peaked component, which is typical of Seyfert galaxies and quasars as well as double-peaked line emitters."831 The fluxes of the narrow lines in the wavelength ranges covering the double-peaked La and LL? lines are measured after the best fitting model is subtracted., The fluxes of the narrow lines in the wavelength ranges covering the double-peaked $\alpha$ and $\beta$ lines are measured after the best fitting model is subtracted.832At other wavelengths. we use the local ‘continuum that includes weak higher order double-peakecl Balmer Ines.,"At other wavelengths, we use the local `continuum' that includes weak higher order double-peaked Balmer lines."833 Narrow lines are modelled using Gaussian functions., Narrow lines are modelled using Gaussian functions.834 Two components are required for [OIL], Two components are required for [OIII]835environment of local Seyfert galaxies 2010).,environment of local Seyfert galaxies .836" That said, if such destruction (Diamond-Stanicwere important, it would mean that the connection between star-formation rate and black hole accretion rate (a central result of this paper, see Section ??)) is actually stronger than we've presented."," That said, if such destruction were important, it would mean that the connection between star-formation rate and black hole accretion rate (a central result of this paper, see Section \ref{sec:results}) ) is actually stronger than we've presented."837 We therefore adopt the 11.3 um aromatic feature as the most robust tracer of the SFR for our sample., We therefore adopt the 11.3 $\mu$ m aromatic feature as the most robust tracer of the SFR for our sample.838 We adopt an uncertainty of 0.2 dex on conversions between the 11.3 wm aromatic feature strength and IR luminosity based on the scatter in this ratio for the SINGS sample and an additional uncertainty of 0.2 dex for conversions between IR luminosity and SFR2009)., We adopt an uncertainty of 0.2 dex on conversions between the 11.3 $\mu$ m aromatic feature strength and IR luminosity based on the scatter in this ratio for the SINGS sample and an additional uncertainty of 0.2 dex for conversions between IR luminosity and SFR.839". Adding these in quadrature, the uncertainty(Rieke on SFRs obtained from equation 2 is 0.28 dex."," Adding these in quadrature, the uncertainty on SFRs obtained from equation \ref{eq:pah} is 0.28 dex."840" In Figure 5,, we show the relationship between BHAR, as traced byrv], and nuclear SFR, as traced by the 11.3 wm aromatic feature."," In Figure \ref{fig:mdotsfr}, we show the relationship between BHAR, as traced by, and nuclear SFR, as traced by the 11.3 $\mu$ m aromatic feature."841 A strong correlation is apparent: Seyferts with larger BHARs tend to have larger nuclear SFRs., A strong correlation is apparent: Seyferts with larger BHARs tend to have larger nuclear SFRs.842 We use the linear regression outlined by to quantify the relationship between nuclear SFR and BHAR: The uncertainties on the regression parameters above correspond to the interval that includes of the posterior distribution for each parameter (see Table 2)., We use the linear regression outlined by to quantify the relationship between nuclear SFR and BHAR: The uncertainties on the regression parameters above correspond to the interval that includes of the posterior distribution for each parameter (see Table 2).843" The best-fit regression line and confidence interval, given the uncertainties in the regression parameters, are shown as solid and dashed lines in Figure 5.."," The best-fit regression line and confidence interval, given the uncertainties in the regression parameters, are shown as solid and dashed lines in Figure \ref{fig:mdotsfr}."844 The observed scatter around this relationship is 0.52 dex, The observed scatter around this relationship is 0.52 dex845or in the form Ea: :)).,"or in the form ) ),."846 For emission in the backward direction is given one formally. needs to replace 0 by «*—B., For emission in the backward direction is given one formally needs to replace $\theta$ by $\pi-\theta$.847 The emissivitv. (3.1)) is analogous to the corresponding formula for svnchrotron enussion (Ginzburg&Svrovalski1965).. with the evelotvon Irequeney. replaced by the frequency. Q. of the LAEW.," The emissivity \ref{eta2}) ) is analogous to the corresponding formula for synchrotron emission \citep{gs65}, with the cyclotron frequency replaced by the frequency, $\Omega$, of the LAEW."848 An important difference is that lor LAE the angular distribution is confined to a small cone about the direction of motion. being zero strictly along this direction. @=0. whereas for svichrotron emission (he maximum occurs where the angle. ϐ of emission is equal to the pitch angle of the particle.," An important difference is that for LAE the angular distribution is confined to a small cone about the direction of motion, being zero strictly along this direction, $\theta=0$, whereas for synchrotron emission the maximum occurs where the angle, $\theta$, of emission is equal to the pitch angle of the particle."849 The dependence on ϐ in 3.1)). (3.1)) causes mathematical difficulties. which we eloss over in the following discussion ancl address explicitly in Sec. 4.4.," The dependence on $\theta$ in \ref{eta2a}) ), \ref{eta2}) ) causes mathematical difficulties, which we gloss over in the following discussion and address explicitly in Sec. \ref{dilemma}."850 The power radiated by an inclividual particle can be evaluated by integrating the enissivitv over frequency anc solid angle., The power radiated by an individual particle can be evaluated by integrating the emissivity over frequency and solid angle.851 Performing the integral over Irequency first gives, Performing the integral over frequency first gives852signal-to-noise of our spectra at 1000 lis about 200.,signal-to-noise of our spectra at 4000 is about 200.853 Appareuth. oue week is too short for observiug recurrent variability or for checking a periodicity of 17.3 d. Iu that respect our results have a preliminary character.," Apparently, one week is too short for observing recurrent variability or for checking a periodicity of 17.3 d. In that respect our results have a preliminary character."854 The positive aspect is. though. that our observations (Fig.," The positive aspect is, though, that our observations (Fig."855 1) do show the presence of rapidly changing absorption compoucuts in profiles of the higher members of the Daliier series., 1) do show the presence of rapidly changing absorption components in profiles of the higher members of the Balmer series.856 The lines showed split absorption cores., The lines showed split absorption cores.857 These are best visible in the profiles of 28/05/99 (solid lues in Fie., These are best visible in the profiles of 28/05/99 (solid lines in Fig.858 1). which have two distinct cores.," 1), which have two distinct cores."859 Closer examination shows that these cores existed also in later spectra., Closer examination shows that these cores existed also in later spectra.860 The cores approached cach other during the next few davs aud nearly merged around 02/06/99., The cores approached each other during the next few days and nearly merged around 02/06/99.861 Thereafter they reappeared., Thereafter they reappeared.862" We measured the waveleugth variation of these absorptio- cores by two methods: directly by determining the central waveleusths of the absorption conponeut. aud also by fitting the profiles by two (απορίαας,"," We measured the wavelength variation of these absorption cores by two methods: directly by determining the central wavelengths of the absorption component, and also by fitting the profiles by two Gaussians."863 This was done in the two lines. so we got. 1 data poiuts for cach componenutat each dav.," This was done in the two lines, so we got 4 data points for each componentat each day."864 The moeasuremenuts agree reasonably well: the mean scatter of the individual velocity nieasureineuts is 3 ans+.," The measurements agree reasonably well; the mean scatter of the individual velocity measurements is 3 ${\rm km{\,}{s^{-1}}}$."865 Table 1 gives the average velocities measured in the two liue components., Table 1 gives the average velocities measured in the two line components.866 We wanted to check if these velocities are related in some way to the photometric period of 17.3 d. reported in Section 1.," We wanted to check if these velocities are related in some way to the photometric period of 17.3 d, reported in Section 1."867 To that cud we replotted the data iu the same diaerauuC» for a period. 17.3 days later (Fig.," To that end we replotted the data in the same diagram for a period, 17.3 days later (Fig."868oO 2)., 2).869 The diagram.C» shows that it isinpossible that the observations have a repetition period of 17.3 d. and it appears moreover that there are possibly two waves that are “swinging in counterphase’. the dotted lines in Fie.," The diagram shows that it is that the observations have a repetition period of 17.3 d, and it appears moreover that there are possibly two waves that are 'swinging in counterphase', the dotted lines in Fig."870 2., 2.871 If this interpretation would be verified by measureimeuts taken over a period of a mouth or longer. then they could be interpreted in terms of a long wave in the wind over the stellar surface having a wavelength of the order of the stars diameter.," If this interpretation would be verified by measurements taken over a period of a month or longer, then they could be interpreted in terms of a long wave in the wind over the stellar surface having a wavelength of the order of the star's diameter."872 This hypothesis will be exaiined in the next section., This hypothesis will be examined in the next section.873 Au interesting aspect of the observations is the hieh acceleration that is interred from the waveleneth variation of the absorption compoucuts., An interesting aspect of the observations is the high acceleration that is inferred from the wavelength variation of the absorption components.874 It is about 10 kmsτα3 , It is about 10 ${\rm km}{\:}{{\rm s}^{-1}}{\:}{{\rm d}^{-1}}$.875"This value stands in contrast to the acceleration of the DACs. which is about LOO times smaller. of the order of 0.1 lausτα tf,"," This value stands in contrast to the acceleration of the long-period DACs, which is about 100 times smaller, of the order of 0.1 ${\rm km}{\,}{{\rm s}^{-1}}{\,}{{\rm d}^{-1}}$ ."876Active Galactic Nuclei (ΔΝ) can be «etected. at. high redshift due to their high luminosities anc hence they can be used to place stringent constraints on the evolution of Large Scale Structure over a wide recshift range (sec Llartwick οἱ Schade 1950].,Active Galactic Nuclei (AGN) can be detected at high redshift due to their high luminosities and hence they can be used to place stringent constraints on the evolution of Large Scale Structure over a wide redshift range (see Hartwick et Schade 1989).877 Our knowledge on the AGN clustering properties comes mainly Crom laree optical UV excess (UWA) surveys lor QSOs (Bovle ο al., Our knowledge on the AGN clustering properties comes mainly from large optical UV excess (UVX) surveys for QSOs (Boyle et al.878 1988. Fane Mo 1993. Shanks Bovle 1994. CrOOLL Shanks 1996. La Franca. Andreani Cristiani 908).," 1988, Fang Mo 1993, Shanks Boyle 1994, Croom Shanks 1996, La Franca, Andreani Cristiani 1998)."879 Croom Shanks (1996) analvsed. the clustering properties. of the Large Might Quasar Survey (LBQS) ancl combined. the results with that obtained from other QSO surveys including the Dirham/AAT UVX sample., Croom Shanks (1996) analysed the clustering properties of the Large Bright Quasar Survey (LBQS) and combined the results with that obtained from other QSO surveys including the Durham/AAT UVX sample.880" They derived a clustering length of ο,=544+1.1h Alpe at:v mean. recdshift of 1.27.", They derived a clustering length of $r_{\circ}=5.4\pm1.1 \;h^{-1}$ Mpc at a mean redshift of 1.27.881 La Franca. Andreani Cristiani (998) derived comparable results (1:=6241.6+ Mpc) by investigating a sample of ~ του uasars in the redshit range 0.3«>< 3.2.," La Franca, Andreani Cristiani (1998) derived comparable results $r_{\circ}=6.2\pm1.6 \;h^{-1}$ Mpc) by investigating a sample of $\sim$ 700 quasars in the redshift range $0.3<z\leq3.2$ ."882 Comparison of these clustering results in dilferent redshifts rather lavours a comoving model for the evolution of clustering., Comparison of these clustering results in different redshifts rather favours a comoving model for the evolution of clustering.883 According to this model the amplitude of the correlation function remains fixed with redshifi πι coordinates as the &galaxy Njpair expands1 together8 with the background mass distribution., According to this model the amplitude of the correlation function remains fixed with redshift in comoving coordinates as the galaxy pair expands together with the background mass distribution.884 In contrast. in the stable model for clustering evolution. AGN trace clumps of mass which have gravitationally collapsed in bound units and have therefore ceased to take part in the general expansion of the universe.," In contrast, in the stable model for clustering evolution, AGN trace clumps of mass which have gravitationally collapsed in bound units and have therefore ceased to take part in the general expansion of the universe."885 However. most of the AGN samples used in the analyses above. come from pencil beam surveys and contain a large fraction of high redshift AGN.," However, most of the AGN samples used in the analyses above, come from pencil beam surveys and contain a large fraction of high redshift AGN."886 Ceorgantopoulos Shanks (1994) studied the correlation function of a sample of about 200 LILAS selected Sevfert galaxies., Georgantopoulos Shanks (1994) studied the correlation function of a sample of about 200 IRAS selected Seyfert galaxies.887 Vhey detect a 2a clustering signal on scales less than 10 f Mpe., They detect a $2\sigma$ clustering signal on scales less than 10 $h^{-1}$ Mpc.888 Although heir statistics are limited. their results are more consisten with a comoving model.," Although their statistics are limited, their results are more consistent with a comoving model."889 It is evident that there is a pressing need for Iarge αλ samples in the local universe in order to dace tight constraints on the clustering evolution in a broa redshift range., It is evident that there is a pressing need for large AGN samples in the local universe in order to place tight constraints on the clustering evolution in a broad redshift range.890 Analogous clustering studies in N-ravs have been scarce out they. are potentially interesting às they provide the opportunity to derive information on how the X-ray selected AGN race the underlying mass distribution., Analogous clustering studies in X-rays have been scarce but they are potentially interesting as they provide the opportunity to derive information on how the X-ray selected AGN trace the underlying mass distribution.891 Phe first direct study (Cusine redshift information) of the correlation function and lustering properties of X-ray selected ΑΝ. is that of Bovle Ao (1993)., The first direct study (using redshift information) of the correlation function and clustering properties of X-ray selected AGN is that of Boyle Mo (1993).892 They studied the local z<0.2 , They studied the local $z<0.2$ 893If. on the other haud. the stacking is orderly. then the situation can be visualized. by cousiderius each laver as a harmonic oscillator resonating at oue of the 7 electrou frequencies (the analog of quantum states) forming a baud.,"If, on the other hand, the stacking is orderly, then the situation can be visualized by considering each layer as a harmonic oscillator resonating at one of the $\pi$ electron frequencies (the analog of quantum states) forming a band."894 If oue ideutical layer is added aud is electrically coupled to the previous ouc. then the frequency degeneracy is lifted by the perturbation. as is well known. and two new resonance frequencies appear instead. bracketing the original one. and separated by an interval Aw proportional to the coupling cocfücient between two adjacent lavers. aud roughly equal to the FWIINI of the π resonance. which is ~1.5 eV for eraphite (Bassani and Parravicini (1973).. Marinopoulos et al. (2001)..," If one identical layer is added and is electrically coupled to the previous one, then the frequency degeneracy is lifted by the perturbation, as is well known, and two new resonance frequencies appear instead, bracketing the original one, and separated by an interval $\Delta\omega$ proportional to the coupling coefficient between two adjacent layers, and roughly equal to the FWHM of the $\pi$ resonance, which is $\sim1.5$ eV for graphite (Bassani and Pastori-Parravicini \cite{bp}, Marinopoulos et al. \cite{mar},"895 IKlinteuberg et al. (2009)))., Klintenberg et al. \cite{kli}) ).896 For N lavers. there are N new resonance frequencies (or electronic quautuu states) but these are all included in this interval (see Cirüuncels et al. (2008).," For N layers, there are N new resonance frequencies (or electronic quantum states) but these are all included in this interval (see Grünneis et al. \cite{gru},"897 Fie., Fig.898 13)., 13).899 Iu eraphite. N is very large aud the interval is densely filled.," In graphite, N is very large and the interval is densely filled."900 If tle full natural width of cach resonance is dw (that of a single laver graphene. which we fouud in Sec.," If the full natural width of each resonance is $\delta\omega$ (that of a single layer graphene, which we found in Sec."901" 2 to be ~0.75 eV) it is apparent that the mimi, ΠΟΟΙ. of oscillators (.0. graphene lavers) necessary to practically “All” that interval is Nay2Aw/dw."," 2 to be $\sim0.75$ eV) it is apparent that the minimum number of oscillators (i.e. graphene layers) necessary to practically “fill"" that interval is $_{m}=2\,\,\Delta\omega/\delta\omega$."902 For the widths adopted above. this turus out to be ~f.," For the widths adopted above, this turns out to be $\sim4$."903 Many measured properties of eraphite are found to be practically recovered by stacking about 5 eraphene lavers (sce Gein and Novoscloy (2007).. Michel id Verberck (2005). Fie.," Many measured properties of graphite are found to be practically recovered by stacking about 5 graphene layers (see Geim and Novoselov \cite{gn}, Michel and Verberck \cite{mv}, Fig."904 (|., 6).905" This iust hold for the dielectric functions as well: as the umber of lavers increases bevoud Land the mumber of quantum states that lie within Aw increases proportionately. the 3-dimensional character of the stack iucreases nearly with N. For a stack of N lavers. if we define a relative umber of lavers as s=(N-1)/N,,. a resonable approximation to the average dielectric function is. therefore. ILowever. this has also to be averaged over alb bricks in a grain. so N ds replaced by its average N. and we cud up again with a weighted average of the dielectric fuuctious of graphene and graphite."," This must hold for the dielectric functions as well: as the number of layers increases beyond 1 and the number of quantum states that lie within $\Delta\omega$ increases proportionately, the 3-dimensional character of the stack increases linearly with N. For a stack of N layers, if we define a relative number of layers as $s$ $_{m}$, a resonable approximation to the average dielectric function is, therefore, However, this has also to be averaged over all bricks in a grain, so N is replaced by its average $\bar{\rm{N}}$, and we end up again with a weighted average of the dielectric functions of graphene and graphite."906 We assume this is true for both polarizations of the electric field., We assume this is true for both polarizations of the electric field.907 Remembering that iu a eiven model IS exain. the geraplütie bricks are randomly oriented. we have now to apply the BrugecimanOO mixing formula to mixtures of parallel aud transverse polarizations. as for PG (see Tutroduction).," Remembering that in a given model IS grain, the graphitic bricks are randomly oriented, we have now to apply the Bruggeman mixing formula to mixtures of parallel and transverse polarizations, as for PG (see Introduction)."908 The dielectric function of eraphite iu the parallel polarization is subject to ercat uncertainty because of the experimental dcifficultv of preseutiug a proper edge face to the light beam (see discussion in Draine aud Lee (198 3)., The dielectric function of graphite in the parallel polarization is subject to great uncertainty because of the experimental difficulty of presenting a proper edge face to the light beam (see discussion in Draine and Lee \cite{dl}) ).909 A further, A further910questions is investigation of structures and kinematics of dense cores around intermediate- protostars.,questions is investigation of structures and kinematics of dense cores around intermediate-mass protostars.911 For these purposes. we have been conducting survey observations toward protostellar cores in the Orion Molecular Cloud -2 and -3 region (OMC-2/3)./ relatively close (d = 450 pe: Genzel and Stutzki 1989) and (he most representative intermecdiate-nass region.," For these purposes, we have been conducting survey observations toward protostellar cores in the Orion Molecular Cloud -2 and -3 region (OMC-2/3), relatively close $d$ = 450 pc; Genzel and Stutzki 1989) and the most representative intermediate-mass star-forming region."912 We. here. report the results of one of the tvpical intermecdiate-1iass protostars. MAS 7 in OMC-3. in the ο (1:0). 2CO(1 0) and 3.3 mm continuum emission with the Nobevama Millimeter Array (NALA) and the Nobevama 45 m telescope.," We, here, report the results of one of the typical intermediate-mass protostars, MMS 7 in OMC-3, in the $^{13}$ $^{+}$ (1–0), $^{12}$ CO(1–0) and 3.3 mm continuum emission with the Nobeyama Millimeter Array (NMA) and the Nobeyama 45 m telescope."913 The 00 (10) molecular line is one of the most appropriate (racers of dense gas. which has a high critical density (regi~LO?em. 3) and is usually detected toward low- to intermediate-mass protostellar envelopes (Takakuwaetal.2000:Saito2001:Fuente2005).," The $^{13}$ $^{+}$ (1–0) molecular line is one of the most appropriate tracers of dense gas, which has a high critical density $n_{\rm{crit}}~{\sim}~10^5~\rm{cm^{-3}}$ ) and is usually detected toward low- to intermediate-mass protostellar envelopes \citep{tak00, sai01, fue05}."914. MAIS 7 is one of the Class 0 candidates identified by the 1.3 min continuum observations toward OAIC-3 (Chinietal.1997)., MMS 7 is one of the Class 0 candidates identified by the 1.3 mm continuum observations toward OMC-3 \citep{chi97}.915.. This object is also identified as CSO12 by the 350 jam continmimum observations (Lisοἱal.1905) and is associated with the IRAS point source 05329-0508., This object is also identified as CSO12 by the 350 $\mu$ m continuum observations \citep{lis98} and is associated with the IRAS point source 05329-0508.916 The dust mass derived from the 1.3 mam continuum emission andthe bolometric luminosity are 8 M. and το L.. respectively (Chinietal.1997).," The dust mass derived from the 1.3 mm continuum emission andthe bolometric luminosity are 8 $_{\odot}$ and 76 $_{\odot}$, respectively \citep{chi97}."917.. This bolometric Iuminosity corresponds to ~ 3 M. or à AO star at the ZAMS., This bolometric luminosity corresponds to $\sim$ 3 $_{\odot}$ or a A0 star at the ZAMS.918 Dense molecular gas around MAIS 7 is also detected by 0Ο (10) observations (Aso.et.al.2000)., Dense molecular gas around MMS 7 is also detected by $^{13}$ $^{+}$ (1–0) observations \citep{aso00}.919. MMS 7 is alsoassociated with a bright rellection nebula. ILaro-5a/6a Claro1953).. and 2\LASS ancl mid-infrared sources are located at the root of the eastern reflection nebula (Nielbocketal. 2003)...," MMS 7 is alsoassociated with a bright reflection nebula, Haro-5a/6a \citep{har53}, and 2MASS and mid-infrared sources are located at the root of the eastern reflection nebula \citep{nie03}. ."920 At MAIS T. a giant molecular outllow (0.86 pe: Aso et al.," At MMS 7, a giant molecular outflow (0.86 pc; Aso et al."921 2000) is also observed along the east-west direction. and 2.12 jan Hy r=1-0 $(1) knots are detected up to 1.45 pe to the western side (Yuetal.1997:Stanke2002).," 2000) is also observed along the east-west direction, and 2.12 $\mu$ m $_2 ~v$ =1-0 $S$ (1) knots are detected up to 1.45 pc to the western side \citep{yu97, sta02}."922. In addition. there is a 3.6 em continuum source elongated along the large-scale outflow. which (races a free-lree jet from the protostar.," In addition, there is a 3.6 cm continuum source elongated along the large-scale outflow, which traces a free-free jet from the protostar."923 The relative isolation of MMS 7 along with the above properties makes (his source one of the most appropriate objects to investigate detailed spatial ancl velocity structures of a dense core around an intermediate-2mnass protostar., The relative isolation of MMS 7 along with the above properties makes this source one of the most appropriate objects to investigate detailed spatial and velocity structures of a dense core around an intermediate-mass protostar.924 We have observed HPCO — (J—10: 86.754 GIIz) and CO (J—10: 115.271 GIIz) lines in MAIS 7 with the six-elements NALA from 2004 November to 2005 March., We have observed $^{13}$ $^{+}$ $J$ =1–0; 86.754 GHz) and CO $J$ =1–0; 115.271 GHz) lines in MMS 7 with the six-elements NMA from 2004 November to 2005 March.925 Detailed NALA observational parameters are summarized in Table 1.., Detailed NMA observational parameters are summarized in Table \ref{obs}.926" Since the minimum projected baseline lengths of the IMCO (1:0) and CO(10) observations were 2.9 and 6.3 kA. our observations were insensitive to structuresmore extended than 57"" (0.13 pe) and 32"" (0.07 pe) at the 10 level. respectively (WilnerandWelch 1994).."," Since the minimum projected baseline lengths of the $^{13}$ $^{+}$ (1–0) and CO(1–0) observations were 2.9 and 6.3 $\lambda$ , our observations were insensitive to structuresmore extended than $''$ (0.13 pc) and $''$ (0.07 pc) at the 10 level, respectively \citep{wil94}. ."927 The overall flux uncertainty was estimated, The overall flux uncertainty was estimated928im detail the iuter-cluster overdensity of the SSC (Drinkwateroetal.1999:Dardell2000:Drinkwateretal. 2001).,"in detail the inter-cluster overdensity of the SSC \citep{dri99,bar00,dri04}."929 Of these. ouly Drinkwateretal.(2001) considers a sinilaur area ou the sky. so we draw a comparison with this study.," Of these, only \citet{dri04} considers a similar area on the sky, so we draw a comparison with this study."930 The SSC intercluster overdensity is 3.3 c 0.1 over 151 deg?. as compared with au overdensity of 2.1 that we fud for the TRS.," The SSC inter-cluster overdensity is 3.3 $\pm$ 0.1 over 151 $^2$, as compared with an overdensity of 2.4 that we find for the HRS."931 A radius of (( 2 Mpe at SSC inean redshift) was excised around ACO clusters within the survey area., A radius of $\sim$ 2 Mpc at SSC mean redshift) was excised around ACO clusters within the survey area.932 This radius iu the SSC corresponds to aat the TRS redshift., This radius in the SSC corresponds to at the HRS redshift.933" C'onsequeuth the TRS siuuple is slightly more restrictive in selecting ouly inter-cluster galaxies. but this is likely a siall differcuce,"," Consequently, the HRS sample is slightly more restrictive in selecting only inter-cluster galaxies, but this is likely a small difference."934 Overall. we find a somewhat smaller. but similar. overdensity in the TRS compared to that iu the SSC.," Overall, we find a somewhat smaller, but similar, overdensity in the HRS compared to that in the SSC."935 Tn addition. we compare the total mass in the TRS intercluster galaxies to that iu the SSC.," In addition, we compare the total mass in the HRS inter-cluster galaxies to that in the SSC."936 Given the differences between the TRS aud SSC studies in overdcusity (2.1/3.5). angular survey reeion (107/151 deg?) aud relative distance (~ 20.000/15.000 +)). we couclude that the total masses of the iuter-cluster regions of the ITRS aud SSC are virtually ideutical.," Given the differences between the HRS and SSC studies in overdensity (2.4/3.3), angular survey region (107/151 $^2$ ), and relative distance $\sim$ 20,000/15,000 ), we conclude that the total masses of the inter-cluster regions of the HRS and SSC are virtually identical."937 Thus our data indeed support the couclusion of previous studies of the distribution of galaxy clusters (Zuccaetal.1993:IIudsonetal.1999:Eimasto2001) that the SSC and IIRS constitute the two largest mass concentrations in the local universe.," Thus our data indeed support the conclusion of previous studies of the distribution of galaxy clusters \citep{zuc93,hud99,ein01} that the SSC and HRS constitute the two largest mass concentrations in the local universe."938 lu 833.2. we found an overall spatialaedshitt trend in the TRS. in that a systematic ducrease in redshift is present with increasing position along a NW axis.," In 3.2, we found an overall spatial-redshift trend in the HRS, in that a systematic increase in redshift is present with increasing position along a $-$ NW axis."939 Bardelli et al. (, Bardelli et al. (9402000) have fitted a plane in (a. ὃν ο) space to their iuter-cluster observations in the SSC.,"2000) have fitted a plane in $\alpha$, $\delta$ , $cz$ ) space to their inter-cluster observations in the SSC."941 Thev note ai «0060 lncrease m average galaxy velocity along the best Bt plane over the ((10. Ape) region. a result reminiscent of the oosition-redshiüft tilt in the IRS.However. when he area ou the sky is expanded (cf.Figure|ofBardellietal.2000) to include the iuter-cluster ealaxies in Drinkwaterotal.(1999).. the iain veal of the ealaxy distribution shifts bv 7 Mpc. ancl the entire distribution is broadened.," They note a $\sim$ 3000 increase in average galaxy velocity along the best fit plane over the (40 Mpc) region, a result reminiscent of the position-redshift tilt in the HRS.However, when the area on the sky is expanded \citep[cf., Figure 4 of][]{bar00} to include the inter-cluster galaxies in \citet{dri99}, the main peak of the galaxy distribution shifts by 7 Mpc, and the entire distribution is broadened."942 Du short. while there appears to be a kinematic eradieut in the SSC. it is not clear whether that feature exteuds over the cutive region of the supercluster.," In short, while there appears to be a kinematic gradient in the SSC, it is not clear whether that feature extends over the entire region of the supercluster."943 As is discussed in §33.3. when the redshitt trend along the aaxis in the URS is fittec and removed. the redshift distribution is biauodal (Fieure 93).," As is discussed in 3.3, when the spatial-redshift trend along the $-$ axis in the HRS is fitted and removed, the redshift distribution is bi-modal (Figure \ref{f9}) )."944 Iu fact. the bianodal signature is observed even in the original vedshift listoervai in Fieure L.," In fact, the bi-modal signature is observed even in the original redshift histogram in Figure \ref{f4}."945 Redshitt binnodality is also strikingly evideut iu the SSC (ch.Figure6ofDrinkwateretal.2001:Quintan," Redshift bi-modality is also strikingly evident in the SSC \citep[cf., Figure 6 of][Figure946 5]{dri04,qui00}."947a There is a lower redshift coniponent to the SSC (at ~&8.000 12.000 ) that is quite distinct from the higher coniponeut at -—1 1.000 18.000j.," There is a lower redshift component to the SSC (at $\sim$ 8,000 $-$ 12,000 ) that is quite distinct from the higher component at $\sim$ 14,000 $-$ 18,000."948. While it was originally though hat the SSC) redshift coniponents are substantially different in size. an extensive follow-up study by Drinkwaterctal.(2001) reveals thetnter-chister populations of the two coniponenuts to be roughly equal (cf.FigurcL3ofDrinkwateretal. 2001).," While it was originally thought that the SSC redshift components are substantially different in size, an extensive follow-up study by \citet{dri04} reveals the populations of the two components to be roughly equal \citep[cf., Figure 5949of][]{dri04}."950. On the other haud. the distribution of thechesters within the SSC is also biauodal and more heavily weighted to the higher redshift comipoueut at 13.000 15.000," On the other hand, the distribution of the within the SSC is also bi-modal and more heavily weighted to the higher redshift component at 13,000 $-$ 18,000."951" Specifically 16 clusters have redshifts above 13.500Ἐν, while oulv 6 lave redshifts below 12.500|."," Specifically, 16 clusters have redshifts above 13,500, while only 6 have redshifts below 12,500."952" The higher redshift componcut coincides with what is designated by BRoeisenegeeretal.(2000) as the collapsing ""Central Region.! centered. ou. the cluster A3nhs."," The higher redshift component coincides with what is designated by \citet{rei00} as the collapsing “Central Region,” centered on the cluster A3558."953 Ax a result. the higher redshift compoucut dominates when both the cluster aud iuter-cluster galaxies are considered.," As a result, the higher redshift component dominates when both the cluster and inter-cluster galaxies are considered."954 More reliable redshitt data for the clusters in the IRS is probably required before a defuitive statement can be made about their redshift distribution. but the available data (cf..," More reliable redshift data for the clusters in the HRS is probably required before a definitive statement can be made about their redshift distribution, but the available data (cf.,"955 Figure 9.. Right panel) iudicate that such a 3:1 Hubalance m cluster uuubers between lower aud hieher redshift components is not present.," Figure \ref{f9}, Right panel) indicate that such a 3:1 imbalance in cluster numbers between lower and higher redshift components is not present."956 Although our observatious reveal a distinct arrangement of field galaxies marking the IIRS. the extent and/or bouudaries of the supercluster are uot casily deteriuued.," Although our observations reveal a distinct arrangement of field galaxies marking the HRS, the extent and/or boundaries of the supercluster are not easily determined."957 Iu fact. percolation and fricuds-oftriends algorithms iuclude other clusters in the ITRS besides those listed im Table 3 (INalinkovetal.1998:Einasto2002).," In fact, percolation and friends-of-friends algorithms include other clusters in the HRS besides those listed in Table \ref{tb3} \citep{kal98,ein02}."958. Recent studies ofthe SSC cover a similar area on the skv and also leave some ambiguityas to the spatial aud. kinematic extent of the supercluster (Quintanaetal.2000:Drinkwater 2001).," Recent studies of the SSC cover a similar area on the sky and also leave some ambiguityas to the spatial and kinematic extent of the supercluster \citep{qui00,dri04}."959. It is quite possible that the boundaries of the IRS extend bevond the region surveyed by us with, It is quite possible that the boundaries of the HRS extend beyond the region surveyed by us with960"Figure 10 shows a colour-magnitude diagram of all WFPC2 sources around NGC 3309/3311, with the confirmed cluster GCs/UCDs highlighted.","Figure \ref{fig:hstcmd} shows a colour-magnitude diagram of all WFPC2 sources around NGC 3309/3311, with the confirmed cluster GCs/UCDs highlighted."961" In Fig. 11,,"," In Fig. \ref{fig:hstucds},"962" we compare the half-light radii and luminosities of the Hydrall GCs/UCDs to globular clusters from the ACS Virgo Cluster Survey (ACSVCS,?),, Milky Way, LMC/SMC and Fornax star clusters from ?,, UCDs from ?,, and the compact object M59cO (?).."," we compare the half-light radii and luminosities of the I GCs/UCDs to globular clusters from the ACS Virgo Cluster Survey \citep[ACSVCS,][]{2009ApJS..180...54J}, Milky Way, LMC/SMC and Fornax star clusters from \citet{2005ApJS..161..304M}, UCDs from \citet{2008A&A...487..921M}, and the compact object M59cO \citep{2008MNRAS.385L..83C}."963" The apparent gAp-band magnitudes of the ACSVCS GCs were transformed into absolute V-band magnitudes using the relation V=gap+0.026—0.307-(gz)4p given in ?,, and a Virgo distance modulus of 31.09 mag (?).."," The apparent $g_{\mathrm{AB}}$ -band magnitudes of the ACSVCS GCs were transformed into absolute $V$ -band magnitudes using the relation $V = g_{\mathrm{AB}} + 0.026-0.307 \cdot (g-z)_{\mathrm{AB}}$ given in \citet{2006ApJ...639..838P}, and a Virgo distance modulus of $31.09$ mag \citep{2007ApJ...655..144M}."964" 7, is the average of the half-light radii measured in the g- and in the z-band.", $r_{\mathrm{h}}$ is the average of the half-light radii measured in the $g$ - and in the $z$ -band.965 My and rj of the star clusters from ? are the King models values., $M_V$ and $r_{\mathrm{h}}$ of the star clusters from \citet{2005ApJS..161..304M} are the King models values.966 The tabulated masses of the UCDs from ? were converted into My with the given M/Ly ratios and a solar absolute magnitude of Myo=4.83 mag (?)..," The tabulated masses of the UCDs from \citet{2008A&A...487..921M} were converted into $M_V$ with the given $M/L_V$ ratios and a solar absolute magnitude of $M_{V,\sun} = 4.83$ mag \citep{1998gaas.book.....B}."967" For M59cO, My was calculated from Mg using B—V=0.96 mag (?).."," For M59cO, $M_V$ was calculated from $M_B$ using $B-V=0.96$ mag \citep{1995PASP..107..945F}."968" The sizes, luminosities and colours of the Hydrall GC/UCDs are fully consistent with the ones of Virgo and Fornax UCDs (e.g.??).. A"," The sizes, luminosities and colours of the I GC/UCDs are fully consistent with the ones of Virgo and Fornax UCDs \citep[e.g.][]{2008AJ....136..461E, 2008A&A...487..921M}."969" size-luminosity relation is visible for objects brighter than My~—10 mag, following the trend observed in other galaxy clusters."," A size-luminosity relation is visible for objects brighter than $M_V \sim -10$ mag, following the trend observed in other galaxy clusters."970" Below this magnitude, we do not regard the size measurements with reliable, since the S/N is smaller than 20 for those objects (cf."," Below this magnitude, we do not regard the size measurements with reliable, since the $S/N$ is smaller than 20 for those objects (cf."971 Table 2))., Table \ref{tab:hstucds}) ).972" With My=—13.37 mag, HUCDI is the brightest object in our sample."," With $M_V=-13.37$ mag, HUCD1 is the brightest object in our sample."973" Its luminosity corresponds to a mass of about 5x107 Mo (assuming M/L= 3), or even ~10? Mo, if assuming a mass-to-light ratio larger than 5, which has been measured for several of the brightest UCDs in Virgo and Fornax (?).."," Its luminosity corresponds to a mass of about $5\times 10^7$ $_{\sun}$ (assuming $M/L=3$ ), or even $\sim 10^8$ $_{\sun}$, if assuming a mass-to-light ratio larger than 5, which has been measured for several of the brightest UCDs in Virgo and Fornax \citep{2008A&A...487..921M}."974" Hence, HUCDI has a luminosity/mass comparable to the most massive UCDs in Virgo and Fornax (VUCD7 and UCD3), but with its half-light radius of rp=25.4 pc, it is the most compact object among the highest luminosity UCDs."," Hence, HUCD1 has a luminosity/mass comparable to the most massive UCDs in Virgo and Fornax (VUCD7 and UCD3), but with its half-light radius of $r_{\mathrm{h}}=25.4$ pc, it is the most compact object among the highest luminosity UCDs."975" VUCD7 and UCD3 are known to feature a two-component light profile with an extended faint envelope and a smaller core component (?),, which is also indicated in Fig. 11.."," VUCD7 and UCD3 are known to feature a two-component light profile with an extended faint envelope and a smaller core component \citep{2007AJ....133.1722E}, which is also indicated in Fig. \ref{fig:hstucds}."976" A weak indication for a faint halo is observed for HUCDI, but not for the other two luminous UCDs."," A weak indication for a faint halo is observed for HUCD1, but not for the other two luminous UCDs."977" This can be seen in Fig. 12,,"," This can be seen in Fig. \ref{fig:hstimages},"978" where from from left to right are shown the three brightest UCDs in our sample, and from top to bottom the according PSF model, the object and the residual (object minus PSF model) image."," where from from left to right are shown the three brightest UCDs in our sample, and from top to bottom the according PSF model, the object and the residual (object minus PSF model) image."979 The residuals for HUCDI are of the order of a few percent of the value of the corresponding science image pixel., The residuals for HUCD1 are of the order of a few percent of the value of the corresponding science image pixel.980" However, King profiles are known not to represent well the outer regions of GCs (e.g.?),, therefore we cannot affirmatively conclude that the observed residual halo of HUCDI is due to a presence of a second component as for VUCD7 and UCD3."," However, King profiles are known not to represent well the outer regions of GCs \citep[e.g.][]{2005ApJS..161..304M}, therefore we cannot affirmatively conclude that the observed residual halo of HUCD1 is due to a presence of a second component as for VUCD7 and UCD3."981" The majority of the confirmed cluster GCs/UCDs is located in the immediate vicinity of NGC 3311, i.e. within a projected distance ofR=5’, or ~70 kpc (see Fig. 7))."," The majority of the confirmed cluster GCs/UCDs is located in the immediate vicinity of NGC 3311, i.e. within a projected distance of $R=5\arcmin$, or $\sim 70$ kpc (see Fig. \ref{fig:map}) )."982 In Fig., In Fig.983 13 we plot the objects apparent magnitudes and radial velocities versus the projected distance from NGC 3311., \ref{fig:radialdist} we plot the objects apparent magnitudes and radial velocities versus the projected distance from NGC 3311.984 There appears to be a trend that brighter objects are in projection located further away from the central galaxy., There appears to be a trend that brighter objects are in projection located further away from the central galaxy.985" This is, however, caused by selection effects,"," This is, however, caused by selection effects,"986matter and radiation epochs there is scaling. ancl these functions may be wrltten as: where .XY represents any pair of superscripts considered above. and why and a=fay’.,"matter and radiation epochs there is scaling, and these functions may be written as: where $XY$ represents any pair of superscripts considered above, and $x=k\eta$ and $x'=k\eta'$."987 Phe purpose of this work is the measurement of the 14 functions of two variables 7NY(rur(eh).," The purpose of this work is the measurement of the 14 functions of two variables $F^{XY}(x,x')$."988 3., 3.989 NUAIERICAL DETERAINATION OF THE CORRELATORS Local strings have an extra complication over global defects. which μαems from the fact that we are unable to simulate the underline field rcory.," NUMERICAL DETERMINATION OF THE CORRELATORS Local strings have an extra complication over global defects, which stems from the fact that we are unable to simulate the underlying field theory."990 Instead. we approximate the true dynamics with line-like relativistic using the Nambu equations of motion.," Instead, we approximate the true dynamics with line-like relativistic strings using the Nambu equations of motion."991 We used a developed stringsimplementation 5. of this algorithm which simulatespreviously the network of neglecting damping7] elfects due to the expansion of the Universe.," We used a previously developed implementation \cite{col,smi} of this algorithm which simulates the network of strings neglecting damping effects due to the expansion of the Universe."992 The stringsvelocities of the strings scements in the network can be constrained to be integers due to the gauge conditions used to cliseretise the equations of motion and this enables the code to obtain very high accuracy with fast computation., The velocities of the strings segments in the network can be constrained to be integers due to the gauge conditions used to discretise the equations of motion and this enables the code to obtain very high accuracy with extremely fast computation.993 ‘To extremelysimulate the extraction of energy [from the system due to the decay of the string into gravitaional radiation. loops of a minimum size are excised from the loopssimulation at cach timestep.," To simulate the extraction of energy from the system due to the decay of the string loops into gravitaional radiation, loops of a minimum size are excised from the simulation at each timestep."994 This also ensures that the network scales with respect to the conformal time 5 which enables us to extend the range of the resulting correlation functions bevond the limited range dynamicalcovered in the simulation., This also ensures that the network scales with respect to the conformal time $\eta$ which enables us to extend the dynamical range of the resulting correlation functions beyond the limited range covered in the simulation.995 We performed simulations with box sizes ranging from 128* to G00H with a cut-olf on the loop size of two links.," We performed simulations with box sizes ranging from $128^3$ to $600^3$, with a cut-off on the loop size of two links."996 Realisation averages were carried out with 256* boxes once it was determined that the general form of the correlators scaled accurately with box size-, Realisation averages were carried out with $256^3$ boxes once it was determined that the general form of the correlators scaled very accurately with box size.997To evaluate the VETCs from the simulations very.we selected times in the range Nd. where N is the box size. when we were sure that the string network was scaling. and when effects. are still excluded: causality.,"To evaluate the UETCs from the simulations we selected times in the range $0.1998N<t<N/4$ , where $N$ is the box size, when we were sure that the string network was scaling, and when boundary effects are still excluded by causality."999" From this we obtained the time boundary.evolution of the networks stress-energvby tensor O,, at cach point on the lattice using where JX(a. is the coordinate of the string with 0 a 9)the string."," From this we obtained the time evolution of the network's stress-energy tensor $\Theta_{\mu\nu}$ at each point on the lattice using where $X(\si,\eta)$ is the spacetime coordinate of the string with $\si$ a parameter along the string."1000 SVT spacetime was carried out on the 10 parameterindependent along of the Fourier decompositiontransform of the stress-energvy tensor at cach timestep., SVT decomposition was carried out on the 10 independent components of the Fourier transform of the stress-energy tensor at each timestep.1001 componentsThis resulted in all the SVT components being from the simulation without assumptions on enerey computedconservationdirectly and on the details of energy. makingdissipation from the network., This resulted in all the SVT components being computed directly from the simulation without making assumptions on energy conservation and on the details of energy dissipation from the string network.1002 The drawback of all the directly in stringsuch à manner is that the process obtainingbecomes computationally componentsintensive for even modestly sized. simulations 256%., The drawback of obtaining all the components directly in such a manner is that the process becomes computationally intensive for even modestly sized simulations e.g. $256^3$.1003 By the decomposed stress-cnergy components e.g.from a central time with eross-correlatingthose from all the stored timesteps the Lt independent ος £73(agg) were Standard computed.," By cross-correlating the decomposed stress-energy components from a central time with those from all the stored timesteps the 14 independent UETC's $f^{XY}(k,\eta,\eta')$ were computed."1004codes for structure formation the root of coherent correlators for sourcing the 10]. requirecillerential squareequations but this is not straightforward in defect. inhomogeneousscenarios because the correlators are incoherent and therefore generalco not factorise., Standard codes for structure formation \cite{cmbfast} require the square root of coherent correlators for sourcing the inhomogeneous differential equations but this is not straightforward in general defect scenarios because the correlators are incoherent and therefore do not factorise.1005 “Lo overcome, To overcome1006Fienre 2 shows the histograms of the fluxes at 350 jan for a stacking box with 0.5<2«0.6 and 0.5Sop«1 mJy.,Figure \ref{fig:HistogramsStackBox2populations} shows the histograms of the fluxes at 350 $\mu$ m for a stacking box with $0.5<z<0.6$ and $0.5<S_{24}<1$ mJy.1007" The iain source of error in the estimate of the fluxes for this case would not be the dispersion in either So, or 2 but the presence of two different. populations. which are imdistinguishable using observations at shorter waveleugths."," The main source of error in the estimate of the fluxes for this case would not be the dispersion in either $S_{24}$ or $z$ but the presence of two different populations, which are indistinguishable using observations at shorter wavelengths."1008 These two populations are the starburst aud he normal (cold) populations described in ?.., These two populations are the starburst and the normal (cold) populations described in \citet{2003MNRAS.338..555L}.1009" Figure 3 shows the nmnuuber of starburst and uormal sources as a function of[ol : for sources with SO<$9,270 wy. 27So,xld ind and So,2 Loaundy."," Figure \ref{fig:SBvsCOLD} shows the number of starburst and normal sources as a function of $z$ for sources with $80<S_{24}<270\,\mu$ Jy, $0.27<S_{24}<1$ mJy, and $S_{24}>1$ mJy."1010 For the τος afore mentioned cases. the cold sources are the dominant population for 5<0.8.20.6. and τς0.5 espeetivelv.," For the three afore mentioned cases, the cold sources are the dominant population for $z<0.8,\, z<0.6$, and $z<0.5$ respectively."1011 There are no effective wavs of separating hese two populations. and this will cause poor estimates of the ποσα colors of each population.," There are no effective ways of separating these two populations, and this will cause poor estimates of the mean colors of each population."1012 This is particukuly inportaunt when the ummber of sources of cach type is approximately equal., This is particularly important when the number of sources of each type is approximately equal.1013" This is recatise we adel ogether wo populations of very different 55539 (cold sources are in general brighter in the submillimeter than starburst sonrces at the same redshifts aud with similar 55,).", This is because we add together two populations of very different $S_{350}$ (cold sources are in general brighter in the submillimeter than starburst sources at the same redshifts and with similar $S_{24}$ ).1014 When one of the populations dominates. this problem becomes jcelieible.," When one of the populations dominates, this problem becomes negligible."1015 Because of the ack of constraints on SEDs at lone wavelengths and their evolution with redshift. the ? model does not ake into account that galaxies of the sale luminosity and redshift could have different values of S4 (apart roni the distinction between normal aux starburst sources).," Because of the lack of constraints on SEDs at long wavelengths and their evolution with redshift, the \citet{2004ApJS..154..112L} model does not take into account that galaxies of the same luminosity and redshift could have different values of $S_{\lambda}$ (apart from the distinction between normal and starburst sources)."1016 To assess the effect of this dispersioL we introduce a random Cassian error into the flix estimated. with the stacking for each of the stacker sources.," To assess the effect of this dispersion, we introduce a random Gaussian error into the flux estimated with the stacking for each of the stacked sources."1017 The errors that we 1iake using tliis x»ocedure are equivalent to those that we would male if we were to use a 1nodel with an intrinsic Gaussian dispersion me ithe Sy of he sources., The errors that we make using this procedure are equivalent to those that we would make if we were to use a model with an intrinsic Gaussian dispersion in the $S_{\lambda}$ of the sources.1018 This type of error does not affect the results ‘or the mean of the sources but the average difference )otween this mean and the fiuxes of the individual sources., This type of error does not affect the results for the mean of the sources but the average difference between this mean and the fluxes of the individual sources.1019 We test the effect on our results for dierent levels of dispersion (nicasured in tenus of the standard deviation in the dispersion compared to the mean flux of the sources)., We test the effect on our results for different levels of dispersion (measured in terms of the standard deviation in the dispersion compared to the mean flux of the sources).1020 In Fig. L.," In Fig. \ref{fig:Histograms-of-the},"1021" we can see the histograms of the errors or a dispersion of 0%. 10:4. 25% for all sources with So,c27044gy.Ax expected. the figure illustrates. how he histograms5o broaden with dispersion."," we can see the histograms of the errors for a dispersion of $0\%$, $10\%$, $25\%$ for all sources with $S_{24}>270\,\mu Jy$.As expected, the figure illustrates how the histograms broaden with dispersion."1022 For a standard deviation in the errors of the fluxes associated with the stacking os; and a standard deviation associated with the Ηχος σαν the final standard deviation in our errors CT; Would be στ=uu|OD," For a standard deviation in the errors of the fluxes associated with the stacking $\sigma_{St}$ and a standard deviation associated with the fluxes $\sigma_{Disp}$ ,the final standard deviation in our errors $\sigma_{Tot}$ would be $\sigma_{Tot}=\sqrt{\sigma_{St}^2+\sigma_{Disp}^2}$."1023 We do not analyze other statistical representations of this effect (16.. non-Gaussian intrinsic dispersion) since we do not have any strong observational constraints.," We do not analyze other statistical representations of this effect (i.e., non-Gaussian intrinsic dispersion) since we do not have any strong observational constraints."1024 The effect of redshift errors are difficult to evaluate., The effect of redshift errors are difficult to evaluate.1025 This is because they combine with the uou-linear k-correction. making the variation in S4 with : complex.," This is because they combine with the non-linear k-correction, making the variation in $S_{\lambda}$ with $z$ complex."1026 In Sect. LL. ," In Sect. \ref{sec:Test}, ,"1027we study the effect of redshift errors for two different relative errors ==:=3 and === 10%., we study the effect of redshift errors for two different relative errors $\frac{\triangle z}{z}=3\%$ and $\frac{\triangle z}{z}=10\%$ .1028The last few vears have witnessed. increased: cllorts in studying the nature anc elfects of the cdillercnt energy injection mechanisms into the interstellar medium. (15M) of galaxies.,The last few years have witnessed increased efforts in studying the nature and effects of the different energy injection mechanisms into the interstellar medium (ISM) of galaxies.1029 Understanding the nature of the turbulence drivers. and their. associated. energy injection rate has important consequences in improving our understanding of the dissipation of turbulence in the SAL and. its impact on star formation., Understanding the nature of the turbulence drivers and their associated energy injection rate has important consequences in improving our understanding of the dissipation of turbulence in the ISM and its impact on star formation.1030 A few recent studies pointed out that kinetic enerev might be injected on large scale (kpc scales) into the SAL of clwarl irregular. galaxies., A few recent studies pointed out that kinetic energy might be injected on large scale (kpc scales) into the ISM of dwarf irregular galaxies.1031 Stanimirovic Lazarian (2001) analyzed the kinetic energy power spectrum in the Small Magellanic. Cloud and did. not. observe. any indication of energy injection up to the largest: considered scale (~ 4) κρο, Stanimirovic Lazarian (2001) analyzed the kinetic energy power spectrum in the Small Magellanic Cloud and did not observe any indication of energy injection up to the largest considered scale $\sim 4$ ) kpc.1032 Dib Burkert (2004.2005) analyzecl the LL gas morphology in Holmberg HE (Ho LL) and a series of numerical simulations of driven turbulence of the laree scale ISAT that include. cooling. heating and the eas self-gravity.," Dib Burkert (2004,2005) analyzed the HI gas morphology in Holmberg II (Ho II) and a series of numerical simulations of driven turbulence of the large scale ISM that include, cooling, heating and the gas self-gravity."1033 They performed. a comparison between the observations of Ho IE and the simulated Lb maps by measuring the autocorrelation length on cillerent scales and concluded that turbulence is injected. into the ISM of the dwarf irregular. low star forming galaxy Ilo Η. on a scale of ~6 kpe.," They performed a comparison between the observations of Ho II and the simulated HI maps by measuring the autocorrelation length on different scales and concluded that turbulence is injected into the ISM of the dwarf irregular, low star forming galaxy Ho II, on a scale of $\sim 6$ kpc."1034 Similar conclusions on the existence of a large scale driving mechanism in the dwarf irregular. galaxy DDO 210 have been reached by Begum et al. (, Similar conclusions on the existence of a large scale driving mechanism in the dwarf irregular galaxy DDO 210 have been reached by Begum et al. (10352006).,2006).1036 Brunt et al. (, Brunt et al. (10372009) recently showed that the size-velocity. clispersion relations of nearby Galactic molecular clouds (MOS). can only be replicated by models in which turbulence is driven on large scales (Le. larger the sizes of the clouds themselves).,"2009) recently showed that the size-velocity dispersion relations of nearby Galactic molecular clouds (MCs) can only be replicated by models in which turbulence is driven on large scales (i.e., larger the sizes of the clouds themselves)."1038 Ixoda et al. (, Koda et al. (10392006) measured the orientations of AICs in. the Galactic Ring region and concluded that the absence of a,2006) measured the orientations of MCs in the Galactic Ring region and concluded that the absence of a1040therefore. the enerex levels quautize aud take the values The time-dependent eigenfuuctious take the form It is important to note that Ais function Αι) has an oscillatory behavior for e«0 (a«IMD ) whiles for o>0 (am Ep decreases monotonically aud is an expouenutiallvy damped function for large « (Fig. 2)).,"therefore, the energy levels quantize and take the values The time-dependent eigenfunctions take the form It is important to note that Airy's function $ \mbox{Ai}(x)$ has an oscillatory behavior for $x<0$ $a<\frac{\displaystyle1041B}{\displaystyle |E_n|}$ ) whiles for $x>0$ $a>\frac{\displaystyle1042B}{\displaystyle |E_n|}$ ) decreases monotonically and is an exponentially damped function for large $x$ (Fig. \ref{fig2}) )."1043 Therefore. the solutions (72)) show a classical behavior for small α aud a quautum behavior for large e.," Therefore, the solutions \ref{radiation-final-1}) ) show a classical behavior for small $a$ and a quantum behavior for large $a$."1044 This is contrary to usually expected results for previous case., This is contrary to usually expected results for previous case.1045 In fact detecting quautum eravitational effects im large Universes is noticeable which has been also observed in FRW. Stephani. and Naluza-Wlein models 10.35.0]..," In fact detecting quantum gravitational effects in large Universes is noticeable which has been also observed in FRW, Stephani, and Kaluza-Klein models \cite{lemos1999,pedramCQG2,Coliteste}."1046 Iu f=1 and w=1/3 (radiation) case. the WD equation (30)) reduces to The above equation can be written as by taking «=αος we have This equation is identical to the time-ndepeudent Schroddinger equation for a simple harmonic oscillator with unit mass ancl energy A whore 2A=(E?|218) aud w=11L.," In $k=1$ and $w=1/3$ (radiation) case, the WD equation \ref{sle2}) ) reduces to The above equation can be written as by taking $x=a- \frac{\displaystyle E}{\displaystyle 12}$ we have This equation is identical to the time-independent Schröddinger equation for a simple harmonic oscillator with unit mass and energy $\lambda$ where $2\lambda = (E^2+24B)$ and $\omega^2=144$."1047" Therefore. the allowed values of A are wt|1/2) aud the possible values of £ are therefore. the stationary solutions are where ZZ, ave Uermite polvnomials."," Therefore, the allowed values of $\lambda$ are $\omega(n+1/2)$ and the possible values of $E$ are therefore, the stationary solutions are where $H_n$ are Hermite polynomials."1048 Towever. neither of the boundary couditious (28)) can be satisfied by wave functions.," However, neither of the boundary conditions \ref{boundary}) ) can be satisfied by these wave functions."1049 theseNow. we present some analytical. solutions. for. the late time Universe.," Now, we present some analytical solutions for the late time Universe."1050 For flat space time (h= 0). dust epoch («= 0). and standard Chaplveiu gas (a= 1). equation CLI)) reduces to," For flat space time $k=0$ ), dust epoch $w=0$ ), and standard Chaplygin gas $\alpha=1$ ), equation \ref{sle2-b}) ) reduces to"1051components on December 6 is consistent with the components being ejected staggered in lime: C1 first ancl C3 last with the spectral steeping occurring Irom radiation losses.,components on December 6 is consistent with the components being ejected staggered in time: C1 first and C3 last with the spectral steeping occurring from radiation losses.1052 II the spectral ageing mechanism is similar for the (μου components then in any nme snapshot Cl will be steeper than C2 and C2 steeper than C3., If the spectral ageing mechanism is similar for the three components then in any time snapshot C1 will be steeper than C2 and C2 steeper than C3.1053 A simple reverse time evolution of the models found in section 5. indicate substantial spectral ageing must have occurred due to svnchrotron losses in a strong internal magnetic fiekl aud inverse Compton losses in a strong ambient X-ray field.," A simple reverse time evolution of the models found in section 5, indicate substantial spectral ageing must have occurred due to synchrotron losses in a strong internal magnetic field and inverse Compton losses in a strong ambient X-ray field."1054 This is a very. plausible explanation of the unusually steep spectral indices of the components C1 and C2 from December 6 to December 14., This is a very plausible explanation of the unusually steep spectral indices of the components C1 and C2 from December 6 to December 14.1055 For example. consider (he existence of a strong X-ray flare curing the Gime that component C1 was launched (hat decaved on a lime scale of hours.," For example, consider the existence of a strong X-ray flare during the time that component C1 was launched that decayed on a time scale of hours."1056 Inverse Compton losses in the first hour alter ejection would make the leptons in Cl cool more than the leptons in C2 and those in C3 might not be cooled very much at all (because (he X-ray. flare was weak when C3 was launched)., Inverse Compton losses in the first hour after ejection would make the leptons in C1 cool more than the leptons in C2 and those in C3 might not be cooled very much at all (because the X-ray flare was weak when C3 was launched).1057 So in principle. for a lime variable source. not only is the elapsed time for spectral ageing important. bul also the strength of the ambient X-ray background al the tme al which the plasma was ejectecl.," So in principle, for a time variable source, not only is the elapsed time for spectral ageing important, but also the strength of the ambient X-ray background at the time at which the plasma was ejected."1058 Similarly. one might have a scenario in which a magnetic flare in the inner accretion [low launches the jet.," Similarly, one might have a scenario in which a magnetic flare in the inner accretion flow launches the jet."1059 The first ejection removes a large fraction of the magnetic flux., The first ejection removes a large fraction of the magnetic flux.1060 Most of ihe remainder is removed bv the ejection of C2 and C3 has a relatively weak magnetic field., Most of the remainder is removed by the ejection of C2 and C3 has a relatively weak magnetic field.1061 This scenario would also create different spectral ageing rates in the spectral components., This scenario would also create different spectral ageing rates in the spectral components.1062 The leptons would cool primarily by svuchrotron self-Compton emission in the early stages and svnchrotron emission later on., The leptons would cool primarily by synchrotron self-Compton emission in the early stages and synchrotron emission later on.1063 The most rapid rate of ageing would also be in Cl., The most rapid rate of ageing would also be in C1.1064 As the plasmoids expand. the rate of spectral ageing would decrease since (he magnetic field is gelling weaker.," As the plasmoids expand, the rate of spectral ageing would decrease since the magnetic field is getting weaker."1065 The constant à. assumption noted above is tantamount to the following scenario., The constant $\alpha$ assumption noted above is tantamount to the following scenario.1066 The plasmoids are ejected aud cooled in different ambient X-ray backgrounds. or magnetic environments.," The plasmoids are ejected and cooled in different ambient X-ray backgrounds, or magnetic environments."1067 In particular. (he X-ray background or magnetic fiekl was strong when Cl was launched and cecaved during the time interval that subsequent ejecta were emitted.," In particular, the X-ray background or magnetic field was strong when C1 was launched and decayed during the time interval that subsequent ejecta were emitted."1068 Thus. one dax alter being ejected. (he leptonic energv spectrum of Cl is steeper than the leptonic energy spectrum of C2 would be one dav after it was ejected.," Thus, one day after being ejected, the leptonic energy spectrum of C1 is steeper than the leptonic energy spectrum of C2 would be one day after it was ejected."1069 By December 6. the spectral ageing rate is comparatively very slow (a much weaker source of cooling ancl most of the hieh energy particles have already been cooled).," By December 6, the spectral ageing rate is comparatively very slow (a much weaker source of cooling and most of the high energy particles have already been cooled)."1070" This conjecture"" is consistent wilh the models found in section 5."," This ""conjecture"" is consistent with the models found in section 5."1071 After December 6. the models are shown in section 5 to have higher svnehrotron cooling rates (han inverse Compton cooling rates.," After December 6, the models are shown in section 5 to have higher synchrotron cooling rates than inverse Compton cooling rates."1072 Thus. after December 6. one would expect that svnchrotron cooling is the dominant driver of spectral ageing.," Thus, after December 6, one would expect that synchrotron cooling is the dominant driver of spectral ageing."1073 The total energy emitted in the svnchrotron spectrum between December 6 and December 14 is z4x10? eres and the total energv stored in (he leptons comprising the ejectecl plasmoids in the models οἱ section 5 is about zz6x101 eres., The total energy emitted in the synchrotron spectrum between December 6 and December 14 is $\approx 4 \times 10^{39}$ ergs and the total energy stored in the leptons comprising the ejected plasmoids in the models of section 5 is about $\approx 6 \times 10^{42}$ ergs.1074 This indicates that the radiation losses are a negligible contributor to the enerey budget of the plasmoids after December 6., This indicates that the radiation losses are a negligible contributor to the energy budget of the plasmoids after December 6.1075 Thus. spectral ageing," Thus, spectral ageing"1076The following section presents the basic analysis of our spectroscopic data. including radial velocity aud rotational velocity meastuements aud line streneth determiuations.,"The following section presents the basic analysis of our spectroscopic data, including radial velocity and rotational velocity measurements and line strength determinations."1077 Since we have uuultiple observatious of mauy stars. we lave used our data to search for velocity variations. aud Section 3) ideutifies several previously-uurecoguised binary systems.," Since we have multiple observations of many stars, we have used our data to search for velocity variations, and Section 3 identifies several previously-unrecognised binary systems."1078 Chromospleric activity is discussed in Section L where we investigate the Ha properties of the sample. describe au relation calibrated with open cluster observations. and use that relation to probe the recent star formation liistory of the Galactic disk.," Chromospheric activity is discussed in Section 4, where we investigate the $\alpha$ properties of the sample, describe an age-activity relation calibrated with open cluster observations, and use that relation to probe the recent star formation history of the Galactic disk."1079 Section 2 suumumarises our mali conclusions., Section 5 summarises our main conclusions.1080 Our primary sample consists of the volume-complete sample of 199 sinele M. dwarls aud M dwarf primaries defined in Paper L. We refer to this as the VC sample., Our primary sample consists of the volume-complete sample of 499 single M dwarfs and M dwarf primaries defined in Paper I. We refer to this as the VC sample.1081 The stars lave absolute uagnuitudes iu the rauge 8xAM€ 16. declilations north of -30° and distances within completeness units rauging from 22 parsecs at My—8.5dE0.5 to 2 parsecs at My=125.50.5.," The stars have absolute magnitudes in the range $8 \le M_V \le 16$ , declinations north of $^o$ and distances within completeness limits ranging from 22 parsecs at $_V = 8.5\pm0.5$ to 5 parsecs at $_V = 15.5\pm0.5$."1082 The latter values were derived. based ou photometric aud trigonometric data available in late 1995. coupled with he (My. TiO») relation derived in paper I. Since the1. Hipparcos-based trigonometric parallaxes Perryinanetal.(1997).. accurate to ~1 milliarcsecod. have become available for approximately wo-thirds of those systems.," The latter values were derived based on photometric and trigonometric data available in late 1995, coupled with the $_V$, TiO5) relation derived in paper I. Since then, Hipparcos-based trigonometric parallaxes \citet{hipparcos}, accurate to $\sim1$ milliarcsecond, have become available for approximately two-thirds of those systems."1083 As will be diseμυWd.sed iu Paper IV of this series (Reidal. in preparatiou). he addition of the new astrometric data allects the iclusionomission of only 155€ of the stars in the Paper I sample.," As will be discussed in Paper IV of this series (Reid, in preparation), the addition of the new astrometric data affects the inclusion/omission of only $\sim15\%$ of the stars in the Paper I sample."1084 While the new cistauce dete10inations should be taken into account in analysis of the kinematics of the local stars (Paper IV they are of little importance for the analysis oL chromospheric activity and age that is uudertakeu 1 this paper., While the new distance determinations should be taken into account in analysis of the kinematics of the local stars (Paper IV) they are of little importance for the analysis of chromospheric activity and age that is undertaken in this paper.1085 Thus. we retain the VC sample as our reference here.," Thus, we retain the VC sample as our reference here."1086 Our data were obtained with the echelle spectrograph (McCarthy1985) on the Palomar 60-inch elescope. which has a 2-pixel resolution of 19.000.," Our data were obtained with the echelle spectrograph \citep{m85}1087 on the Palomar 60-inch telescope, which has a 2-pixel resolution of 19,000."1088 Observatious between May 199f aud February 1995 used the original quartz 607 cross-dispersing prisius. giviug waveleugth coverage [rom 3700A to 9500A..," Observations between May 1994 and February 1995 used the original quartz $60\arcdeg$ cross-dispersing prisms, giving wavelength coverage from $3700 $ to $9500 $."1089 Beyond. 7000. there are gaps in waveleugtli coverage where the orders extend olf the CCD., Beyond $7000 $ there are gaps in wavelength coverage where the orders extend off the CCD.1090 Because our targets are all red. stars. very few counts were obtained in the blue (A< 1800À)) except for the very wieghtest stars.," Because our targets are all red stars, very few counts were obtained in the blue $\lambda < 4800$ ) except for the very brightest stars."1091 The resu was that the useful part of the spectrum was squeezed into the lower quarter of the CCD. leading to partial overlap between the reddest. adjacent orders.," The result was that the useful part of the spectrum was squeezed into the lower quarter of the CCD, leading to partial overlap between the reddest adjacent orders."1092 To solve this oxoblem. begiuning in June 1995. a uew set of cross-dispersiug prisins (with SE3 glass and a 127 apex) was used.," To solve this problem, beginning in June 1995, a new set of cross-dispersing prisms (with SF3 glass and a $42\arcdeg$ apex) was used."1093 This setup gives waveleugtli coverage [rom1800 to 0500A. aad has complete order separation. although the gaps in wavelength coverage remain.," This setup gives wavelength coverage from$4800$ to $9500 $ and has complete order separation, although the gaps in wavelength coverage remain."1094 For both, For both1095such objects among brown dwarls 2 out of 6. Le. 32223.,"such objects among brown dwarfs – 2 out of 6, i.e. $33\pm 23$."1096 IE we consider all objects. there are three stars (64:35) with ratios >5%.. and one brown dwarf. i.e. 11d:11% (IC348-613 [rom Ixleinetal. (2003))).," If we consider all objects, there are three stars $6 \pm 3$ ) with ratios $>5$, and one brown dwarf, i.e. $11\pm 11$ (IC348-613 from \citet{kap03}) )."1097 Here we neglect upper limits close to5%.. because these objects are unlikely to have disk masses >5%.," Here we neglect upper limits close to, because these objects are unlikely to have disk masses $>5$."1098. Taking into account the large uncertainties in disk and object masses. (here is no statistical basis (ο claim that brown cdwarls lack (relative) massive disks.," Taking into account the large uncertainties in disk and object masses, there is no statistical basis to claim that brown dwarfs lack (relative) massive disks."1099 There is still the possibility. (hat such an effect is hidden in the non-detections., There is still the possibility that such an effect is hidden in the non-detections.1100 The average disk to object mass ratio is 1.9% in the stellar regime. which is an upper limit. because it does not take into account non-cletections.," The average disk to object mass ratio is $1.9$ in the stellar regime, which is an upper limit, because it does not take into account non-detections."1101 Assuming no trend with mass. we expect half of the brown chwarls to have values below1.," Assuming no trend with mass, we expect half of the brown dwarfs to have values below."1102954.. For nine of our brown dwarls we can rule owl that thev have ratios >1.956.. and Dor two of them we know that the ratio is >1.9%.," For nine of our brown dwarfs we can rule out that they have ratios $>1.9$, and for two of them we know that the ratio is $>1.9$."1103.. Thus. if most of the upper limits with ratios >1.9% are in fact values lower than1.9%.. there would be reason to believe that brown clwarls have more often very little circumstellar malerial (han stars. which would in (urn be indirect evidence for truncated disks.," Thus, if most of the upper limits with ratios $>1.9$ are in fact values lower than, there would be reason to believe that brown dwarfs have more often very little circumstellar material than stars, which would in turn be indirect evidence for truncated disks."1104 Dased on the available observational data. this possibility cannot be definitely excluded.," Based on the available observational data, this possibility cannot be definitely excluded."1105 The next eeneration of submmn telescopes will hopefully provide the means to verily this hypothesis., The next generation of submm telescopes will hopefully provide the means to verify this hypothesis.1106 Recapitulating. we do not see anv overall trend of relative disk mass with object mass.," Recapitulating, we do not see any overall trend of relative disk mass with object mass."1107 The dominant feature seen in Fig., The dominant feature seen in Fig.1108 3. is a large scatter over the entire mass range. which appears to be more significant (han any possible mass dependency.," \ref{f3} is a large scatter over the entire mass range, which appears to be more significant than any possible mass dependency."1109 Previous studies of this problem are inconclusive: Andrews&Williams(2005) find that (absolute) disk masses in Daurus-Aurga scatter over three orders of magnitude and do not show anv trend with stellar masses (which range from 0.1 to 2.5.M. in their sample).," Previous studies of this problem are inconclusive: \citet{aw05} find that (absolute) disk masses in Taurus-Auriga scatter over three orders of magnitude and do not show any trend with stellar masses (which range from 0.1 to $\,M\,_{\odot}$ in their sample)."1110 The same result has been obtained by Mannings&Sargent(2000).. who include datapoints for ILerbig Ae stars with masses between 1 and 4...," The same result has been obtained by \citet{ms00}, who include datapoints for Herbig Ae stars with masses between 1 and $\,M_{\odot}$."1111" If however. the range of disk masses is constant over such a large mass range. (his implies an increase of the relative disk mass with decreasing, stellar mass by about one order of magnitude."," If, however, the range of disk masses is constant over such a large mass range, this implies an increase of the relative disk mass with decreasing stellar mass by about one order of magnitude."1112 This is in agreement with the studies of Nuernberger.(1997). ancl Nuernbergeretal.(1998)... who find a weak correlation of relative disk mass and stellar mass in Lupus and p OOph. in the sense that stars with low masses («0.7M. ) tend to have more massive disks.," This is in agreement with the studies of \citet{ncz97} and \citet{nby98}, who find a weak correlation of relative disk mass and stellar mass in Lupus and $\rho$ Oph, in the sense that stars with low masses $<0.7\,M_{\odot}$ ) tend to have more massive disks."1113 On the other hand. the review paper by Natta.CGrinin.&Mannings(2000) does not report any evidence for such a trend.," On the other hand, the review paper by \citet{ngm00} does not report any evidence for such a trend."1114 They compile disk masses for T Tori aid Herbig Ae/Be stars. and find a positive correlation of (absolute) disk mass with stellar mass.," They compile disk masses for T Tauri and Herbig Ae/Be stars, and find a positive correlation of (absolute) disk mass with stellar mass."1115 Converted to relative disk mass. the correlation clisappears.," Converted to relative disk mass, the correlation disappears."1116 Consequently. (hey claim Chat the ratio of disk to stellar mass is roughly constant in the mass range from 4 to M...," Consequently, they claim that the ratio of disk to stellar mass is roughly constant in the mass range from 4 to $\,M_{\odot}$."1117 This is confirmed by our Fig. 3..," This is confirmed by our Fig. \ref{f3},"1118 which does not show any clear sign of a trend., which does not show any clear sign of a trend.1119 In anv ease. if (here is a correlation. it is much weaker than the scatter in disk masses.," In any case, if there is a correlation, it is much weaker than the scatter in disk masses."1120where the He luminosity term accounts for the ionizing radiation (?) and Lys for the UV radiation absorbed by grains and re-emitted in the IR.,where the $\alpha$ luminosity term accounts for the ionizing radiation \citep{1999ApJS..123....3L} and $_{TIR}$ for the UV radiation absorbed by grains and re-emitted in the IR.1121 We have not considered the continuum radiation longward of 2800 which becomes important when young clusters have luminosities lower than 10 ere s'., We have not considered the continuum radiation longward of 2800 which becomes important when young clusters have luminosities lower than $^{38}$ erg $^{-1}$ .1122" The estimated TIR luminosity. Ly;5. has been computed following ?:: and the luminosity function from the 8 and 24 ffluxes as vL,."," The estimated TIR luminosity, $L_{TIR}$, has been computed following \citet{2009A&A...493..453V}: and the luminosity function from the 8 and 24 fluxes as $\nu L_\nu$."1123 The above expression is correct for HII regions whose IR emission peaks longward of 24 .Evvolved stars 51.52.83 have no emisson in the FIR and the above formula overestimates the TIR.," The above expression is correct for HII regions whose IR emission peaks longward of 24 \\.Evvolved stars s1,s2,s3 have no emisson in the FIR and the above formula overestimates the TIR."1124 For these sources we compute the luminosity according to?) who have shown that the dust emission from evolved stars closely follows that of a black body at temperatures of 400-600 K. In particular. using the 5 and M ffluxes given in Table 1. we estimate the luminosity in. solar luminosity units as: For s4 we give the luminosities and stellar masses using the formulae for HII regions but put the values in brackets since it might be an evolved star.," For these sources we compute the luminosity according to \citet{2007MNRAS.376..313G} who have shown that the dust emission from evolved stars closely follows that of a black body at temperatures of 400-600 K. In particular, using the 8 and 24 fluxes given in Table 1, we estimate the luminosity in solar luminosity units as: For s4 we give the luminosities and stellar masses using the formulae for HII regions but put the values in brackets since it might be an evolved star."1125" Extinction corrections for Ha fluxes in HII regions are 0.83 Ay where Ay. the visual extinction, is given as In Table 3 we list extinction values and cluster masses."," Extinction corrections for $\alpha$ fluxes in HII regions are 0.83 $_V$ where $_V$, the visual extinction, is given as In Table 3 we list extinction values and cluster masses."1126 Following ?.. we set 0.1 M. asthe lower limit of the IMF and use a Salpeter slope of2.3 down to 0.5 M... and |3 between 0.5 and 0.1 M...," Following \citet{2009A&A...495..479C}, we set 0.1 $_\odot$ asthe lower limit of the IMF and use a Salpeter slope of 2.3 down to 0.5 $_\odot$, and 1.3 between 0.5 and 0.1 $_\odot$."1127 When bolometric luminosities are small. as in the case of these clusters. the upper end of the IMF is not fully populated and one has to take ito account stochastic effects.," When bolometric luminosities are small, as in the case of these clusters, the upper end of the IMF is not fully populated and one has to take into account stochastic effects."1128 In the stochastic regime. uncertainties are large and a simple scaling law between luminosity and cluster mass does not apply (and it would underestimate the cluster mass).," In the stochastic regime, uncertainties are large and a simple scaling law between luminosity and cluster mass does not apply (and it would underestimate the cluster mass)."1129" For Lay10! erg s7! the possible range of cluster naàsses for a given Lj, is up to one order of magnitude (?).. ", For $_{bol}<10^{40}$ erg $^{-1}$ the possible range of cluster masses for a given $_{bol}$ is up to one order of magnitude \citep{2010arXiv1011.1097C}. .1130"In Table 3 we show the median cluster mass for the given bolometric luminosity. the expected median He luminosity. Li,” (corresponding to the observed Lj) and the relative uncertainties."," In Table 3 we show the median cluster mass for the given bolometric luminosity, the expected median $\alpha$ luminosity, $^{exp}_{H\alpha}$ (corresponding to the observed $_{bol}$ ) and the relative uncertainties."1131 The observed values corrected for extinction. Liclearv. are also given.," The observed values corrected for extinction, $^{obs}_{H\alpha}$, are also given."1132 Finally. we comment when there is not a peak u the images at the location of the 24 ssource.," Finally, we comment when there is not a clear peak in the images at the location of the 24 source."1133 In the next Section and in the rest of this Section we discuss possible correlations between cloud properties., In the next Section and in the rest of this Section we discuss possible correlations between cloud properties.1134 We will quote in parenthesis the Pearson linear correlation coetlicient rp for data relative to star forming regions. excluding AGBs.," We will quote in parenthesis the Pearson linear correlation coefficient $r_P$ for data relative to star forming regions, excluding AGBs."1135" The ratio of line intensities {ο/7,o in Figure | correlates with the ratio of line-widths (rp=0.74 for log Ws,/Wi.5 — log ντ1ας ).", The ratio of line intensities $I_{2-1}/I_{1-0}$ in Figure 1 correlates with the ratio of line-widths $r_P=0.74$ for log $W_{2-1}/W_{1-0}$ – log $I_{2-1}/I_{1-0}$ ).1136" If we exclude $19. the source with the largest W_,/W)_o. the Pearson linear correlation coefficient is 0.91and the slope is 1.940.2."," If we exclude s19, the source with the largest $W_{2-1}/W_{1-0}$, the Pearson linear correlation coefficient is 0.91and the slope is $\pm 0.2$ ."1137 The observed line ratio shows amarginal decrease as the galactocentrie radius increases (rp= —0.46) and it does not correlate with the cluster bolometric, The observed line ratio shows amarginal decrease as the galactocentric radius increases $r_P=-0.46$ ) and it does not correlate with the cluster bolometric1138concluded hat the system is a semi-detached Aleo-iype binary with a mass ratio of y= Q.1. ali orbital iiclination of 7=s7T° .9. and a temperature dillereuce between the compouents of AT=2.600 ]x. BRecettly. székkely (2003) aud Dvorak (2009) performed CCD observatious iu order to locate ὃ sculi-type »ulsations but failed to detect them.,"concluded that the system is a semi-detached Algol-type binary with a mass ratio of $q$ =0.4, an orbital inclination of $i$ $^\circ$ .9, and a temperature difference between the components of $\Delta T$ =2,600 K. Recently, Székkely (2003) and Dvorak (2009) performed CCD observations in order to locate $\delta$ Scuti-type pulsations but failed to detect them."1139 Although he orbital period of SZ Her has been exziniued several times (Ixreiner 1971: Mallama 1980: Zavala et al., Although the orbital period of SZ Her has been examined several times (Kreiner 1971; Mallama 1980; Zavala et al.1140 2002). a detailed stily of its orbital period was made by Székkely (2003) and Sovdugan (2008).," 2002), a detailed study of its orbital period was made by Székkely (2003) and Soydugan (2008)."1141 They reported that he period chaugeOm can be described usinge either a sine curve or a sinele liglt-time ellect (LITE) epemeris due to a third body with implied periods of 66 ancl Tl vr. respectively.," They reported that the period change can be described using either a sine curve or a single light-time effect (LITE) ephemeris due to a third body with implied periods of 66 and 71 yr, respectively."1142 Sovdugan (2008) aso suggested that the tiring residuals from the LITE fit inclicate au additionalcdit short-erm oscilation with a period witliin about 20 vyr., Soydugan (2008) also suggested that the timing residuals from the LITE fit indicate an additional short-term oscillation with a period within about 20 yr.1143 More thau oue housaud eclipse imiugs. spalming ~ 110v r. should be sufficient o resolve the confusion regarding he orbital beliavior of SZ Her.," More than one thousand eclipse timings, spanning $\sim$ 110 yr, should be sufficient to resolve the confusion regarding the orbital behavior of SZ Her."1144 Noueteless. the period variation oL this system has uot vet been studied as coneltsively as recuired.," Nonetheless, the period variation of this system has not yet been studied as conclusively as required."1145 m this article. a new photornetric study of SZ Her based on noderni observations and anaΝΟ is p'esented. aud it is demonsrated that the SZ Her system is ikely a quadruple one containing two low-1nass companions.," In this article, a new photometric study of SZ Her based on modern observations and analyses is presented, and it is demonstrated that the SZ Her system is likely a quadruple one containing two low-mass companions."1146 We performed uew CCD photometry of SZ Her on 13 nights [rom 20JOS February 28 through May 17., We performed new CCD photometry of SZ Her on 13 nights from 2008 February 28 through May 17.1147 The observations were taken with a SITe 2I CCD camera aud à BVH filter set attached to the 61-cim reflector at Sobaeksan Optical Astronomy Observatory (SOAQ) in Ixo‘ea., The observations were taken with a SITe 2K CCD camera and a $BVRI$ filter set attached to the 61-cm reflector at Sobaeksan Optical Astronomy Observatory (SOAO) in Korea.1148 The instrument aud reduction method used were he same as those clescribed by Lee οἱ al. (, The instrument and reduction method used were the same as those described by Lee et al. (11493XJOT. 2010b).,"2007, 2010b)."1150 GSC (BD+3392925. TYC 2610-1116-1) audGSC 2610-0821. imaeect ou the chip at the same time as the program target. were selecled as comparison and check stars. respectively.," GSC 2610-1116 $\rm BD+33^{o} 2925$, TYC 2610-1116-1) and GSC 2610-0821, imaged on the chip at the same time as the program target, were selected as comparison and check stars, respectively."1151 The lo-values of the clispersious of the imaguitice differences between these stars are withit +£0.01 mae [or all bandpasses., The $\sigma$ -values of the dispersions of the magnitude differences between these stars are within $\pm$ 0.01 mag for all bandpasses.1152 The coordinates ai| Tycho magnitudes for the three stars of ilerest are given in Table 1., The coordinates and Tycho magnitudes for the three stars of interest are given in Table 1.1153 A total of 1.728 individial observations were obtained ainone the four bandpasses (135 in B. [37 in V. 139 in AZ. aud [17 in £) amd a sample of them ts lised in Table 2.," A total of 1,728 individual observations were obtained among the four bandpasses (435 in $B$, 437 in $V$, 439 in $R$, and 417 in $I$ ) and a sample of them is listed in Table 2."1154 The light. curves of SZ Her delined by the SOAO observations are plotted in Figwe Las the V-C) differential maguitucesversus orbital phase. which was computed according o the ephemeris for our binary inodel determined later in this a‘ticle with the Wilsou-Devinuey svithesis code (Wilson Devinuey 1971. hereafter W-D).," The light curves of SZ Her defined by the SOAO observations are plotted in Figure 1 as the $-$ C) differential magnitudes orbital phase, which was computed according to the ephemeris for our binary model determined later in this article with the Wilson-Devinney synthesis code (Wilson Devinney 1971, hereafter W-D)."1155 Iu addition to these compele light curves. two eclipse timies were observed in 2001 June aud 2011: May using the same telescope.," In addition to these complete light curves, two eclipse timings were observed in 2004 June and 2011 May using the same telescope."1156 The 20(dL data were ¢¢)lectecl using the SITe 3 CCD camera aid B filter. aud the 2(11 ones using an ELI IMGEX OLECCD camera aidBVB filters.," The 2004 data were collected using the SITe 2K CCD camera and $B$ filter, and the 2011 ones using an FLI IMG4301E CCD camera and $BV$ filters."1157 GSC 2610-1116 also served as the coriparison star for these data collectious., GSC 2610-1116 also served as the comparison star for these data collections.1158 Details of the new detector have been giveu previously by Lee et al. (, Details of the new detector have been given previously by Lee et al. (11593(11).,2011).1160Ar and Ca abundances ave relatively insensitive to small changes in the temperature and Fe abundance.,Ar and Ca abundances are relatively insensitive to small changes in the temperature and Fe abundance.1161 “Pherefore. systematic uncertainties on these measurements are likely small.," Therefore, systematic uncertainties on these measurements are likely small."1162 The Ar and Ca abundance profiles (Fig., The Ar and Ca abundance profiles (Fig.1163 3aa-b) are also centrally peaked., \ref{fig:ArCa}a a-b) are also centrally peaked.1164" The Ar and Ca abunances peak ab Za,~1.2 and Zo,~LA Solar. respectively."," The Ar and Ca abundances peak at $Z_{\rm Ar}\sim1.2$ and $Z_{\rm Ca}\sim1.4$ Solar, respectively."1165" They decrease to Za,0.5 and. Zon0.6 Solar respectively by r~35 κρο", They decrease to $Z_{\rm Ar}\sim0.5$ and $Z_{\rm Ca}\sim0.6$ Solar respectively by $r\sim35$ kpc.1166 As with Fe and οἱ. both. proiles show a marginal.but plausible increase in abundance at rà30 kpc that may be due to the uplift of cool. metal-rich material. as described in Paper I. Both the Ar/Fe and Ca/Fe abundance ratio profiles (Fig.," As with Fe and Si, both profiles show a marginal,but plausible increase in abundance at $r\sim30$ kpc that may be due to the uplift of cool, metal-rich material, as described in Paper I. Both the Ar/Fe and Ca/Fe abundance ratio profiles (Fig."1167" A3cc-d) are consistent with constant values of Zr,fZyo=094cx0.02 and ZounfZp.=1.25+0.03 Solar. respectively."," \ref{fig:ArCa}c c-d) are consistent with constant values of $Z_{\rm Ar}/Z_{\rm Fe}=0.94\pm0.02$ and $Z_{\rm Ca}/Z_{\rm Fe}=1.25\pm0.03$ Solar, respectively."1168 Similarly to twe οή axindance ratio profile. the Ar/Ca ixiuidance ratio profile (big.," Similarly to the Si/S abundance ratio profile, the Ar/Ca abundance ratio profile (Fig."1169e Φου) is also an interestingὃν cross-check upon the determination of these abundances., \ref{fig:ArCa}e e) is also an interesting cross-check upon the determination of these abundances.1170 Like Si and S. Ar and Ca are also created in similar quantities by 5ος and SNla.," Like Si and S, Ar and Ca are also created in similar quantities by ${\rm SN_{CC}}$ and ${\rm SN\;Ia}$."1171 Vherelore. the Ar/Ca abundance. ratio profile is insensitive to the relative number of supernovae that explode as Όλος and SNla.," Therefore, the Ar/Ca abundance ratio profile is insensitive to the relative number of supernovae that explode as ${\rm SN_{CC}}$ and ${\rm SN\;Ia}$."1172 Instead. the Ar/Ca »indance ratio is primarily sensitive to changes in the average vields of SNla.," Instead, the Ar/Ca abundance ratio is primarily sensitive to changes in the average yields of ${\rm SN\;Ia}$."1173 The Ar/Ca abundance ratio is consistent with being flat as a function of radius and its οerved mean value is Ζγρζο=0.71£0.02 Solar., The Ar/Ca abundance ratio is consistent with being flat as a function of radius and its observed mean value is $Z_{\rm Ar}/Z_{\rm Ca}=0.71\pm0.02$ Solar.1174 The top row of Fig., The top row of Fig.1175 4. shows the abundance profiles of Ne (left). Me (middle). and Ni (right). respectively.," \ref{fig:NeMgNi} shows the abundance profiles of Ne (left), Mg (middle), and Ni (right), respectively."1176 The bottom row shows the abundance ratio profiles of Ne/Fe (Left). Meg/Ee (middle). and Nifke (right).," The bottom row shows the abundance ratio profiles of Ne/Fe (left), Mg/Fe (middle), and Ni/Fe (right)."1177 There are significant systematic uncertainties present in the determination of the Ne. Meg. and Ni abundances (see Sect. 3.3.1)).," There are significant systematic uncertainties present in the determination of the Ne, Mg, and Ni abundances (see Sect. \ref{section:bias}) )."1178 The inferred abundance profiles of Ne. Ale. ancl Ni are all centrally peaked at Zxe2.3 Solar. Ζω1.2 Solar. and Zy;2.8 Solar. respectively.," The inferred abundance profiles of Ne, Mg, and Ni are all centrally peaked at $Z_{\rm Ne}\sim 2.3$ Solar, $Z_{\rm Mg}\sim1.2$ Solar, and $Z_{\rm Ni}\sim2.8$ Solar, respectively."1179 Lhe profiles decline to ZwxeLO Solar. Zain0.6 Solar. and Zyi~1.0 Solar by r~25 kpc.," The profiles decline to $Z_{\rm Ne}\sim1.0$ Solar, $Z_{\rm Mg}\sim0.6$ Solar, and $Z_{\rm Ni}\sim1.0$ Solar by $r\sim25$ kpc."1180 Although here are differences between the abundances measured wih andMEL. strong central peaks in the abundance profiles of Ne. Mg. and Ni are found with both codes.," Although there are differences between the abundances measured with and, strong central peaks in the abundance profiles of Ne, Mg, and Ni are found with both codes."1181 These profiles also show the enhancements at 30 kpc., These profiles also show the enhancements at $\sim$ 30 kpc.1182 Phe Ne/Fe abundance ratio determined by both plasma codes is centrally. peakect., The Ne/Fe abundance ratio determined by both plasma codes is centrally peaked.1183 The Mg/Fe ratio is marginally consistent with being Lat as a function of radius and its mean values is Zi/Zy.=0.988+0.012 Solar., The Mg/Fe ratio is marginally consistent with being flat as a function of radius and its mean values is $Z_{\rm Mg}/Z_{\rm Fe}=0.988\pm0.012$ Solar.1184 This is in agreement with previous results once the dillerences between plasma codes are taken into account. (, This is in agreement with previous results once the differences between plasma codes are taken into account. (1185see Matsushita 2003. Simionescu 2010).,"see Matsushita 2003, Simionescu 2010)."1186 The Ni/Ee results derived using both plasma codes are poorly fit to a linear model., The Ni/Fe results derived using both plasma codes are poorly fit to a linear model.1187 Phe reduced. 47s are. high. vp 60/34.," The reduced $\chi^2$ s are high, $\chi^2/\nu\sim60/34$ ."1188 Primarily this is due to increased scatter in the Ni measurements at large radii (ον30 kpe)., Primarily this is due to increased scatter in the Ni measurements at large radii $r>30$ kpc).1189 Within the radial range of 5<r«30 kpc and using the code the, Within the radial range of $5<r<30$ kpc and using the code the1190Ix.,.1191(r) Ehe thermal condition T;/T.>1 for the formation of CR-NTVG leads to a family of critical lines on the P— diagram (see Fig., The thermal condition $T_i/T_s>1$ for the formation of CR-NTVG leads to a family of critical lines on the $P-\dot{P}$ diagram (see Fig.1192 1 in GMOL) 23 where A=(2.7x10*)/?52 for ASOL case (eq. [, 1 in GM01) where ${\cal A}=(2.7\times 10^3)^{1/2}=52$ for AS91 case (eq. [1193"3]) and A=(3.96xLO"")?1990 for JSG case (eq. ",3]) and ${\cal A}=(3.96\times 10^6)^{1/2}=1990$ for J86 case (eq. [11944).,4]).1195" Alternatively. one can find a minimum required surface magnetic field D,=bD, expressed by the coefficient b in the formΑπρ27.9) In this case the gap height —ος is determined by the condition /=/,,~{ων where f, is (he mean free path of the electron to emit a photon with energy fiw=25heD./ime (Zhang GAIOL).."," Alternatively, one can find a minimum required surface magnetic field $B_s=bB_d$ expressed by the coefficient $b$ in the form In this case the gap height $h=h_{ICS}$ is determined by the condition $h=l_{ph}\sim l_e$, where $l_e$ is the mean free path of the electron to emit a photon with energy $\hbar\omega=2\gamma\hbar eB_s/mc$ \citep[][ GM01]{zhm00}."1196 The ICS-NTVG model is described by the folowing parameters: the heieht of a quasi steady gap (10)the gap potential drop )6:0.72hο... and the surface temperatureN, The ICS-NTVG model is described by the following parameters: the height of a quasi steady gap the gap potential drop ) and the surface temperature.1197"12) The thermal condition 7;/1T,>1 for the formation of ICS-NTWVG leads to a family of critical lines on the J?—P diagram (see Fig.", The thermal condition $T_i/T_s>1$ for the formation of ICS-NTVG leads to a family of critical lines on the $P-\dot{P}$ diagram (see Fig.1198 1 in GAIOL)ο»... where B=(2x102)?14 for AS91 case (eq. [, 1 in GM01) where ${\cal B}=(2\times 10^2)^{1/2}=14$ for AS91 case (eq. [11993]) and B—(1.69x107)?=130 for J36 case (eq. |,3]) and ${\cal B}=(1.69\times 10^4)^{1/2}=130$ for J86 case (eq. [12004]).,4]).

Showing the first 1,200 of 10280 lines. Download the file for the rest.