ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Ii such cases; we need to solve the liuear system(10).," In such cases, we need to solve the linear system."3". The direct method described. above las several advantages with respect to the ""kernel method aud to the over-relaxation method: 0) The method is fast in the case where an orthouormal set of fictions cau be found.", The direct method described above has several advantages with respect to the “kernel” method and to the over-relaxation method: (i) The method is fast in the case where an orthonormal set of functions can be found.4" Tn fact. we need oulv to evaluate one iuteeral for cach coefficient e, that we want to caleulate. ("," In fact, we need only to evaluate one integral for each coefficient $c_\alpha$ that we want to calculate. ("50) The method does not require a large amount of memory: we necd to retain onlv the » values of the coefficieuts ος. ,ii) The method does not require a large amount of memory: we need to retain only the $n$ values of the coefficients $c_\alpha$. (6i) The precision ofthe aversion is driven in a natural wav bv the value of o.,iii) The precision ofthe inversion is driven in a natural way by the value of $n$.7" Typically. the larger a is. the smaller the leneth scale of f£, (see below). ("," Typically, the larger $\alpha$ is, the smaller the length scale of $f_\alpha$ (see below). ("8"iv) Tn some cases. the decomposition of the mass deusitv &(0) im terms of the functions f, cau be useful.","iv) In some cases, the decomposition of the mass density $\tilde\kappa(\vec\theta)$ in terms of the functions $f_\alpha$ can be useful."9 When the field Q is rectangular. an orthonormal set of fictions can be written casily.," When the field $\Omega$ is rectangular, an orthonormal set of functions can be written easily."10 Tere we consider the special case when @ is a square of leugth z (iu some suitable units) auv rectangular field cau be haudlecd iu a simular mauner., Here we consider the special case when $\Omega$ is a square of length $\pi$ (in some suitable units); any rectangular field can be handled in a similar manner.11 In the case considered. an orthonormal set of functions 1s given by with (n.3)C(0.0).," In the case considered, an orthonormal set of functions is given by with $(\alpha, \beta) \in \mathbb{N}^2 \setminus (0,0)$."12 The normalization nm; is defined as The function fgg ix not defined., The normalization $n_{\alpha\beta}$ is defined as The function $f_{00}$ is not defined.13 Note that here we use two indices for the set., Note that here we use two indices for the set.14 Cosimes must be used in order to have a set (see Eqs.(," Cosines must be used in order to have a set (see Eqs.,"157)..(12).. and Appendix A)., and Appendix A).16" Our problemi is solved in terms of the coefficients 6,5: We now observe that the particular choice of the orthonormal set {fii} allows us to use fast Fourier transform (FFT) techniques to evaluate Eqs.", Our problem is solved in terms of the coefficients $c_{\alpha\beta}$: We now observe that the particular choice of the orthonormal set $\{ f_{\alpha\beta} \}$ allows us to use fast Fourier transform (FFT) techniques to evaluate Eqs.17 aud(18)., and.18 The use of FFT aakes the direct method very cfücieut: in articular the method becomes of order O(N?log N)., The use of FFT makes the direct method very efficient: in particular the method becomes of order $\O\bigl( N^2 \log N \bigr)$ .19 Moreover. several optimized FFT libraries are available.," Moreover, several optimized FFT libraries are available."20 The optimal fruncation for the series is determined bv the adopted erid uunbers for a erid of NSAL points. à should run from Oto No1. aud > from Oto AL1 (this is standard practice for FET libraries).," The optimal truncation for the series is determined by the adopted grid numbers: for a grid of $N \times M$ points, $\alpha$ should run from $0$ to $N-1$, and $\beta$ from $0$ to $M-1$ (this is standard practice for FFT libraries)."21 Our iethod has been implemented iu aud inIDL., Our method has been implemented in and in.22" The version uses the library (""Fastest Fourier Transform in the West.” version 2.0.1) to perform discrete Fourier transforms (DET)."," The version uses the library (“Fastest Fourier Transform in the West,” version $2.0.1$ ) to perform discrete Fourier transforms (DFT)."23 This library. written by Matteo Frigo and Steven €. Jolusou. is considered the quickest DET library publicly available.," This library, written by Matteo Frigo and Steven G. Johnson, is considered the quickest DFT library publicly available."24 The performance of our direct method is compared ith that of the over-relaxation method. also implemented in C.," The performance of our direct method is compared with that of the over-relaxation method, also implemented in ."25 The procedure used in the tests is suniunarized in the following points: The results obtainedin the tests are the followiug:, The procedure used in the tests is summarized in the following points: The results obtainedin the tests are the following:26and it is surrounded by a complex morphology of cores and filaments.,and it is surrounded by a complex morphology of cores and filaments.27" In conclusion, the radial behaviour of the collapsing cores is in good agreement with previous models."," In conclusion, the radial behaviour of the collapsing cores is in good agreement with previous models."28 The core structure can therefore be thought of as a superposition of two effects., The core structure can therefore be thought of as a superposition of two effects.29" Firstly, there is the high density peak and steep radial decrease in mean density due to gravitational collapse."," Firstly, there is the high density peak and steep radial decrease in mean density due to gravitational collapse."30" Secondly, there is the effect of the density field of the molecular cloud which is highly filamentary and irregular."," Secondly, there is the effect of the density field of the molecular cloud which is highly filamentary and irregular."31" It is this medium which has rapidly undergone collapse, and there is no time or mechanism for this structure to be homogenised (?).."," It is this medium which has rapidly undergone collapse, and there is no time or mechanism for this structure to be homogenised \citep{Klessen00}."32" For example, the irregular cores seen by ? will have their non-axisymetric geometries enhanced during collapse, as denser regions collapse more swiftly than less dense ones."," For example, the irregular cores seen by \citet{Stutz09} will have their non-axisymetric geometries enhanced during collapse, as denser regions collapse more swiftly than less dense ones."33 Generally low mass and high mass sinks have different accretion histories., Generally low mass and high mass sinks have different accretion histories.34 Typically low mass sinks only have accretion from within their core as outlined in the previous section., Typically low mass sinks only have accretion from within their core as outlined in the previous section.35" However, the growth of higher mass sinks is dominated by a secondary stage of accretion; the accretion of additional infalling material from outside the original core radius."," However, the growth of higher mass sinks is dominated by a secondary stage of accretion; the accretion of additional infalling material from outside the original core radius."36 We shall now examine whether this additional material is also inhomogeneously distributed., We shall now examine whether this additional material is also inhomogeneously distributed.37" To answer this question, we use a modification of the previous Hammer projection technique."," To answer this question, we use a modification of the previous Hammer projection technique."38" As the SPH method is Lagrangian, it is possible to flag each individual gas particle that becomes accreted by a sink."," As the SPH method is Lagrangian, it is possible to flag each individual gas particle that becomes accreted by a sink."39" We take a shell at the boundary of the core, and track the positions at which flagged particles from the external environment pass through the shell."," We take a shell at the boundary of the core, and track the positions at which flagged particles from the external environment pass through the shell."40 We do this over a period of 2x104 years (which is about the mean dynamical time of the bound pre-stellar cores found in SCB09) and then make Hammer projections of the integrated density on the shell in a similar manner to before., We do this over a period of $2\E^4$ years (which is about the mean dynamical time of the bound pre-stellar cores found in SCB09) and then make Hammer projections of the integrated density on the shell in a similar manner to before.41 Figures 9 and 10 show the resulting ‘accretion surfaces’ of material which passes through shells at r=Re0.01 pc and r=0.1 pc around the representative cores over the time interval., Figures \ref{CDaccinner} and \ref{CDaccouter} show the resulting `accretion surfaces' of material which passes through shells at $r=R_c=0.01$ pc and $r=0.1$ pc around the representative cores over the time interval.42 We do not specify the absolute values on this surface as it is merely a way of visualising the direction at which accreted material approaches the core., We do not specify the absolute values on this surface as it is merely a way of visualising the direction at which accreted material approaches the core.43 It is not a real density., It is not a real density.44 The cores in 9 show differing amounts of additional accreted material passing through the core surface at r=R.., The cores in \ref{CDaccinner} show differing amounts of additional accreted material passing through the core surface at $r=R_c$.45" For instance, core (a) has almost no additional accretion from its external environment, but core (b) has a large amount."," For instance, core (a) has almost no additional accretion from its external environment, but core (b) has a large amount."46" The distribution of accreted material is again highly anisotropic, and shows a resemblance to the column density distribution in 1.."," The distribution of accreted material is again highly anisotropic, and shows a resemblance to the column density distribution in \ref{CDprojections}."47 Accreted material from the environment is therefore entering the core along the high column density filaments., Accreted material from the environment is therefore entering the core along the high column density filaments.48 10 shows how the accreted material enters, \ref{CDaccouter} shows how the accreted material enters49iu the PNLE for the spectroscopically coufiiied ICPNe oe row Virgo fields.,than the PNLF for the spectroscopically confirmed ICPNe in our Virgo fields.502. Caven the well-studied xoperties of the PNLF in galaxies. this Figure shows that jese chuission line candidates are not a population of PNe at the distance of LO Mpc. and are not associated with the II cloud in the Leo intragroup region.," Given the well-studied properties of the PNLF in galaxies, this Figure shows that these emission line candidates are not a population of PNe at the distance of 10 Mpc, and are not associated with the HI cloud in the Leo intragroup region."51 The bright edge of the LF for the point-like cinission line candidates in the Leo field at συ=26.7 rules out PNe at a distance of LO Ape as a possible explanation for these sources., The bright edge of the LF for the point-like emission line candidates in the Leo field at $m_{5007} = 26.7$ rules out PNe at a distance of 10 Mpc as a possible explanation for these sources.52" These objects must therefore be background clnitters, either [OTT] emitters at νε 0.31 or Evo cutters at 2ozXd."," These objects must therefore be background emitters, either [OII] emitters at $z\simeq530.34$ , or $\alpha$ emitters at $z\simeq 3.1$."54 Our criteria will select preferentially Lya galaxies at z23.1. while ΟΤΙ cinitters at 20.31 are iulikelv because of the EW >100A requirement (Πας et 11997: Ποσο et 11995).," Our criteria will select preferentially $\alpha$ galaxies at $z \simeq 3.1$, while [OII] emitters at $ z\simeq 0.34$ are unlikely because of the EW $ >55100 $ requirement (Hammer et 1997; Hogg et 1998)."56 We must then compare the LF of our selected. emission line objects with the LF of Lvo galaxies at :=3.1 , We must then compare the LF of our selected emission line objects with the LF of $\alpha$ galaxies at $ z = 3.1$ .57There is no well-:uupled Ίσα LF of such objects at >=Xd in the literature., There is no well-sampled $\alpha$ LF of such objects at $ z = 3.1$ in the literature.58 Very recent results for the Subaru deep feld were obtained for the Lya population at 2=Lad (Ouchi ot 22003). but Arnaboldi et ((2002) derived it from the work by Steidel et ((2000).," Very recent results for the Subaru deep field were obtained for the $\alpha$ population at $z=4.86$ (Ouchi et 2003), but Arnaboldi et (2002) derived it from the work by Steidel et (2000)."59" The resulting LF. scaled to our effective volumelore"". is shown as full line in Figure 6.."," The resulting LF, scaled to our effective volume, is shown as full line in Figure \ref{Fig:lfla}."60 We have also added data points coustructed from the Lya blauk field search done by Cowie IIu (]998). and the spectroscopically confirmed Lyra sample from Wudritzki et ((2000).," We have also added data points constructed from the $\alpha$ blank field search done by Cowie Hu (1998), and the spectroscopically confirmed $\alpha$ sample from Kudritzki et (2000)."61.6. The resulting LEs for the field Exo cutters at zoXd agree very well ith that computed from Steidel et (20001., The resulting LFs for the field $\alpha$ emitters at $z \sim 3.1$ agree very well with that computed from Steidel et (2000).62 Tn these works (Cowie IIu 1988: Steidel et 22000: I&uditzki et 22000) Ίνα eunitters are ideutified via their uarrow line excess. regardless of whether theirenissjon is poiut-like or resolved.," In these works (Cowie Hu 1988; Steidel et 2000; Kudritzki et 2000) $\alpha$ emitters are identified via their narrow line excess, regardless of whether theiremission is point-like or resolved."63 Therefore we have produceda catalog of all objects with liue excess, Therefore we have produceda catalog of all objects with line excess64NV334.,.65. Because of the low brightness of the system (V=18.83 mag) we were able to obtain only a very rough estimate of the velocity amplitude., Because of the low brightness of the system $V=18.83$ mag) we were able to obtain only a very rough estimate of the velocity amplitude.66" We therefore modifed our analysis: we estimated the inclination based on the light curve. and compared the calculated velocity amplitude K,. to the observed one."," We therefore modifed our analysis: we estimated the inclination based on the light curve, and compared the calculated velocity amplitude $K_c$ to the observed one."67" Assuming that the spectrum originates in the secondary leads to Κι. being much larger than &,5. regardless of the value of g."," Assuming that the spectrum originates in the secondary leads to $K_c$ being much larger than $K_{obs}$, regardless of the value of $q$."68 The same analysis applied to the configuration with the primary generating the spectrum produces results shown in Table 6.. which seem to indicate that NV334 is similar to V251.," The same analysis applied to the configuration with the primary generating the spectrum produces results shown in Table \ref{tab: NV334p}, which seem to indicate that NV334 is similar to V251."69 However. the similarity is superticial: the effective temperatures of the two systems differ by ~4000 K. causing the cooler and dimmer NV334 to occupy an entirely different position in the color-magnitude diagram of c) Cen (see Fig. 1)).NV360..," However, the similarity is superficial: the effective temperatures of the two systems differ by $\sim4000$ K, causing the cooler and dimmer NV334 to occupy an entirely different position in the color-magnitude diagram of $\omega$ Cen (see Fig. \ref{fig:70omegacmd})."71 Our standard analysis fails for this system (iterations of the inclination 7 based on the velocity curve are either converging extremely slowly or entirely diverging). so that we again have to apply a modified version based on the light curve.," Our standard analysis fails for this system (iterations of the inclination $i$ based on the velocity curve are either converging extremely slowly or entirely diverging), so that we again have to apply a modified version based on the light curve."72" As for the case of NV334. assuming that the spectrum originates in the secondary leads to K. being much larger than K,5,. regardless of the value of q."," As for the case of NV334, assuming that the spectrum originates in the secondary leads to $K_c$ being much larger than $K_{obs}$, regardless of the value of $q$."73 Systems with the primary being responsible for the spectrum fare much better (see Table 7)). and we conclude that NV360. unlike NV334. is truly similar toNV400.," Systems with the primary being responsible for the spectrum fare much better (see Table \ref{tab: NV360p}) ), and we conclude that NV360, unlike NV334, is truly similar to."74. When the secondary is assumed to generate the spectrum. the calculated light-curve amplitude ts for all values of q much higher than the observed one.," When the secondary is assumed to generate the spectrum, the calculated light-curve amplitude is for all values of $q$ much higher than the observed one."75 Assuming that the spectrum originates in the primary we get the results displayed in Table 8.., Assuming that the spectrum originates in the primary we get the results displayed in Table \ref{tab: NV400p}.76 Apparently. NV400 is yet another system similar to V240.," Apparently, NV400 is yet another system similar to ."77. Assuming that the spectrum originates from the primary we obtain the results shown in Table 9:: computed light-curve amplitudes are far too low (1f the secondary were the source of the spectrum they would be even lower)., Assuming that the spectrum originates from the primary we obtain the results shown in Table \ref{tab: V240p}: computed light-curve amplitudes are far too low (if the secondary were the source of the spectrum they would be even lower).78 Moreover. the total mass of the system is unacceptably high (as we argued in the case of V214. it should not exceed 1.84 M).," Moreover, the total mass of the system is unacceptably high (as we argued in the case of V214, it should not exceed 1.84 $M_\odot$ )."79 In other words. the color and apparent magnitude of V240 are incompatible with its observed velocity amplitude. suggesting that this system does not belong to w Cen.," In other words, the color and apparent magnitude of V240 are incompatible with its observed velocity amplitude, suggesting that this system does not belong to $\omega$ Cen."80 Another possibility is that the light of V240 is dominated by a tertiary component. and the recorded spectra are unrelated to the photometric," Another possibility is that the light of V240 is dominated by a tertiary component, and the recorded spectra are unrelated to the photometric"81Wilkineetal.(2005) compiled a list of 316 voung stus in LI688. the main cloud of the p Ophiuchi complex.,"\citet{wil08} compiled a list of 316 young stars in L1688, the main cloud of the $\rho$ Ophiuchi complex."82 About 46 HILT objects Darsony 2004)... 119 [ls emission features (Grossoetal.2001:Gómez2003:INhanzadyan 2004).. 16 high velocity CO outflows (Dontemps1996:Bussmannetal.2007) have been observed in the p Ophiuchi complex.," About 46 HH objects \citep{wilking97,wu02,phe04}, , 119 $_2$ emission features \citep{gro01,gom03,kha04}, 16 high velocity CO outflows \citep{bontemps96,bussmann07} have been observed in the $\rho$ Ophiuchi complex."83 For a summary of the p Ophiuchi complex we reler to the recent review by Wilkingetal.(2003) The IRAC observations of p Ophiuchi were conducted on 2004 March 28-30 and 2005 september 15 and 18.," For a summary of the $\rho$ Ophiuchi complex we refer to the recent review by \citet{wil08}84 The IRAC observations of $\rho$ Ophiuchi were conducted on 2004 March 28-30 and 2005 September 15 and 18."85 These observations are part of the Spitzer legacy. programced and the program ID is 177., These observations are part of the Spitzer legacy program \citep{eva03} and the program ID is 177.86 Fig., Fig.87 1 shows the Ciree-color image of (he area mapped with IRAC which is about 8.0 ?deg., \ref{fig1} shows the three-color image of the area mapped with IRAC which is about 8.0 $^2$.88 This area is included in the region mapped with MIPS (see Fig.l of Paclgettetal.2008) ), This area is included in the region mapped with MIPS (see Fig.1 of \citealt{pad08}) ).89 Each region in Fig., Each region in Fig.90 1. was individually mapped in two epochs with an exposure lime of 12 seconds (24 seconds in total) ancl one map was macle in the high dynamic range (IID) mode. which added a short exposure time of 0.6 seconds (Evansetal.2003)..," \ref{fig1} was individually mapped in two epochs with an exposure time of 12 seconds (24 seconds in total) and one map was made in the high dynamic range (HDR) mode, which added a short exposure time of 0.6 seconds \citep{eva03}."91 Six off-cloud fields were observed in order to sample the backeround source counts., Six off-cloud fields were observed in order to sample the background source counts.92 We used the final mosaic images and source catalogs of the fourth delivery of IRAC and ALIPS data., We used the final mosaic images and source catalogs of the fourth delivery of IRAC and MIPS data.93 The data were processed withSpitzer Science Center pipeline S13 and calibrated bye2d team (details can be found in the Delivery Document!))., The data were processed with Science Center pipeline S13 and calibrated by team (details can be found in the Delivery ).94 In the final mosaic images bright source artifacts have been removed., In the final mosaic images bright source artifacts have been removed.95 ILowever. some instrumental effects still survive.," However, some instrumental effects still survive."96" For example. optical ghosts produced by internal reflections within the filters (labelled with ""eh in the subsequent figures) and strav lights produced by ilhuninating sources olf the eclees of the detector axravs (labelled in the subsequent figures with “Stray light) remain in the final mosaic images (see the [TRAC data for details)."," For example, optical ghosts produced by internal reflections within the filters (labelled with “gh"" in the subsequent figures) and stray lights produced by illuminating sources off the edges of the detector arrays (labelled in the subsequent figures with “Stray light"") remain in the final mosaic images (see the IRAC data for details)."97 A catalog of voung stellar object candidates in the p Ophiuchi. YSOc. is also delivered bv thecd team.," A catalog of young stellar object candidates in the $\rho$ Ophiuchi, YSOc, is also delivered by the team."98 Evansetal.(2009). improved over this catalog., \citet{eva09} improved over this catalog.99 Thev removed some suspect sources. added known sources and data obtained at other wavelengths. and calculated additional quantities such as bolometric temperature and bolometricIuminositv.," They removed some suspect sources, added known sources and data obtained at other wavelengths, and calculated additional quantities such as bolometric temperature and bolometricluminosity."100 Their list, Their list101then we cross-correlated them with. the ρα images.,then we cross-correlated them with the $R$ -band images.102" The depth of the ραπ images is rather homogeneous. providing a S/Nzc3 for objects with /?z25 (with a 3"" diameter photometric aperture)."," The depth of the $R$ -band images is rather homogeneous, providing a $S/N \approx 3$ for objects with $R \approx 25$ (with a $^{\prime \prime}$ diameter photometric aperture)."103 The magnitudes were measured in 3” diameter apertures anc converted to 6 magnitudes by applying the aperture corrections estimated from field stars., The magnitudes were measured in $^{\prime \prime}$ diameter apertures and converted to $^{\prime \prime}$ magnitudes by applying the aperture corrections estimated from field stars.104 Such corrections are typically in the range of 0.08-0.15 and 0.04-0.10 magnitudes in AY and. £2 bands respectively., Such corrections are typically in the range of 0.08-0.15 and 0.04-0.10 magnitudes in $K^{\prime}$ - and $R$ - bands respectively.105 Photometric errors were estimated according to the poissonian noise from the objects and [rom the background., Photometric errors were estimated according to the poissonian noise from the objects and from the background.106 Finally. the magnitudes. were corrected. for Galactic extinction using the ele values derived. from. the extinction maps of Burnstein LHeiles (1982) and adopting ely=0.56.15 and cla=0:07. (see Table 1).," Finally, the magnitudes were corrected for Galactic extinction using the $A_{B}$ values derived from the extinction maps of Burnstein Heiles (1982) and adopting $A_{R}=0.56 A_{B}$ and $A_{K}=0.07 A_{B}$ (see Table 1)."107 The completeness in cach A field was estimated: by constructing a background image for each field and adding random simulated objects with Mollfat. profiles consistent with the observed seeing., The completeness in each $K^{\prime}$ field was estimated by constructing a background image for each field and adding random simulated objects with Moffat profiles consistent with the observed seeing.108 The simulated objects were then detected with SExtractor as a function of magnitude., The simulated objects were then detected with SExtractor as a function of magnitude.109 The completeness is rather homogeneous among all the fields., The completeness is rather homogeneous among all the fields.110" The global completeness of the whole A""- selected. sample is for ἐν=19.0E0.1 and it decreases to about [or A=19.240.1.", The global completeness of the whole $K^{\prime}$ - selected sample is for $K=19.0 \pm 0.1$ and it decreases to about for $K=19.2 \pm 0.1$.111 The common sky area covered by /?- ancl A-band imaging is 40.0 arcmin., The common sky area covered by $R$ - and $K$ -band imaging is 40.0 $^2$.112 At A<19.2. only 3 objects are undetected in {ρα (see Figure 3).," At $K<19.2$, only 3 objects are undetected in $R$ -band (see Figure 3)."113 The main uncertainty in galaxy counts comes from the stargalaxy separation., The main uncertainty in galaxy counts comes from the star–galaxy separation.114 Because of the poor pixcl sampling in the IRAC?2B images. the morphological classifier used by SExtractor is more reliable. at moderately bright magnitudes.," Because of the poor pixel sampling in the IRAC2B images, the morphological classifier used by SExtractor is more reliable at moderately bright magnitudes."115 Such classifier. called ~stellarity inclex” (9). is by definition equal to 1.0 for stars and to 0.0 for galaxies.," Such classifier, called “stellarity index” $S$ ), is by definition equal to 1.0 for stars and to 0.0 for galaxies."116 We found that S>0.9 provided. a reliable classification of stellar objects at AC «17.5., We found that $S>0.9$ provided a reliable classification of stellar objects at $K<$ 17.5.117 We then decided to statistically subtract the star counts using the method. outlined. by Saraceo et al. (, We then decided to statistically subtract the star counts using the method outlined by Saracco et al. (118"1997): taking advantage of the good secing and better sampling of the SUSI images. we extracted a subsample of objects that were reliably classified as stellar in the ρα, ","1997): taking advantage of the good seeing and better sampling of the SUSI images, we extracted a subsample of objects that were reliably classified as stellar in the $R$ -band."119In order to assess the fraction of objects that are misclassified in the near-Hi images. we computed the ratio Re=INSp where Nog and Nop are the number of objects that SlExtractor classifies as stars (ic. with S2 0.9) respectively in the LRAC2B and in the SUSL images.," In order to assess the fraction of objects that are misclassified in the near-IR images, we computed the ratio $R_s = N_{s,K}/N_{s,R}$, where $N_{s,K}$ and $N_{s,R}$ are the number of objects that SExtractor classifies as stars (i.e. with $S>0.9$ ) respectively in the IRAC2B and in the SUSI images."120 The ratio 2. was computed as a function of the A'- magnitude over a range 15.0fv «19.0., The ratio $R_s$ was computed as a function of the $K$ -band magnitude over a range $<K<$ 19.0.121 As expected. Haod for 15.0.A «17.0. indicating that SLExtractor correctly classifies the stellar objects at moderately. bright {να magnitudes.," As expected, $R_s\sim1$ for $<K<$ 17.0, indicating that SExtractor correctly classifies the stellar objects at moderately bright $K$ -band magnitudes."122 However. the classification. elliciency drops to Ry~0.5 for 11.0ἐν «19.0. showing that at fainter A-band magnitudes SExtractor underestimates the number of stars in LRAC2B images.," However, the classification efficiency drops to $R_s\sim0.5$ for $<K<$ 19.0, showing that at fainter $K$ -band magnitudes SExtractor underestimates the number of stars in IRAC2B images."123 The galaxy counts were then corrected adopting the Ry values computed for cach {να magnitude bin., The galaxy counts were then corrected adopting the $R_s$ values computed for each $K$ -band magnitude bin.124" We excluded from the counts the objects lving at «3"" from the edges of the images in order to avoid. photometric uncertainties due to edge elfects.", We excluded from the counts the objects lying at $<3^{\prime \prime}$ from the edges of the images in order to avoid photometric uncertainties due to edge effects.125 This implied a reduction of the useful total area of the 14 fields from 46.4 arcmin? to 43.4 arcmin?., This implied a reduction of the useful total area of the 14 fields from 46.4 $^{2}$ to 43.4 $^{2}$.126 ‘Table 2 and Figure 2 show the galaxy counts in. A- as derived from the observed fields compared with the literature averaged. counts of Hall Green (1998) ane the counts of the wide-field (181. aremin?) survey of Minezaki et al. (, Table 2 and Figure 2 show the galaxy counts in $K$ -band as derived from the observed fields compared with the literature averaged counts of Hall Green (1998) and the counts of the wide-field (181 $^{2}$ ) survey of Minezaki et al. (1271998).,1998).128 Por A.«17.5. our counts are consistent. with the average literature counts.," For $K<17.5$, our counts are consistent with the average literature counts."129 For 17.5«A19 the counts seem to have a systematic excess of a factor of 21.2-1.5 respect to average literature counts., For $17.5<K<19$ the counts seem to have a systematic excess of a factor of $\approx$ 1.2-1.5 respect to average literature counts.130 For A719 the counts are incomplete in most of the fields (see Tab 2)., For $K>19$ the counts are incomplete in most of the fields (see Tab 2).131 In. order ο provide an estimate of the field-to-Bield variations of the counts. we do not show their poissonian uncertainties. but he their rms derived from the different fields.," In order to provide an estimate of the field-to-field variations of the counts, we do not show their poissonian uncertainties, but the their rms derived from the different fields."132" We also tried an alternative approach deriving the A- counts by selecting objects with band: ""stellarity incex” «0.9 from the total area in common between he A- and the B-bancd images (40 arcmin?).", We also tried an alternative approach deriving the $K$ -band counts by selecting objects with $R$ -band “stellarity index” $<0.9$ from the total area in common between the $K$ - and the $R$ -band images (40 $^2$ ).133 Within the »oissonian uncertainties. we obtained results consistent with he previous method. thus independently: confirming the," Within the poissonian uncertainties, we obtained results consistent with the previous method, thus independently confirming the"134as anomalies before (Stetson1991) aud may be due o the composite colors of binary stars. or clance superpositions which are [i yw only one star in the reduction procedure.,"as anomalies before \citep{stet94} and may be due to the composite colors of binary stars, or chance superpositions which are fit by only one star in the reduction procedure."135 Both of these Seeestlous cani be well sl«liec with simulations of the c'ompleteuess aid crowding effects. aud will be discussed in a future paper.," Both of these suggestions can be well studied with simulations of the completeness and crowding effects, and will be discussed in a future paper."136 Thi‘ly. the theory precicts too many bright blue stragelers.," Thirdly, the theory predicts too many bright blue stragglers."137 Since ile first two problems cainol JO adressed in the context of our theoretical moclels. we focus here ou the third point. and exylore what is required to prodwe a theoretical blue stragelerMD distribution Wwlich terminates at the s:une uaglLE.udes as the observed blue stragelers.," Since the first two problems cannot be addressed in the context of our theoretical models, we focus here on the third point, and explore what is required to produce a theoretical blue straggler distribution which terminates at the same magnitudes as the observed blue stragglers."138 As discussed above. he brigh iue stragelers have ligh masses. aud do not live very long.," As discussed above, the bright blue stragglers have high masses, and do not live very long."139" Tjerefore. in order to have a poplation of blue strag[n]oOers which lacks bright stars. the blue st""Qooassers must have stopped Lormilg SOLDe time ago."," Therefore, in order to have a population of blue stragglers which lacks bright stars, the blue stragglers must have stopped forming some time ago."140 In the cortext of the models described above. we [iuc| that a blue stragelerMD [orijalion ‘ate which teriiated 3 Cyr ago reprocduc‘es the upper art o ‘the observed. blue strageler distribution quite wel (Fieure 9).," In the context of the models described above, we find that a blue straggler formation rate which terminated 3 Gyr ago reproduces the upper part of the observed blue straggler distribution quite well (Figure 9)."141 However. we caution that e precise date of the termiuatio of blie strageler forulaio slould not be taken too seriously.," However, we caution that the precise date of the termination of blue straggler formation should not be taken too seriously."142 Virst. he formation rates use here are ιοί realistic.," First, the formation rates used here are not realistic."143 A hl« Ivnamical model of tle evolution of e cluster wotld be required to jxoduce accurate time-cdepedent rates., A full dynamical model of the evolution of the cluster would be required to produce accurate time-dependent rates.144 Second. «our resuls ale [Iuenced by our choice of bitary parameters and mass hictions. although tlie saue qualialive ects will appv regardless of he choice of these paramete‘s.," Second, our results are influenced by our choice of binary parameters and mass functions, although the same qualitative effects will apply regardless of the choice of these parameters."145 ThiX. the observed sauiple is Diasec two importauo ways.," Third, the observed sample is biased in two important ways."146 Incoimupleteness due to crowding wi allect the distribution. particularly at e [aint eud.," Incompleteness due to crowding will affect the distribution, particularly at the faint end."147 However this shoud not affect the lack of bright blue stragelers. which is the observe Βατο we are rving to reproduce.," However this should not affect the lack of bright blue stragglers, which is the observed feature we are trying to reproduce."148 Νore inuportantly. we do uo have complee spatial coveage of e cluster.," More importantly, we do not have complete spatial coverage of the cluster."149 HST results sugges that he blue stragelerMOD distributiou extends to brigler lijtsiuthle ster core (Cillilaxletal.19O8)., HST results suggest that the blue straggler distribution extends to brighter limits in the cluster core \citep{Getal98}.150. It is possible that the blue str:welerMD clistributio1 ds clille'ent it e core because the contribution of lue stragelers created by Ην mergers. raljer {1an stellar lisious. is larger 1 ithe outer regions.," It is possible that the blue straggler distribution is different in the core because the contribution of blue stragglers created by binary mergers, rather than stellar collisions, is larger in the outer regions."151 If so. the lack of b‘ight bue stragelers in this region may be ye Closely relate to the characteristies of the binary »opulaion iu this region han he stellar lision rate.," If so, the lack of bright blue stragglers in this region may be more closely related to the characteristics of the binary population in this region than the stellar collision rate."152 Detailed inodels of binary merger evolutionlary tracks. coujned witL preclictious of he binary populatiοι5 Luerger rate. will be necessary O μίαςe this degeneracy.," Detailed models of binary merger evolutionary tracks, combined with predictions of the binary population's merger rate, will be necessary to untangle this degeneracy."153 Tje. lacς of right blue stragelers outside tie core could be explainec| by mass segregation. eitjer becatse the nore uassive blue stragelers sitk to the cluster center (al10wel this effect shoid iοἱ be domiuaut since he mass ΠΟΝΟ between the bright aud faint bh esrageleMODs would be relaively sinall). or N driving tie [ew remaining binaries toward the center of re cluser.," The lack of bright blue stragglers outside the core could be explained by mass segregation, either because the more massive blue stragglers sink to the cluster center (although this effect should not be dominant since the mass difference between the bright and faint blue stragglers would be relatively small), or by driving the few remaining binaries toward the center of the cluster."154 However. we do not believe hat mass segregation alone coud account [or the sharp ¢oll in due strageleree hu-—jnosities 1hat we observe in the absence of a sieuilicant change iu the ιο strageler [ormatiou rae. since it is ard to yelieve the upper part of he blue strageler distribjon coud be lost from our observatOLS eiven that tiere are large nube's of observed blue stragelers in 11e L.1—L1AM. lalee.," However, we do not believe that mass segregation alone could account for the sharp cutoff in blue straggler luminosities that we observe in the absence of a significant change in the blue straggler formation rate, since it is hard to believe the upper part of the blue straggler distribution could be lost from our observations given that there are large numbers of observed blue stragglers in the $1.1-1.4 M_{\sun}$ range."155 Thi stle data appear to suggest that 17 Tuc lias passe through a stage similar to the cur'ent state of M80 at some point in tie. past., Thus the data appear to suggest that 47 Tuc has passed through a stage similar to the current state of M80 at some point in the past.156 The large extent of the blue straMDooler seqence in ALSO observec by Ferraroetal.(1999) teuds to support this interpretation. suce a cluster whose blue strageler formation rate is uuusially high at the present. tite should tend to appea: [ike that iu," The large extent of the blue straggler sequence in M80 observed by \cite{fer99} tends to support this interpretation, since a cluster whose blue straggler formation rate is unusually high at the present time should tend to appear like that in"157"and n,=[fF(5)ds is the particle number density.",and $n_e\equiv\int F(\gamma){\rm d}\gamma$ is the particle number density.158" Obviously. the radio beam intensity I, decreases because of the photon scattering to thestate with 04>0 and the maximum scatlerine probability. corresponds to £4—s and νι=p/4«pr."," Obviously, the radio beam intensity $I_\nu$ decreases because of the photon scattering to thestate with $\theta_1159>\theta$ and the maximum scattering probability corresponds to $\theta_{1_{\rm max}}=\pi$ and $\nu_1=\nu\theta^2/4\ll\nu$."160 Note that in this situation the azimuthal angle ©; is of no interest., Note that in this situation the azimuthal angle $\phi_1$ is of no interest.161 The svstem (4) has the following solution (see.e.g...Petrova2004b):: I is (he first integral.," The system (4) has the following solution \citep[see, e.g.,162][]{p04b}: where is the first integral."163 Thus. in our approximate consideration the induced scattering results in the net intensitv transfer of the radio beam intensity into the background.," Thus, in our approximate consideration the induced scattering results in the net intensity transfer of the radio beam intensity into the background."164 Of course. an exact (reatment of the problem taking into account (he complete rather than approximate cross-sections would show that the total intensity. J somewhat decreases. (he energy being deposited to the scattering particles. and the number of photons »+n4 is conserved instead of the intensity. nz+mp.," Of course, an exact treatment of the problem taking into account the complete rather than approximate cross-sections would show that the total intensity $I$ somewhat decreases, the energy being deposited to the scattering particles, and the number of photons $n+n_1$ is conserved instead of the intensity $n\nu+n_1\nu_1$."165" One can see that as a significant. Traction of photons comes to the backeround state. nycn<n. the corresponding intensity is nip,en!pg?<n VU."," One can see that as a significant fraction of photons comes to the background state, $n_1\sim n^{(0)}\ll n$, the corresponding intensity is $n_1\nu_1\sim n^{(0)}\nu\theta^2 <n^{(0)}\nu$ ."166 Hence. of the original energy of the radio beam. nr. about 0?!0)p is transferred to the backeround state and ~(1—67?)0'p is gained bv the particles.," Hence, of the original energy of the radio beam, $n^{(0)}\nu$, about $\theta^2n^{(0)}\nu$ is transferred to the background state and $\sim (1-\theta^2)n^{(0)}\nu$ is gained by the particles."167 Although @<1. the background intensity grows drastically (ef.," Although $\theta\la 1$, the background intensity grows drastically (cf."168 eq.|13] below). and the intensity transfer between the two states greatlv dominates the evolution of the total intensity 7.," eq.[13] below), and the intensity transfer between the two states greatly dominates the evolution of the total intensity $I$."169 In our treatment. the latter is ignored and the intensity of the efficiently. growing component is intended to be roughly comparable with the original radio beam intensity or at least be above the detection level.," In our treatment, the latter is ignored and the intensity of the efficiently growing component is intended to be roughly comparable with the original radio beam intensity or at least be above the detection level."170According to equation (6). the efficiency of intensity transfer is characterized by the quantity P.—agr.,"According to equation (6), the efficiency of intensity transfer is characterized by the quantity $\Gamma =Iag^{ij}r$."171" Provided that P.>1 the background intensity grows exponentially. IEEEnuexp(I). but still remains much less than the original radio beam intensity. and. correspondingly. Z, IU."," Provided that $\Gamma\ga 1$ the background intensity grows exponentially, $I_{\nu_1}\approx I_{\nu_1}^{(0)}\exp (\Gamma)$, but still remains much less than the original radio beam intensity, and, correspondingly, $I_\nu\approx I_\nu^{(0)}$ ."172" A significant part of the beam intensity is (transferred (o the backeround on a more stringent. condition: At still larger D. the background intensity increases verv weakly. slowly approaching (he initialradio beam intensity. Z,,2f)(01(1— 1/£). whereas the beamintensity decreases"," A significant part of the beam intensity is transferred to the background on a more stringent condition: At still larger $\Gamma$ , the background intensity increases very weakly, slowly approaching the initialradio beam intensity, $I_{\nu_1}\approx I_\nu^{(0)}(1-1/\xi)$ , whereas the beamintensity decreases"173~ 53000-53500). while it was relatively large at around 2000 (MJD - 51800).,"$\sim$ 53000–53500), while it was relatively large at around 2000 (MJD $\sim$ 51800)."174 To quantify the level of activity. we defined all surface elements cooler than Ty=4200 K as spots.," To quantify the level of activity, we defined all surface elements cooler than $T_\mathrm{spot} = 4200$ K as spots."175 The spot coverage of each Doppler image was then estimated by calculating the percentage of the surface covered by elements with temperatures lower than Typo., The spot coverage of each Doppler image was then estimated by calculating the percentage of the surface covered by elements with temperatures lower than $T_\mathrm{spot}$.176 There is a clear decrease in the spot coverage. which explains the increase in the mean temperature occurring after 2002 (Fig. 5)).," There is a clear decrease in the spot coverage, which explains the increase in the mean temperature occurring after 2002 (Fig. \ref{activity}) )."177 We note that the photometric amplitudes derived by Jetsu et al. (, We note that the photometric amplitudes derived by Jetsu et al. (178in prep.),in prep.)179 also support the notion that the spot activity has decreased after 2002., also support the notion that the spot activity has decreased after 2002.180 A comparison with our earlier Doppler images (?) also shows that the spots were more dominant in the images from 1994-2002 than during 2004-2010., A comparison with our earlier Doppler images \citep{lindborg2011} also shows that the spots were more dominant in the images from 1994–2002 than during 2004–2010.181" In a recent paper (?).. we published 16 temperature maps for the star during 1994-2002. revealing short-term. irregular, “fip-flop’-type events and a systematic. drift of the active regions in the orbital reference frame of the binary system."," In a recent paper \citep{lindborg2011}, we published 16 temperature maps for the star during 1994–2002, revealing short-term, irregular, 'flip-flop'-type events and a systematic drift of the active regions in the orbital reference frame of the binary system."182 The 12 new temperature maps for 2004-2010 clearly show that the behaviour of the star is quite different from the earlier epoch: Several studies have found cyclic behaviour in, The 12 new temperature maps for 2004–2010 clearly show that the behaviour of the star is quite different from the earlier epoch: Several studies have found cyclic behaviour in.183Peg.. ? analysed 25 years of photometry and reported cycles of 13.5. 9.5. and 6.8 years in the spot activity.," \cite{rodono2000} analysed 25 years of photometry and reported cycles of 13.5, 9.5, and 6.8 years in the spot activity."184 Furthermore. ? reported à 4.65 year cycle in the “flip-flop” events.," Furthermore, \cite{berdyugina1999} reported a 4.65 year cycle in the 'flip-flop' events."185 We cannot confirm any regular flip-flops., We cannot confirm any regular flip-flops.186 The same conclusion was also drawn in the analysis of photometry by ?.., The same conclusion was also drawn in the analysis of photometry by \cite{roettenbacher2011}.187 We also see no evidence of a drift in. the spot-generating mechanism during 2004—2010 in the Doppler images., We also see no evidence of a drift in the spot-generating mechanism during 2004–2010 in the Doppler images.188 However. indications of a drift can be seen in the photometric minima during 2006-2010.," However, indications of a drift can be seen in the photometric minima during 2006–2010."189 We interpret the drift itself as a dynamo wave migrating in the azimuthal direction., We interpret the drift itself as a dynamo wave migrating in the azimuthal direction.190 The spots are thus generated by an underlying structure with a higher angular velocity than the surface of the star., The spots are thus generated by an underlying structure with a higher angular velocity than the surface of the star.191 It is clear that the star has entered a state of weaker activity than during 1994-2002., It is clear that the star has entered a state of weaker activity than during 1994–2002.192 The spot evolution seems fast and random. which could mean that a dynamo wave cannot be clearly detected.," The spot evolution seems fast and random, which could mean that a dynamo wave cannot be clearly detected."193 This could be related to à minimum im the star's eycle., This could be related to a minimum in the star's cycle.194 In this respect. the 13.5 year cycle found by ? i5 plausible.," In this respect, the 13.5 year cycle found by \cite{rodono2000} is plausible."195 A period of higher activity and a clearly detectable drifting active longitude. alternating with a period of lower activity. could then constitute the activity cycle.," A period of higher activity and a clearly detectable drifting active longitude, alternating with a period of lower activity, could then constitute the activity cycle."196 The collected time series of Doppler images. however. is still too short to make a decisive conclusion about the existence of such a stellar cycle.," The collected time series of Doppler images, however, is still too short to make a decisive conclusion about the existence of such a stellar cycle."197rapidly increasing impact of the (1|z)* cosmological dimming US Lq445445 exceeds unity.,rapidly increasing impact of the $(1+z)^4$ cosmological dimming as $z_{quasar}$ exceeds unity.198 The /v-band observations of Falomoetal.(2004) and. others involve exposure times equivalent to a factor ~+ larger than those presented here and peak surface brightnesses of the detected host galaxies are typically found to be fry20 7 for quasars that are not strong radio sources.," The $K$ -band observations of \citet{2004ApJ...604..495F} and others involve exposure times equivalent to a factor $\sim$ 4 larger than those presented here and peak surface brightnesses of the detected host galaxies are typically found to be $\mu_K \ga 20\,$ $^{-2}$ for quasars that are not strong radio sources."199 The almost complete lack of extended. approximately symmetric sources surrounding the quasar images is thus no surprising given the much higher surface brightness thresholds. Hasc21d 7. to which our imaging reaches.," The almost complete lack of extended, approximately symmetric sources surrounding the quasar images is thus not surprising given the much higher surface brightness thresholds, $\mu_K \simeq20021.1\,$ $^{-2}$, to which our imaging reaches."201 I is certainly possible that a number of small. faint. fy=20.0. residuals centred on the PSF-subtracted quasar images could be the higher surface brightness cores of host galaxies but the presence of any such hosts does not result in contamination of the catalogue of dy=20.0 objects that forms the basis for the statistical analysis of the absorber hosts.," It is certainly possible that a number of small, faint, $K > 20.0$, residuals centred on the PSF-subtracted quasar images could be the higher surface brightness cores of host galaxies but the presence of any such hosts does not result in contamination of the catalogue of $K \le 20.0$ objects that forms the basis for the statistical analysis of the absorber hosts."202 The discussion in the preceding sub-sections indicate that the census of galaxies with A'x20.0 is essentially complete to within 100 of the quasar images., The discussion in the preceding sub-sections indicate that the census of galaxies with $K \le 20.0$ is essentially complete to within 0 of the quasar images.203 Fainter than A~19.0. the census is substantially incomplete for galaxies with very small z05 separations from the quasar.," Fainter than $K \simeq 19.0$, the census is substantially incomplete for galaxies with very small $\la0\farcs5$ separations from the quasar."204 The corresponding spatial separations are small however. corresponding to scales Z4 Kkkpe.," The corresponding spatial separations are small however, corresponding to scales $\la$ kpc."205 The galaxy catalogue is also expected to suffer from little contamination by quasar host galaxies or other galaxies at the quasar redshifts. based both on prior expectations from larger. Av-band surveys and on he lack of any empirical evidence from the imaging observations hemselves.," The galaxy catalogue is also expected to suffer from little contamination by quasar host galaxies or other galaxies at the quasar redshifts, based both on prior expectations from larger $K$ -band surveys and on the lack of any empirical evidence from the imaging observations themselves."206 The statistics of the sample of galaxies presented in Table 2. should thus provide direct constraints on any population of galaxies associated with the absorbers that are bright enough ο be included in the Ax20.0 sample.," The statistics of the sample of galaxies presented in Table \ref{tab:galcat} should thus provide direct constraints on any population of galaxies associated with the absorbers that are bright enough to be included in the $K \le20720.0$ sample."208 The effectiveness of an imaging survey for constraining the oresenee of associated galaxies depends on: i) the number of argets. i) the surface density of unrelated galaxies and ii) the number of associated galaxies per target.," The effectiveness of an imaging survey for constraining the presence of associated galaxies depends on: i) the number of targets, ii) the surface density of unrelated galaxies and iii) the number of associated galaxies per target."209 Figure 2. presents the number counts of galaxies detected in our survey. as a function of separation from the absorber.," Figure \ref{fig:galex} presents the number counts of galaxies detected in our survey, as a function of separation from the absorber."210 On the right we show the observed number of galaxies within a given radius. divided by that predicted: the grey shading shows the associated. Poissonian errors for the magnitude limit of 20.," On the right we show the observed number of galaxies within a given radius, divided by that predicted; the grey shading shows the associated Poissonian errors for the magnitude limit of 20."211" For our observations. the angular scale on which such an excess population could be identified with any significance is 75-7"" and we chose a radius of 6700 for the presentation of the associated galaxy statistics."," For our observations, the angular scale on which such an excess population could be identified with any significance is $\simeq$ $^{\prime\prime}$ and we chose a radius of 0 for the presentation of the associated galaxy statistics."212 We have not attempted to maximise the significance of the excess by optimising the choice of magnitude limit and separation (e.g. quoting statistics for the excess at « 3700 with a magnitude limit of A'< 19.0)., We have not attempted to maximise the significance of the excess by optimising the choice of magnitude limit and separation (e.g. quoting statistics for the excess at $<$ 0 with a magnitude limit of $K\le19.0$ ).213 The existence of a galaxy overdensity close to the quasars with intervening absorbers is clear., The existence of a galaxy overdensity close to the quasars with intervening absorbers is clear.214 Table 5. summarises the number-counts within 6700 for the absorber and control quasars down to magnitude limits of A4.=19.0 and 20.0.," Table \ref{tab:galden}215 summarises the number-counts within 0 for the absorber and control quasars down to magnitude limits of $K=19.0$ and $20.0$."216 The predictec numbers of galaxies are calculated from the number counts of galaxies at 296/0 from the target quasar in all targe fields (Section ??))., The predicted numbers of galaxies are calculated from the number counts of galaxies at $\Delta\theta > 6\farcs0$ from the target quasar in all target fields (Section \ref{sec:nummag}) ).217 The numerical excess of galaxies observec is high. with observed:predicted numbers of 30:9.6 and 16:42 for magnitude limits of 46=20.0 and 19.0 respectively.," The numerical excess of galaxies observed is high, with observed:predicted numbers of 30:9.6 and 16:4.2 for magnitude limits of $K=20.0$ and $19.0$ respectively."218 Figure 2 illustrates the form of the observed excess as a function of separation., Figure \ref{fig:galex} illustrates the form of the observed excess as a function of separation.219 The quasars are widely separated on the sky and large-scale structure is not an issue but the small-scale clustering of galaxies may be expected to enhance the apparent overdensity., The quasars are widely separated on the sky and large-scale structure is not an issue but the small-scale clustering of galaxies may be expected to enhance the apparent overdensity.220 However. allowing onlyone galaxy within a 6700 radius of each absorber quasar to contribute to our number count of observec galaxies. gives observed:predicted numbers of 21:9.6 and |2:4.2 for magnitude limits of 47=20.0 and 19.0 respectively.," However, allowing only galaxy within a 0 radius of each absorber quasar to contribute to our number count of observed galaxies, gives observed:predicted numbers of 21:9.6 and 12:4.2 for magnitude limits of $K=20.0$ and $19.0$ respectively."221 The significance of the excess at both magnitudes is 23., The significance of the excess at both magnitudes is $\simeq$ $\sigma$.222 The statistics for the small control sample of quasars are entirely consistent with the predictions from the field A-band number counts., The statistics for the small control sample of quasars are entirely consistent with the predictions from the field $K$ -band number counts.223" The form of the local A-band galaxy luminosity function (GLF) is now well-established (Table4ofKochaneketal.2001) and we adopt the Schechter fit presented in this paper. with AJ;=23.39 (for H4,=100kkmss landa =—1.09. as our fiducial reference when considering the absolute magnitude distribution of the galaxies apparently associated with the absorbers."," The form of the local $K$ -band galaxy luminosity function (GLF) is now well-established \citep[Table 4 of][]{2001ApJ...560..566K} and we adopt the Schechter fit presented in this paper, with $M_K^*=-23.39$ (for $H_0=100$ $^{-1}$ ) and $\alpha=-1.09$, as our fiducial reference when considering the absolute magnitude distribution of the galaxies apparently associated with the absorbers."224 The A-band observations probe the rest-frame wavelength region 10 ας= 0.7) to ~13 (2= 1.2) where the galaxy light is dominated by stars., The $K$ -band observations probe the rest-frame wavelength region $\sim$ $z=0.7$ ) to $\sim$ $z=1.2$ ) where the galaxy light is dominated by K-giant stars.225 As a result. the K-corrections in the A-band are largely insensitive to the spectral type of the galaxies.," As a result, the K-corrections in the $K$ -band are largely insensitive to the spectral type of the galaxies."226 Galaxies observed at increasing redshift are expected to brighten. as their stellar populations become younger. and there is broad agreement between the predictions of theoretical models (e.g.Pozzetti.Bruzual.&Zamorani1996) and observations (Droryetal.2003:Saraccoetal. 20060).," Galaxies observed at increasing redshift are expected to brighten, as their stellar populations become younger, and there is broad agreement between the predictions of theoretical models \citep[e.g.][]{1996MNRAS.281..953P} and observations \citep{2003ApJ...595..698D, 2006MNRAS.367..349S}."227 To characterise the evolution we adopt a simple pure luminosity evolution model for the GLF. with the characteristic absolute magnitude evolving as ALi(ο.=ALR(s) 0.532. as determined by Droryetal.(2003)..," To characterise the evolution we adopt a simple pure luminosity evolution model for the GLF, with the characteristic absolute magnitude evolving as $M_K^*(z)=M_K^*(z=0)-0.53z$ , as determined by \citet{2003ApJ...595..698D}."228 Our observations oobe a magnitude range typically within £1.5 mag of Mj. and xossible variations in the slope of the faint end slope of the ΟΤΕ do not significantly affect our results.," Our observations probe a magnitude range typically within $\pm1.5\,$ mag of $M_K^*$ and possible variations in the slope of the faint end slope of the GLF do not significantly affect our results."229 Cols., Cols.230 2 and 3 of Table 4 oresents K-corrections as a function of redshift for galaxies of two spectral types.," 2 and 3 of Table \ref{tab:galkcorr}231 presents K-corrections as a function of redshift for galaxies of two spectral types."232 Absolute magnitudes before and after accounting or redshift evolution are given for objects with fy=18.0 and 20.0 in Cols., Absolute magnitudes before and after accounting for redshift evolution are given for objects with $K=18.0$ and $20.0$ in Cols.233 4. 5. 7 and 8.," 4, 5, 7 and 8."234 Finally. luminosities relative to £ are orovided in Cols.," Finally, luminosities relative to $L^*$ are provided in Cols."235 6 and 9., 6 and 9.236 The A -band at 2-1 approximates the -/- at 220.0 and the K-corrections and luminosity evolution are well-detined at the relatively long wavelengths probed by the A- observations., The $K$ -band at $z$$\simeq$ 1 approximates the $J$ -band at $z$ =0.0 and the K-corrections and luminosity evolution are well-defined at the relatively long wavelengths probed by the $K$ -band observations.237 Beginning with the ./-band luminosity function of Coleetal.(2001) and using the observed A-band ΟΤΕ at 21 rom Pozzettietal.(2003) results in luminosities for the galaxies. relative to L. that are essentially indistinguishable from those calculated above.," Beginning with the $J$ -band luminosity function of \citet{2001MNRAS.326..255C} and using the observed $K$ -band GLF at $z$$\simeq$ 1 from \citet{2003A&A...402..837P} results in luminosities for the galaxies, relative to $L^*$ , that are essentially indistinguishable from those calculated above."238eascous or conipact objects (Cary 1991. Gerhard Sills could explain. at least iu part. the dark ealactic haloes.,"gaseous or compact objects (Carr 1994, Gerhard Silk \nocite{Carr:1994,Gerhard:1996} could explain, at least in part, the dark galactic haloes."239 Tu 1986. proposed microleusing techuiques for measuriug the abundance of compact objects iu galactic haloes.," In 1986, \nocite{Paczynski:1986} proposed microlensing techniques for measuring the abundance of compact objects in galactic haloes."240 The LMC stars are favourable targets for nücroleusing events searches., The LMC stars are favourable targets for microlensing events searches.241 Since 1990 and 1992. the EROS (Aubourg et al.," Since 1990 and 1992, the EROS (Aubourg et al."242 aud ATACTIO (Alcock et al., \nocite{Aubourg:1993} and MACHO (Alcock et al.243 eroups have studied this line of sight., \nocite{Alcock:1993} groups have studied this line of sight.244 The detection of LO mucrolensing events has been claimed iu he huge mass rauge 0.05LAL. (Auboure et al., The detection of 10 microlensing events has been claimed in the large mass range $0.05 - 1 M_{\odot}$ (Aubourg et al.245 1993. Alcock et al. 1996).," 1993, Alcock et al. \nocite{Aubourg:1993,Alcock:1996}."246. This detection vate. smaller than expected with a full halo. indicates hat the most likely raction of conrpact «jects in the dark halo is f=0.5 (Alcock et al. 1996).," This detection rate, smaller than expected with a full halo, indicates that the most likely fraction of compact objects in the dark halo is $f = 0.5$ (Alcock et al. \nocite{Alcock:1996}."247. Coucurrenth. the siuill mass range jas been excluded for a wide range of ealactic models by he EROS and MACTIO eroups.," Concurrently, the small mass range has been excluded for a wide range of galactic models by the EROS and MACHO groups."248 Objects iu he nass range (510TAL.cM5<ΤΑ. could not account for nore than of the standard halo mass (Alcock et al. 1998).," Objects in the mass range $5\times 10^{-7} M_\odot < M249< 5 \times 10^{-4} M_\odot$ ) could not account for more than of the standard halo mass (Alcock et al. \nocite{Alcock:1998}."250. In the meantime. the DUO (Alard et al. 1995).," In the meantime, the DUO (Alard et al. \nocite{Alard:1995},"251 AFACTIO (Alcock al., MACHO (Alcock et al.252 and OCLE (Udalski et al., \nocite{Alcock:1995a} and OGLE (Udalski et al.253 eroups look- towards the galactic bulge where star- events are expected., \nocite{Udalski:1995b} groups look towards the galactic bulge where star-star events are expected.254 The detection rate is higher than expected from galactic models (see for iustauce Evans 1991. Alcock et al.," The detection rate is higher than expected from galactic models (see for instance Evans 1994, Alcock et al."255 1995. Stanek et al. 1997).," 1995, Stanek et al. \nocite{Evans:1994,Alcock:1995a,Stanek:1997}."256. The events detected in these two directious demonstrate the efficacy of the nicroleusing techniques based on the iionitoriug of several iillious of stars., The events detected in these two directions demonstrate the efficacy of the microlensing techniques based on the monitoring of several millions of stars.257 The detection of a larger nuniber of eveuts is one of the big challeuges in nicroleunsing searches., \nocite{Ansari:1997a} The detection of a larger number of events is one of the big challenges in microlensing searches.258 This basically requires the monitoring of a larecr nuuuber of stars., This basically requires the monitoring of a larger number of stars.259 The Pixel Method. initially presented by Baillou et al. (1993).," The Pixel Method, initially presented by Baillon et al. \nocite{Baillon:1993},"260. gives a new auswer to this problem: monitoriug pixel fluxes., gives a new answer to this problem: monitoring pixel fluxes.261 On nuages of galaxies. imost of the pixel fluxes come from unresolved stars. which contribute to the background fux.," On images of galaxies, most of the pixel fluxes come from unresolved stars, which contribute to the background flux."262 If oue of these stars is maguified by unicrolensing. the pixel fux will vary proportionally.," If one of these stars is magnified by microlensing, the pixel flux will vary proportionally."263 Such a huninosity variation can be detected above a given threshold. provided the magnification is large enough.," Such a luminosity variation can be detected above a given threshold, provided the magnification is large enough."264 Unlike other approaches (uamely star imouitoriug aud Differential huage Photometry. see below). the Pixel Methodnof perforin a photometry of the stars but is designed to achieve a hieh cfitcicncy for the detection of luminosity variations aecting umesolved stars.," Unlike other approaches (namely star monitoring and Differential Image Photometry, see below), the Pixel Method perform a photometry of the stars but is designed to achieve a high efficiency for the detection of luminosity variations affecting unresolved stars."265 This caus that we will work with pixel fluxes andnot with star fluxes., This means that we will work with pixel fluxes and with star fluxes.266 A theoretical study of the pixel lensing 1iethod has been published by Cotda (€1996b)., A theoretical study of the pixel lensing method has been published by Gould \nocite{Gould:1996}.267. This pixel mouitoring approach has two types of application., This pixel monitoring approach has two types of application.268 Firstly it allows us to investigate more distant galaxies aud thus to study other lines of sight.," Firstly, it allows us to investigate more distant galaxies and thus to study other lines of sight."269 This has led to observations of the M21 ealaxyv., This has led to observations of the M31 galaxy.270 The ACGAPE team (Ausari et al., The AGAPE team (Ansari et al.271 has shown that this method works on M21. data. aud huuinositv variations conipatible with the expected microlensing eveuts have been detected but the complete analysis is still in progress (Cüraud-Iérraud 1997).," \nocite{Ansari:1997a} has shown that this method works on M31 data, and luminosity variations compatible with the expected microlensing events have been detected but the complete analysis is still in progress (Giraud-Hérraud \nocite{YGH:1997}."272 Ao simular approach. though technically different. called Differcutial huage Photometry is also investigated by the VATT/Colmubia collaboration (Crotts 1992. Tomaney Crotts 1996).," A similar approach, though technically different, called Differential Image Photometry is also investigated by the VATT/Columbia collaboration (Crotts 1992, Tomaney Crotts \nocite{Crotts:1992,Tomaney:1996}."273. Some prospective work has also been doue towards ALS? (Gould 1995).., Some prospective work has also been done towards M87 (Gould \nocite{Gould:1995b}.274 The second possibility is to apply pixel mucrolensing on existing data. thus exteudiug the sensitivity of previous analyses to unresolved stars.," The second possibility is to apply pixel microlensing on existing data, thus extending the sensitivity of previous analyses to unresolved stars."275 This is precisely the subject of this paper aud of the two which will follow: we present the implementation of the Pixel Method on CCD images of the LMC., This is precisely the subject of this paper and of the two which will follow: we present the implementation of the Pixel Method on CCD images of the LMC.276 We have applied for the first time a comprehensive pixel analysis ou existing LAIC images collected by the EROS collaboration., We have applied for the first time a comprehensive pixel analysis on existing LMC images collected by the EROS collaboration.277 With respect to previous analyses (Queinnec 1991. Aubourg et al.," With respect to previous analyses (Queinnec 1994, Aubourg et al."278 1995. Renault 1996).. our analysis of the saue data using pixel monitonrus allows us to extend the mass range of iutercst up to LAL... and to increase the sensitivity of microlensing searches.," 1995, Renault \nocite{Queinnec:1994,Aubourg:1995,Renault:1996}, our analysis of the same data using pixel monitoring allows us to extend the mass range of interest up to $1\, M_\odot$ and to increase the sensitivity of microlensing searches."279 On these nuages. a large fraction of ae stars retains unresolved: typically 5 to 10 stars conladusibute to of the pixel flux iu one square arc-second.," On these images, a large fraction of the stars remains unresolved: typically 5 to 10 stars contribute to of the pixel flux in one square arc-second."280 Since this approach potentially uses all the inage coutent (aud not ouly the resolved stars}. the volume of the data to haudle is wich lareer.," Since this approach potentially uses all the image content (and not only the resolved stars), the volume of the data to handle is much larger."281 Weuce we perform this first exploratory analysis- on a relativelyH stall data set: 0.25E dee?JD covering: a period of observation of 120 davs. which correspouds to ofthe LAIC CCD data (91-9D).," Hence we perform this first exploratory analysis on a relatively small data set: 0.25 $^2$ covering a period of observation of 120 days, which corresponds to of the LMC CCD data (91-94)."282 This paper is the first of a series of three. describing the data treatineut (this paper). the microlensing search (Melchior et al.," This paper is the first of a series of three, describing the data treatment (this paper), the microlensing search (Melchior et al."283 1998a. hereafter Paper IT) aud a catalogue of variable stars (Melchior ct al.," 1998a, hereafter Paper II) and a catalogue of variable stars (Melchior et al."284 1998b. hereafter Paper IIT.," 1998b, hereafter Paper III)."285 Iu the companion papers (Papers IT aud WT). we show how the data treatinent described here to produce pixel helt curves allows us to perform analyses that iucrease the scusitivity to microlensiug eveuts and variable stars with respect to the star monitoring analysis applied ou the same field: an order of magnitude in the πο of detectable luminosity variations ix gained.," In the companion papers (Papers II and III), we show how the data treatment described here to produce pixel light curves allows us to perform analyses that increase the sensitivity to microlensing events and variable stars with respect to the star monitoring analysis applied on the same field: an order of magnitude in the number of detectable luminosity variations is gained."286 To discover real variations. the nuages and lieht curves have to be corrected for various sources of fake variabilities. such as geometrical aud plotometric nüsmatehn. or secine changes between successive Haiages.," To discover real variations, the images and light curves have to be corrected for various sources of fake variabilities, such as geometrical and photometric mismatch, or seeing changes between successive images."287 The construction of light curves cleaned frou these effects is the subject of this first paper., The construction of light curves cleaned from these effects is the subject of this first paper.288 If the flux of a eiven star contributes to the pixel fiux. the latter cau be expressed as follows: where oque ds the flux of the given star. f the fraction of the star flux that enters the pixel. hereafter called," If the flux of a given star contributes to the pixel flux, the latter can be expressed as follows: where $\phi_{\rm star}$ is the flux of the given star, $f$ the fraction of the star flux that enters the pixel, hereafter called"289a conipact steep-spectrmm (CSS) quasar.,a compact steep-spectrum (CSS) quasar.290" The CSSs are defined to be of subealactic dimensions. with an overall size S20 kpc and a steep ligh-frequency spectrum having a 20.5. where Sx»""."," The CSSs are defined to be of subgalactic dimensions, with an overall size $\lapp$ 20 kpc and a steep high-frequency spectrum having $\alpha\geq$ 0.5, where $\propto\nu^{-\alpha}$."291 They are believed to be voune sources at an carly stage of evolution (Fanti et al., They are believed to be young sources at an early stage of evolution (Fanti et al.292 1995: Readhead et al., 1995; Readhead et al.293 1996a.b: ODea 1998: Owsiauik Couway 1998).," 1996a,b; O'Dea 1998; Owsianik Conway 1998)."294 We present below the results of our observations of the CSS quasar 136 and discuss their implications., We present below the results of our observations of the CSS quasar $-$ 136 and discuss their implications.295" The radio source 1 136 is associated with a quasar at aredshift of 1.687. so that 1"" correspouds to δ.LL kpc in a Universe with IL,2 50 kins ! loud qo=0.5."," The radio source $-$ 136 is associated with a quasar at a redshift of 1.687, so that $^{\prime\prime}$ corresponds to 8.44 kpc in a Universe with $_o$ = 50 km $^{-1}$ $^{-1}$ and $_o$ =0.5."296 It has a steep radio spectrum (o299 zz 0.61) which appears to flatten below a few hundred MITZz (cf., It has a steep radio spectrum $\alpha^{5000}_{750}$ $\approx$ 0.64) which appears to flatten below a few hundred MHz (cf.297 Steppe et al., Steppe et al.298 1995)., 1995).299 Tt is a well-known ow-frequency variable source (Bondi et al., It is a well-known low-frequency variable source (Bondi et al.300 1996 aud references therein). although it does nof exhibit significant variations at cni waveleugths (Pacdriclli et al.," 1996 and references therein), although it does not exhibit significant variations at cm wavelengths (Padrielli et al."301 1987)., 1987).302 VLA observations at 5 GIIz find 136 to be slightly exteude while at 15 GIIz it is resolved iuto a double source (Mantovani et al.," VLA observations at 5 GHz find $-$ 136 to be slightly extended, while at 15 GHz it is resolved into a double source (Mantovani et al."303 1991)., 1994).304 Interplauctary sclutillation observations at 327 ΛΠΙ show that about 50 per cent of the total flux density is from a compact structure with a size of Z100 millarcsee (Jevakuuiar ct al., Interplanetary scintillation observations at 327 MHz show that about 50 per cent of the total flux density is from a compact structure with a size of $\lapp$ 100 milliarcsec (Jeyakumar et al.305 2000)., 2000).306 The source was observec with MERLIN onu 1995 May 06 at 5 CIIz for zSN hr using a bandwidth of 28 AIIIz., The source was observed with MERLIN on 1995 May 06 at 5 GHz for $\approx$ 8 hr using a bandwidth of 28 MHz.307 Initial values for the telescope aud correlator gaius were determined from a short observation of the bright unresolved calibrator BO552|398 (DA193)., Initial values for the telescope and correlator gains were determined from a short observation of the bright unresolved calibrator B0552+398 (DA193).308 The primary flux deusity and polarization calibrator was 3€286 aud the secondary calibrator was | 100., The primary flux density and polarization calibrator was 3C286 and the secondary calibrator was $+$ 100.309 The global VLBI MI2 observations were made on 1992 September 21 with a ceutral frequeney of 1663.99 NITz aud a 2-MIIzZ baucdsvidth., The global VLBI MK2 observations were made on 1992 September 21 with a central frequency of 1663.99 MHz and a 2-MHz bandwidth.310 The array consisted of auteunas at Ousala. Effelsberg. Westerbork. Jodrell Dauk (the Lovell telescope). Medicina. Creeu Bank. Tavstack. VLBA-KP. VLBA-PT. VLBA-LA. VLBA-NL. VLBA-BR and VLBA-OV.," The array consisted of antennas at Onsala, Effelsberg, Westerbork, Jodrell Bank (the Lovell telescope), Medicina, Green Bank, Haystack, VLBA-KP, VLBA-PT, VLBA-LA, VLBA-NL, VLBA-BR and VLBA-OV."311 The source was tracked for a total of 7.5 hr in a shap-shot observing progranuuc., The source was tracked for a total of 7.5 hr in a snap-shot observing programme.312 The data were processed at the Caltech Block? correlator., The data were processed at the Caltech Block2 correlator.313 These observations were carried out on 1996 Feb 01 with the Verv Long Baseline Árrav (VLBA) aud one antenna of the VLA recording both right- aud left-circular polarization (RCP. LCP) using one-hit sampling.," These observations were carried out on 1996 Feb 04 with the Very Long Baseline Array (VLBA) and one antenna of the VLA recording both right- and left-circular polarization (RCP, LCP) using one-bit sampling."314 The source was observed at 8.1 CGIIz for 1.3 hr with a bandwidth of 32 ΠΕ., The source was observed at 8.4 GHz for 4.3 hr with a bandwidth of 32 MHz.315 Amplitude calibration was derived from ineastrements of the svstemi temperature during the observations aud the knowu anteuna eain of each clement of the array., Amplitude calibration was derived from measurements of the system temperature during the observations and the known antenna gain of each element of the array.316 Complex correlation coefficients were recovered at the Array Operations Center in Socorro., Complex correlation coefficients were recovered at the Array Operations Center in Socorro.317 Phase eradieuts iu frequency and time were corrected by the use of global fringe fitting (Schwab Cotton 1983)., Phase gradients in frequency and time were corrected by the use of global fringe fitting (Schwab Cotton 1983).318 Earlicr VLA A-array observations of 136 at 5 6112 show the source to be sheltly extended. while the 15-CGIIz observations show it to be clearly double-Iobed with an angular size of 0.32 arcsec along a PA of 307.," Earlier VLA A-array observations of $-$ 136 at 5 GHz show the source to be slightly extended, while the 15-GHz observations show it to be clearly double-lobed with an angular size of 0.32 arcsec along a PA of $^\circ$."319 The weaker southern compoucut is 1.3 por cent polarized at 15 CGIIz compared to 2.5 per cout for the northern component (Mautovani ct al., The weaker southern component is 4.3 per cent polarized at 15 GHz compared to 2.5 per cent for the northern component (Mantovani et al.320 1991)., 1994).321 Beceut polarization observations with the VLBA show the northern jet and hot-spot to have very hieh rotation measures (RAIS) in the range of 3000 το 10.000 τας 7 in the rest frame of the source (Mantovani ot al.," Recent polarization observations with the VLBA show the northern jet and hot-spot to have very high rotation measures (RMs) in the range of $-$ 3000 to $-$ 10,000 rad $^{-2}$ in the rest frame of the source (Mantovani et al."322 2002)., 2002).323 The MERLIN 5-CGITz mage with au angular resolution of 11 nas along a PA of 167 and an rus noise of 0.3 11Jv/beaa is shown in Fie., The MERLIN 5-GHz image with an angular resolution of $\times$ 41 mas along a PA of $^\circ$ and an rms noise of 0.3 mJy/beam is shown in Fig.324 1., 1.325 This image is of higher resolution than the VLA images aud reveals ereater details of the structure., This image is of higher resolution than the VLA images and reveals greater details of the structure.326 The two main componcuts. northern aud souther are clearly seen and are Ll and 7.9 per cent polarized respectively.," The two main components, northern and southern, are clearly seen and are 4.4 and 7.9 per cent polarized respectively."327 The peak brightuess of the northeru, The peak brightness of the northern328"The MFGalactic magnetic field has two components: a regular component B, and a random component b, so that the total Galactic magnetic field is given as B=B,+6 (forreviewsee??)..","The Galactic magnetic field has two components: a regular component $\vec{B_{r}}$ and a random component $\vec{b}$, so that the total Galactic magnetic field is given as $\vec{B}=\vec{B_r}+\vec{b}$ \citep[for review see][]{beck96, han02}."329" The regular component is usually simulated as a combination of a disk and a halo field, whereas the random field component is simulated as a Gaussian random field, Gaussian random field, (for ??).."," The regular component is usually simulated as a combination of a disk and a halo field, whereas the random field component is simulated as a Gaussian random field, Gaussian random field, \citep[for details see][]{sun08, sun09}."330" Note that for our calculations, we split Bina component parallel (Bj) and perpendicular (11) to the LOS, so that Faraday rotation is defined by By and synchrotron emission by B."," Note that for our calculations, we split $\vec{B}$ in a component parallel $B_{\parallel}$ ) and perpendicular $B_{\perp}$ ) to the LOS, so that Faraday rotation is defined by $B_{\parallel}$ and synchrotron emission by $B_{\perp}$."331" Considering the aim of our effort to simulate the Galactic emission for a small patch of the sky, we treat the regular field component in a simplified way."," Considering the aim of our effort to simulate the Galactic emission for a small patch of the sky, we treat the regular field component in a simplified way."332 The regular field component is assumed to be uniform in the xy-plane and to have an exponential decrease in the z direction., The regular field component is assumed to be uniform in the $xy$ -plane and to have an exponential decrease in the $z$ direction.333 The typical value of the regular field component is a few µία (forreviewsee??)..," The typical value of the regular field component is a few $\mu334G$ \citep[for review see][]{beck96, han02}."335 For the random field component we follow ?? and simulate it as a Gaussian random field.," For the random field component we follow \citet{sun08, sun09} and simulate it as a Gaussian random field."336" The power spectrum of the field follows a power law, with spectral index —8/3."," The power spectrum of the field follows a power law, with spectral index $-8/3$."337 This spectral index is commonly used for a Kolmogorov-like turbulence spectrum., This spectral index is commonly used for a Kolmogorov-like turbulence spectrum.338" In our simulation, the realization. of the random field component is done in the following way."," In our simulation, the realization of the random field component is done in the following way."339" First we generate three different Gaussian random fields for the bz, ὂν and b; component."," First we generate three different Gaussian random fields for the $b_x$, $b_y$ and $b_z$ component."340" From those three fields, we then calculatethe amplitude of b and normalize it to the desired value."," From those three fields, we then calculatethe amplitude of $\vec{b}$ and normalize it to the desired value."341 A typical value for the mean random field strength is b=3wG (?).., A typical value for the mean random field strength is $b=3~{\rm \mu G}$ \citep{sun08}.342" At high Galactic latitudes, the warm ionized medium consists mostly of diffuse ionized gas (DIG) with total emission measure of -5pccm? and T,=8000K (?).."," At high Galactic latitudes, the warm ionized medium consists mostly of diffuse ionized gas (DIG) with total emission measure of $\sim3435~{\rm pc~cm^{-6}}$ and $T_e=8000~{\rm K}$ \citep{reynolds90}."344 The properties of the DIG can be traced by its free-free emission and dispersion (DM) of pulsars (e.g.?).., The properties of the DIG can be traced by its free-free emission and dispersion (DM) of pulsars \citep[e.g.][]{gaensler08}.345 Recent simulations of the Galactic emission (???) used the ? model for the thermal electron distribution.," Recent simulations of the Galactic emission \citep{sun08,346waelkens09, sun09} used the \citet{cordes02} model for the thermal electron distribution."347" That model simulates the Galaxy as several large-scale (e.g. thin and thick disk, and spiral arms) and small-scale (e.g. supernovae bubbles) structures."," That model simulates the Galaxy as several large-scale (e.g. thin and thick disk, and spiral arms) and small-scale (e.g. supernovae bubbles) structures."348" In our simulation, we follow our previous model of the Galactic free-free emission (?) and simulate the thermal electron density distribution as a Gaussian random field with the power law type of the spectrum."," In our simulation, we follow our previous model of the Galactic free-free emission \citep{jelic08} and simulate the thermal electron density distribution as a Gaussian random field with the power law type of the spectrum."349 The spectral index of the power law is —3., The spectral index of the power law is $-3$.350 The amplitude of the Gaussian random field (thermal electron density) is normalized in a way to match the typical observed EM of the quasars at high Galactic latitudes (EMvaluesaretakenfrom?)..," The amplitude of the Gaussian random field (thermal electron density) is normalized in a way to match the typical observed EM of the quasars at high Galactic latitudes \citep[EM values are taken351from][]{berkhuijsen06}."352" It is important to note that our model is flexible to include additional features of the thermal electron distribution, e.g. dense bubbles or clumpy distribution."," It is important to note that our model is flexible to include additional features of the thermal electron distribution, e.g. dense bubbles or clumpy distribution."353 Some of these features are presented in Sec. ??.., Some of these features are presented in Sec. \ref{sec:ex}.354 Here we summarize the steps we follow to obtain maps of the Galactic emission at a desired frequency., Here we summarize the steps we follow to obtain maps of the Galactic emission at a desired frequency.355 The flow chart of the algorithm is presented in Fig. 1.., The flow chart of the algorithm is presented in Fig. \ref{fig:fc}. .356 In the following section we will show some examples of the Galactic emission maps obtained by this algorithm., In the following section we will show some examples of the Galactic emission maps obtained by this algorithm.357 7.. ὃν and οι.,"\cite{barlow}, , \cite{anupama}, , and \cite{evans}."358 A high helium abundance would suggest tha he secondary in U Sco is a highly evolved star. however the value derived. here is close to the solar value NOLO)ZN(LI) = 0.085 (?)..," A high helium abundance would suggest that the secondary in U Sco is a highly evolved star, however the value derived here is close to the solar value N(He)/N(H) = 0.085 \citep{asplund}."359 Our estimate suggests that assumptions abou he helium-rich nature of the secondary are unfouncded arc hat the secondary did not acerete helium-rich materia significantly in the post common envelope phase (2).., Our estimate suggests that assumptions about the helium-rich nature of the secondary are unfounded and that the secondary did not accrete helium-rich material significantly in the post common envelope phase \citep{hachisukato}.360 In the evolutionary calculations of 2.. U Sco is best. represente ον their sequence. D. with a hvdrogen-rich. secondary.," In the evolutionary calculations of \cite{sarna}, U Sco is best represented by their sequence B, with a hydrogen-rich secondary."361 For our helium abundance analysis we have considered the data (day S) and the lines for which we can be confident tha case D recombination applies., For our helium abundance analysis we have considered the data $\ge$ day 8) and the lines for which we can be confident that case B recombination applies.362 Other estimates have been arecr than our value. especially those based only on the ines.," Other estimates have been larger than our value, especially those based only on the lines."363 These large abundances result from the use of the Saha equation. which is highly sensitive to the assumed electron emperature.," These large abundances result from the use of the Saha equation, which is highly sensitive to the assumed electron temperature."364 However. we use both and lines in our analvsis. avoiding the use of the equation to derive the otal helium abundance.," However, we use both and lines in our analysis, avoiding the use of the equation to derive the total helium abundance."365 The temperature dependencies in his analysis are in the recombination cocllicients. which are relatively insensitive to the value chosen. and in the collisional correction for lines.," The temperature dependencies in this analysis are in the recombination coefficients, which are relatively insensitive to the value chosen, and in the collisional correction for lines."366 Using a range of lines rom days Slo and 9.439 and a range of recombination coellicients we determine several ratios. cach of which suggest electron temperatures of 2.3.2.43.10 Ih. herefore we use recombination coellicients for the closes empoerature from ? of 1.=2107 KK. The hyerogen line profiles have broad. and narrow components.," Using a range of lines from days 8.81 and 9.43 and a range of recombination coefficients we determine several $^{++}$ $^{+}$ ratios, each of which suggest electron temperatures of $2.3-2.4 \times 10^4$ K, therefore we use recombination coefficients for the closest temperature from \cite{hummer} of $T_e =2 \times 10^4$ K. The hydrogen line profiles have broad and narrow components."367 The narrow components become prominen arouncl day S and arise from slow moving material at speeds of ~1000 kkmis. !., The narrow components become prominent around day 8 and arise from slow moving material at speeds of $\sim1000$ $^{-1}$ .368" ""Phe broad. components are composed of material travelling at. velocities of up to 10.000 "," The broad components are composed of material travelling at velocities of up to $\sim10,000$ $^{-1}$."369The rise in the strength of the narrow components coulc be explained by optical depth ellects., The rise in the strength of the narrow components could be explained by optical depth effects.370 “Phe ejecta begins to become optically thin around day S: before this time slower moving material would be concealed by the shell of optically thick fast moving material., The ejecta begins to become optically thin around day 8; before this time slower moving material would be concealed by the shell of optically thick fast moving material.371 As the faster moving material expands anc becomes optically thin the Dux from the slow moving material becomes visible., As the faster moving material expands and becomes optically thin the flux from the slow moving material becomes visible.372 lines also begin to dominate the helium emission around dav S. however this is most likely due to photo-ionisation of helium rather than an optical depth elfect.," lines also begin to dominate the helium emission around day 8, however this is most likely due to photo-ionisation of helium rather than an optical depth effect."373 Phe measured line widths are generally higher for the SAAQO data than for near-simultaneous LT data: possible reasons for this are the dillerence in resolution and signal to noise issues., The measured line widths are generally higher for the SAAO data than for near-simultaneous LT data; possible reasons for this are the difference in resolution and signal to noise issues.374 “Phere also appears to be clear structure to the LL? line in the LT spectra which is not present in the SAAQO spectra., There also appears to be clear structure to the $\beta$ line in the LT spectra which is not present in the SAAO spectra.375 Ao possible source. of the narrow components is an inclination cllect: since U Seo is a high inclination svstem most of the velocity is in the plane of the sky. therefore we may be observing a small racial component.," A possible source of the narrow components is an inclination effect; since U Sco is a high inclination system most of the velocity is in the plane of the sky, therefore we may be observing a small radial component."376 This is consistent with the asymmetric ejecta models of ?.., This is consistent with the asymmetric ejecta models of \cite{drake}.377 Data over a longer time range ancl with better orbital phase coverage is required to explore this further., Data over a longer time range and with better orbital phase coverage is required to explore this further.378 We present optical and near-infrared observations of the 2010 outburst of U Scorpii., We present optical and near-infrared observations of the 2010 outburst of U Scorpii.379 We find the helium abundance of U Scorpii to be IN(IIG)/ZN(IID)2.0.073+0.031., We find the helium abundance of U Scorpii to be $N$ $N$ $ = 0.073\pm0.031$.380 This estimate is lower than most previous studies ancl does not support their conclusions which suggest that the secondary in this svstem is helium-rich., This estimate is lower than most previous studies and does not support their conclusions which suggest that the secondary in this system is helium-rich.381 The velocities seen in the 2010 outburst are consistent with those seen in previous outbursts. with some hydrogen and lines being seen to have both narrow and broad components in our later spectra.," The velocities seen in the 2010 outburst are consistent with those seen in previous outbursts, with some hydrogen and lines being seen to have both narrow and broad components in our later spectra."382 Further observations are planned to investigate the nature and metallicity of the secondary in this svstem., Further observations are planned to investigate the nature and metallicity of the secondary in this system.383 The Liverpool Telescope is operated. on the island. of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias with financial support from the Ulx Science and. Technology. Facilities Council., The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias with financial support from the UK Science and Technology Facilities Council.384" ALBININ. was supported by the Peter and Patricia Gruber Foundation through the LAU-PPGE fellowship. by the Peking University One όνος, Talent. Fund. (985). and by the National Natural Science Foundation of China (grants 11010237 and 11043007)."," M.B.N.K. was supported by the Peter and Patricia Gruber Foundation through the IAU-PPGF fellowship, by the Peking University One Hundred Talent Fund (985), and by the National Natural Science Foundation of China (grants 11010237 and 11043007)."385 This paper uses observations mace at the South African Astronomical Observatory (SAAQ)., This paper uses observations made at the South African Astronomical Observatory (SAAO).386 Pheauthors would like to thank theanonymousreferee for their helpful comments., Theauthors would like to thank theanonymousreferee for their helpful comments.387scale.,scale.388 y-values between -0.2 and -0.3 were found., $\gamma$ -values between -0.2 and -0.3 were found.389 This agrees with Sakaietal.(2004)... who related the difference between TRGB and Cepheid distances to the galactic region metallicities (not taking into account metallicity gradients. though) and derived a similar PLR dependence on metallicity.," This agrees with \citet{sakai04}, who related the difference between TRGB and Cepheid distances to the galactic region metallicities (not taking into account metallicity gradients, though) and derived a similar PLR dependence on metallicity."390 However. these results are highly uncertain.," However, these results are highly uncertain."391 Rizzietal.(2007) have argued that many of the TRGB distances used by Sakaietal.(2004) need to be revised., \citet{rizzi07} have argued that many of the TRGB distances used by \citet{sakai04} need to be revised.392 With the Rizzietal.(2007) TRGB distances the Sakaetal.(2004) dependence of the PLR on metallicity disappears and the results are in much closer agreement with stellar pulsation theory (Fiorentinoetal.2002.2007:Marconi2005:Bonoetal. 2008).. which predicts just the opposite. namely thatmetallicity.," With the \citet{rizzi07} TRGB distances the \citet{sakai04} dependence of the PLR on metallicity disappears and the results are in much closer agreement with stellar pulsation theory \citep{fiorentino02, fiorentino07, marconi05, bono08}, which predicts just the opposite, namely that."393 It is very important to note that all the region (oxygen) metallicities adopted. when comparing Cepheid distances with TRGB distances or when using metallicity gradients. are highly uncertain.," It is very important to note that all the region (oxygen) metallicities adopted, when comparing Cepheid distances with TRGB distances or when using metallicity gradients, are highly uncertain."394" They result from the of a simplified analysis method. the so called ""strong-line method"". which uses only the fluxes of the strongest forbidden lines of (most commonly) [OIL] and [OIL] relative to Hg."," They result from the application of a simplified analysis method, the so called ``strong-line method”, which uses only the fluxes of the strongest forbidden lines of (most commonly) [OII] and [OIII] relative to $_{\beta}$."395 Unfortunately. abundances obtained with the strong-line method depend heavily on the calibration used.," Unfortunately, abundances obtained with the strong-line method depend heavily on the calibration used."396 As a striking example. Kewley&Ellison(2008) have demonstratedll that the quantitative shape of the mass-metallicity relationshipli of galaxies can change from very steep to almost flat depending on the calibration used reffig+)).," As a striking example, \citet{kewley08}397 have demonstrated that the quantitative shape of the mass-metallicity relationship of galaxies can change from very steep to almost flat depending on the calibration used \\ref{fig4}) )."398 In the same way. as shown by Kudritzkietal.(2008) and Bresolinetal.(2009).. metallicity gradients of spiral galaxies can change from steep to flat and absolute values of metallicity can shift by as much as 0.6 dex. again as the result of different calibrations of the strong line method reftig5)).," In the same way, as shown by \citet{kud08} and \citet{bresolin09}, metallicity gradients of spiral galaxies can change from steep to flat and absolute values of metallicity can shift by as much as 0.6 dex, again as the result of different calibrations of the strong line method \\ref{fig5}) )."399 The Cepheid work on spiral galaxies usually uses the calibration by Zaritskyetal.(1994)., The Cepheid work on spiral galaxies usually uses the calibration by \citet{zaritsky94}.400.. As shown by Kudritzkietal.(2008).. this calibration gives far too high metallicities.," As shown by \citet{kud08}, this calibration gives far too high metallicities."401 The results displayed in reffigt and 5 are shocking., The results displayed in \\ref{fig4} and \ref{fig5} are shocking.402 They reveal that at present galaxy metallicities based on strong-line methods are uncertain by 0.6 to 0.8 dex because of the systematic uncertainties applicationnherent in the strong line methods used., They reveal that at present galaxy metallicities based on strong-line methods are uncertain by 0.6 to 0.8 dex because of the systematic uncertainties inherent in the strong line methods used.403 Even with a strictly differential approach in the use of Cepheids as distance indicators as applied by Riessetal.(2011). this PAtill creates residual uncertainties. which become important in view of the ambitious goal to determine Hy with an accuracy of a few percent.," Even with a strictly differential approach in the use of Cepheids as distance indicators as applied by \citet{riess11} this still creates residual uncertainties, which become important in view of the ambitious goal to determine $_{0}$ with an accuracy of a few percent."404" This major problem requires a c,resh approach and is begging for the development of a ,ew and independent method less affected by systematic uncertainties.", This major problem requires a fresh approach and is begging for the development of a new and independent method less affected by systematic uncertainties.405" Very obviously. in order to better assess the systematic. uncertainties of the determination. of distances to star-forming spiral and irregular galaxies in the local universe an independent and complementary method is ""Sesirable. which can overcome the problems of interstellar extinction and variations of chemical composition."," Very obviously, in order to better assess the systematic uncertainties of the determination of distances to star-forming spiral and irregular galaxies in the local universe an independent and complementary method is desirable, which can overcome the problems of interstellar extinction and variations of chemical composition."406 In the ext section we introduce such a method. which is based on the quantitative spectroscopy of BSGs.," In the next section we introduce such a method, which is based on the quantitative spectroscopy of BSGs."407 The method makes use of the enormous intrinsic brightness of these objects., The method makes use of the enormous intrinsic brightness of these objects.408processes could dramatically chauee he mass. orbital parameters. composition aud nmuuber of the planets.,"processes could dramatically change the mass, orbital parameters, composition and number of the planets."409 We remark that a related. bii ess radical suggestion has been made by Ribas (2007)... who propose that ile largest extra-solar planets (exoplauets) may have formed by direct gravitational collapse.," We remark that a related, but less radical suggestion has been made by \cite{Rib+07}, who propose that the largest extra-solar planets (exoplanets) may have formed by direct gravitational collapse."410 TIev poiut out that this mechanism cau explain the large masses aud high eccentricities of many ofd ie. recently discovered. exoplauets (Butler. 2006)., They point out that this mechanism can explain the large masses and high eccentricities of many of the recently discovered exoplanets \citep{But+06}.411. By contrast. we propose tla lie concurrent collapse mechauisin could be the origin of all planetary systems.," By contrast, we propose that the concurrent collapse mechanism could be the origin of all planetary systems."412 Other theorles LLvolving direct raviatioial collapse of planets have been proposed in tle past: we meution 11ese in section 3.. contrastug tliem with our own hypothesis.," Other theories involving direct gravitational collapse of planets have been proposed in the past: we mention these in section \ref{sec: 3}, contrasting them with our own hypothesis."413 The consequeuces of our new theory are uucertaiu. but we argue tiat d has the potential to resolve all of the difficulties with the standard «ust-aMDooregatiou model.," The consequences of our new theory are uncertain, but we argue that it has the potential to resolve all of the difficulties with the standard dust-aggregation model."414 Becaise we liypotliesise that planetary systems developed from very diffe'eut. initial. couclitious. qtantitative modelling of the processes which we cousider is an open-euced problem.," Because we hypothesise that planetary systems developed from very different initial conditions, quantitative modelling of the processes which we consider is an open-ended problem."415 However. prelimiiary quantitative estimates of the processes which we describe using dlausible initial data viek verv encouraging results.," However, preliminary quantitative estimates of the processes which we describe using plausible initial data yield very encouraging results."416 2.1exoplanets The discovery ofsubstantia uumbers of extra-solar planeS (ButleWright&Marey2006) has produced observations whicl Challenge the standard moel., 2.1 The discovery ofsubstantial numbers of extra-solar planets \citep{But+06} has produced observations which challenge the standard model.417 Significant numbers of these exoplanets have large orbit eccentricity., Significant numbers of these exoplanets have large orbital eccentricity.418 Various models have been proposed to account [o mis (Zakaiinka&Tremaine2()0D.., Various models have been proposed to account for this \citep{Zak+04}.419 The most plausible of tlese Is a slow-actine three-body instabilityab resuling iu a drift of orl)ital paraueters. leadiug to a between two planets.," The most plausible of these is a slow-acting three-body instability resulting in a drift of orbital parameters, leading to a near-collision between two planets."420 TThis could «'ause escape of one mjanet aud scattering of the other to ali eccentric Guxl probably 1on-equatoria) orbit (Ford.Basiο&Yu2003)., This could cause escape of one planet and scattering of the other to an eccentric (and probably non-equatorial) orbit \citep{For+03}.421. It is as yet not clear whetier the large proportion ol exoplauets with eccentric orbis can be explaited by this moclel., It is as yet not clear whether the large proportion of exoplanets with eccentric orbits can be explained by this model.422 The uodel would suggest that large pleatels are less likely to be scattered iuto highly eccentric orbits., The model would suggest that large planets are less likely to be scattered into highly eccentric orbits.423 Tiere seetus. Lowever. to be a posltive correlation between ecceutricity aud mass 2007).," There seems, however, to be a positive correlation between eccentricity and mass \citep{Rib+07}."424. There are also exauples of plauets where the orbital plaue is at a very large augle of iuclinatiou to the equator of the star: see for example Jolius-Ixrulefal(2008)., There are also examples of planets where the orbital plane is at a very large angle of inclination to the equator of the star: see for example \cite{Joh+08}.425. lauy of the exoplauets are found to have very high masses. which suggests that the circuustellar accretion disk woild have to have a much hieher mass than is usually assumed.," Many of the exoplanets are found to have very high masses, which suggests that the circumstellar accretion disk would have to have a much higher mass than is usually assumed."426 A related difficulty is the observation of substautial uumbers of exoplauets which are gas-giaus orbiting close to their star. at raclii whe'e the temperature would| be too high to allow them to formsitu.," A related difficulty is the observation of substantial numbers of exoplanets which are gas-giants orbiting close to their star, at radii where the temperature would be too high to allow them to form."427 It is inferred that these have crifted closer o their star after their formation. but the mechanisin for this is not fully uuderstood.," It is inferred that these have drifted closer to their star after their formation, but the mechanism for this is not fully understood."428 It bas been suggested that the inward imigratiou of the orbits of 101 Jupiters is relatec to the their gravitational interaction with a circumstellar dise 1996).., It has been suggested that the inward migration of the orbits of hot Jupiters is related to the their gravitational interaction with a circumstellar disc \citep{Lin+96}. .429 This also appears to require larger quautities of gas in the circumstellar ¢isc, This also appears to require larger quantities of gas in the circumstellar disc430the Blanudford-AMcIvec adiabatic selfsimuilarity solution with the sinularity variable at auv radius r. where 59—το) is the Lorentz factor of the shocked iiediun just behind the forward shock whose radius is denoted as R. aud f;=(G0)11(1653)] is the time measured iu the local mediuu’s rest frame.,"the Blandford-McKee adiabatic self-similarity solution with the similarity variable at any radius $r$, where $\gamma_2\equiv \gamma_2(R)$ is the Lorentz factor of the shocked medium just behind the forward shock whose radius is denoted as $R$, and $t_l=(R/c)[1+1/(16\gamma_2^2)]$ is the time measured in the local medium's rest frame."431 The radius r can thus be expressed as function of \ by which implies r=Ro as long as xcx1653 for au ultrarelativistic forward shock., The radius $r$ can thus be expressed as function of $\chi$ by which implies $r\simeq R$ as long as $\chi\ll 16\gamma_2^2$ for an ultrarelativistic forward shock.432" This justifies theepprocination, in which the width of region 2 is insiguificaut as compared to the shock radius Z7."," This justifies the, in which the width of region 2 is insignificant as compared to the shock radius $R$."433" According to Blandford Meee (1976). the pressure and Loreutz factor of the shocked medium at radius à are eiven bv where js, is the proton mass."," According to Blandford McKee (1976), the pressure and Lorentz factor of the shocked medium at radius $r$ are given by where $m_p$ is the proton mass."434 Aloug the contact discontinuity. 55=22(c0) aud P=Por). which vield where the thin-shell approximation rz£7 has been considered.," Along the contact discontinuity, $\gamma_3=\gamma_2(r)$ and $P'_3=P'_2(r)$, which yield where the thin-shell approximation $r\simeq R$ has been considered."435 For an ultrarclativistic. adiabatic forward shock. Blandford Melzee (1976) found its total cherey. Iu deriving the temporal laws of the similarity variable at the location ofthe contact discontinutyv and the Loreutz factors of regions 2 and 3. we should note one crucial effect that the photous that are radiated frou regions 2 and 3 at the same time in the local πουπας rest frame will be detected at different observer times.," For an ultrarelativistic, adiabatic forward shock, Blandford McKee (1976) found its total energy, In deriving the temporal laws of the similarity variable at the location of the contact discontinuity and the Lorentz factors of regions 2 and 3, we should note one crucial effect that the photons that are radiated from regions 2 and 3 at the same time in the local medium's rest frame will be detected at different observer times."436" This is because the Loreutz factor of region 3 is snaller than σος) by a factor of 4.1/7 so that for a same time interval in the local medius rest frame. fe. the emission from region 3 will reach the observer at time. and the emission from region 2 will reach the observer at tinc. Using equatious (12) aud (11). the Lorentz factor of the shocked medium G.c.. region 2) just behind the forward shock is found to evolve with time as From equations (10)-(13). we have the similarity variable at the location of the contact ciscoutinuity. aud the Loreutz factor of region 3. where Ly.prLE,.flO""-oresἩ, Eso=Ey/ans»10°?eres. and 03 aud f£ are in uuits of Lem7 aud 1 day respectively."," This is because the Lorentz factor of region 3 is smaller than $\gamma_2(R)$ by a factor of $\chi^{-1/2}$ so that for a same time interval in the local medium's rest frame, $R/c$, the emission from region 3 will reach the observer at time, and the emission from region 2 will reach the observer at time, Using equations (12) and (14), the Lorentz factor of the shocked medium (i.e., region 2) just behind the forward shock is found to evolve with time as From equations (10)-(13), we have the similarity variable at the location of the contact discontinuity, and the Lorentz factor of region 3, where $L_{w,47}=L_w/10^{47}\,{\rm erg}\,{\rm s}^{-1}$, $E_{52}=E_0/10^{52}\,{\rm ergs}$, and $n_1$ and $t$ are in units of $1\,{\rm cm}^{-3}$ and 1 day respectively."437 Letting 4=1. we define a critical time At this time. the injection enerev to the fireball sienificantly exceeds its initial cnerey.," Letting $\chi=1$, we define a critical time At this time, the injection energy to the fireball significantly exceeds its initial energy."438 For f<fq. the similarity variable 4>1.," For $t<t_{\rm cr}$, the similarity variable $\chi>1$."439 Tt should be eiuipliasized that the dynamics denoted by equation (17) is simply calculated by equatiug the pressures of the two-sided shocked fluids at the contact discontinuity., It should be emphasized that the dynamics denoted by equation (17) is simply calculated by equating the pressures of the two-sided shocked fluids at the contact discontinuity.440 In this derivation. we have ucelected auv work done on region 2 by region 3 because the pressure of region 3 is much less than PG at t<fa.," In this derivation, we have neglected any work done on region 2 by region 3 because the pressure of region 3 is much less than $P'_2(R)$ at $t<t_{\rm cr}$."441 Once the observer’s time exceeds fü. the sinularity variable 4=1.," Once the observer's time exceeds $t_{\rm cr}$, the similarity variable $\chi=1$."442 At this stage. the total kinetic cnerey of region 2 is approximated as Zu»(33LALye. where ALS(li/3) is the swopt-up medium mass.," At this stage, the total kinetic energy of region 2 is approximated as $E_{{\rm kin},2}=(\gamma_2^2-1)M_{\rm443sw}c^2$, where $M_{\rm sw}=(4\pi/3)R^3n_1m_p$ is the swept-up medium mass."444" Encrey conservationMain, requires that any increase of kinetic enerev of region 2 should be equal to work done by PELION 3. where d!H=[7R2dR’qnPugR23) is the volume change of region 23 in its own rest frame."," Energy conservation requires that any increase of kinetic energy of region 2 should be equal to work done by region 3, where $dV'_3=4\pi R^2dR'=4\pi R^2(dR/\gamma_3)$ is the volume change of region 3 in its own rest frame."445 Since reeious 2 and 3 should keep velocity equality aloug the coutact discontinuity (viz.," Since regions 2 and 3 should keep velocity equality along the contact discontinuity (viz.,"446 5»= 53). we rewrite equation (19) as Considering equation (5) aud the shock radius D 55ct. the solutiou. of equation. (20)κ beconies where the dependence of 2» on £ is consistent with the oue derived for a pure electromagnetic energy injection by Dai Lu (1998).," $\gamma_2=\gamma_3$ ), we rewrite equation (19) as Considering equation (5) and the shock radius $R\simeq4474\gamma_2^2ct$ , the solution of equation (20) becomes where the dependence of $\gamma_2$ on $t$ is consistent with the one derived for a pure electromagnetic energy injection by Dai Lu (1998)."448 We next discuss the dyvuaudes of a relativistic wind bubble: Iu the case of fy=—Tipo (viz.," We next discuss the dynamics of a relativistic wind bubble: In the case of $t_{\rm cr}>T_{{\rm M},0}$ (viz.,"449 Eso> the wind bubble should evolve based on equations (15) aud (17): for fa<Tipy (viz.," $E_{52}>0.46I_{45}P_{0,{\rm ms}}^{-2}$ ), the wind bubble should evolve based on equations (15) and (17); for $t_{\rm cr}<T_{{\rm450M},0}$ (viz.,"451 Es»« qu however. the Loreutz factors of the wine bubble decay initially as 59Xf77 and ryX£777058 at t<fu (stage T). subsequently as το=55xfPl at CauTypo) (stage ID. and finally again as 55X£78 at fTip (stage IIT).," $E_{52}<0.46I_{45}P_{0,{\rm ms}}^{-2}$ ), however, the Lorentz factors of the wind bubble decay initially as $\gamma_2\propto t^{-3/8}$ and $\gamma_3\propto t^{-39/136}$ at $t<t_{\rm cr}$ (stage I), subsequently as $\gamma_2=\gamma_3\propto t^{-1/4}$ at $t\in (t_{\rm cr},T_{{\rm452M},0})$ (stage II), and finally again as $\gamma_2\propto453t^{-3/8}$ at $t>T_{{\rm M},0}$ (stage III)."454 It should be pointed out that this, It should be pointed out that this455the results obtained for each of the sources. All X-ray fluxes and luminosities are corrected for absorption.,"the results obtained for each of the sources, All X-ray fluxes and luminosities are corrected for absorption."456comparison.. ? suggested an association of the IGR source with PKS 0352-686. a blazar of BL Lac type at 7=0.087. based on its location inside the IBIS error circle (2) as well as the fact that these objects are known to be strong emitters of X- and gamma-rays.," \citet{masetti06a} suggested an association of the IGR source with PKS $-$ 686, a blazar of BL Lac type at $z$ =0.087, based on its location inside the IBIS error circle \citep{gotz06} as well as the fact that these objects are known to be strong emitters of X- and gamma-rays."457 The source detected by Swift//XRT is from the position of reported in NED. The extended 2MASX source that lies within the XRT error circle (Table 3)) has already been associated with the BL Lac.," The source detected by /XRT is from the position of reported in NED, The extended 2MASX source that lies within the XRT error circle (Table \ref{tab:ircounterparts}) ) has already been associated with the BL Lac."458source.. The ssource Is coincident with which is classified as being a cluster of galaxies in SIMBAD., The source is coincident with which is classified as being a cluster of galaxies in SIMBAD.459 An absorbed power-law represents the Swift//XRT spectrum well with y7=00.98 for 63 degrees of freedom (dof)., An absorbed power-law represents the /XRT spectrum well with 0.98 for 63 degrees of freedom (dof).460 The value of the absorption (Table 5)) is compatible with the value of Galactic absorption along the line of sight., The value of the absorption (Table \ref{tab:spectral}) ) is compatible with the value of Galactic absorption along the line of sight.461 At z20.087. the 2-10 keV luminosity of the source is ~2.5x107 erg/s. We note that the extrapolated 20-40 keV flux of the sspectrum ts about twice as high as the fflux reported in ?..," At $z$ =0.087, the 2–10 keV luminosity of the source is $\sim2.5\times 10^{44}$ erg/s. We note that the extrapolated 20–40 keV flux of the spectrum is about twice as high as the flux reported in \citet{gotz06}."462 If the extrapolation of the power-law is valid. ther this indicates variability. as expected ina BL Lac.," If the extrapolation of the power-law is valid, then this indicates variability, as expected in a BL Lac."463 Based on positional comeidence and the good agreement between the aandROSAT spectral «παρε. 9 suggested that IGR J05346—-5759 is the hard X-ray counterpart to TW Pic. a Cataclysmic Variable (CV).," Based on positional coincidence and the good agreement between the and spectral shape, \citet{gotz06} suggested that IGR $-$ 5759 is the hard X-ray counterpart to TW Pic, a Cataclysmic Variable (CV)."464 There is a unique and quite bright XRT source within the IBIS error circle., There is a unique and quite bright XRT source within the IBIS error circle.465 We first checked the XRT count rates for variability between the different pointings., We first checked the XRT count rates for variability between the different pointings.466 The source shows some variability between high flux states (up to ~O.45 cts/s) and lower flux states (down to ~0.11 ets/s)., The source shows some variability between high flux states (up to $\sim 0.45$ cts/s) and lower flux states (down to $\sim0.11$ cts/s).467 We extracted a single spectrum from one of each of the three pointings., We extracted a single spectrum from one of each of the three pointings.468 An absorbed fits the data well in all cases (yz-11.19 for 89 dof. 1.29 for 14 dof and 1.26 for 98 dof. for pointings #11. 2 and 3. respectively).," An absorbed fits the data well in all cases 1.19 for 89 dof, 1.29 for 14 dof and 1.26 for 98 dof, for pointings 1, 2 and 3, respectively)."469 The best spectral parameters of all three pointings are reported in Table 5.. and they are in good agreement with This variability has been used by ? to refute the Intermediate Polar (IP) type for this source.," The best spectral parameters of all three pointings are reported in Table \ref{tab:spectral}, , and they are in good agreement with This variability has been used by \citet{norton00} to refute the Intermediate Polar (IP) type for this source."470 A refined position for this object has recently been, A refined position for this object has recently been471excluded from further analysis. since data were affected by the strong telluric OI features.,"excluded from further analysis, since data were affected by the strong telluric OI features."472 In order to speed up the data analysis. we co-added spectra into sufficiently narrow phase bins in such à way that phase smearing of the stellar absorption lines originating. from stellar RV. variation. sub-pixel shifts and the effect of barycentric motion of the Earth remained negligible.," In order to speed up the data analysis, we co-added spectra into sufficiently narrow phase bins in such a way that phase smearing of the stellar absorption lines originating from stellar RV variation, sub-pixel shifts and the effect of barycentric motion of the Earth remained negligible."473 The main criterion for the size of the phase bins was that the unseen planetary lines did not suffer from smearing in excess of 2kms!, The main criterion for the size of the phase bins was that the unseen planetary lines did not suffer from smearing in excess of $2~{\rm km~s^{-1}}$.474 With that step. we were able to reduce the number of spectra from 406 to 58.," With that step, we were able to reduce the number of spectra from 406 to 58."475 We model the starlight reflected from the planet as a copy of the stellar spectrum. strongly scaled down in brightness and Doppler-shifted according to the orbital motion of the planet.," We model the starlight reflected from the planet as a copy of the stellar spectrum, strongly scaled down in brightness and Doppler-shifted according to the orbital motion of the planet."476 The target spectra have an average S/N of 300 to 600 per spectral bin., The target spectra have an average S/N of $300$ to $600$ per spectral bin.477 With expected planet-to-star flux ratios of the order of a few times 107. it is clear that the reflected spectrum from the planet is deeply buried in the noise of the stellar spectrum.," With expected planet-to-star flux ratios of the order of a few times $10^{-5}$, it is clear that the reflected spectrum from the planet is deeply buried in the noise of the stellar spectrum."478 The weak planetary signal is boosted by the large number of spectra. and more importantly. by the combination of the approximately 1500 absorption lines. achieved using the data-synthesis method (cf.," The weak planetary signal is boosted by the large number of spectra, and more importantly, by the combination of the approximately $1500$ absorption lines, achieved using the data-synthesis method (cf."479 Charbonneau et al., Charbonneau et al.480 1999) described below in a nutshell (for a detailed description see Rodler et al., 1999) described below in a nutshell (for a detailed description see Rodler et al.481 2008)., 2008).482 In the first step. a high S/N. virtually planet-free superspectrum Is computed by co-adding the observed spectra after correcting their wavelength values for the barycentric motion of the Earth and for the stellar orbital motion.," In the first step, a high S/N, virtually planet-free superspectrum is computed by co-adding the observed spectra after correcting their wavelength values for the barycentric motion of the Earth and for the stellar orbital motion."483 Then. we construct a model to describe each of the original. unmodified object spectra: the dominant stellar signal is represented by the superspectrum. shifted to the position. of the observed spectrum.," Then, we construct a model to describe each of the original, unmodified object spectra: the dominant stellar signal is represented by the superspectrum, shifted to the position of the observed spectrum."484 Imperfections in flux and wavelength are corrected by using the approach deseribed in Rodler et al. (, Imperfections in flux and wavelength are corrected by using the approach described in Rodler et al. (4852008).,2008).486 Furthermore. the contribution to the spectrum of the planetary signal is created as follows: we adopt the spectrum of the slowly rotating F6 IV star HD 136351 to mimic the expected sharp reflected lines of the planet (Section 3.2)).," Furthermore, the contribution to the spectrum of the planetary signal is created as follows: we adopt the spectrum of the slowly rotating F6 IV star HD 136351 to mimic the expected sharp reflected lines of the planet (Section \ref{rotbro}) )."487 For each observed object spectrum. we need to co-align that spectrum of HD 136351 with the stellar model spectrum (shifted superspectrum).," For each observed object spectrum, we need to co-align that spectrum of HD 136351 with the stellar model spectrum (shifted superspectrum)."488 The co-aligned spectrum of HD 136351 is then scaled down by the factor eC)µίώ.Ὀ and shifted by velocity V.C.6) with respect to the stellar spectrum.," The co-aligned spectrum of HD 136351 is then scaled down by the factor $\epsilon(\lambda)~\mu(\phi, i)$ and shifted by velocity $V_{\rm{p}}(K_{\rm{p}},\phi)$ with respect to the stellar spectrum."489 Hence. the two free parameters of the planetary model are the planet-to-star flux ratio for the fully-illuminated planet εί). and the orbital inclination /.. which corresponds to an RV semi-amplitude of the planet Ay=Kysini155.6sinikmsl.," Hence, the two free parameters of the planetary model are the planet-to-star flux ratio for the fully-illuminated planet $\epsilon(\lambda)$ , and the orbital inclination $i$, which corresponds to an RV semi-amplitude of the planet $K_{\rm{p}}=K_{\rm{p,max}} \sin i=155.6490 \sin i~~{\rm km~s^{-1}}$."491 Concerning the albedo function pGU. we adopt (1) a grey albedo assumption (Le. pCGlU=constant for all wavelengths 2) and (it) the irradiated atmospheric Class IV. model by Sudarsky et al. (," Concerning the albedo function $p(\lambda)$, we adopt (i) a grey albedo assumption (i.e. $p(\lambda)={\rm constant}$ for all wavelengths $\lambda$ ) and (ii) the irradiated atmospheric Class IV model by Sudarsky et al. ("4922000). which predicts higher reflectivity at shorter wavelengths (Fig. 2)).,"2000), which predicts higher reflectivity at shorter wavelengths (Fig. \ref{sudarsk}) )."493" The search range for the RV semi-amplitude comprises Kp=40 to 180kms! (corresponding to orbital inclinations i=15° to 907. plus twice the error of Ky, (Le. 18.3kms! with a step width of 3kms""."," The search range for the RV semi-amplitude comprises $K_{\rm p} = 40$ to $180~{\rm km~s^{-1}}$ (corresponding to orbital inclinations $i=15^\circ$ to $90^\circ$, plus twice the error of $K_{\rm p,max}$ (i.e. $18.3~{\rm km~s^{-1}}$ with a step width of $3~{\rm km~s^{-1}}$."494 This is a good compromise between computing time and sampling the expected average reflected absorption line profile with the FWHM of =12kms!," This is a good compromise between computing time and sampling the expected average reflected absorption line profile with the FWHM of $\approx495 12~{\rm km~s^{-1}}$."496 Since the tidal lock hypothesis and a spin-orbit alignment ts very likely. low-inelination orbits can be furthermore excluded: the observed vsini=15kms?! woulc imply an improbable true rotational velocity above 50kms! for an F7 IV-V star at very low orbital inclinations (c.f.," Since the tidal lock hypothesis and a spin-orbit alignment is very likely, low-inclination orbits can be furthermore excluded: the observed $v \sin i = 15~{\rm km~s^{-1}}$ would imply an improbable true rotational velocity above $50~{\rm km~s^{-1}}$ for an F7 IV-V star at very low orbital inclinations (c.f."497 rotational velocities of late F-type stars: Glebocki Stawikowski 2000)., rotational velocities of late F-type stars; Glebocki Stawikowski 2000).498 Using simulations. we found that for small inclinations of the planetary orbit. where the planets appear only slightly illuminated. and the method would fail to detect Jupiter-size objects even with very highalbedos.," Using simulations, we found that for small inclinations of the planetary orbit, where the planets appear only slightly illuminated, and the method would fail to detect Jupiter-size objects even with very highalbedos."499This equation is general as long as à. 6. and F are in the same coordinate system (e.g.. Earth mean equator J2000) and the object located at (0. 4) is infinitely far away.,"This equation is general as long as $\alpha$ , $\delta$ , and $\vec{r}$ are in the same coordinate system (e.g., Earth mean equator J2000) and the object located at $\alpha$, $\delta$ ) is infinitely far away."500 Other forms of this equation in the literature assume that we have the angular coordinates of the new origin or that the Earth and the new origin are in the same plane (e.g..BinnendiJk1960;Henden&Kaitchuck1982;HirshfeldSinnott 1997).. but we explain in refsec:calculating why this form is most. practical. for calculating the delay.," Other forms of this equation in the literature assume that we have the angular coordinates of the new origin or that the Earth and the new origin are in the same plane \citep[e.g.,][]{binnendijk60,501henden82, hirshfeld97}, but we explain in \\ref{sec:calculating}502 why this form is most practical for calculating the delay."503 The HJD. which uses the Sun as the origin of the new reference frame. is only accurate to 8 s because of the acceleration of the Sun due primarily to Jupiter and Saturn (Fig. 2)).," The HJD, which uses the Sun as the origin of the new reference frame, is only accurate to 8 s because of the acceleration of the Sun due primarily to Jupiter and Saturn (Fig. \ref{fig:bjdvhjd}) )."504 It was popular when people first began considering this effect because it is relatively simple to calculate from tables without a computer (Landolt&Blondeau1972).. and remains popular because self-contained algorithms exist to approximate it without any external tables (e.g..Duffett-Smith 1989).," It was popular when people first began considering this effect because it is relatively simple to calculate from tables without a computer \citep{landolt72}, and remains popular because self-contained algorithms exist to approximate it without any external tables \citep[e.g.,][]{duffet89}."505.. However. because the HJD is not useful when accuracies of better than 8 s are needed. most of the algorithms im use today use approximations that are only precise at the I s level. and it becomes impossible to back out the original JD from the HJD unless we know the exact algorithm used.," However, because the HJD is not useful when accuracies of better than 8 s are needed, most of the algorithms in use today use approximations that are only precise at the 1 s level, and it becomes impossible to back out the original JD from the HJD unless we know the exact algorithm used."506 Because of these problems. the HJD was formally deprecated by International Astronomical Union (AU) Resolution A4 in 1991. in favor of the BJD. a time referenced to the Solar System Barycenter (SSB).," Because of these problems, the HJD was formally deprecated by International Astronomical Union (IAU) Resolution A4 in 1991, in favor of the BJD, a time referenced to the Solar System Barycenter (SSB)."507 The analogous correction to the Roemer delay 1n our Solar System can also be significant in the target system., The analogous correction to the mer delay in our Solar System can also be significant in the target system.508 We refer to this as Ay., We refer to this as $\Delta_{R}$.509 For example. for transiting planets with cz:0.06 AU. Ap can be as large as 30 s. In general. the position of the planet during primary transit has become the unspoken standard reference frame for transiting planets. while the host star's photosphere is the unspoken standard for Radial Velocity (RV) planets.," For example, for transiting planets with $a \approx5100.06$ AU, $\Delta_{R}$ can be as large as 30 s. In general, the position of the planet during primary transit has become the unspoken standard reference frame for transiting planets, while the host star's photosphere is the unspoken standard for Radial Velocity (RV) planets."511 In theory. the timing would be much more stable in the targets barycentric reference frame. but the accuracy with which we can convert to this frame depends on the measurements of the system.," In theory, the timing would be much more stable in the target's barycentric reference frame, but the accuracy with which we can convert to this frame depends on the measurements of the system."512 Since different observers may use different values as measurements improve. quoting the JD in the frame of the target’s barycenter may obfuscate the long term reliability of timing.," Since different observers may use different values as measurements improve, quoting the JD in the frame of the target's barycenter may obfuscate the long term reliability of timing."513 Therefore. we argue it 1s better to quote Julian Date in the SSB reference frame. and correct for Ay only when comparing observations at different phases in the planet's orbit.," Therefore, we argue it is better to quote Julian Date in the SSB reference frame, and correct for $\Delta_{R}$ only when comparing observations at different phases in the planet's orbit."514 This A correction is not necessary for TTVs of the primary transit. since the planet is always in the same phase.," This $\Delta_{R}$ correction is not necessary for TTVs of the primary transit, since the planet is always in the same phase."515 Nevertheless. we should explicitly state the objects reference frame to avoid any potential ambiguity. particularly when comparing any combination of primary transits. secondary transits. RVs. and another primary transit of a different planet in the same system. when it may not be obvious which origin is being used.," Nevertheless, we should explicitly state the object's reference frame to avoid any potential ambiguity, particularly when comparing any combination of primary transits, secondary transits, RVs, and another primary transit of a different planet in the same system, when it may not be obvious which origin is being used."516 For RV measurements. which are taken at many different phases. the effect is much smaller and can generally be ignored because the star's orbit around the barycenter is small.," For RV measurements, which are taken at many different phases, the effect is much smaller and can generally be ignored because the star's orbit around the barycenter is small."517 For a typical hot Jupiter. (1.e.. a Jupiter mass planet in à 3 day orbit around à solar mass star). the maximum time difference in the RV signal (for an edge-on orbit) is 20 ms. which would change the measured RV by ~50jm s'.," For a typical hot Jupiter, (i.e., a Jupiter mass planet in a 3 day orbit around a solar mass star), the maximum time difference in the RV signal (for an edge-on orbit) is 20 ms, which would change the measured RV by $\sim 50 \518\mu$ m $^{-1}$."519 While planets farther out will cause a larger timing offset. the difference in the measured RV is even smaller.," While planets farther out will cause a larger timing offset, the difference in the measured RV is even smaller."520 To be clear. the JD can be specified in many time standards (Seidelmann1992).. and while the [AU has made no explicit statement regarding the allowed time standards of the GJD. HJD. or BJD. their meaning in any given time standard is unambiguous.," To be clear, the JD can be specified in many time standards \citep{seidelmann92}, and while the IAU has made no explicit statement regarding the allowed time standards of the GJD, HJD, or BJD, their meaning in any given time standard is unambiguous."521 Unfortunately. they have been specified in many standards. usually implicitly.," Unfortunately, they have been specified in many standards, usually implicitly."522 However. the particular time standard used affects how useful the time stamp is as an absolute reference.," However, the particular time standard used affects how useful the time stamp is as an absolute reference."523 We must be careful not to directly compare BJDs or HJDs in different time standards. as each has different offsets. periodic terms. and/or rates. which can introduce systematic errors of over |] minute.," We must be careful not to directly compare BJDs or HJDs in different time standards, as each has different offsets, periodic terms, and/or rates, which can introduce systematic errors of over 1 minute."524 For this reason. it 15 critical that any stated BJD or HJD also specify the time standard used when one-minute accuracies are important. and the uncertainty of a time that is quoted without a standard should beassumedto be at least | minute.," For this reason, it is critical that any stated BJD or HJD also specify the time standard used when one-minute accuracies are important, and the uncertainty of a time that is quoted without a standard should beassumedto be at least 1 minute."525 First. it may be useful to summarize the relevant standards of time:," First, it may be useful to summarize the relevant standards of time:"526artificially inflating the 47 value of the fit.,artificially inflating the $\chi^2$ value of the fit.527 Such calibration errors probably dominate the \7 of the best-lit model for Abell 1795., Such calibration errors probably dominate the $\chi^2$ of the best-fit model for Abell 1795.528 The fit is unacceptable (42=3.85). but the residual pattern in Figure 5 demonstrates the effect of a probable gain offset below 0.5 keV (Prieto. coupled with small statistical errors derived [rom the large number of counts (6x 10).," The fit is unacceptable $\chi^2_r = 3.85$ ), but the residual pattern in Figure \ref{f5} demonstrates the effect of a probable gain offset below 0.5 keV \citep{phs} coupled with small statistical errors derived from the large number of counts $6\tenup{5}$ )."529 The cooling flow model fit of 1795 is of equally poor quality. but the resulting excess column is closer to the Galactic value for the cooling flow versus [or the two-thermal component model).," The cooling flow model fit of A1795 is of equally poor quality, but the resulting excess column is closer to the Galactic value for the cooling flow versus for the two-thermal component model)."530 Two-component model fits to the remaining 5 clusters are poor. but if they are physically senilicant thev show the same behavior as the rest of the sample.," Two-component model fits to the remaining 5 clusters are poor, but if they are physically significant they show the same behavior as the rest of the sample."531 Allowing each of their Galactic columns to vary does reduce the 4? of the fit. vielding a model with a higher column. alihough the significance of this is difficult to ascertain due to the poor significance of the resulting model.," Allowing each of their Galactic columns to vary does reduce the $\chi^2$ of the fit, yielding a model with a higher column, although the significance of this is difficult to ascertain due to the poor significance of the resulting model."532 If we assume that these fits are physically significant. the columnis exceed their Galactic values by <3856 [for the cooling flow models). which is twpically an order of magnitude smaller than the excesses found by WEJMA (see Table 3)).," If we assume that these fits are physically significant, the columns exceed their Galactic values by $\leq38$ for the cooling flow models), which is typically an order of magnitude smaller than the excesses found by WFJMA (see Table \ref{tab:comp}) )."533 For example. clusters displaving absorption above the Galactic value in both studies (Abell 35. 1795. 2029. and 2199). but otherwise do not seem to be unusual in Ay. show an excess that is 40 times ereater in WFEJMA than in the present work.," For example, clusters displaying absorption above the Galactic value in both studies (Abell 85, 1795, 2029, and 2199), but otherwise do not seem to be unusual in $\NHoo$, show an excess that is 40 times greater in WFJMA than in the present work."534 The models of WJEMA all contain cooling lows. so for completeness we ran cooling flow models with variable absorption (both Galactic and proximal to the cooling flow) for ihe entire sample.," The models of WJFMA all contain cooling flows, so for completeness we ran cooling flow models with variable absorption (both Galactic and proximal to the cooling flow) for the entire sample."535 For (hose models which contain only a Galactic absorption component (as a [ree parameter). in no case was a substantial excess absorption required to model the Mnission.," For those models which contain only a Galactic absorption component (as a free parameter), in no case was a substantial excess absorption required to model the emission."536 We cannot rule out the presence of a significant quantity of cool gas al the center of cooling flows. bul we stress (hat a significant excess absorption is not arequired leature ol these spectra.," We cannot rule out the presence of a significant quantity of cool gas at the center of cooling flows, but we stress that a significant excess absorption is not a feature of these spectra."537 Of the internally absorbed cooling flow clusters common to both WFJMA and (his study. only one quarter show significant excess absorption.," Of the internally absorbed cooling flow clusters common to both WFJMA and this study, only one quarter show significant excess absorption."538 The fact (hat any show significant. absorplion is not surprising. since absorption can be invoked (o obscure any inount of cooling flow emission: that onlv a quarter actually display (his behavior suggests that excess internal absorption is probably not a ubiquitous feature of these svstenis.," The fact that any show significant absorption is not surprising, since absorption can be invoked to obscure any amount of cooling flow emission; that only a quarter actually display this behavior suggests that excess internal absorption is probably not a ubiquitous feature of these systems."539 Tt is difficult to compare these results with those of Allen&Fabian(1997) since the methods differ significantlv: however. one point is worth mentioning.," It is difficult to compare these results with those of \citet{af} since the methods differ significantly; however, one point is worth mentioning."540" The “color profile"" approach that they adopted used data [rom 0.4 keV through 2 keV. In low Galactic column clusters most of the absorption is manifest from 0.2 to 0.4 keV. where our technique is «quite sensitive."," The “color profile” approach that they adopted used data from 0.4 keV through 2 keV. In low Galactic column clusters most of the absorption is manifest from 0.2 to 0.4 keV, where our technique is quite sensitive."541 For example. thev compute an excess «πι of almost lor 2029. whereas our (wo-component model is only larger than the Galactic value (and lower still for our externally absorbed cooling flow moclel: see Table 3)).," For example, they compute an excess $\NHoo$ of almost for A2029, whereas our two-component model is only larger than the Galactic value (and lower still for our externally absorbed cooling flow model; see Table \ref{tab:comp}) )."542 Figures 6. and 7 show a direct comparison between between our cooling flow moclels of, Figures \ref{f6} and \ref{f7} show a direct comparison between between our cooling flow models of543Luminous optical/UV. emission-lines ancl Lar-infrarecl dust reraciation emission than observed.,luminous optical/UV emission-lines and far-infrared dust reradiation emission than observed.544 ltellection from a partially ionized aceretion disk could emerge in this energv range., Reflection from a partially ionized accretion disk could emerge in this energy range.545 However. the soft X-ray variation which is larger than that of the primary continuum is hard to explain by this hypothesis.," However, the soft X-ray variation which is larger than that of the primary continuum is hard to explain by this hypothesis."546 A possible solution is to introduce a warm absorber which varies in response to the continuum source., A possible solution is to introduce a warm absorber which varies in response to the continuum source.547 This moclel needs no extra N-rav. source but does necd. cülferent physical. conditions for the X-ray absorber., This model needs no extra X-ray source but does need different physical conditions for the X-ray absorber.548 Unlike. many higher luminosity Sevfert 1: galaxies in which OVIL (0.74 keV) ancl OVILE (0.85 keV) are major features of the warn absorption. the primary absorption occurs around 2 keV in NGOC4395.," Unlike many higher luminosity Seyfert 1 galaxies in which OVII (0.74 keV) and OVII (0.85 keV) are major features of the warm absorption, the primary absorption occurs around 2 keV in NGC4395."549" It requires a higher ionization parameter and the absorption is due to highly ionized O. Ne. Mg. οἱ, 5 and I'e-L. Introducing a warm absorber reduces the opacity of the cold absorption significantly."," It requires a higher ionization parameter and the absorption is due to highly ionized O, Ne, Mg, Si, S and Fe-L. Introducing a warm absorber reduces the opacity of the cold absorption significantly."550 This is more consistent. with the observed. optical/UV. properties. e.g.. non-stellar UV continuum. broad. permitted line emission. high excitation lines usually seen in Sevfert 1: nuclei.," This is more consistent with the observed optical/UV properties, e.g., non-stellar UV continuum, broad permitted line emission, high excitation lines usually seen in Seyfert 1 nuclei."551 The column density implied. from cold absorption alone fitted to the 210 keV spectrum corresponds to zly~10 when the Galactic gas-to-dust ratio is used., The column density implied from cold absorption alone fitted to the 2–10 keV spectrum corresponds to $A_{\rm V}\sim 10$ when the Galactic gas-to-dust ratio is used.552 The Sevfert-1 like optical/UV. properties would not be observed if the central source and the broad-line region (BLY are behind such a heavy. obscuration. although some reddening to the BLE and the central source is still required. given the ionizing photon deficit. (e.g... [or narrow 113). pointed out by. Lira et al (1999). INraemer ct al (1999) and Moran et al (1999). the Sevfert 1.8/1.9 nature in the optical emission-line spectrum (e.e.. Ho et al 1997) and the laree Ho fll? ratio (~5.1. Ixraemer et al 1999) for the broad component.," The Seyfert-1 like optical/UV properties would not be observed if the central source and the broad-line region (BLR) are behind such a heavy obscuration, although some reddening to the BLR and the central source is still required given the ionizing photon deficit (e.g., for narrow $\beta$ ), pointed out by Lira et al (1999), Kraemer et al (1999) and Moran et al (1999), the Seyfert 1.8/1.9 nature in the optical emission-line spectrum (e.g., Ho et al 1997) and the large $\alpha$ $\beta$ ratio $\sim 5.1$, Kraemer et al 1999) for the broad component."553 As the soft X-ray part of the spectrum. has such a variable nature. a spectral variability study is more useful in modelling the soft X-ray spectrum than using the time-averaged spectrum alone.," As the soft X-ray part of the spectrum has such a variable nature, a spectral variability study is more useful in modelling the soft X-ray spectrum than using the time-averaged spectrum alone."554 The warm absorber hypothesis is then explored to account for the spectral variability between the active and quiescent states in the next section., The warm absorber hypothesis is then explored to account for the spectral variability between the active and quiescent states in the next section.555 We try to model the spectral change in the soft X-ray band between the quiescent ancl active states with a warm absorber., We try to model the spectral change in the soft X-ray band between the quiescent and active states with a warm absorber.556 “Phe continuum source is assumed. to have a constant spectral shape of a power-law with photon-index P=1.72. obtained from the integrated 210 keV cata. but cillerent luminosities in the two states.," The continuum source is assumed to have a constant spectral shape of a power-law with photon-index $\Gamma = 1.72$, obtained from the integrated 2–10 keV data, but different luminosities in the two states."557 Since it is unlikely. that column density changes in response to the continuum source. the ionization parameter. £=L(n.BR?) of absorbing matter is a primary driver of the spectral change.," Since it is unlikely that column density changes in response to the continuum source, the ionization parameter, $\xi=L/(nR^2)$ of absorbing matter is a primary driver of the spectral change."558 We used the warm absorber model by Done et al (1992) in NSPEC., We used the warm absorber model by Done et al (1992) in XSPEC.559 Cold absorption is also included in the model., Cold absorption is also included in the model.560 The small aperture SIS spectra of the quiescent ancl active states are fitted. jointly to obtain parameters which are. shared between the two datasets. such as a column density of an absorber.," The small aperture SIS spectra of the quiescent and active states are fitted jointly to obtain parameters which are shared between the two datasets, such as a column density of an absorber."561" The mean absorption-corrected 210 keV. Duxes of the power-law continuum during the quiescent and active states are 3.2""190.12""erg aand 6.1.10Perg1. respectively."," The mean absorption-corrected 2–10 keV fluxes of the power-law continuum during the quiescent and active states are $3.2\times 10^{-12}$ and $6.1\times 10^{-12}$, respectively."562 A model of a single warm absorber with variable £ plus a constant cold absorber is first fitted., A model of a single warm absorber with variable $\xi$ plus a constant cold absorber is first fitted.563" Lhe temperature of the absorbing gas is assumed to be 3LO"" Ix. Column densities of warm and cold absorbers aandNy. respectively) are free parameters but. set. to vary in unison between the two datasets."," The temperature of the absorbing gas is assumed to be $3\times 10^5$ K. Column densities of warm and cold absorbers and, respectively) are free parameters but set to vary in unison between the two datasets."564 The ionization xwameter (£) of the absorber is allowed to take different values in the two datasets as well as the normalization of power-law., The ionization parameter $\xi$ ) of the absorber is allowed to take different values in the two datasets as well as the normalization of power-law.565 This parameter setting enables the model to acconmmodate a physically reasonable change in parameters in à variable warm absorber hypothesis., This parameter setting enables the model to accommodate a physically reasonable change in parameters in a variable warm absorber hypothesis.566 This nioclel. jowever. gives a poor fit to the absorption band. (below 2 keV).," This model, however, gives a poor fit to the absorption band (below 2 keV)."567 Despite the Dux change of more than a factor of two. he ionization parameters obtained from the fit dilfer very ittle (Moclel-1 in Table 2).," Despite the flux change of more than a factor of two, the ionization parameters obtained from the fit differ very little (Model-1 in Table 2)."568 When the values of Nyy for the two datasets are also allowed to vary independently. the fit improved significantly (Mocdel-2 in Table 2).," When the values of $N_{\rm W}$ for the two datasets are also allowed to vary independently, the fit improved significantly (Model-2 in Table 2)."569 However. not only has Wy decreased in the active state. but the lower £ in the active state is in the opposite sense to the photoionization hypothesis.," However, not only has $N_{\rm W}$ decreased in the active state, but the lower $\xi$ in the active state is in the opposite sense to the photoionization hypothesis."570 Therefore the single warm absorber model fails. but the latter fit instead. suggests the existence of two physically distinct. absorbers. Le. the absorption feature in the two spectra are dominated by clillerent absorbers.," Therefore the single warm absorber model fails, but the latter fit instead suggests the existence of two physically distinct absorbers, i.e, the absorption feature in the two spectra are dominated by different absorbers."571 Next we try a multi-laver warm absorber moclel., Next we try a multi-layer warm absorber model.572with respect to MyvLD (Lable 1)).,with respect to MvLD (Table \ref{SourcesTable}) ).573 Phe computed ὅτι flux excesses have also changed. with the result that LELD 42044 no longer has a significant (30)excesst..," The computed $\mu$ m flux excesses have also changed, with the result that LEID 42044 no longer has a significant $>$ $\sigma$ )."574 The extraction and calibration of these spectra followed the standard sequence (for details. see Sargentetal. 2010)).," The extraction and calibration of these spectra followed the standard sequence (for details, see \citealt{SSM+10}) )."575 Spectra. were extracted after. images ha been cdilferenced to remove background. radiation and ceaned to. replace bad. pixels., Spectra were extracted after images had been differenced to remove background radiation and cleaned to replace bad pixels.576 We extracted: spectra from ineliviclual images. coaddec them. anc calibrated. them with LR 6348 (for Short-Low: SL) and LR 6348 and LID 73511 (for Long-Low: LL).," We extracted spectra from individual images, coadded them, and calibrated them with HR 6348 (for Short-Low; SL) and HR 6348 and HD 173511 (for Long-Low; LL)."577 The SL and LL spectra wer| joined. stitehed. and trimmed.," The SL and LL spectra were joined, stitched, and trimmed."578 Stitching means thata| segments were normalized multiplicatively (upward. by : 4)) to remove discontinuities between orders caused by pointing-induced errors in radiation passing through the s»ectroscopie. slits.," Stitching means thatall segments were normalized multiplicatively (upward, by $\leq$ ) to remove discontinuities between orders caused by pointing-induced errors in radiation passing through the spectroscopic slits."579 ‘Trimming means removing extraneous cd:ua from the ends of the segments., Trimming means removing extraneous data from the ends of the segments.580 The final spectra are shown in reflBRSEig.. 3 4. with Dux plotted in ATE. to better show deviations from a blackbocdy.," The final spectra are shown in \\ref{IRSFig}, \ref{IRS2Fig} \ref{IRS3Fig}, with flux plotted in $\lambda^2 F_\nu$, to better show deviations from a blackbody."581 We have obtained optical spectra of 12 of our 14 targets in order to compare gas mass loss tracers to the amount and type of dust. produced., We have obtained optical spectra of 12 of our 14 targets in order to compare gas mass loss tracers to the amount and type of dust produced.582 These spectra were taken with the Magellan Inamori Kyocera Eecholle (MEE) clouble-ecchelle spectograph mounted on the Magellan: 6.5-m Clay telescope at Las Campanas observatory using a 0.75 5 arcsec2 slit.. giving. a resolution. of. 240000.," These spectra were taken with the Magellan Inamori Kyocera Écchelle (MIKE) double-écchelle spectograph mounted on the Magellan 6.5-m Clay telescope at Las Campanas observatory using a 0.75 $\times$ 5 $^2$ slit, giving a resolution of $R \approx 40\,000$."583 Phe spectra were reduced with bias subtraction. Dlat-fielding and sky subtraction using the MEIIZ-IDL pipeline updated by PProchaska.," The spectra were reduced with bias subtraction, flat-fielding and sky subtraction using the MIKE-IDL pipeline updated by ."584. VhAr are exposures bracketting the stellar targets were used to determine the wavelength scale., Th–Ar arc exposures bracketting the stellar targets were used to determine the wavelength scale.585?..,\cite{Kuntschner2000}.586 We do not increase the extraction width to obtain a better S/N. because the velocity dispersion can be strongly aperture-dependent (as shown in Fig. 13. ," We do not increase the extraction width to obtain a better S/N, because the velocity dispersion can be strongly aperture-dependent (as shown in Fig. \ref{dispersionextractionwindow}) ),"587and it is important to be well-matched to ?.., and it is important to be well-matched to \cite{Kuntschner2000}.588 As a tinal step. the spectra were rebinned to a common wavelength increment (2.43 A//pix). approximately doubling the original step size.," As a final step, the spectra were rebinned to a common wavelength increment (2.43 /pix), approximately doubling the original step size."589 To illustrate the effect of the different data reduction steps. we refer to Fig. 2.. ," To illustrate the effect of the different data reduction steps, we refer to Fig. \ref{fig:data_reduction}, ,"590"where a central 3.7"" extraction of NCG 1381 is shown for the different data reduction steps: the ISAAC pipeline result. after removing the telluric lines and after rebinning to the À//pix wavelength increment."," where a central $3.7''$ extraction of NCG 1381 is shown for the different data reduction steps: the ISAAC pipeline result, after removing the telluric lines and after rebinning to the /pix wavelength increment."591" For each extracted spectrum. we have derived an empirical S/Nfollowing the method described by ὃν,"," For each extracted spectrum, we have derived an empirical S/Nfollowing the method described by \cite{Stoehr2007}."592 The resulting S/N per rebinned element ranges from ~7 for the faintest galaxies to more than 100 for the brightest ones and is listed in Table I.., The resulting S/N per rebinned element ranges from $\sim7$ for the faintest galaxies to more than 100 for the brightest ones and is listed in Table \ref{galaxies}. .593" The instrumental set-up. as described in Table 2.. gives a FWHM resolution of approximately 107 km + . producing an instrumental contribution to the dispersion of ~45 km """," The instrumental set-up, as described in Table \ref{tabinstrsetup}, , gives a FWHM resolution of approximately 107 km $^{-1}$ , producing an instrumental contribution to the dispersion of $\sim45$ km $^{-1}$."594 The sharp blue edge of the CO band head allows us to measure the Kinematics accurately (2)..., The sharp blue edge of the CO band head allows us to measure the kinematics accurately \citep{Silge2003}.595 There are several techniques to obtain the internal kinematical information. e.g. the Fourier correlationquotient (FCQomethod. developedby ? and usedby ?..," There are several techniques to obtain the internal kinematical information, e.g. the Fourier correlationquotient (FCQ)method, developedby \cite{Bender1994} and usedby \cite{Kuntschner2000}. ."596 We use the, We use the597igh mass stars are very short lived. no more than a few million years.,"high mass stars are very short lived, no more than a few million years."598 There is no evidence of a deuse stellar association near NCC 5los ΝΤ aud origin in the closest super-star cluster would require a transit tine to the preseut location on the order of 30 Myr. (vaaretctal. 2003)., There is no evidence of a dense stellar association near NGC 5408 X-1 and origin in the closest super-star cluster would require a transit time to the present location on the order of 30 Myr \citep{phil03}.599. Thus. the companion mass is likely significantly lower. near 20 AZ.. or less simular to fouud frou studies of the stellar euviroumieuts of other ULAs (Caiséetal.2008. 2011).," Thus, the companion mass is likely significantly lower, near 20 $M_{\odot}$ or less similar to found from studies of the stellar environments of other ULXs \citep{fabien08,fabien11}."600. Figure 3., Figure 3.601 shows the upper bound on the conrpact object mass for donors of 120 M... and 20 M. as a function of orbital period., shows the upper bound on the compact object mass for donors of 120 $_{\odot}$ and 20 $_{\odot}$ as a function of orbital period.602 Wiel black hole masses are excluded. except for very short periods.," High black hole masses are excluded, except for very short periods."603 We note that the liue shift was the same in ΟΡΟ versus OBS aud OBG. taken one day apart. suggesting that the period is longer than few davs.," We note that the line shift was the same in OB3 versus OB5 and OB6, taken one day apart, suggesting that the period is longer than few days."604 Thus. the black hole mass is likely," Thus, the black hole mass is likely"605"ISM perturbations lead to 3-dimensional effects that alters the density distribution patterns observed in Fig.l and, in addition, cause significant warping of the disk.","ISM perturbations lead to 3–dimensional effects that alters the density distribution patterns observed in \ref{f1} and, in addition, cause significant warping of the disk."606" However, if the optical depth 7 is of the order of 10:53, the collisional disruption prevents the building up of significant inclination, as it does for the eccentricity, and the warping of the disk is negligible."," However, if the optical depth $\tau$ is of the order of $10^{-3}$, the collisional disruption prevents the building up of significant inclination, as it does for the eccentricity, and the warping of the disk is negligible."607 This is shown in Fig.4 where the outcome of a simulation with irs=60° and 7=107° is illustrated as a density plot in the (ᾳ—2) plane., This is shown in \ref{f4} where the outcome of a simulation with $i_{ISM}=60^o$ and $\tau = 10^{-3}$ is illustrated as a density plot in the $(x-z)$ plane.608 Only small out-of-plane features can be detected in the figure which cannot be interpreted neither as warping or clumping due to their limited extension along the z-axis (the scale along the axis is 1/10 of the scale along the x-axis)., Only small out--of--plane features can be detected in the figure which cannot be interpreted neither as warping or clumping due to their limited extension along the z–axis (the scale along the z--axis is 1/10 of the scale along the x–axis).609 A large optical depth prevents then a disk to develop structures not only in the parent body plane but also out-of-plane., A large optical depth prevents then a disk to develop structures not only in the parent body plane but also out–of–plane.610 This finding is confirmed for any value of irs., This finding is confirmed for any value of $i_{ISM}$.611 In Fig.5 the median inclination of the grains in the disk is plotted for different values of 4;s. with T=10? (top plot)., In \ref{f5} the median inclination of the grains in the disk is plotted for different values of $i_{ISM}$ with $\tau = 10^{-3}$ (top plot).612 The value of the median of the particle inclinations is always smaller than 5? since the Stark cycle is quickly interrupted by a collision., The value of the median of the particle inclinations is always smaller than $5^o$ since the Stark cycle is quickly interrupted by a collision.613" This confirms that if the disk is optically thick, the ISM"," This confirms that if the disk is optically thick, the ISM"614offsets remained coustaut relative to the levels found from the overscan pixels (the latter determined separately for the four quadrants).,offsets remained constant relative to the levels found from the overscan pixels (the latter determined separately for the four quadrants).615 For bias subtraction. therefore. we subtracted both the levels from the overscan regious iu individual frames. as well as au average of the bias frames.," For bias subtraction, therefore, we subtracted both the levels from the overscan regions in individual frames, as well as an average of the bias frames."616 The frames were corrected for sensitivitv variations using flat fields coustructed from images of the sky taken at dusk aud dawn., The frames were corrected for sensitivity variations using flat fields constructed from images of the sky taken at dusk and dawn.617 Since the conditious were not photometric diming either of our two runs. we cannot reliably calibrate our," Since the conditions were not photometric during either of our two runs, we cannot reliably calibrate our"618missing XRB (Vignali et al.,missing XRB (Vignali et al.619 in preparation)., in preparation).620" In the GOODS-S and -N fields, the population of CT AGN at z1 can be tracked further down to intrinsicLy of 10“erg.."," In the GOODS-S and -N fields, the population of CT AGN at $z\sim 1$ can be tracked further down to intrinsic$L_X$ of $10^{42}$."621" Furthermore, since the CT samples obtained using the X/NeV diagnostic appear to be relatively free from contamination by less obscured objects, by using spectroscopic surveys with well defined selection functions it would be possible to estimate the space density of this missing AGN population."," Furthermore, since the CT samples obtained using the X/NeV diagnostic appear to be relatively free from contamination by less obscured objects, by using spectroscopic surveys with well defined selection functions it would be possible to estimate the space density of this missing AGN population."622" It has to be noted that selection of obscured AGN through the [Ne V]3426 line is most likely a lower limit, since it misses objects with dusty Narrow Line Regions, which can instead be picked up by mid-IR selection(??)."," It has to be noted that selection of obscured AGN through the [Ne V]3426 line is most likely a lower limit, since it misses objects with dusty Narrow Line Regions, which can instead be picked up by mid-IR selection."623". However, when sufficiently deep mid-IR coverage is not available for spectroscopic survey fields, [Ne V]-selection may then represent a promising and ready-to-use method to get large samples of z~1 CT AGN."," However, when sufficiently deep mid-IR coverage is not available for spectroscopic survey fields, [Ne V]-selection may then represent a promising and ready-to-use method to get large samples of $z\sim 1$ CT AGN."624" In addition, the comparison between the space density of [Ne V] and mid-IR selected objects with similar redshifts and bolometric luminosities in fields with full multiwavelength coverage, may indicate the fraction of objects in which the NLR is free from obscuration, thus constraining the physical scale on which absorption arises."," In addition, the comparison between the space density of [Ne V] and mid-IR selected objects with similar redshifts and bolometric luminosities in fields with full multiwavelength coverage, may indicate the fraction of objects in which the NLR is free from obscuration, thus constraining the physical scale on which absorption arises."625" Popular semi-analytic models of galaxy formation and evolution propose that, at least for the most massive and luminous objects, nuclear activity and star formation are both triggered by major mergers of gas rich galaxies, and that at the early stages of the merger, when star formation is more vigorous, the AGN is embedded within optically thick gas shrouds."," Popular semi-analytic models of galaxy formation and evolution propose that, at least for the most massive and luminous objects, nuclear activity and star formation are both triggered by major mergers of gas rich galaxies, and that at the early stages of the merger, when star formation is more vigorous, the AGN is embedded within optically thick gas shrouds."626 According to this scenario one can therefore expect to observe strong star formation in obscured QSOs., According to this scenario one can therefore expect to observe strong star formation in obscured QSOs.627" At redshift z~2, the concurrent obscured black hole growth and star formation has been observed in the population of bright submillimeter sources ?)."," At redshift $z\sim 2$, the concurrent obscured black hole growth and star formation has been observed in the population of bright submillimeter sources ."628". At low redshifts (z< 0.3), noted that the ratio between the luminosity of the [O II]3727 line and the [O III]5007 line is a factor of ~4 higher in the type- QSO composite spectrum of than in the average spectrum of type-1 QSOs."," At low redshifts $z<0.3$ ), noted that the ratio between the luminosity of the [O II]3727 line and the [O III]5007 line is a factor of $\sim 4$ higher in the type-2 QSO composite spectrum of than in the average spectrum of type-1 QSOs."629" Because of the similar [O III] luminosity of the two populations, the excess of low-ionized oxygen in type-2 QSOs could not be explained in terms of different ionization parameters of the Narrow Line Regions, and was then interpreted by as due to enhanced star formation in type- QSOs."," Because of the similar [O III] luminosity of the two populations, the excess of low-ionized oxygen in type-2 QSOs could not be explained in terms of different ionization parameters of the Narrow Line Regions, and was then interpreted by as due to enhanced star formation in type-2 QSOs."630 We then investigated the star formation in the samples presented in this work by measuring the [O ΠΙ2727 flux on the SDSS spectra and considered the [O II] over [Ne V] ratio as a function of the nuclear obscuration., We then investigated the star formation in the samples presented in this work by measuring the [O II]3727 flux on the SDSS spectra and considered the [O II] over [Ne V] ratio as a function of the nuclear obscuration.631 To compareobjects in the same redshift range we restricted our analysis to z> 0.4., To compareobjects in the same redshift range we restricted our analysis to $z>0.4$ .632" As shown in Fig.4,, this cut removes only low-redshifts"," As shown in \ref{lnexz_sdss}, , this cut removes only low-redshifts"633linear array of lavee ellipticals to the west of NGC 1275 (sce Paper II).,linear array of large ellipticals to the west of NGC 1275 (see Paper II).634 To calibrate the photometry. we fitted zero poiut offsets. adiinass and color terms using UDR Lanclolt standard star photometiv acquired during cach welt (Paper ID).," To calibrate the photometry, we fitted zero point offsets, airmass and color terms using UBR Landolt standard star photometry, acquired during each night (Paper II)."635 Photometry for each galaxy was done in two wars., Photometry for each galaxy was done in two ways.636" \lagnitudes representative of the cutire galaxy were measured within au aperture. BR. equal to three times au inverted Petrosian (1976) radius, or Ry, = 3 \ r(j = 0.5)."," Magnitudes representative of the entire galaxy were measured within an aperture, $_p$, equal to three times an inverted Petrosian (1976) radius, or $_{p}$ = 3 $\times$ $\eta$ = 0.5)."637 We also meastred core magnitudes aud colors within the ceutral of cach galaxy., We also measured core magnitudes and colors within the central of each galaxy.638 We also correct the plotometry for Galactic extinction. and for a shelt k-correction.," We also correct the photometry for Galactic extinction, and for a slight k-correction."639 For more details on the sample sclection and the data themselves. see Paper IL.," For more details on the sample selection and the data themselves, see Paper II."640 Monte. Carlo simulations. aud comparisons between the S2kD aud Minui-Mosaic plotometry. which is further discussed in Paper IL also argues that our photometry is reliable down to D = 21 or Mp—10.5 at the adopted distance of the Perseus cluster.," Monte Carlo simulations, and comparisons between the S2kB and Mini-Mosaic photometry, which is further discussed in Paper II, also argues that our photometry is reliable down to B = 24 or $_{\rm B} = -10.5$ at the adopted distance of the Perseus cluster."641 This section briefly discusses how we are able to confidently identify LAICCs in the Perseus cluster from other cluster meuber aud backeround systems in a reliable manner., This section briefly discusses how we are able to confidently identify LMCGs in the Perseus cluster from other cluster member and background systems in a reliable manner.642 See Paper II for a more detailed discussion aud description of these issues., See Paper II for a more detailed discussion and description of these issues.643 The πα of our sclection criteria is the following: (i)., The summary of our selection criteria is the following: (i).644 Total (B8...R)y colors of < 2., Total $(B-R)_{0}$ colors of $<$ 2.645 Galaxies redder than this are alinost always in the backgrouud., Galaxies redder than this are almost always in the background.646 Frou stellar population svuthesis models (Worthey 1991). no normal passively evolviug galaxy with metallicity [Fe/II] = 10.5 is redder than (B Rjy ~ 2 after an 18 Gyr initial burst of star formation.," From stellar population synthesis models (Worthey 1994), no normal passively evolving galaxy with metallicity [Fe/H] = +0.5 is redder than $(B-R)_{0}$ $\sim$ 2 after an 18 Gyr initial burst of star formation."647 We further onlv investigate in detail Perseus LAICGs with maenitudes Mp<<12.5 to be more certain of cluster mcuibership and to have very accurately measured photometric aud structural parameters. G, We further only investigate in detail Perseus LMCGs with magnitudes $_{\rm B} < -12.5$ to be more certain of cluster membership and to have very accurately measured photometric and structural parameters. (648i}.,ii).649 Svuumetric. round or elliptical shapes. without evidence for internal structures that might be due to star formation. spiral arius. or other internal features.," Symmetric, round or elliptical shapes, without evidence for internal structures that might be due to star formation, spiral arms, or other internal features."650 Dwiuf cllipticals/spheroidals appear this way. whereas backerouud galaxies are often morphologically disturbed and can be identified as such in high resolution images (o.@.. Conselice 2001). (," Dwarf ellipticals/spheroidals appear this way, whereas background galaxies are often morphologically disturbed and can be identified as such in high resolution images (e.g., Conselice 2001). ("651i).,iii).652 A central surface brightuess in a aperture fainter than µ = 21.0 mae > and a nouw-ceutrally concentrated light profile that is close to exponcutial (see Paper IT and 833.1)., A central surface brightness in a aperture fainter than $\mu_{{\rm B}}$ = 24.0 mag $^{-2}$ and a non-centrally concentrated light profile that is close to exponential (see Paper II and 3.1).653 These are properties of uearby dwarf ellipticals aud cau be used to distinguish LAICGs from eiaut ellipticals aud. backerouud systems. (, These are properties of nearby dwarf ellipticals and can be used to distinguish LMCGs from giant ellipticals and background systems. (654iv) As discussed in Paper Π. we also limit our study to objects that are detected above a So coutidence.,"iv) As discussed in Paper II, we also limit our study to objects that are detected above a $\sigma$ confidence."655 Figure 1l shows Uarris R-baud nuages of the LAICGs which we study in this paper aud which constitute the ealaxies that make up the low-Iunünositv scatter in the Perseus coloranagnuitude relationship (Paper IL. Figure 2 8.1).," Figure 1 shows Harris R-band images of the LMCGs which we study in this paper and which constitute the galaxies that make up the low-luminosity scatter in the Perseus color-magnitude relationship (Paper II, Figure 2 3.4)."656 We also somewhat arbitrarily separate LAICCs iuto blue aud red types for discussion purposes., We also some-what arbitrarily separate LMCGs into blue and red types for discussion purposes.657 We define blue Perseus galaxies as those which are within |20 of the coloranaguitude relationship scatter at Mp=17 ou the red side and everything bluer than this (see Cousclice 2002)., We define blue Perseus galaxies as those which are within $+2 \sigma$ of the color-magnitude relationship scatter at $_{\rm B} = -17$ on the red side and everything bluer than this (see Conselice 2002).658 LAICCs which are >26 redder than the CAIR are denoted as red systems., LMCGs which are $>2\sigma$ redder than the CMR are denoted as red systems.659 Figure 3 show the (5.R)y color histogram for Perseus LMCCs. down to Mp=11. with the red and blue svsteuis shaded differently (see Paper IT).," Figure 3 show the $(B-R)_{0}$ color histogram for Perseus LMCGs, down to $_{\rm B} = -11$, with the red and blue systems shaded differently (see Paper II)."660 Tn addition. to Figure 1. which displavs the R-baud inages of the 53 ealaxies studied in this paper. Table 1 lists their positions. photometric aud structural properties.," In addition to Figure 1, which displays the R-band images of the 53 galaxies studied in this paper, Table 1 lists their positions, photometric and structural properties."661 For reference. the absolute imiaguitude Mp. (08.R99 color. and identification number are also listed iu Table 1 aud printed on Figure 1 for each. galaxy.," For reference, the absolute magnitude $_{{\rm B}}$, $(B-R)_{0}$ color, and identification number are also listed in Table 1 and printed on Figure 1 for each galaxy."662 Iun Paper ID we studied the basic plotometric aud structural properties of the Perseus LAICCs. including siuple global morphological indexes such as asvuunuetiv and concentration (Bershady. Jaugreun Conselice 2000: Cousclice. Bershady Jaueren 2000).," In Paper II we studied the basic photometric and structural properties of the Perseus LMCGs, including simple global morphological indexes such as asymmetry and concentration (Bershady, Jangren Conselice 2000; Conselice, Bershady Jangren 2000)."663 We found that the concentration meexes aud asviletrics of the LAICGs are consistent with dwart clliptical-like objects., We found that the concentration indexes and asymmetries of the LMCGs are consistent with dwarf elliptical-like objects.664 We also found a strong correlation between the central surface brightucss aud magnitudes of these objects. with the same relationship seeu for nearby dwarfs (ec... Binegech Cameron 1991).," We also found a strong correlation between the central surface brightness and magnitudes of these objects, with the same relationship seen for nearby dwarfs (e.g., Binggeli Cameron 1991)."665 This sugeests that all Perseus LAICGs are objects that would be classified. if nearby. as dwarf elliptical or spheroidals and are also thus likely cluster ienibers.," This suggests that all Perseus LMCGs are objects that would be classified, if nearby, as dwarf elliptical or spheroidals and are also thus likely cluster members."666 We go bevoud this basic examination in this paper to deteriunue the detailed structural properties of the Perseus LMCC5z., We go beyond this basic examination in this paper to determine the detailed structural properties of the Perseus LMCGs.667 Previous studies lave found that LAICGs in other clusters. as well as Local Croup dSplis. have exponeutial profiles. while large ellipticals lave steeper rb/! de Vaucouleur profiles aud are found to be more rouuder than LAICCs (Ryden Terndrup 1991).," Previous studies have found that LMCGs in other clusters, as well as Local Group dSphs, have exponential profiles, while large ellipticals have steeper $^{1/4}$ de Vaucouleur profiles and are found to be more rounder than LMCGs (Ryden Terndrup 1994)."668 What is the case for the Perseus LAICCs?, What is the case for the Perseus LMCGs?669 To auswer this. we perform surface plotometry on all 53 early-type LAICCs aud all the elliptical galaxies used in this paper.," To answer this, we perform surface photometry on all 53 early-type LMCGs and all the elliptical galaxies used in this paper."670 We fit elliptical isophotes to cach galaxy image using the Fourier prescription of Jedrzejewsiki (1987). naploimieuted by the ELLIPSE routine in STSDAS.," We fit elliptical isophotes to each galaxy image using the Fourier prescription of Jedrzejewski (1987), implemented by the ELLIPSE routine in STSDAS."671 We allow the cllipticity. position angle. and center to chauge while fitting.," We allow the ellipticity, position angle, and center to change while fitting."672 The general form of the Fourier decomposition of intensity distributions is with «e; aud b; the fitted Fourier components. aud o the position angele at cach fitted semi-major leugth r.," The general form of the Fourier decomposition of intensity distributions is with $a_{i}$ and $b_{i}$ the fitted Fourier components, and $\phi$ the position angle at each fitted semi-major length$r$ ."673 For each surface brightuess profile. Γεν a general Sérrsic profile of the form: is fit. where Z(r) is the intensity of au isoplote at r. Ly is the central iutensity. aud rj is the scale leneth.," For each surface brightness profile, $I(r)$, a general Sérrsic profile of the form: is fit, where $I(r)$ is the intensity of an isophote at $r$, $I_{0}$ is the central intensity, and $_{0}$ is the scale length."674 The value n = Lis for pure exponential profiles. and n = Lis the de Vaucouleurs profile that fits eiut elliptical galaxies fairly well (ee. CGavazzi ct al.," The value n = 1 is for pure exponential profiles, and n = 4 is the de Vaucouleurs profile that fits giant elliptical galaxies fairly well (e.g., Gavazzi et al."675 2000)., 2000).676 Table 1 lists the fitted profile parameters. n aud ry. forcach LAICG.," Table 1 lists the fitted profile parameters, n and $_{0}$ , foreach LMCG."677Argentina. a field containing the reflection nebula NGC 6726 was imaged twice. with roughly one month between observations.,"Argentina, a field containing the reflection nebula NGC 6726 was imaged twice, with roughly one month between observations."678 The observations were made by means of our large field-of-view 0.45m [2.5 remotely-operated telescope on top of Cerro Burek. through a broad (390 to 700 nm) filter.," The observations were made by means of our large field-of-view 0.45m f/2.8 remotely-operated telescope on top of Cerro Burek, through a broad (390 to 700 nm) filter."679 The detector was a 4008 x 2672 pixel camera based on the Kodak KAFI1000M CCD chip., The detector was a 4008 $\times$ 2672 pixel camera based on the Kodak KAF11000M CCD chip.680 The pixel scale of the setup was 1.47 arcsec/pixel., The pixel scale of the setup was 1.47 arcsec/pixel.681 In the course of standard reductio and analysts. the very young stars S CrA and R CrA experienced significant brightness changes.," In the course of standard reduction and analysis, the very young stars S CrA and R CrA experienced significant brightness changes."682 Closer inspection of the data clearly revealed material that appeared to be expanding away from the two stars at superluminal speeds for all plausible values of the distances to the stars., Closer inspection of the data clearly revealed material that appeared to be expanding away from the two stars at superluminal speeds for all plausible values of the distances to the stars.683 The appearance of such superluminal motion is well understood as the scattering of a light pulse off of stationary. dust near the line of sight. rather than physical expansion of material ejected by the star (Coudere1939).," The appearance of such superluminal motion is well understood as the scattering of a light pulse off of stationary dust near the line of sight, rather than physical expansion of material ejected by the star \citep{Couderc1939}."684. Light echoes not only illuminate otherwise hidden dust. but also directly reveal that material's. 3-dimensional structure.," Light echoes not only illuminate otherwise hidden dust, but also directly reveal that material's 3-dimensional structure."685 Since mapping the circumstellar environments of such young stars has the potential to reveal critical details of their formation processes. we re-imaged the field of NGC 6726 over the next three months at one-to-two week intervals. as listed in. Table 1.," Since mapping the circumstellar environments of such young stars has the potential to reveal critical details of their formation processes, we re-imaged the field of NGC 6726 over the next three months at one-to-two week intervals, as listed in Table 1."686 Each image was reduced in the standard manner bbias subtraction. flat fielding). then geometrically registered to a common orientation using field stars as described in Sugerman (2005).," Each image was reduced in the standard manner bias subtraction, flat fielding), then geometrically registered to a common orientation using field stars as described in \citet{Suger2005}."687. RMS residuals of geometric registration are less than 0.1 pixel for all images., RMS residuals of geometric registration are less than 0.1 pixel for all images.688 Finally. data were flux calibrated to the V-band (+0.1 mags) using field stars from the UCAC? catalog (Zachariasetal.2004).," Finally, data were flux calibrated to the $V$ -band $\pm 0.1$ mags) using field stars from the UCAC2 catalog \citep{Zacha2004}."689. Fig.l shows the reflection nebula NGC 6726 and surrounding field. with the locations of S and κ CrA indicated.," \ref{fig1} shows the reflection nebula NGC 6726 and surrounding field, with the locations of S and R CrA indicated."690 A full color animation of that field is provided as supplementary online material., A full color animation of that field is provided as supplementary online material.691 The time-variable light echoes are immediately obvious i1 the animation as they move away from their sources., The time-variable light echoes are immediately obvious in the animation as they move away from their sources.692 A temporal sequence of the field immediately surrounding S CrA is shown in reffig2.., A temporal sequence of the field immediately surrounding S CrA is shown in \\ref{fig2}.693 The echoes can also be faintly seen in the direct images (left column of reftig2)). and are slightly more apparent when the function (PSF) of S CrA has been removed (Stetson 1987).. as shown in the middle column.," The echoes can also be faintly seen in the direct images (left column of \\ref{fig2}) ), and are slightly more apparent when the point-spread function (PSF) of S CrA has been removed \citep{Stet1987}, as shown in the middle column."694 Since light echoes are transient phenomena. we employed the PSF-matching and difference-imaging techniques of Tomaney&Crotts(1996) and Sugerman(2005).. combined with the NN2 algorithm of Newman&Rest(2006).. to remove all sources of constant flux and produce echo-only images for each epoch.," Since light echoes are transient phenomena, we employed the PSF-matching and difference-imaging techniques of \citet{Tomaney1996} and \citet{Suger2005}, combined with the NN2 algorithm of \cite{Newman2006}, to remove all sources of constant flux and produce echo-only images for each epoch."695 These are shown in the right column of reffig2.., These are shown in the right column of \\ref{fig2}.696 An animated close-up view of S CrA subtracted images in which the residuals from the subtractions have been removed by interpolation is shown as online material., An animated close-up view of S CrA subtracted images in which the residuals from the subtractions have been removed by interpolation is shown as online material.697 Expanding concentric rings or “quasi-rings” are clearly visible in. this animation.," Expanding concentric rings or “quasi-rings"" are clearly visible in this animation."698 As revealed by the animation and in reffig2.. the echoes have an overall circular morphology but the circles are incomplete: that is why the term “quasi-rings” would be more appropriate to describe them.," As revealed by the animation and in \\ref{fig2}, the echoes have an overall circular morphology but the circles are incomplete; that is why the term “quasi-rings"" would be more appropriate to describe them."699 These expanding echoes are analyzed in the next section., These expanding echoes are analyzed in the next section.700 In the typical approach to model nova or supernovae echoes. the light curve and distance of the source are used to deduce the 3-dimensional structure of the dust producing the echoes (Sugerman2003).," In the typical approach to model nova or supernovae echoes, the light curve and distance of the source are used to deduce the 3-dimensional structure of the dust producing the echoes \citep{Suger2003}."701. Such an approach is not feasible with our S CrA data., Such an approach is not feasible with our S CrA data.702 S ΟΓΑ has been observed as part of the All Sky Automated Survey (Pojmanski1997).. however the photometry (Fig.3)) does not have homogenous time sampling. making the epochs of maximum light and the pulse durations moderately uncertain.," S CrA has been observed as part of the All Sky Automated Survey \citep{Pojman1997}, however the photometry \ref{fig3}) ) does not have homogenous time sampling, making the epochs of maximum light and the pulse durations moderately uncertain."703 Additionally. there are few distance measurements to the star. with the best estimates ranging from 120 to 140 pe (Marraco&Rydgren1981:Carmonaetal.2007).. as there are no reliable parallax measurements from Hipparcos (seethediscussiononthistopicinNeuháuser&Forbrich 2008).," Additionally, there are few distance measurements to the star, with the best estimates ranging from 120 to 140 pc \citep{Marraco1981,Carmona2007}, as there are no reliable parallax measurements from Hipparcos \citep[see the704 discussion on this topic in][]{Neu2008}."705.. In such circumstances. we instead constrain the above properties by comparing the echoes that would arise from a test dust geometry to those actually observed.," In such circumstances, we instead constrain the above properties by comparing the echoes that would arise from a test dust geometry to those actually observed."706 We initially use guesses for the dates and times of four outbursts and later on we refine such estimates based on the observed light curve., We initially use guesses for the dates and times of four outbursts and later on we refine such estimates based on the observed light curve.707 A minimum of 4 outbursts are needed for the following reason., A minimum of 4 outbursts are needed for the following reason.708 While we see echoes expanding in the first set of observations. a month later on we see echoes of smaller angular size. which means that new echoes have been formed: thus another pulse is needed and so on.," While we see echoes expanding in the first set of observations, a month later on we see echoes of smaller angular size, which means that new echoes have been formed; thus another pulse is needed and so on."709 As a first-order model to explain expanding rings. the two simplest dust geometries are a spherical shell and a planar slab. which are representative of circumstellar and interstellar media. respectively.," As a first-order model to explain expanding rings, the two simplest dust geometries are a spherical shell and a planar slab, which are representative of circumstellar and interstellar media, respectively."710 Echoes from à spherical shell always appear concentric and can only grow as large as the shells radius., Echoes from a spherical shell always appear concentric and can only grow as large as the shell's radius.711 For a slab. on the other hand. echoes should be observable to an arbitrary size. and will only appear as concentric. circular rings tf the slab is aligned with the plane of the sky (an unlikely constraint).," For a slab, on the other hand, echoes should be observable to an arbitrary size, and will only appear as concentric, circular rings if the slab is aligned with the plane of the sky (an unlikely constraint)."712 Qualitatively. the echoes we observe are concentric at all epochs and end abruptly around 80 aresee from S. CrA. therefore the spherical shell model is already favored from these arguments. but for completeness we explore both the sphertcal-shell and the planar-slab geometries.," Qualitatively, the echoes we observe are concentric at all epochs and end abruptly around 80 arcsec from S CrA, therefore the spherical shell model is already favored from these arguments, but for completeness we explore both the spherical-shell and the planar-slab geometries."713" We identified the angular distance of each echo from 5 CrA (see Table 1). by iteratively fitting Gaussian functions to radial surface-brightness profiles taken at various position angles (Sugerman 2009). avoiding ai ""ear-like filament which shows brightness variations over time."," We identified the angular distance of each echo from S CrA (see Table 1), by iteratively fitting Gaussian functions to radial surface-brightness profiles taken at various position angles \citep{Suger2005}, , avoiding an “ear”-like filament which shows brightness variations over time."714 This “ear” feature is stationary and can be explained as an isolated knot with thickness smaller than cAr. where Af is a typical pulse duration.," This “ear"" feature is stationary and can be explained as an isolated knot with thickness smaller than $c \Delta t$, where $\Delta t$ is a typical pulse duration."715 An isolated knot in the direction perpendicular to the line of sight would only brighten and fade and no apparent motion would be seen., An isolated knot in the direction perpendicular to the line of sight would only brighten and fade and no apparent motion would be seen.716 Therefore it seems that the ear-like feature is an isolated. overdensity of muchsmaller size than the dust structure causing the expanding quasi-rings., Therefore it seems that the ear-like feature is an isolated overdensity of muchsmaller size than the dust structure causing the expanding quasi-rings.717 Hence we excluded the ear-like signal from the analysis of the angular size of the expanding rings., Hence we excluded the ear-like signal from the analysis of the angular size of the expanding rings.718 , 719Iu the collapsar scenario for GRBs. pulses iu GRD light curves are thought to be produced w collisions between relativistic shells. ejected roni a central eugime (see Zhaue&MOszaros(2001) for a recent review).,"In the collapsar scenario for GRBs, pulses in GRB light curves are thought to be produced by collisions between relativistic shells ejected from a central engine (see \citet{zm04} for a recent review)."720 The interception of a iore slowly moving shell by a second sheLB hat is ejected at a later tine. bu with faster speed aud larger Lorentz factor. produces a shoc- hat dissipates internal enerev to cnereize the articles that enmüt the CRB radiation.," The interception of a more slowly moving shell by a second shell that is ejected at a later time, but with faster speed and larger Lorentz factor, produces a shock that dissipates internal energy to energize the particles that emit the GRB radiation."721 This scOnarbo d:8 Wadely considered to explain pulses i- GRD light cuves (IKobavashi.Piran.&Sari1997:Daigue&Mochkovitch 1998).," This scenario is widely considered to explain pulses in GRB light curves \citep{kps97,dm98}."722. Stulos of pulses aro muportait fo decide if CRB sources require engines thatare long-lastiug or impulsive (απο&Mitinan2 03)... with inportant implications for the nature «ft the central chee. which is often argued to bx| à newly formed black hole powered by the accreion of a massive dense torus.," Studies of pulses are important to decide if GRB sources require engines thatare long-lasting or impulsive \citep{dm03}, with important implications for the nature of the central engine, which is often argued to be a newly formed black hole powered by the accretion of a massive dense torus."723 Tere we construct an elementarv kinematic model for colliding shells. assumed. spliexical aud uniform within jet opening augle 0;.," Here we construct an elementary kinematic model for colliding shells, assumed spherical and uniform within jet opening angle $\theta_j$."724 This is the sort of jet tha Frailetal.(2001) discuss regarding the sαπατά cherev reservoir result. where jet opening angeles aro inferred from. the time of achromatic spectral breaks iu optical afterglow light curves.," This is the sort of jet that \citet{fra01} discuss regarding the standard energy reservoir result, where jet opening angles are inferred from the time of achromatic spectral breaks in optical afterglow light curves."725" Weiso perform this study in order to quautity the curvature coustraiut of a spherically οιτας shell traveling with bulk Loreutz factor EP. which implies hat the shell radius in order to produce variability ou timescale f, (Rybicki&Liglhtinan1979:Fenimore.Madras.Navakshiun 1996)."," We also perform this study in order to quantify the curvature constraint of a spherically emitting shell traveling with bulk Lorentz factor $\Gamma$, which implies that the shell radius in order to produce variability on timescale $t_{var}$ \citep{rl79,fmn96}."726. This study also quantifies the rate at which fux decays at a given energy due to curvature effects; aud the rauge of validity of the approximate relation between internal comoving cnerey density aj and observed enerev fux Be. where dy is," This study also quantifies the rate at which flux decays at a given energy due to curvature effects, and the range of validity of the approximate relation between internal comoving energy density $u_0^\prime$ and observed energy flux $\Phi_E$ , where $d_L$ is"727We therefore conclude that the ULXs are most likely powered by super-Eddington accretion onto LOAL. black holes.,We therefore conclude that the ULXs are most likely powered by super-Eddington accretion onto $\sim 10 \msun$ black holes.728 The two different types of ULXs result from their different inclination angles., The two different types of ULXs result from their different inclination angles.729 The permanently super-Eddington power-law type ULXs might be viewed at low inclination angles., The permanently super-Eddington power-law type ULXs might be viewed at low inclination angles.730 The low-luminosity ULXs with non-power-law spectra are consistent of being super-Eddington sources viewed at high inclinations (NGC 253 X-4. NGC 253 X-9) or they are sub-Eddington sources (NGC 253 X-2).," The low-luminosity ULXs with non-power-law spectra are consistent of being super-Eddington sources viewed at high inclinations (NGC 253 X-4, NGC 253 X-9) or they are sub-Eddington sources (NGC 253 X-2)."731 The authors thank Dr. Hua Feng and the anonymous referee for useful comments., The authors thank Dr. Hua Feng and the anonymous referee for useful comments.732 This work was supported by the Finnish Graduate School in Astronomy and Space Physics (JTEK) and the Academy of Finland grant 127512 (IP)., This work was supported by the Finnish Graduate School in Astronomy and Space Physics (JJEK) and the Academy of Finland grant 127512 (JP).733 This research has made use of data obtained from the Chandra Data Archive and software provided by the Chandra X-ray Center (CXC) in the application packages CIAO. ChIPS. and Sherpa.," This research has made use of data obtained from the Chandra Data Archive and software provided by the Chandra X-ray Center (CXC) in the application packages CIAO, ChIPS, and Sherpa."734 This research was based on observations obtained with XMM-Newton. an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA.," This research was based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA."735(his value. are not available to us.,"this value, are not available to us."736 We however point out that (he maximunr acceleration of the CME occurred at or less than 215. in height., We however point out that the maximum acceleration of the CME occurred at or less than $2\Rsun$ in height.737 In the higher corona.i.e... in the COR? ΕΟΝ we see that the acceleration once again rises by a small amount to reach I8ms7 at a heieht of around 8H... whieh can be attributed to the residual acceleration phase consistent with the model proposed by Chen&Ixrall(2003).," In the higher corona, in the COR2 FOV we see that the acceleration once again rises by a small amount to reach $18\mpss$ at a height of around $8\Rsun$, which can be attributed to the residual acceleration phase consistent with the model proposed by \citet{Chen.Krall2003}."738. This was a bright CME with a well-defined sviimetriceal LE on the south-east limb ol the Sun. as seen in Figure 2..," This was a bright CME with a well-defined symmetrical LE on the south-east limb of the Sun, as seen in Figure \ref{F:img31dec}."739 The CAIE was associated with a C8 class flare which originated in NOAA active region (AR) 10980., The CME was associated with a C8 class flare which originated in NOAA active region (AR) 10980.740 The CME appeared in the DD FOV at UUT. and at UUT in AA ΕΟΝ.," The CME appeared in the B FOV at UT, and at UT in A FOV."741 The CME crossed the CORI FOV in about one hour. indicating that it was a relatively fast CME. and appeared in (he (Ομ FOV at UUT in the two spacecralt.," The CME crossed the COR1 FOV in about one hour, indicating that it was a relatively fast CME, and appeared in the COR2 FOV at UT in the two spacecraft."742 This CME showed an unusual cusp in its LE. whieh was distinctly visible in the COR2 images.," This CME showed an unusual cusp in its LE, which was distinctly visible in the COR2 images."743 Thernisienetal.(2009) have used the eracduated evlindrical shell model (Thernisienetal.2006) to fit (vo shells flanking the cusp for this CALE. aud emploved the forward modelling technique to determine its (rue direction of propagation and speed.," \citet{Thernisien.etal2009} have used the graduated cylindrical shell model \citep{Thernisien.etal2006} to fit two shells flanking the cusp for this CME, and employed the forward modelling technique to determine its true direction of propagation and speed."744 We have used this feature for the purpose of reconstruction., We have used this feature for the purpose of reconstruction.745 On the eastern limb of the image Irom EUVIDD. a flare can be seen at CUT. followed by opening up of the field lines which can be clearly seen in and images from DD. From the GOES soft X-ray flix data. we find that the flare started al ULT. and peaked at UUT.," On the eastern limb of the image from B, a flare can be seen at UT, followed by opening up of the field lines which can be clearly seen in and images from B. From the GOES soft X-ray flux data, we find that the flare started at UT, and peaked at UT."746 The CME speed increased in the lower corona to reach S12kms.| at a height of 2.8RR... and showed a little dip before attaining a constant value of around 870kms|! (Figure 8)).," The CME speed increased in the lower corona to reach $812\kmps$ at a height of $2.8\Rsun$, and showed a little dip before attaining a constant value of around $870\kmps$ (Figure \ref{F:res31dec}) )."747"on a time scale of 2/4,=11.3692 dd. It is this that gives rise to the subharmonie and its harmonic in the frequency spectrum,",on a time scale of $2 P_{\rm rot} = 11.3692$ d. It is this that gives rise to the subharmonic and its harmonic in the frequency spectrum.748" Ellipsoidal variables often have frequency spectra where the highest peak occurs at 274,4, as a consequence of a double-wave light curve with different maxima and minima.", Ellipsoidal variables often have frequency spectra where the highest peak occurs at $2\nu_{\rm rot}$ as a consequence of a double-wave light curve with different maxima and minima.749 If the highest peak in the frequency spectrum were mistaken for μις then there would appear to be a subharmonic present.," If the highest peak in the frequency spectrum were mistaken for $\nu_{\rm rot}$, then there would appear to be a subharmonic present."750" This is most unlikely to be the case for 110195926,", This is most unlikely to be the case for 10195926.751 The light curve seen in reftig:le. bears little resemblance to an ellipsoidal variable. and the incidence of short period binaries in magnetic Ap stars is very low.," The light curve seen in \\ref{fig:lc} bears little resemblance to an ellipsoidal variable, and the incidence of short period binaries in magnetic Ap stars is very low."752 We ure gathering high resolution spectra of this star to study ifs rotational variations., We are gathering high resolution spectra of this star to study its rotational variations.753 Those will be used to test for binary radial velocity variations., Those will be used to test for binary radial velocity variations.754 It is clear that the dominant rotational light variation is of the a7 CCVn type. ie. that the star is a spotted magnetic variable.," It is clear that the dominant rotational light variation is of the $\alpha^2$ CVn type, i.e. that the star is a spotted magnetic variable."755 If the variation with 2/7.) were to be interpreted as orbital. then the rotation period would have to be half that of the orbital period. which is again unlikely.," If the variation with $2P_{\rm rot}$ were to be interpreted as orbital, then the rotation period would have to be half that of the orbital period, which is again unlikely."756 Subharmonics are seen in stellar pulsation. most strikingly in theKepler data for LLyr stars. including LLyrae itself (2:: ?)).," Subharmonics are seen in stellar pulsation, most strikingly in the data for Lyr stars, including Lyrae itself \citealt{szaboetal10}; \citealt{kolenbergetal10}) )."757 This is interpreted in terms of nonlinear effects in the pulsation. but this explanation does not apply here because we are dealing with rotation. not pulsation.," This is interpreted in terms of nonlinear effects in the pulsation, but this explanation does not apply here because we are dealing with rotation, not pulsation."758 We interpret the rotational light variations in reftig:le. to be caused by surface spots and we conclude that the rotation period of the star is αρ=5.68459 dd. It is unlikely for an a? CCVn spotted rotator to have two sets of spots that are as asymmetric as those of 110195926. but nearly identical on two opposite hemispheres. as would be required if the rotation period were twice the value given above.," We interpret the rotational light variations in \\ref{fig:lc} to be caused by surface spots and we conclude that the rotation period of the star is $P_{\rm rot} 759= 5.68459$ d. It is unlikely for an $\alpha^2$ CVn spotted rotator to have two sets of spots that are as asymmetric as those of 10195926, but nearly identical on two opposite hemispheres, as would be required if the rotation period were twice the value given above."760 There is no known case of this amongst the a= CCVn variables., There is no known case of this amongst the $\alpha^2$ CVn variables.761 We fitted by least-squares a harmonie series for Sloat, We fitted by least-squares a harmonic series for $\frac{1}{2}\nu_{\rm rot}$.762" Other than #4, and SIZE only the harmonies of £4, were significant."," Other than $\nu_{\rm rot}$ and $\frac{3}{2}\nu_{\rm rot}$, only the harmonics of $\nu_{\rm rot}$ were significant."763 Thus if the subharmonie were the rotation frequency. we would have a harmonic series describing the data that included only every other harmonic.," Thus if the subharmonic were the rotation frequency, we would have a harmonic series describing the data that included only every other harmonic."764 This. too. is unprecedented and unlikely.," This, too, is unprecedented and unlikely."765 reffig:tt-lowfs shows an amplitude spectrum of the short cadence data seen in the middle panel of reftig:le.., \\ref{fig:ft-lowfs} shows an amplitude spectrum of the short cadence data seen in the middle panel of \\ref{fig:lc}.766 To study the pulsation frequencies we ran a high pass filter to remove completely the rotational light variations. and any instrumental drift.," To study the pulsation frequencies we ran a high pass filter to remove completely the rotational light variations, and any instrumental drift."767 This was done by sequential automatic prewhitening of sinusoids until the noise level at low frequencies matches that at higher frequencies., This was done by sequential automatic prewhitening of sinusoids until the noise level at low frequencies matches that at higher frequencies.768 The highest amplitude noise peaks are less than j/mag. so all peaks with amplitudes greater than this at frequencies less than mmHz were removed from," The highest amplitude noise peaks are less than $\mu$ mag, so all peaks with amplitudes greater than this at frequencies less than mHz were removed from"769calculations with N=64 and N=128 produce virtually indistinguishable results. so that the former mode number ts sufficient for all practical purposes.,"calculations with $N=64$ and $N=128$ produce virtually indistinguishable results, so that the former mode number is sufficient for all practical purposes."770 In general. the steepness of the power determines which maximum wave number N is necessary: spectrumInserting the wave numbers from Eq. (7))," In general, the steepness of the power spectrum determines which maximum wave number $N$ is necessary: Inserting the wave numbers from Eq. \ref{eq:pxi_derivation_kn}) )"771 into Eq. (23)).," into Eq. \ref{eq:pxi_derivation_gausspwrspec}) ),"772 we see that N~ Lop/?1s sufficient for Gaussian power spectra.," we see that $N \sim L \, \sigma_P / 2$ is sufficient for Gaussian power spectra."773 In this section. we calculate the moments of the distribution.," In this section, we calculate the moments of the distribution."774 Apart from possible use in future applications. this 1s also useful as a check for the distribution function derived above. since we can derive the moments in two independent ways and compare results.," Apart from possible use in future applications, this is also useful as a check for the distribution function derived above, since we can derive the moments in two independent ways and compare results."775 First. we can get the moments Μι from the derivatives of the characteristic function (2.p.63): The first derivative yields the mean of the distribution. or the expectation value of &: For the variance and other higher-order quantities. we use the central moments. which are the moments of the distribution of ἕ--ἕ.," First, we can get the moments $M_k$ from the derivatives of the characteristic function \citep[p. 63]{Kendall1977}: The first derivative yields the mean of the distribution, or the expectation value of $\xi$: For the variance and other higher-order quantities, we use the central moments, which are the moments of the distribution of $\xi-\overline{\xi}$."776 The centralised characteristic function is simply and it yields the central moments as (?.pp.57.63) The first six non-zero central moments are then From these moments. we can also obtain some conventional statistical quantities: the variance V(£)=Mi». the standard deviation e=/V(é). the skewness S(€)=Mao? and the kurtosis K(£)=Mo?--3.," The centralised characteristic function is simply and it yields the central moments as \citep[pp. 57, 63]{Kendall1977}777 The first six non-zero central moments are then From these moments, we can also obtain some conventional statistical quantities: the variance $V(\xi) = M_{\mathrm{c}2}$, the standard deviation $\sigma = \sqrt{V(\xi)}$, the skewness $S(\xi) = M_{\mathrm{c}3}/\sigma^3$ and the kurtosis $K(\xi)~=~M_{\mathrm{c}4}/\sigma^4 - 3$."778 Alternatively. we can also calculate the moments from the probability distribution function by the integrals An important check for the sanity of the distribution function will be to re-obtain the normalisation as unity by this approach.," Alternatively, we can also calculate the moments from the probability distribution function by the integrals An important check for the sanity of the distribution function will be to re-obtain the normalisation as unity by this approach."779 We can compute it as the moment of order zero:, We can compute it as the moment of order zero:780keV. respectively (the width of the line core. as well as tha of the added lines. held fixed to Ll eV for simplicity).,"keV, respectively (the width of the line core, as well as that of the added lines, held fixed to 1 eV for simplicity)."781 Tho fi improves signilicantIv (A 2157115 clo).," The fit improves significantly $\chi^2$ =157/115 d.o.f.),"782 and no strong systematic residuals are left (Eig. 3))., and no strong systematic residuals are left (Fig. \ref{good_sp}) ).783 The significance of each Line is 99.99‘a. according to the F-test.," The significance of each line is $>$, according to the F-test."784 Phe Compton shoulder is actually expected. to have a finite width. bu unfortunately its profile (Sunvaev Churazov 1996. Mat 2002) is different [rom any model.," The Compton shoulder is actually expected to have a finite width, but unfortunately its profile (Sunyaev Churazov 1996, Matt 2002) is different from any model."785 As a check. we tried for the Compton Shoulder also gaussians with a=40 and 70 eV. (the latter value corresponding to a ENIM of the same order of the total width of the Shoulder).," As a check, we tried for the Compton Shoulder also gaussians with $\sigma$ =40 and 70 eV (the latter value corresponding to a FWHM of the same order of the total width of the Shoulder)."786 A worse fit (\7 of 162/115 and 1797115. respectively) is obtained. but the [Luxes of both the Shoulder ane the line core remain almost unchanged.," A worse fit $\chi^2$ of 162/115 and 179/115, respectively) is obtained, but the fluxes of both the Shoulder and the line core remain almost unchanged."787 The best fit results are summarized. in Table 1.., The best fit results are summarized in Table \ref{bestfit}.788 All errors refer to confidence. level for one interesting xuwameter., All errors refer to confidence level for one interesting parameter.789 Lois worth noting that the measured ion hea ine cnerey constrains. at the above confidence level. the ionization state of iron (before the photoionization) to be ναοσα and.XUL (Llouse 1969). or to be exactly if the MOSI value is considered.," It is worth noting that the measured iron $\alpha$ line energy constrains, at the above confidence level, the ionization state of iron (before the photoionization) to be between and (House 1969), or to be exactly if the MOS1 value is considered."790 While of course. given he present uncertainties on the instruments calibrations. hese results cannot be taken too literally. they highlight he potentially extraordinary precision of the measure.," While of course, given the present uncertainties on the instruments calibrations, these results cannot be taken too literally, they highlight the potentially extraordinary precision of the measure."791 More solid constraints on the ionization states can be derived fron he Fe /Ixeo lines ratio. as ciiscussed in Sec.," More solid constraints on the ionization states can be derived from the Fe $\beta$ $\alpha$ lines ratio, as discussed in Sec."792 3.3., 3.3.793 The iron abundance is directly. measured. by the depth of the iron edge in the Compton reflection continuum. and it is measured with respect to the clements responsible for the photoabsorption below the edge. mainly. oxvecn and neon," The iron abundance is directly measured by the depth of the iron edge in the Compton reflection continuum, and it is measured with respect to the elements responsible for the photoabsorption below the edge, mainly oxygen and neon"794We intend to expand this work to higher redshifts usingCOMBO-17+4. the NIR extension of COMBO-17. for which observations are currently being carried out using the 24°2 Omega2000 camera at the 3.5m-telescope on Calar Alto. Spain.,"We intend to expand this work to higher redshifts using, the NIR extension of COMBO-17, for which observations are currently being carried out using the $2k\times2k$ Omega2000 camera at the 3.5m-telescope on Calar Alto, Spain."795 Combining the existing optical data base from COMBO-17 with NIR observations in one broad and three medium band filters (covering the wavelength range from 1040 to nnm). we expect to obtain ~1200 galaxy redshifts with an accuracy οσα|2)=0.02 up to :=2.," Combining the existing optical data base from COMBO-17 with NIR observations in one broad and three medium band filters (covering the wavelength range from 1040 to nm), we expect to obtain $\simeq 4200$ galaxy redshifts with an accuracy of $\sigma_z/(1+z)=0.02$ up to $z=2$."796 This much longer baseline in cosmic time will allow us to observe much larger evolution of halo masses. testing the idea of hierarchical growth back to a time close to the formation of luminous galaxies.," This much longer baseline in cosmic time will allow us to observe much larger evolution of halo masses, testing the idea of hierarchical growth back to a time close to the formation of luminous galaxies."797FL5400 and a power-law index a.,$F_{5100}$ and a power-law index $\alpha$.798 We chose three fitting ranges. 2200-2230A.. 4180-4220 aane 5050-5100A.. as continuum windows. since these area have little emission lines (see Figure 5)).," We chose three fitting ranges, 2200-2230, 4180-4220 and 5050-5100, as continuum windows, since these area have little emission lines (see Figure \ref{fig:cont_fit}) )."799 Phere are. however. the Balmer continuum and the ppseuco-continuum underneath these regions. requiring some corrections.," There are, however, the Balmer continuum and the pseudo-continuum underneath these regions, requiring some corrections."800 Tsuzukietal.(2006). gives 14 quasar spectra covering a wide wavelength range and measured accurately their continuum levels., \citet{tzk} gives 14 quasar spectra covering a wide wavelength range and measured accurately their continuum levels.801 We fitted power-law continuum models to their spectra in the continuum windows. and compared the continuum levels with those given by Tsuzukictal.(2006).," We fitted power-law continuum models to their spectra in the continuum windows, and compared the continuum levels with those given by \citet{tzk}."802. We found that our method systematically overestimates the continuum levels. at 2200-2230A. at 4180-4220 aand at 5050-5100A.," We found that our method systematically overestimates the continuum levels, at 2200-2230, at 4180-4220 and at 5050-5100."803 According to these results. we first reduced the (ux densities of the object by these amount ab each continuum window. then fitted the power-LIaw continuum.," According to these results, we first reduced the flux densities of the object by these amount at each continuum window, then fitted the power-law continuum."804 n example of the fitted power-law continuum is indicated as the dashed line in Figure 5.., An example of the fitted power-law continuum is indicated as the dashed line in Figure \ref{fig:cont_fit}.805 Phe measurement error of the continuum levels is estimated to be less than., The measurement error of the continuum levels is estimated to be less than.806. Grandi.(1982) eives à formula describing the Balmer continuum produced. by a uniform. temperature. partiallv optically thick cloud: where Ον) is the Planck function at the electron temperature 7). and ree is the optical depth at the Balmer edge at A=3646A.," \citet{gra} gives a formula describing the Balmer continuum produced by a uniform temperature, partially optically thick cloud: where $B_\lambda(T_e)$ is the Planck function at the electron temperature $T_e$, and $\tau_{BE}$ is the optical depth at the Balmer edge at $\lambda = 3646$."807. Ixurk.etal.(2007). assumed gas clouds of uniform.temperature (7;;=15.000 Ix) and the optical depth fixed to Tee=I. and fit equation (8)) to their sample quasar spectra to estimate the strength of the Balmer continuum (see also Dietrichetal. 20033).," \citet{kurk} assumed gas clouds of uniformtemperature $T_e=15,000$ K) and the optical depth fixed to $\tau_{BE}=1$, and fit equation \ref{eq:bac}) ) to their sample quasar spectra to estimate the strength of the Balmer continuum (see also \citealt{die03}) )."808 We followed their method and assumed 7;=15.000 Ix. zzz=1.," We followed their method and assumed $T_e=15,000$ K, $\tau_{BE}=1$."809 The only one parameter. namely the scale factor foe. is set free and is decided by fitting equation (8)) to the power-law subtracted spectrum at 3645A.," The only one parameter, namely the scale factor $F_{BaC}$, is set free and is decided by fitting equation \ref{eq:bac}) ) to the power-law subtracted spectrum at $-$ 3645."810 An example of the fitted Balmer continuum is indicated as the dotted line in Figure 5.., An example of the fitted Balmer continuum is indicated as the dotted line in Figure \ref{fig:cont_fit}.811 Since hhas enormous energy levels. neighboring enmüssion lines contaminate heavily with each other. which makes it cillicult to measure the eemission lines.," Since has enormous energy levels, neighboring emission lines contaminate heavily with each other, which makes it difficult to measure the emission lines."812 One approach to measure the cemission lines is to use templates., One approach to measure the emission lines is to use templates.813 So far. several templates are derived from the narrow-line Sevfert 1 galaxy. 1 Zw I.," So far, several templates are derived from the narrow-line Seyfert 1 galaxy, I Zw 1."814 In the UV. Vestergaard&Wilkes(2001). ane etal.(2006). give their templates.," In the UV, \citet{vw} and \citet{tzk} give their templates."815" The template given by Vestergaard.&Wilkes(2001) do not cover around. πο, F", The template given by \citet{vw} do not cover around line.816suzukiοἱal.(2006) used a synthetic. spectrum calculated with the Cloudy photoionization code in order ο separate the comission from the line. and derived semiempiricallv the template which covers around the line.," \citet{tzk} used a synthetic spectrum calculated with the Cloudy photoionization code in order to separate the emission from the line, and derived semiempirically the template which covers around the line."817 Since we want to measure the comission line. we decided to use the UV template given by Tsuzukietal.(2006).," Since we want to measure the emission line, we decided to use the UV template given by \citet{tzk}."818. In the optical. Véron-C'etty.Joly&Veron(2004). and Tsuzukietal.(2006). open their templates to the publie.," In the optical, \citet*{vjv} and \citet{tzk} open their templates to the public."819 Véron-C'ettyietal.(2004) carefully analyzed the comission lines in 1 Zw 1. finding that the lines are emitted. from both BLR and Narrow Line Region LIU.," \citet{vjv} carefully analyzed the emission lines in I Zw 1, finding that the lines are emitted from both BLR and Narrow Line Region (NLR)."820 Phev succeeded. to separate them and called the woad line svstem L1 and the narrow Dine system N3. respectively.," They succeeded to separate them and called the broad line system L1 and the narrow line system N3, respectively."821 “Psuzukietal.(2006)— also. analyzed the μα»ectrum of L Zw 1 and derived the optical emplate. which was however not separated into the BL and the NLR components.," \citet{tzk} also analyzed the spectrum of I Zw 1 and derived the optical template, which was however not separated into the BLR and the NLR components."822 We applied both the broad line system L1 template given by Véron-C'ettyetal.(2004). and 16 optical emplate given. by Fsuzukictal.(2006) to. all of our samples. finding that the latter has a slightly smaller average vo value (median xzLAS for suzukietal.2006... while meclian M~1.58 for Véron-C'ettvetal. 20049).," We applied both the broad line system L1 template given by \citet{vjv} and the optical template given by \citet{tzk} to all of our samples, finding that the latter has a slightly smaller average $\chi_\nu^2$ value (median $\chi_\nu^2 \sim 1.48$ for \citealt{tzk}, while median $\chi_\nu^2823\sim 1.58$ for \citealt{vjv}) )."824 Here we adopt to use the optical template given by Tsuzukietal. (2006)., Here we adopt to use the optical template given by \citet{tzk}. .825 Prior to applying the template. to each quasar. broadening of the template spectrum is needed.," Prior to applying the template to each quasar, broadening of the template spectrum is needed."826 Thus we modeled the flux clensity as follows:, Thus we modeled the flux density as follows:827period at 91 s in one section (see Table 2)).,period at 91 s in one section (see Table \ref{dno7tab2}) ).828 The 45 s and 42 s periodicities are svnodic and sidereal DNOs. for which we would expect a QPO at ~ 460 s. and there is indeed a strong QPO at 420 s (see Fig. 9)).," The 45 s and 42 s periodicities are synodic and sidereal DNOs, for which we would expect a QPO at $\sim$ 460 s, and there is indeed a strong QPO at 420 s (see Fig. \ref{SALT005omc1}) )."829 Also seen in Fig., Also seen in Fig.830 9. is the switch from svnodic to sidercal period and back. as seen before in VW LEVE (see figure 12 of Paper IV which is switching of the second harmonic anc seen in PY PsA in figures 1 and 2 in Paper VI).," \ref{SALT005omc1} is the switch from synodic to sidereal period and back, as seen before in VW Hyi (see figure 12 of Paper IV – which is switching of the second harmonic – and seen in TY PsA in figures 1 and 2 in Paper VI)."831 One section of the light curve contains the 29 s and 45 s modulations present simultaneously., One section of the light curve contains the 29 s and 45 s modulations present simultaneously.832 1n the various runs we have identified many of the DNO and QPO components seen in previous studies — namely κο WOPO. DNOBoro. 2 WDNO- 2 WOO aud VIpDNO-," In the various runs we have identified many of the DNO and QPO components seen in previous studies – namely $\omega_{\rm DNO}$, $\omega_{\rm QPO}$ , $\omega_{\rm DNO} - \omega_{\rm QPO}$, 2 $\omega_{\rm DNO}$, 2 $\omega_{\rm QPO}$ and $\omega_{\rm lpDNO}$."833" 1n addition. for the first time we have seen the harmonic 2 (wpwo wore) early in the evolution. when the fundamental »eriod. was  22 s. and the harmonics 2(Gopyo were) and A(wpxo— woro) at the very end of an outburst. when the undamental period was  90 s. We will not attempt to explain these components beyond. what has already. been ooposed in Papers LE and IV. but we note the richness of he E""Es. which is a challenge for any model."," In addition, for the first time we have seen the harmonic 2 $\omega_{\rm DNO} - \omega_{\rm QPO}$ ) early in the evolution, when the fundamental period was $\sim$ 22 s, and the harmonics $\omega_{\rm DNO} - \omega_{\rm QPO}$ ) and $\omega_{\rm DNO} - \omega_{\rm QPO}$ ) at the very end of an outburst, when the fundamental period was $\sim$ 90 s. We will not attempt to explain these components beyond what has already been proposed in Papers II and IV, but we note the richness of the FTs, which is a challenge for any model."834 Our new observations add an interesting extension to he evolution of DNO periods and harmonics in VW Lyi., Our new observations add an interesting extension to the evolution of DNO periods and harmonics in VW Hyi.835 Fie., Fig.836 10. is based on figure |. of Paper IV. with addition of our new observations., \ref{DNOevolSALT} is based on figure 1 of Paper IV with addition of our new observations.837 We note first that our identification in Section 3.1 of the 24.78 s modulation as yy rather than Poyo results in better agreement with the Poxe evolution., We note first that our identification in Section 3.1 of the 24.78 s modulation as $P_{\rm SYN}$ rather than $P_{\rm DNO}$ results in better agreement with the $P_{\rm DNO}$ evolution.838 The DNOs observed at Z7 — LS d. and the later DNO mentioned in Sect.," The DNOs observed at $T$ $\sim$ 1.8 d, and the later DNO mentioned in Sect."839 2.2.4. demonstrate that ater Po Ld d the DNOs cease to increase. as was suspectecLin Paper IV.," 2.2.4, demonstrate that after $T$ $\sim$ 1.1 d the DNOs cease to increase, as was suspected in Paper IV."840 The maximum value of the implied. Lone is  90 s. which is essentially the same as the maximum vaue ol Ppxo in VW νι (Paper HI) the IpDNOs show little evolution in VW Lyi. increasing in. period by about through an outburst. ultimately converging on the same »eriod. as that of the fundamental DNO.," The maximum value of the implied $P_{\rm DNO}$ is $\sim$ 90 s, which is essentially the same as the maximum value of $P_{\rm lpDNO}$ in VW Hyi (Paper III) – the lpDNOs show little evolution in VW Hyi, increasing in period by about through an outburst, ultimately converging on the same period as that of the fundamental DNO."841 We thank Albert Jones for providing alerts of outbursts of VW. Livi. and Pauline Loader. for supplying a file of the archived. observations of VW νι collected. by the Roval Astronomical Society of New Zealand.," We thank Albert Jones for providing alerts of outbursts of VW Hyi, and Pauline Loader for supplying a file of the archived observations of VW Hyi collected by the Royal Astronomical Society of New Zealand."842 We kindly acknowledge the use of the AAWSO. observational archive of VN Livi for determining the shape of the 2008 January outburst light curve., We kindly acknowledge the use of the AAVSO observational archive of VW Hyi for determining the shape of the 2008 January outburst light curve.843 PAWs and DW's rescare ris supported by the University of Cape Town and the Naticnal Research Foundation of South Africa., PAW's and BW's research is supported by the University of Cape Town and the National Research Foundation of South Africa.844 AIF of the observaions reported in this paper were obtained with the Souhern. African Large Telescope (SALT). a consortium consisting of the National Research Foundation of South Africa. Nicholas Copernicus Astronomical Center of the Polis1 Academy of Sciences. Hobby. Eberly Telescope Founding 1nstitutions. Rutgers University. Georg-Xugust-Università| Gotttingen. University of Wisconsin. -. Madison. Carnegie Mellon University. University of Canterbury. United. Ixingdom SALT Consortium. University of North Carolina- Chapel Hil. Dartmouth College. American Museum. of Natural llistory and the InterUniversity Centre for Astronomy and Astrophysics. India.," All of the observations reported in this paper were obtained with the Southern African Large Telescope (SALT), a consortium consisting of the National Research Foundation of South Africa, Nicholas Copernicus Astronomical Center of the Polish Academy of Sciences, Hobby Eberly Telescope Founding Institutions, Rutgers University, Georg-August-Universitätt Götttingen, University of Wisconsin - Madison, Carnegie Mellon University, University of Canterbury, United Kingdom SALT Consortium, University of North Carolina- Chapel Hill, Dartmouth College, American Museum of Natural History and the Inter-University Centre for Astronomy and Astrophysics, India."845at a distance of 2170 pc and an age 10 Myr (?)..,at a distance of 2170 pc and an age 10 Myr \citep{bhatt94}.846 The nebulous cloud in the central region is in the background as inferred by ? based on the reddening determinations in different regions of the cluster., The nebulous cloud in the central region is in the background as inferred by \cite{pandey86} based on the reddening determinations in different regions of the cluster.847 NGC 637 has been studied by ? in the RGU photographic system., NGC 637 has been studied by \cite{grub75} in the $RGU$ photographic system.848 Photoelectric observations in the (DV. system were made by ? to obtain a distance of 2500 pc. reddening 0.66 and age 15 Myr.," Photoelectric observations in the $UBV$ system were made by \cite{huest91} to obtain a distance of 2500 pc, reddening 0.66 and age 15 Myr."849 ? also observed this cluster and obtained a younger age of 0-4 Myr., \cite{phelps94} also observed this cluster and obtained a younger age of 0–4 Myr.850" A conspicuous gap was found in its color-magnitude diagram, which is not a result of incompleteness of data."," A conspicuous gap was found in its color–magnitude diagram, which is not a result of incompleteness of data."851 ? presented V/ photometry of this cluster and monitored the cluster for variables., \cite{piet06} presented $VI$ photometry of this cluster and monitored the cluster for variables.852 NGC 189 is a young compact cluster in the vicinity of Stock 24 and Do 12., NGC 189 is a young compact cluster in the vicinity of Stock 24 and Do 12.853 It has been studied by ? and the distance to this cluster was found to be τοῦ pc., It has been studied by \cite{balazs61} and the distance to this cluster was found to be 790 pc.854" In this paper, we have used the 2MASS database."," In this paper, we have used the 2MASS database."855" The point-source 5/.N=10 limit is acheived at or fainter than /=15.8"", 17=15.1"" and N=14.35"" for virtually the entire sky and hence we have used the above criteria to extract the 2MASS data using Vizier!."," The point-source $S/N =10$ limit is acheived at or fainter than $J=15.8^{m}$, $H=15.1^{m} $ and $K =14.3^{m}$ for virtually the entire sky and hence we have used the above criteria to extract the 2MASS data using Vizier."856". Further, we have also added the constraint that photometric errors in each band are <0.2""."," Further, we have also added the constraint that photometric errors in each band are $\leq 0.2^{m}$."857" For accurate determination of the cluster parameters, it is essential to have the knowledge of the radial extent of the clusters."," 						 For accurate determination of the cluster parameters, it is essential to have the knowledge of the radial extent of the clusters."858 Mass segregation might lead to a larger ‘true’ cluster size than stated in the ? catalogue., Mass segregation might lead to a larger `true' cluster size than stated in the \cite{dias07} catalogue.859 As the 2MASS, As the 2MASS860mass per unit star formation rate) that is highest at early times.,mass per unit star formation rate) that is highest at early times.861 This scaling has been shown to give a much better match to the luminosity function of the Milky Wavy satellites iui mocdels that use a constant wind velocity (2).., This scaling has been shown to give a much better match to the luminosity function of the Milky Way satellites than models that use a constant wind velocity \citep{Okamoto2010a}.862 One further addition to the model is needed to ensure that SNe driven. winds act as intended., One further addition to the model is needed to ensure that SNe driven winds act as intended.863 ?/— showed. that standard. kinetic feedback is more ellective in low mass galaxies. where wind particles tend. to crag neighbouring eas out with them.," \citet{DVecchiaSchaye2008} showed that standard kinetic feedback is more effective in low mass galaxies, where wind particles tend to drag neighbouring gas out with them."864 In high mass galaxies on the other hand. 10 pressure of the ISM can be sullicient to prevent much of 10 mass in the wind from escaping.," In high mass galaxies on the other hand, the pressure of the ISM can be sufficient to prevent much of the mass in the wind from escaping."865 Since we wish to be able o prescribe the mass loading and wind velocity directly. we 'hoose to decouple wind particles from the hyvdrodsnamic alculation for a short time in order to allow them to escape 1e high density star forming regions.," Since we wish to be able to prescribe the mass loading and wind velocity directly, we choose to decouple wind particles from the hydrodynamic calculation for a short time in order to allow them to escape the high density star forming regions."866 When the density jas [allen to ny=O.OLeEm the particles feel the usual ivdrodyvnamic force again.," When the density has fallen to $n_{H}=0.01cm^{-3}$, the particles feel the usual hydrodynamic force again."867 Lo they do not reach sullicientlv ow densities after a time IOkpe/vasia. they are recoupled ΗΝΜΑΝ.," If they do not reach sufficiently low densities after a time $\rm 10kpc/v_{wind}$, they are recoupled anyway."868" When a gas particle receives SNe energy from a neighbouring star particle. the wind speed (ο) is obtained from the local velocity dispersion and then the particle is assigned a probability to be added to the wind: where AQ is the total feedback. energy. received by the gas particle and. miu, is the current mass of the SPILL particle."," When a gas particle receives SNe energy from a neighbouring star particle, the wind speed $v_{w}$ ) is obtained from the local velocity dispersion and then the particle is assigned a probability to be added to the wind: where $\Delta Q$ is the total feedback energy received by the gas particle and $m_{sph}$ is the current mass of the SPH particle."869 Note that an SPIEL particle's mass may increase if it receives mass from SNe or AGB stars in. neighbouring star particles. or decrease if it spawns a new star particle. which has a mass of half the original gas particle mass.," Note that an SPH particle's mass may increase if it receives mass from SNe or AGB stars in neighbouring star particles, or decrease if it spawns a new star particle, which has a mass of half the original gas particle mass."870 Hf p; exceeds unity. that is. if there is energy available in excess of that needed. to add the particle to the wind. then the extra energy. is distributed to the gas particles neighbours as an increase in internal energy.," If $p_{w}$ exceeds unity, that is, if there is energy available in excess of that needed to add the particle to the wind, then the extra energy is distributed to the gas particle's neighbours as an increase in internal energy."871 The direction in which wind particles are propelled is chosen at random to be parallel or anti-parallel to the vector (66—0)ο@yrae where (o is the velocity of the gas particle before it receives feedback energy. @groe is the gravitational acceleration. vector. pointing approximately to the local potential minimum (halo centre) and v9 is the bulk velocity of the halo. which we take to be the mean velocity of the eas particle’s dark matter neighbours.," The direction in which wind particles are propelled is chosen at random to be parallel or anti-parallel to the vector $(\vec{v_{0}}-\vec{\overline{v}})\times872\vec{a}_{grav}$ where $v_{0}$ is the velocity of the gas particle before it receives feedback energy, $\vec{a}_{grav}$ is the gravitational acceleration vector, pointing approximately to the local potential minimum (halo centre) and $\overline{v}$ is the bulk velocity of the halo, which we take to be the mean velocity of the gas particle's dark matter neighbours."873 The result of this treatment is a wind launched preferentially along an object’s rotation axis (7).., The result of this treatment is a wind launched preferentially along an object's rotation axis \citep{SpringelHernquist2003}.874 Our model for SNe winds cillers from that described bv 7. in two ways., Our model for SNe winds differs from that described by \citet{Okamoto2010a} in two ways.875 Firstly. we allow all gas particles. not. just. those above the star formation density threshold. to be added to the wind if they receive feedback energy.," Firstly, we allow all gas particles, not just those above the star formation density threshold, to be added to the wind if they receive feedback energy."876 The original prescription can result in à variable wind mass loading depending on how well the star forming region is resolved., The original prescription can result in a variable wind mass loading depending on how well the star forming region is resolved.877 Secondly. only type LE SNe contribute to the winds. type la SNe energy is added to the eas as thermal energy.," Secondly, only type II SNe contribute to the winds, type Ia SNe energy is added to the gas as thermal energy."878 Galaxies are identified using a version of the aalgorithm (7). adapted by 2.. which identifies self-bound structures ancl includes the internal energy. of gas when computing particle binding energies.," Galaxies are identified using a version of the algorithm \citep{Springel2001} adapted by \citet{Dolag2009}, which identifies self-bound structures and includes the internal energy of gas when computing particle binding energies."879 From the ~S5Alpe high resolution region. we select all galaxies within 280kpc of the centre of the most massive (central) galaxy.," From the $\rm \sim5Mpc$ high resolution region, we select all galaxies within 280kpc of the centre of the most massive (central) galaxy."880 ‘This distance was chosen to match the limiting maenitucle of the completeness-corrected. satellite Luminosity function constructed by 2.., This distance was chosen to match the limiting magnitude of the completeness-corrected satellite luminosity function constructed by \citet{Koposov2008}.881" Phe largest satellite in our Aq-C-4 run is resolved with about 1.510"" particles in total. ~3107 of which are star particles."," The largest satellite in our Aq-C-4 run is resolved with about $1.5\times10^5$ particles in total, $\sim3\times10^4$ of which are star particles."882 In the following. we consider all galaxies with more than ten star particles. which. taking into account the tvpical mass fraction lost. through stellar evolution for our choice of IME. implies a stellar mass limit of ~1.2.107M. for Aq-C-4.," In the following, we consider all galaxies with more than ten star particles, which, taking into account the typical mass fraction lost through stellar evolution for our choice of IMF, implies a stellar mass limit of $\rm883\sim1.2\times10^{5}M_{\odot}$ for Aq-C-4."884 Using a dark matter only CDMO) counterpart of our Aq-C-4 run. simulated as part of the Aquarius project. (7)... we have examined the extent to which the dynamics of the barvons alter the structure of clark matter (sub)halos. of satellite galaxies over the course of their formation.," Using a dark matter only (DMO) counterpart of our Aq-C-4 run, simulated as part of the Aquarius project \citep{Springel2008}, we have examined the extent to which the dynamics of the baryons alter the structure of dark matter (sub)halos of satellite galaxies over the course of their formation."885" The DALO run had. identical initial conditions to our Aq-C-4. but for the absence of barvons and a corresponcinely higher dark matter particle mass by a factor L/(OO,ο)."," The DMO run had identical initial conditions to our Aq-C-4, but for the absence of baryons and a correspondingly higher dark matter particle mass by a factor $\sim8861/(1-\Omega_{b}/\Omega_{m})$."887 Naively. one might simply compare cach subhalo with its DMO equivalent at 2=0. but this turns out to be problematic.," Naively, one might simply compare each subhalo with its DMO equivalent at $z=0$, but this turns out to be problematic."888 As has been noted in previous N-bocdy simulations at cillerent resolutions. small phase deviations in subhalo orbits get amplified over time. such that subhalos can be in quite different positions at z=0 (e.g..?2?7)..," As has been noted in previous N-body simulations at different resolutions, small phase deviations in subhalo orbits get amplified over time, such that subhalos can be in quite different positions at $z=0$ \citep[e.g.,][]{Frenk1999,Springel2008}."889 We see similar dillerences between Aq-C-4 and the DMO run., We see similar differences between Aq-C-4 and the DMO run.890 Subhalo orbits are also allected by other factors such as subhalo-subhalo scattering ancl variations in the potential due to small dillerences in the growth history of the main halo., Subhalo orbits are also affected by other factors such as subhalo-subhalo scattering and variations in the potential due to small differences in the growth history of the main halo.891 Since the streneth of tical shocking is strongly dependent on. pericentric distance. (77)... small orbital deviations can cause large dillerences in subhalo structure. which are entirely unrelated. to the presence or. absence of barvons.," Since the strength of tidal shocking is strongly dependent on pericentric distance \citep{Gnedin1999,Mayer2001}, small orbital deviations can cause large differences in subhalo structure, which are entirely unrelated to the presence or absence of baryons."892 This complication can be avoided. either. by choosing subhalos with no close pericentre. or by making the comparison at the epoch when the satellite is first accreted into the halo of the main galaxy. before the orbits have had a chance to diverge.," This complication can be avoided, either by choosing subhalos with no close pericentre, or by making the comparison at the epoch when the satellite is first accreted into the halo of the main galaxy, before the orbits have had a chance to diverge."893 We choose the latter option. since the former restricts us to a very small number of cases. although we note that one massive halo in a low eccentricity ( 0.2) orbit with a distant pericentre (200kpc) shows comparatively small dillerences in its dark matter density profile at 2=0 relative to the DMO case.," We choose the latter option, since the former restricts us to a very small number of cases, although we note that one massive halo in a low eccentricity $\sim0.2$ ) orbit with a distant pericentre $\rm \sim200kpc$ ) shows comparatively small differences in its dark matter density profile at $z=0$ relative to the DMO case."894 In the [ew instances where the accretion times of the subhalo ciller slightly between the hyerodyvnamical and DAIO runs. we choose the earlier of the two epochs.," In the few instances where the accretion times of the subhalo differ slightly between the hydrodynamical and DMO runs, we choose the earlier of the two epochs."895" In we show spherically averaged profiles for the dark matter density. (plotted as pr? to emphasise small differences) of our most massive satellites. for Aq-C-+4 and he DMO run at the output time when each satellite first joins the main [riend-of-friends. (FOL:?) ""The dillerences in the subhalo density. profiles with and without xwvons Clearly exceed the uncertainties. associated with inite sampling. indicated by the error bars."," In we show spherically averaged profiles for the dark matter density (plotted as $\rho r^{2}$ to emphasise small differences) of our most massive satellites, for Aq-C-4 and the DMO run at the output time when each satellite first joins the main friend-of-friends \citep[FOF;][]{Davis1985} The differences in the subhalo density profiles with and without baryons clearly exceed the uncertainties associated with finite sampling, indicated by the error bars."896 They. are also, They are also897hat caused event OGLE-2005-BLO-390. which is known to host a cool Super-Earth (2)..,"that caused event OGLE-2005-BLG-390, which is known to host a cool Super-Earth \citep{PLANET:planet}."898 It is easier to detect a planet than being able o claim that there are no other planets orbiting the same star. and if one aims for quantifying multiplicity. this needs to be addressed.," It is easier to detect a planet than being able to claim that there are no other planets orbiting the same star, and if one aims for quantifying multiplicity, this needs to be addressed."899 Planetary multiplicity however becomes an obvious johenomenon with the detection of respective systems. such as the pair of gas-giant planets orbiting OGLE-2006-BLG-|09L (2). which resemble a half-scale version of the Jupiter-Saturn yart of the Solar system.," Planetary multiplicity however becomes an obvious phenomenon with the detection of respective systems, such as the pair of gas-giant planets orbiting OGLE-2006-BLG-109L \citep{DoubleCatch}, which resemble a half-scale version of the Jupiter-Saturn part of the Solar system."900 Interestingly. the planets found to orbit HR 8799 look like the complementary double-scale version (2)..," Interestingly, the planets found to orbit HR 8799 look like the complementary double-scale version \citep{Marois:planet}."901 It is particularly striking that very early opportunities to detect such systems by gravitational microlensing or direct imaging. respectively. were successful. while one needs to keep in mind hat planetsi] with an orbital period similar to Saturn cannot be detected from radial-velocity surveys so far (given a 10-15 year iistory of respective campaigns). and observing planetary transits is further disfavoured by the small transit probability.," It is particularly striking that very early opportunities to detect such systems by gravitational microlensing or direct imaging, respectively, were successful, while one needs to keep in mind that planets with an orbital period similar to Saturn cannot be detected from radial-velocity surveys so far (given a 10–15 year history of respective campaigns), and observing planetary transits is further disfavoured by the small transit probability."902 However. ‘or Super-Earths and planets with Neptune-class masses in closer orbits. radial-velocity surveys find a very high level of multiplicity as well (22).," However, for Super-Earths and planets with Neptune-class masses in closer orbits, radial-velocity surveys find a very high level of multiplicity as well \citep{Mayor:abundance,HARPS:abundance2}."903. The detection of the pair of Jupiter- and Saturn-like planets orbiting OGLE-2006-BLG-109L. (2). is often hailed because of the striking similarity with the Solar system. albeit that there is basically nothing that can be said about potential inner rocky planets other than that such cannot be excluded.," The detection of the pair of Jupiter- and Saturn-like planets orbiting OGLE-2006-BLG-109L \citep{DoubleCatch} is often hailed because of the striking similarity with the Solar system, albeit that there is basically nothing that can be said about potential inner rocky planets other than that such cannot be excluded."904 There is however another important result arising from this discovery: outer giant planets are not of the lonesome tvpe., There is however another important result arising from this discovery: outer gas-giant planets are not of the lonesome type.905 How does one arrive at such a conclusion?, How does one arrive at such a conclusion?906 Regardless of the large detection efficiency for such planets in events with a peak magnification as large as that of OGLE-2006-BLG-109 (ly~ 290). the planetary abundance is moderate or small.," Regardless of the large detection efficiency for such planets in events with a peak magnification as large as that of OGLE-2006-BLG-109 $A_0 \sim 290$ ), the planetary abundance is moderate or small."907 Tf we consider an abundance of 5 per cent. the probability for a double catch would be just 0.25 per cent if the planets were drawn independently from the population.," If we consider an abundance of 5 per cent, the probability for a double catch would be just 0.25 per cent if the planets were drawn independently from the population."908 This would mean an expected detection of ~1/30 systems amongst the 13 events comprising the systematic sample reported by ?.. so that we would have been very lucky to find the detected pair.," This would mean an expected detection of $\sim\,1/30$ systems amongst the 13 events comprising the systematic sample reported by \citet{Gould:abundance}, so that we would have been very lucky to find the detected pair."909 Therefore. it appears the more likely assumption that the two detections were not the result of independent draws. but instead the probabilitv for a planet to orbit a star is larger if one considers a star that is known to host planets as compared to an arbitrarily chosen star that might host planets or not.," Therefore, it appears the more likely assumption that the two detections were not the result of independent draws, but instead the probability for a planet to orbit a star is larger if one considers a star that is known to host planets as compared to an arbitrarily chosen star that might host planets or not."910 This however means that it is not appropriate to consider a planetary mass function with planets randomly drawn from it. but instead one needs to distinguish between stars with or without planets. as the formalism suggested in the previous section does.," This however means that it is not appropriate to consider a planetary mass function with planets randomly drawn from it, but instead one needs to distinguish between stars with or without planets, as the formalism suggested in the previous section does."911 These arguments however get weaker if the planetary abundance was as large as 20 per cent. because this would mean a probability of + per cent for a pair. or 1/2 expected to be detected for 13 events as compared to the one found.," These arguments however get weaker if the planetary abundance was as large as 20 per cent, because this would mean a probability of 4 per cent for a pair, or 1/2 expected to be detected for 13 events as compared to the one found."912 The statistical analysis of microlensing events in order to derive planet abundance estimates provides an illustrative example of the challenges one is facing., The statistical analysis of microlensing events in order to derive planet abundance estimates provides an illustrative example of the challenges one is facing.913 If a planet orbits the lens star. a detectable signal will only arise with a finite probability.," If a planet orbits the lens star, a detectable signal will only arise with a finite probability."914 This finite detection efficiency for planets of given mass and orbital separation from their host star is of relevance not only for assessing abundances by means of detections but also for drawing conclusions from the absence of planetary signals., This finite detection efficiency for planets of given mass and orbital separation from their host star is of relevance not only for assessing abundances by means of detections but also for drawing conclusions from the absence of planetary signals.915 Moreover. the host stars of planets detected by gravitational microlensing arise stochastically from the underlying population of stars that intervene the observed targets. with current experiments most of the masses of the lens stars are only known up to a broad probability distribution. although the mass of the lens star is frequently known for events in which jxanetary signals have been detected (2)...," Moreover, the host stars of planets detected by gravitational microlensing arise stochastically from the underlying population of stars that intervene the observed targets, with current experiments most of the masses of the lens stars are only known up to a broad probability distribution, although the mass of the lens star is frequently known for events in which planetary signals have been detected \citep{Gould:abundance}."916 The lack of information about the planet's host star is troublesome. since it is important to distinguish planet population statistics in the range of stellar masses between 0.1 and 0.5 AZ... covered by mierolensing. given that current planet-formation models predict substantial differences. in xirticular for the abundance of gas-giant planets (2)..," The lack of information about the planet's host star is troublesome, since it is important to distinguish planet population statistics in the range of stellar masses between 0.1 and 0.8 $M_\odot$, covered by microlensing, given that current planet-formation models predict substantial differences, in particular for the abundance of gas-giant planets \citep{IdaLin}."917 For the rather small planet samples acquired so far. let us yjowever neglect this issue for the time being. and just compare jxanet abundance estimates that refer to the sample of probed ens stars.," For the rather small planet samples acquired so far, let us however neglect this issue for the time being, and just compare planet abundance estimates that refer to the sample of probed lens stars."918 Recently. there have been some discussions about jxanetary mass functions that can be extracted from microlensing observations.," Recently, there have been some discussions about planetary mass functions that can be extracted from microlensing observations."919 While the discussion by ?. is not based on a well- criterion for selecting the considered 160 planet detections rom the so far published 24 candidates towards the Galactic bulge (2).. and moreover no relation has been given between these ‘detections’ and the efficiency of the full observing campaigns. ? in contrast adopted selection criteria that lead to a well-detined event sample. and evaluated the detection efficiencies properly.," While the discussion by \citet{Sumi:planet} is not based on a well-defined criterion for selecting the considered 10 planet detections from the so far published 24 candidates towards the Galactic bulge \citep{Dominik:review}, and moreover no relation has been given between these `detections' and the efficiency of the full observing campaigns, \citet{Gould:abundance} in contrast adopted selection criteria that lead to a well-defined event sample, and evaluated the detection efficiencies properly."920 However. hey refer to a planetary mass function described by means of the jxlanet-to-star mass ratio. whose value is highly questionable. given hat rather obviously one does not expect the same number of half-massive planets to form around half-massive stars.," However, they refer to a planetary mass function described by means of the planet-to-star mass ratio, whose value is highly questionable, given that rather obviously one does not expect the same number of half-massive planets to form around half-massive stars."921 In particular. coagulation and accretion processes depend on the masses of the bodies involved and their spatial density. but not on the mass of the star.," In particular, coagulation and accretion processes depend on the masses of the bodies involved and their spatial density, but not on the mass of the star."922 Nevertheless. the sample drawn by ? allows for an insightful urther look.," Nevertheless, the sample drawn by \citet{Gould:abundance} allows for an insightful further look."923 ? refer to 13 events with a peak magnification lo 200. densely monitored by (and other campaigns) rom 2005 to 2008.," \citet{Gould:abundance} refer to 13 events with a peak magnification $A_0 > 200$ , densely monitored by (and other campaigns) from 2005 to 2008."924" Amongst those events. 2 provided a signal hat indicates the presence of a massive gas-giant planet above 150 M, (0.5 AZ. namely OGLE-2006-BLG-109 and MOA-"," Amongst those events, 2 provided a signal that indicates the presence of a massive gas-giant planet above 150 $M_\oplus$ $0.5~M_\rmn{jup}$ ), namely OGLE-2006-BLG-109 and MOA-2007-BLG-400."925 For such planets. the detection efficiency for orbital separations that correspond to the “lensing can broadly be assumed to be of the order of 100 per cent (?)..," For such planets, the detection efficiency for orbital separations that correspond to the `lensing can broadly be assumed to be of the order of 100 per cent \citep{GS98}. ."926 One would therefore estimate the abundance of such planets to be about 15 per cent., One would therefore estimate the abundance of such planets to be about 15 per cent.927 Rather than just focussing on events with large peak magnifications. the PLANET collaboration (22). has acquired data on a much larger sample of about 50 events per year with li from 2002 to 2007. with sampling intervals of around 2 hrs or better. where the average detection efficiency for Jupiter-mass planets in the ‘lensing zone’ for such a sample is about [5 to 20 per cent (2).," Rather than just focussing on events with large peak magnifications, the PLANET collaboration \citep{PLANET:first,PLANET:EGS} has acquired data on a much larger sample of about 50 events per year with $A_0 > 2$ from 2002 to 2007, with sampling intervals of around 2 hrs or better, where the average detection efficiency for Jupiter-mass planets in the `lensing zone' for such a sample is about 15 to 20 per cent \citep{GL92}."928. Only one respective planet has been reported: OGLE-2005-BLG-071Lb (22).. as compared to expected 45-60 if those reside around each of the lens stars.," Only one respective planet has been reported: OGLE-2005-BLG-071Lb \citep{OB71,OB71:Dong}, as compared to expected 45–60 if those reside around each of the lens stars."929 This gives a rough abundance estimate of [.5—2 per cent. which looks substantially smaller than what one guesses from the densely monitored events with ly 200.," This gives a rough abundance estimate of 1.5–2 per cent, which looks substantially smaller than what one guesses from the densely monitored events with $A_0 > 200$ ."930 The PLANET team earlier claimed an upper abundance limit (at 95 per cent confidence) of 33 per cent on Jupiter-mass planets in the sameorbital range based on the absence of any detection amongst 42 events well-covered from 1995 to 1999 (??)..," The PLANET team earlier claimed an upper abundance limit (at 95 per cent confidence) of 33 per cent on Jupiter-mass planets in the sameorbital range based on the absence of any detection amongst 42 events well-covered from 1995 to 1999 \citep{PLANET:fiveshort,PLANET:fivelong}. ."931its spectral type aud hIuimositv class. aud its V-maenitude of 8.91 (Wrielt et al.,"its spectral type and luminosity class, and its V-magnitude of 8.91 (Wright et al."932 2003). the spectroscopic distance of IID 81032 is 110415 pe. iuplviug an N-rav luninosity of 9.2+L2«1000 Cresyes ! (based ou the RASS count rate aud aLass πιο conversion factor of 6«1012 erg ? PSPC 1: a self-consistent N-rav huninosity based ou the RASS data is calculated 1 88). and a radio Inniuositv of 1.640.7«1016 LT d where tre (stance is the «In13uimant contributor to the large errors in the eresΗννο," 2003), the spectroscopic distance of HD 81032 is $140 \pm 45$ pc, implying an X-ray luminosity of $9.2 \pm 4.2 \times 10^{30}$ erg $^{-1}$ (based on the RASS count rate and an assumed conversion factor of $6 \times 10^{-12}$ erg $^{-2}$ PSPC $^{-1}$: a self-consistent X-ray luminosity based on the RASS data is calculated in 8), and a radio luminosity of $1.6 \pm 0.7 \times 10^{16}$ erg $^{-1}$ $^{-1}$, where the distance uncertainty is the dominant contributor to the large errors in the luminosities."933 These valies clearly showmuacertainty that the star ΠΟ S1032 has au active corona., These values clearly show that the star HD 81032 has an active corona.934 Iu this paper. we preseut extensive optical photometryic and spectroscopic observations. as well as an analysis of arcTival Xaw data. of the star ΠΟ 51032.," In this paper, we present extensive optical photometric and spectroscopic observations, as well as an analysis of archival X-ray data, of the star HD 81032."935 Tjs ds the first detailed optical photomietric aud spectroscopic study of tus star. although a brief discussion of some of the optical photometry of this star was even in Paudey ct al. (," This is the first detailed optical photometric and spectroscopic study of this star, although a brief discussion of some of the optical photometry of this star was given in Pandey et al. ("9362002).,2002).937 Based on our loug-teru optical studs. the star TID 81032 is shown to be a new ienmber of the (loug-period) RS CVu class.," Based on our long-term optical study, the star HD 81032 is shown to be a new member of the (long-period) RS CVn class."938 The organization of his paper is as follows., The organization of this paper is as follows.939 The obxyvations and data reduction are preseuted in 2, The observations and data reduction are presented in \ref{obsdat.sec}.940 Tn we an the data for the periodicity.," In \ref{ligper.sec}, we analyse the data for the periodicity."941 In stbed.wO discuss the leh curve aud phase of uit 83.," In \ref{phopha.sec}, we discuss the light curve and phase of minimum light."942Chromosphlericuvse οwission features are descri 1n &h.., Chromospheric emission features are described in \ref{halcai.sec}.943 Tn &6 and 87.. the spectral type aud. IR helt.excess of ITD 81032 are discussed. while the N-vay spectra are discussed i18a.," In \ref{physical.sec} and \ref{sed.sec}, the spectral type and IR excess of HD 81032 are discussed, while the X-ray spectra are discussed in \ref{xspec.sec}."944 Finally. 9. stuumarizes the main results of this paper.," Finally, \ref{con.sec} summarizes the main results of this paper."945 IID 81032 was observed in the Johuson D. V aud Cousins R filter for 122 nights during four observing runs - vear 2000-2001 (19 nights). vear 2001-2002 (51 uightsj. vear 2002-2003 (21 uiehts) and vear 2003-2001 (25 nights) - at the Arvabhatta Research Institute o: Olervatioual Scicuces (ARTES).," HD 81032 was observed in the Johnson B, V and Cousins R filter for 122 nights during four observing runs - year 2000-2001 (19 nights), year 2001-2002 (54 nights), year 2002-2003 (24 nights) and year 2003-2004 (25 nights) - at the Aryabhatta Research Institute of Observational Sciences (ARIES)."946" The observatious were made with the LObem Sampurnanand telescop Caud using a 2h«2k CCD camera in vears 2000 - WwD.03, and a LhSTA CCD camera he vears 243 - 5001."," The observations were made with the 104-cm Sampurnanand telescope and using a $2k\times2k$ CCD camera in years 2000 - 2003, and a $1k\times1k$ CCD camera during the years 2003 - 2004."947 A few CCD frames were taken in the V and R filters on every might. with duringex)osure fines ranelic from 2 to 6) secs. epending upon the xcing conditions aud the filter used.," A few CCD frames were taken in the B, V and R filters on every night, with exposure times ranging from 2 to 60 secs, depending upon the seeing conditions and the filter used."948 Several bias aud twilight flat frames were also taken during cach run., Several bias and twilight flat frames were also taken during each observing run.949 Dias subtraction. flat and aperture photometry were performed usingIRAF*.," Bias subtraction, flat fielding and aperture photometry were performed using."950. For the observingstar IID 81032 the aud fieldingcheck stars wore TYC 5171 1315 1 and USNO-À2.0 Yes]ectivelv.," For the star HD 81032 the comparison and check stars were TYC 5471 1345 1 and USNO-A2.0 0750-06845737, respectively."951 Differential comparisonphotometry. in the SCLISC οf variable minus the comparison star. was done. since program. comparison aud check stars were ¢ul in tje sale CCD frame.," Differential photometry, in the sense of variable minus the comparison star, was done, since program, comparison and check stars were all in the same CCD frame."952 UBVRE observations ο “TD 81032 along with some of the Landol (1992) standard reeion SA 98 woYe obtained on 22 February 2OL for photometric calibration., UBVRI observations of HD 81032 along with some of the Landolt (1992) standard region SA 98 were obtained on 22 February 2004 for photometric calibration.953" The average maguitucdes UD 81032 duriie the observing vears 2003 2001in the U. D. VOR and I filters were 10.2[20.007b.9.of335d0.005.8.61£0,005.8.101+0.003. and b.73640.003, respectively."," The average magnitudes of HD 81032 during the observing years 2003 - 2004 in the U, B, V, R and I filters were $10.242 \pm 0.007, 9.635 \pm 0.005, 8.614 \pm 0.005, 8.401 \pm 0.003,$ and $7.736 \pm 0.003$, respectively."954 Spectroscopic observations were carried out durius 2003 Jaunary 20 to 21 at the Vainn Bappu Observatory. Kavalur with the OMR spectrograph fed by the 231-12. Vainn Bappu Telescope.," Spectroscopic observations were carried out during 2003 January 20 to 24 at the Vainu Bappu Observatory, Kavalur with the OMR spectrograph fed by the 234-cm Vainu Bappu Telescope."955 The data were acquired with a 1021« CCD camera of 21&2Lyi sqtare pixel size covering a range of 120041 aud having a dis]sion of 1.254/prec., The data were acquired with a $1024\times1024$ CCD camera of $ 24\times 24 \mu m$ square pixel size covering a range of $1200 \AA$ and having a dispersion of $1.25 \AA/pixel$.956 The star was observed in the waveleueth ranges of 3500 - L700 and 5700 - 6900., The star was observed in the wavelength ranges of 3500 - 4700 and 5700 - 6900.957À.. Four of ITD 81032 in the Πα 1 ale three iu the Call IT aud Is were obtained., Four spectra of HD 81032 in the $\alpha$ region and three in the CaII H and K region were obtained.958 asienal-to-noise of spectraratio between 20 to. LO was regioacliievec in these spectra., a signal-to-noise of ratio between 20 to 40 was achieved in these spectra.959 IID 71952. a regionIKOIV. type star. also ο a reference 81032.," HD 71952, a K0IV type star, was also observed as a reference star for HD 81032."960 wasThe spectra servedwereas extracted starusing forthe ITDstaudar doreuction procedures in the IRAF packages (bias subraction. flat fielding. extraction of the spectrum aid waveleneth calibration using arc laps).," The spectra were extracted using the standard reduction procedures in the IRAF packages (bias subtraction, flat fielding, extraction of the spectrum and wavelength calibration using arc lamps)."961 The spec‘tral τεolutiou was determined bv using enission ines of arc lamps taken on the same nights., The spectral resolution was determined by using emission lines of arc lamps taken on the same nights.962 A spectral resolution (6A) of 2.74 at GBOOA aud 3 Aat 10OA was achieved., A spectral resolution $\delta \lambda$ ) of $2.7 \AA$ at $6300 \AA$ and $3.7 \AA$ at $4000 \AA$ was achieved.963 All the spectra were normalized to the counntun and equivalent widths for the euissio1 lines were computed using the IRAF tasksplot., All the spectra were normalized to the continuum and equivalent widths for the emission lines were computed using the IRAF task.964 The errors iu he iieasureimenut of the withs of the emissiou lines were determuned by mcasirine the equivaleu widths of some moderate equivalent lines for cach , The errors in the measurement of the equivalent widths of the emission lines were determined by measuring the equivalent widths of some moderate absorption lines for each spectra.965We the staudard deviation or the equivalent width of cach feature. absorptionand finally determined spectra.the mean computedstandard deviation to be (0.02.4.," We computed the standard deviation for the equivalent width of each feature, and finally determined the mean standard deviation to be $ 0.02 \AA$ ."966of reionization.,of reionization.967 Clearly the process is much more rapid in f250C than in the T-QSO simulation. which in turn is more rapid than the T-star simulation.," Clearly the process is much more rapid in f250C than in the T-QSO simulation, which in turn is more rapid than the T-star simulation."968 In Fig., In Fig.969 3 a dip in the skewness is discernible in T-QSO. but is much less obvious than the dip in f250C because of fluctuations at high redshift. and because it spans a wider range in redshift due to the more extended reionization.," \ref{fig:skewandmean} a dip in the skewness is discernible in T-QSO, but is much less obvious than the dip in f250C because of fluctuations at high redshift, and because it spans a wider range in redshift due to the more extended reionization."970 This may be an indication that a more extended reionization process will be harder to detect using the skewness., This may be an indication that a more extended reionization process will be harder to detect using the skewness.971 The T-star simulation provides some hope. however.," The T-star simulation provides some hope, however."972 Despite an even more gradual reduction in the mean differential brightness temperature. the skewness is negative for quite a large range in redshift.," Despite an even more gradual reduction in the mean differential brightness temperature, the skewness is negative for quite a large range in redshift."973 Because the density field is positively skewed. a negative skewness is a clear signature of reionization.," Because the density field is positively skewed, a negative skewness is a clear signature of reionization."974 It is therefore possible that the skewness could provide a detection even in the case of very extended reionization., It is therefore possible that the skewness could provide a detection even in the case of very extended reionization.975 We now proceed to test the possibilities for signal extraction using the skewness. starting with a rather more optimistic case than will be encountered with the actual LOFAR EoR experiment.," We now proceed to test the possibilities for signal extraction using the skewness, starting with a rather more optimistic case than will be encountered with the actual LOFAR EoR experiment."976 We first note that the 5°ϱ field of one LOFAR EoR window corresponds to a distance of approximately SOOAlpe.500 (comoving) at >=10 in our assumed cosmology.," We first note that the $5^{\circ}\times 5^{\circ}$ field of one LOFAR EoR window corresponds to a distance of approximately $800\ \mathrm{Mpc}\times977800\ \mathrm{Mpc}$ (comoving) at $z=10$ in our assumed cosmology."978 At each redshif we therefore tile copies of the simulation to produce a slice of the correct size. then interpolate this onto a 2567 mesh.," At each redshift we therefore tile copies of the simulation to produce a slice of the correct size, then interpolate this onto a $256^2$ mesh."979 Since each pixel will be affected differently by foregrounds and noise. and since we consider only one-point statistics. we do not anticipate tha this will strongly affect our conclusions.," Since each pixel will be affected differently by foregrounds and noise, and since we consider only one-point statistics, we do not anticipate that this will strongly affect our conclusions."980 To produce our datacube. we add the simulated foregrounds described in Section 1ὁ to the CS. then smooth each slice using a Gaussian kernel to the estimated LOFAR resolution of =4arcmin.," To produce our datacube, we add the simulated foregrounds described in Section \ref{subsec:fg} to the CS, then smooth each slice using a Gaussian kernel to the estimated LOFAR resolution of $\approx 4\ \mathrm{arcmin}$."981 We then add uncorrelatec random noise as described by ?.., We then add uncorrelated random noise as described by\citet{JEL08}. .982 The noise has an of 52mlx at 150MlIz and has two contributions: a frequency dependen component coming from the sky. which scales as sv2. and a frequency-independent part from the receivers.," The noise has an of $52\983\mathrm{mK}$ at $150\ \mathrm{MHz}$ and has two contributions: a frequency dependent component coming from the sky, which scales as $\nu^{-2.55}$, and a frequency-independent part from the receivers."984 The noise on each image pixel is independent., The noise on each image pixel is independent.985 In reality. this will not be the case: rather. the noise on individual visibilities will be independent.," In reality, this will not be the case: rather, the noise on individual visibilities will be independent."986 We will tackle this more difficult case with realistic noise and a non-Gaussian point spread function below., We will tackle this more difficult case with realistic noise and a non-Gaussian point spread function below.987 Spatial slices are separated by 0.5MlIz in frequency. p.," Spatial slices are separated by $0.5\ \mathrm{MHz}$ in frequency, $\nu$."988 At 150MllIz this corresponds to a difference in redshift. A>=0.03. the slices having a comoving hickness of approximately 7+ Alpe.," At $150\ \mathrm{MHz}$ this corresponds to a difference in redshift, $\Delta z\approx 0.03$, the slices having a comoving thickness of approximately $7\ h^{-1}\ \mathrm{Mpc}$ ."989 Once we have a datacube with EoR signal. foregrounds and noise. we fit a third-order polynomial in log£ to each pixel.," Once we have a datacube with EoR signal, foregrounds and noise, we fit a third-order polynomial in $\log\nu$ to each pixel."990 We lave experimented with using different functional forms. but find hat so long as we obtain a visually reasonable fit. our results or the skewness do not change enough to affect our conclusions.," We have experimented with using different functional forms, but find that so long as we obtain a visually reasonable fit, our results for the skewness do not change enough to affect our conclusions."991 Teasurements of the variance are rather more sensitive to under- and over-fitting. which demonstrates the importance of understanding the foregrounds well. and of using robust statistics.," Measurements of the variance are rather more sensitive to under- and over-fitting, which demonstrates the importance of understanding the foregrounds well, and of using robust statistics."992 It is also possible to estimate the foregrounds by removing noise using a filtering procedure., It is also possible to estimate the foregrounds by removing noise using a filtering procedure.993 While this requires fewer assumptions about the nature of the foregrounds. it tends to over-tit.," While this requires fewer assumptions about the nature of the foregrounds, it tends to over-fit."994 After a fit has been obtained. it is subtracted from the total. leaving residuals which are an estimate of the CS plus the noise.," After a fit has been obtained, it is subtracted from the total, leaving residuals which are an estimate of the CS plus the noise."995 The skewness of these residuals as a function of redshift is shown in Fig. 4.., The skewness of these residuals as a function of redshift is shown in Fig. \ref{fig:rskew3}.996 While the CS which goes into the datacube is different for each panel of the plot. the noise and foregrounds are the same.," While the CS which goes into the datacube is different for each panel of the plot, the noise and foregrounds are the same."997 This accounts for the fact that the skewness as a function of redshift shows very similar features in each panel., This accounts for the fact that the skewness as a function of redshift shows very similar features in each panel.998 At first sight this seems rather discouraging. with the desired signal totally dominated by fitting errors (FEs) and noise.," At first sight this seems rather discouraging, with the desired signal totally dominated by fitting errors (FEs) and noise."999 The situation can be improved. however.," The situation can be improved, however."1000 At each frequency. we may denoise the residual image by smoothing.," At each frequency, we may denoise the residual image by smoothing."1001 This is possible because our images are oversampled. with more than one pixel per resolution element. and because there are no pixel-to-pixel correlations in the noise (by construction).," This is possible because our images are oversampled, with more than one pixel per resolution element, and because there are no pixel-to-pixel correlations in the noise (by construction)."1002 While this is clearly unrealistic. if serves as a prototype for the more difficult denoising step when we consider the dirty maps.," While this is clearly unrealistic, it serves as a prototype for the more difficult denoising step when we consider the dirty maps."1003 The effect of smoothing on the different components of the residual maps — these components being the CS. FEs and noise — is illustrated in Fig. 5..," The effect of smoothing on the different components of the residual maps – these components being the CS, FEs and noise – is illustrated in Fig. \ref{fig:scaledep},"1004 in which we show how the absolute value of the third moment of the one-point distribution of these components. [r/5|. is affected when they are smoothed with windows of different size.," in which we show how the absolute value of the third moment of the one-point distribution of these components, $|\mu_3|$, is affected when they are smoothed with windows of different size."1005 We show the result for the slice of the T-star datacube at 115MllIz. corresponding to a redshift of 11.35.," We show the result for the slice of the T-star datacube at $115\1006\mathrm{MHz}$, corresponding to a redshift of $11.35$."1007 This slice is chosen because the skewness of the T-star simulation is significantly negative here: we get similar results with the other simulations if we choose an appropriate slice in which the skewness is significantly different from zero., This slice is chosen because the skewness of the T-star simulation is significantly negative here; we get similar results with the other simulations if we choose an appropriate slice in which the skewness is significantly different from zero.1008 When the smoothing window is very narrow. so that there is almost no smoothing. [αἱ for the noise exceeds the value for the CS.," When the smoothing window is very narrow, so that there is almost no smoothing, $|\mu_3|$ for the noise exceeds the value for the CS."1009 This occurs even though (£775=0 (where the expectation is taken over different noise realizations). simply because the noise is so much larger than that of the (significantly skewed) cosmological signal.," This occurs even though $\langle\mu_3^\mathrm{noise}\rangle=0$ (where the expectation is taken over different noise realizations), simply because the noise is so much larger than that of the (significantly skewed) cosmological signal."1010 As the size of the smoothing window is increased. pes”liols| drops much more quickly than RAI since the smoothing averages together uncorrelated noise pixels. but correlated signal pixels.," As the size of the smoothing window is increased, $|\mu_3^\mathrm{noise}|$ drops much more quickly than $|\mu_3^\mathrm{CS}|$ since the smoothing averages together uncorrelated noise pixels, but correlated signal pixels."1011 At large scales. the signal also becomes uncorrelated. so its small means that [οςΣΤnelse once more.," At large scales, the signal also becomes uncorrelated, so its small means that $|\mu_3^\mathrm{CS}|<|\mu_3^\mathrm{noise}|$ once more."1012" The scale at which: [j/$C exceeds |yoke by the greatest amount in this residual map is ""|approximately 3— 4 aremin.", The scale at which $|\mu_3^\mathrm{CS}|$ exceeds $|\mu_3^\mathrm{noise}|$ by the greatest amount in this residual map is approximately $3$ $4$ $\mathrm{arcmin}$ .1013 Inthe case of the fitting errors. |i55|shows less variation as the smoothing scale is changed than either of theother components.," Inthe case of the fitting errors, $|\mu_3^\mathrm{FE}|$shows less variation as the smoothing scale is changed than either of theother components."1014 If the foreground fitting were completely unbiased.," If the foreground fitting were completely unbiased,"1015We have discussed previously PA-00-SL (PaulinIeuriksson et al.,We have discussed previously PA-00-S4 (Paulin-Henriksson et al.1016 2002)., 2002).1017" Tt lies 22/31"" from the centre of XBL. but uulike PA-99-N9. there ig a concentration of stellar lenses along the line of sieht."," It lies $22'31''$ from the centre of M31, but unlike PA-99-N2, there is a concentration of stellar lenses along the line of sight."1018" Tudeed. it is oulv 2/51"" from the ceutre of M32,"," Indeed, it is only $2'54''$ from the centre of M32."1019 The blue. color (RofF=0.04 0.101) of this event argues strouely for a source that lies in the disk of M31., The blue color $R-I=0.0\pm 0.14$ ) of this event argues strongly for a source that lies in the disk of M31.1020 Because M2 is believed to be iu the foreground (Byrd1976:Ford.JacobyandJeuuer 1978).. the proximity of the leus to the line of seht of M32 suggests that it is most naturally interpreted ἂν star in M32.," Because M32 is believed to be in the foreground \cite{byrd,ford}, the proximity of the lens to the line of sight of M32 suggests that it is most naturally interpreted as a star in M32."1021 However. this arguienut is not conclusive and a MACTIO interpretation is still possible.," However, this argument is not conclusive and a MACHO interpretation is still possible."1022for ReeCC is about 0.5kpe.,for $R_{GC}$ is about 0.5kpc.1023 Note that the ageo indicator in Friel’s work is based on the MAI (Morphological Age Index). whieh was only intended. to provide a relative age ranking of clusters. therefore. it is not fully comparable.," Note that the age indicator in Friel's work is based on the MAI (Morphological Age Index), which was only intended to provide a relative age ranking of clusters, therefore, it is not fully comparable."1024 But there still has σου overall correlation between Friels catalogue ancl ours., But there still has good overall correlation between Friel's catalogue and ours.1025 Our following analvsis will be mainly based on those two catalogues. but excluding cluster Berkeley 29.," Our following analysis will be mainly based on those two catalogues, but excluding cluster Berkeley 29."1026 In CAT 1. Berkeley 29 has the galactocentrie radius of 23 kpc. the E(B—V) of 0.15 and a metallicity of -0.18 dex. [rom the compilation of Diasetal.(2002).," In CAT 1, Berkeley 29 has the galactocentric radius of 23 kpc, the $E(B-V)$ of 0.15 and a metallicity of -0.18 dex, from the compilation of \citet{dia02}."1027. llowever. different values lor (hese parameters of Berkeley 29 were published in the literature.," However, different values for these parameters of Berkeley 29 were published in the literature."1028 Kaluzny(1994) gave a much smaller galactocentrie clistance of about 19 kpc. the reddening E(B—V) lager than 0.21 ancl based on the CMD morphology and comparison to other old clusters. he also deduced a [Fe/IH] value of lower than -1.," \citet{kal94} gave a much smaller galactocentric distance of about 19 kpc, the reddening $E(B-V)$ larger than 0.21 and based on the CMD morphology and comparison to other old clusters, he also deduced a [Fe/H] value of lower than -1."1029" In the work of Noriega-Mendoza (1997).. who applied a new technique for simultaneous determination of [Fe/II] and ΓΗ)—V). a [Fe/II]—-0.30 and a E(B—V)=0.01 were given for Berkeley 29,"," In the work of \citet{nor97}, who applied a new technique for simultaneous determination of [Fe/H] and $E(B-V)$, a [Fe/H]=-0.30 and a $E(B-V)=0.01$ were given for Berkeley 29."1030 As the properties of Berkeley 29 are quite uncertain we do not include this object in the following calculations., As the properties of Berkeley 29 are quite uncertain we do not include this object in the following calculations.1031 In general. (he uncertainty for the metallicity determinations in open clusters is about. 0.1dex.," In general, the uncertainty for the metallicity determinations in open clusters is about 0.1dex."1032 Usine data trom Diasοἱal.(2002). [or 571 open clusters with distance and age data. we plotted (he cluster positions on an (X.Y) coordinate svstem. with the zero point in X al the ealactic center (the Sun is assunied to be al 8.5 kpe) as Figure 1 shows.," Using data from \citet{dia02} for 571 open clusters with distance and age data, we plotted the cluster positions on an (X,Y) coordinate system, with the zero point in X at the galactic center (the Sun is assumed to be at 8.5 kpc) as Figure 1 shows."1033 IIere the full line are represents the solar circle about the ealactic center., Here the full line arc represents the solar circle about the galactic center.1034 One will [ind from this figure that in the ealactic plane. voung clusters (wilh ages vounger than that of Hyades. 0.8 Gyr. see Phelpsetal. (1994))) distributed cuite uniformly. around the Sun. while roughly only 20%. of the old clusters are inside the solar circle. most of the old ones are located further awav [rom the ealactie center than the Sun.," One will find from this figure that in the galactic plane, young clusters (with ages younger than that of Hyades, 0.8 Gyr, see \citet{phe94}) ) distributed quite uniformly around the Sun, while roughly only $20\%$ of the old clusters are inside the solar circle, most of the old ones are located further away from the galactic center than the Sun."1035 This result is quite consistent. with the early comprehensive study of Phelpsetal.(1994)., This result is quite consistent with the early comprehensive study of \citet{phe94}.1036. The deficiency. of older clusters in the inner part of the disk has been ascribed to the preferential destruction of the clusters when they encountered with eiut molecular clouds. which were primarily found in (he inner Galaxy.," The deficiency of older clusters in the inner part of the disk has been ascribed to the preferential destruction of the clusters when they encountered with giant molecular clouds, which were primarily found in the inner Galaxy."1037 The distribution of either old or voung clusters perpendicular to the galactic plane could be fitted by a simple exponential law. which are plotted in Figure 2.," The distribution of either old or young clusters perpendicular to the galactic plane could be fitted by a simple exponential law, which are plotted in Figure 2."1038 The vounger clusters are distributed on the galactic plane almost svnuuetrically about (he Sun. with a perpendicular scale height of approximately 57 pc.," The younger clusters are distributed on the galactic plane almost symmetrically about the Sun, with a perpendicular scale height of approximately 57 pc."1039 In contrast. about 80% old clusters are in the outer," In contrast, about $80\%$ old clusters are in the outer"1040wiltlis of the helini-like aud bydrogen-like thermal iron lines frou theaccretion flows.,widths of the helium-like and hydrogen-like thermal iron lines from the accretion flows.1041 From Table and Figure Ll. we can see iat the thera A-orayv line emission from the ADAFs without «mtfows is always auch lower iu ehussion from the ADAFs with outflows.," From Table 1 and Figure 4, we can see that the thermal X-ray line emission from the ADAFs without outflows is always much lower than emission from the ADAFs with outflows."1042 Moreover. the eqivaleut widths of he Te-like aud II-like thermal τοι1 lines siguificautly increase with 10 nass accretio Lrate Pas at the outer radius of ic ADAF and the outflow strenetL parameter p.," Moreover, the equivalent widths of the He-like and H-like thermal iron lines significantly increase with the mass accretion rate $\dot{m}_{\rm out}$ at the outer radius of the ADAF and the outflow strength parameter $p$."1043 Iu fact. these two ]aralucters are not independent. oees the outflow strength paramcter p is required ο increase with the mass accretion rate (mi at 1ο outer radius te| reproduce the observed. N-rav continluu Spectruni.," In fact, these two parameters are not independent, i.e., the outflow strength parameter $p$ is required to increase with the mass accretion rate $\dot{m}_{\rm out}$ at the outer radius to reproduce the observed X-ray continuum spectrum."1044 The sensitive ependence of aud EW» ou the two parameters can be explained bv the VOYV «ifereut deusity profi (seo Figure 1). lxCATS he temerature profi of the ulasmias are| quite simularor ADAFs with different αταλοςqs adoος (see.Figure2.aNaravan&Yi 1995).," The sensitive dependence of $_{1}$ and $_{2}$ on the two parameters can be explained by the very different density profiles (see Figure 1), because the temperature profiles of the plasmas are quite similarfor ADAFs with different parameters adopted \citep*[see Figure 2, and][]{ny95}."1045. As showed in Eqation (2)). the equivaleut width of a thera line is the rati of the line luminosity anc the spectral huuinosity at the enerev of the line.," As showed in Equation \ref{EW}) ), the equivalent width of a thermal line is the ratio of the line luminosity and the spectral luminosity at the energy of the line."1046 All our derived spec are required to fit the observed. N-rav. continuni ato6.83 keV. Thus. the οuvaleut widths of hermal iron lines will mainly depend on he iue bhuuinosities.," All our derived spectra are required to fit the observed X-ray continuum at $\sim 6.83$ keV. Thus, the equivalent widths of thermal iron lines will mainly depend on the line luminosities."1047 From Equation (3)). the line ο cutter| from a certain annulus of the accretion flow is determuned by the deusity and the μηine enissivitv. which is a function of the clectron cluperature.," From Equation \ref{l_line}) ), the line luminosity emitted from a certain annulus of the accretion flow is determined by the density and the line emissivity, which is a function of the electron temperature."1048 As tl1¢ temperature profiles are quite simular for the uxxlels with different parameters. he line hnünositv is mostly deermuned by the deusitv distributiπι of the eas iu the accretion flow.," As the temperature profiles are quite similar for the models with different parameters, the line luminosity is mostly determined by the density distribution of the gas in the accretion flow."1049 For the LOCels with larger p. ic. stronger outflows. he deusity profiles are flatter. which neans that mich nore gases are blown away frou he accretion How.," For the models with larger $p$, i.e., stronger outflows, the density profiles are flatter, which means that much more gases are blown away from the accretion flow."1050 This leads to higher elecron densities im the o1ter region of the accretion flow (rim 09) with t1ο electron. temperature X1o Ix. where the thernal IHe-Iike aud U-like iron lines enmuütted.," This leads to higher electron densities in the outer region of the accretion flow $r\gtrsim 10^{3}$ ) with the electron temperature $\lesssim10^8$ K, where the thermal He-like and H-like iron lines emitted."1051 Therefore. the lue luninosity aud the equivalent witlis of he thermal ion lines are hieher iu hese nxdels.," Therefore, the line luminosity and the equivalent widths of the thermal iron lines are higher in these models."1052interest for future study is whether the central engine is the same in the radio-Ioud quasars as it is in the racio-quiet quasars.,interest for future study is whether the central engine is the same in the radio-loud quasars as it is in the radio-quiet quasars.1053 For instance the structure of the accretion disk may be different in the raclio-loucl quasars: sensitive stuclies of the rellection component and iron line in the RLOs (ic. with XMM) may help to determine this., For instance the structure of the accretion disk may be different in the radio-loud quasars; sensitive studies of the reflection component and iron line in the RLQs (i.e. with XMM) may help to determine this.1054 The properties of theredio-quict quasars on the whole seem more complex., The properties of the quasars on the whole seem more complex.1055 As has been seen in this paper. there Is little or no dependence on the N-ray. continua of. quasers on luminosity and therefore presumably the black hole mass.," As has been seen in this paper, there is little or no dependence on the X-ray continua of quasars on luminosity and therefore presumably the black hole mass."1056 Llowever perhaps the one driving factor responsible for the individual properties of quasars may be the m ol the central engine (i.e. the ratio of mass accretion to the Iddington rate - or the Iddington ratio)., However perhaps the one driving factor responsible for the individual properties of quasars may be the $\dot{m}$ of the central engine (i.e. the ratio of mass accretion to the Eddington rate - or the Eddington ratio).1057 A high fractional accretion rate can result in the surface lavers of the disk becoming highly: photoionised. which subsequently. can have several effects on the X-ray spectra.," A high fractional accretion rate can result in the surface layers of the disk becoming highly photoionised, which subsequently can have several effects on the X-ray spectra."1058 Depending on the degree of ionisation. ionisec rather than neutral iron Ix. emission lines can dominate the disk reflection spectrum. as observed.," Depending on the degree of ionisation, ionised rather than neutral iron K emission lines can dominate the disk reflection spectrum, as observed."1059 Furthermore at even higher ionisations. the neutral reflection. component. (anc iron line emission) can appear to be weaker. particularly if the disk is Cully ionised to several Thomson depths (e.g. Navakshin 1999. Ross 1999). also in agreemen with the apparent properties of the more luminous raclio-uiet quasars.," Furthermore at even higher ionisations, the neutral reflection component (and iron line emission) can appear to be weaker, particularly if the disk is fully ionised to several Thomson depths (e.g. Nayakshin 1999, Ross 1999), also in agreement with the apparent properties of the more luminous radio-quiet quasars."1060 A further effect is that. for high. Eddington ratios. stronger soft. N-rav. emission can be produced.," A further effect is that, for high Eddington ratios, stronger soft X-ray emission can be produced."1061 This may parth arise as the intrinsic thermal emission. from the disk can become stronger (Ross. Fabian Mineshinge 1992).," This may partly arise as the intrinsic thermal emission from the disk can become stronger (Ross, Fabian Mineshinge 1992)."1062 Another possibility is that as the disk is more highlv ionised. it can become more rellective at. soft. X-rav energies. producing a steepening ofthe X-ray. spectrum. ab low energies (Le. as à result of the ionised disk reflection component).," Another possibility is that as the disk is more highly ionised, it can become more reflective at soft X-ray energies, producing a steepening of the X-ray spectrum at low energies (i.e. as a result of the ionised disk reflection component)."1063 Thus a high accretion rate (relative to Eddington) may explain the strong soft excesses in some of the objects considered. earlier., Thus a high accretion rate (relative to Eddington) may explain the strong soft excesses in some of the objects considered earlier.1064 Lt is also interesting to return to the question of the dichotomy between the broad. ancl narrow Line quasers ju was considered. earlier., It is also interesting to return to the question of the dichotomy between the broad and narrow line quasars that was considered earlier.1065 It has been postulated in the iterature that the narrow optical LL? lines may be an indicator of a high accretion rate (Pounds 1995. Laor 1997).," It has been postulated in the literature that the narrow optical $\beta$ lines may be an indicator of a high accretion rate (Pounds 1995, Laor 1997)."1066 Wo this is correct this could. indeed: account or the dilferences between the 2 types of objects., If this is correct this could indeed account for the differences between the 2 types of objects.1067 A high accretion rate may explain the strong soft excesses observed roth in the narrow-line quasars (6 out of S objects in this sample) and also the lower luminosity narrow-line Sevfert Is (NLSIs) in other samples (Vaughan 1999. Leighly 1999).," A high accretion rate may explain the strong soft excesses observed both in the narrow-line quasars (6 out of 8 objects in this sample) and also the lower luminosity narrow-line Seyfert 1s (NLS1s) in other samples (Vaughan 1999, Leighly 1999)."1068 Xs explained this can be caused by increased intrinsic disk emission or an increase in the rellectiveness of the disk in the soft. N-rav. band., As explained this can be caused by increased intrinsic disk emission or an increase in the reflectiveness of the disk in the soft X-ray band.1069 Also if the intrinsic clisk emission is stronger. this can also account for the steeper 2-10. X-ray slopes. via increased C'ompton cooling (Pounds 1995).," Also if the intrinsic disk emission is stronger, this can also account for the steeper 2-10 X-ray slopes, via increased Compton cooling (Pounds 1995)."1070 Additionally some of the narrow-line quasars also show evidence for ionised iron. Ix. emission. also indicative of a high ionisation disk and thus a high accretion. rate. although the evidence is tentative so far (also see Vaughan 1999).," Additionally some of the narrow-line quasars also show evidence for ionised iron K emission, also indicative of a high ionisation disk and thus a high accretion rate, although the evidence is tentative so far (also see Vaughan 1999)."1071 So the narrow-line quasars (as well as the NLSIs) may radiate at a relatively. high. fraction of the Ecelington rate. whereas in general the broad-lined: quasars may. be sub-ISddington. perhaps similar to the normal Sevfert 15. but with more massive central black holes.," So the narrow-line quasars (as well as the NLS1s) may radiate at a relatively high fraction of the Eddington rate, whereas in general the broad-lined quasars may be sub-Eddington, perhaps similar to the normal Seyfert 1s, but with more massive central black holes."1072 n important question to ask is whether there are any narrow-line quasars at higher recshilts (z>1)., An important question to ask is whether there are any narrow-line quasars at higher redshifts $>1$ ).1073 Finally the amount of soft. X-ray absorption towards quasars was found to inerease with redshift (also see Fiore 1998 for a similar analysis of qquasars)., Finally the amount of soft X-ray absorption towards quasars was found to increase with redshift (also see Fiore 1998 for a similar analysis of quasars).1074 This correlation is apparently robust. even when calibration cllects ancl uncertainties in the amount of local absorption are taken into account.," This correlation is apparently robust, even when calibration effects and uncertainties in the amount of local absorption are taken into account."1075 The main question that has arisen from this. is whether this absorption is intrinsic to the quasars or whether it originates from line-obsight matter.," The main question that has arisen from this, is whether this absorption is intrinsic to the quasars or whether it originates from line-of-sight matter."1076 Given the low-number density of high column systems (such as clamped Ly-a svstems) that could cause appreciable X-ray absorption (O'Flaherty Jakobsen 1997). the most likely scenario is that the bulk of this absorbing material is local to the quasars or the host galaxy environment.," Given the low-number density of high column systems (such as damped $\alpha$ systems) that could cause appreciable X-ray absorption (O'Flaherty Jakobsen 1997), the most likely scenario is that the bulk of this absorbing material is local to the quasars or the host galaxy environment."1077 As there seems to be a comparative lack of absorption in some racio-quict quasars (see Fiore 1998). this excess numav be associated with radio-loud quasars (also see Sambruna 1999).," As there seems to be a comparative lack of absorption in some radio-quiet quasars (see Fiore 1998), this excess may be associated with radio-loud quasars (also see Sambruna 1999)."1078 However further data (with. NMM and Chandra) is required to determine the exact. location and. cause of this absorption., However further data (with XMM and Chandra) is required to determine the exact location and cause of this absorption.1079 We thank the ssupport teams. ab GskC and ISAS. for their help.," We thank the support teams, at GSFC and ISAS, for their help."1080 1n xwtieular. thanks to Ken Pounds and Simon Vaughan for rool reading the paper and providing useful discussions.," In particular, thanks to Ken Pounds and Simon Vaughan for proof reading the paper and providing useful discussions."1081 We also thank the anonymous referee for providing sugecstions o improve the paper., We also thank the anonymous referee for providing suggestions to improve the paper.1082 This research made use of data obtained. from the Leicester Database ancl Archive Service (LEDAS) at the Department of Physics and Astronomy. Leicester University. Ulx. and the High Enerey Astrophysics Science Archive Researeh Center (IIEASARC). provided by ολοδν Goddard Space Flight Center.," This research made use of data obtained from the Leicester Database and Archive Service (LEDAS) at the Department of Physics and Astronomy, Leicester University, UK, and the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center."1083In the last two cases. the gas density decreases. smoothly with ealactocentric distance. without any sharp clrop in the gas density.,"In the last two cases, the gas density decreases smoothly with galactocentric distance, without any sharp drop in the gas density."1084 In that case. the outer limit of the OLL might characterize a gas density threshold below which the formation of bound. clusters is either inhibited. or delaved.," In that case, the outer limit of the OH might characterize a gas density threshold below which the formation of bound clusters is either inhibited or delayed."1085 1n the first case. the 5σας density drops sharply at. closer ealactocentric distance (1.0... 2=25 kkpce).," In the first case, the gas density drops sharply at closer galactocentric distance (i.e., $D \simeq 25$ kpc)."1086 Nevertheless. we emphasize that all four curves Gt well the observed. OLL profile in the range kkpe. thus accounting for the vast majority of the OIL GC's (less than ten per cent of the OLL GCs are located bevond. 2= 20kkpc).," Nevertheless, we emphasize that all four curves fit well the observed OH profile in the range kpc, thus accounting for the vast majority of the OH GCs (less than ten per cent of the OH GCs are located beyond $D \simeq 20$ kpc)."1087 This strengthens our hypothesis following which the mass density profile of the OLL GCS is tracing the cold barvonic material available to the star formation process some GGvr ago., This strengthens our hypothesis following which the mass density profile of the OH GCS is tracing the cold baryonic material available to the star formation process some Gyr ago.1088 An aclelitional issue of potential concern is that a small. but non negligeable. fraction of the ΟΙ may itself be an accreted component.," An additional issue of potential concern is that a small, but non negligeable, fraction of the OH may itself be an accreted component."1089 Several of the GCs in nearby dwarf ealaxies have colour-magnitucde diagram characteristics that are inelistineuishable from OLL GC's (see. es. Grebel Gallagher 2004 anc references. therein).," Several of the GCs in nearby dwarf galaxies have colour-magnitude diagram characteristics that are indistinguishable from OH GCs (see, e.g., Grebel Gallagher 2004 and references therein)."1090 Furthermore. Alackev Gilmore (2004) emphasize that. although the overwhelming contribution of clusters from external galaxies would be to the Galactic YII subsystem. there may also be a non-zero contribution to the ΟΙ subsystem. which they estimate of order 15 per cent or. equivalently. a dozen of GCs.," Furthermore, Mackey Gilmore (2004) emphasize that, although the overwhelming contribution of clusters from external galaxies would be to the Galactic YH subsystem, there may also be a non-zero contribution to the OH subsystem, which they estimate of order 15 per cent or, equivalently, a dozen of GCs."1091 Such accreted OLL cluster candidates may be identified on the basis of either an extended core radius or spatial motions more typical of YII objects. although these alone are not definite indicators of an extra-Galactic origin.," Such accreted OH cluster candidates may be identified on the basis of either an extended core radius or spatial motions more typical of YH objects, although these alone are not definite indicators of an extra-Galactic origin."1092 \lackey Gilmore (2004) find live OIL GCs in the first. category (NGC 6809. 6101. 7492. 5897 and Pal 15) and 6 GCs in the second (NGC 1904. 2298.5024.5904.6205.7089).," Mackey Gilmore (2004) find five OH GCs in the first category (NGC 6809, 6101, 7492, 5897 and Pal 15) and 6 GCs in the second (NGC 1904, 2298,5024,5904,6205,7089)."1093 We have thus relittecl ecuation (S)) to the observed OLL cistribution. after rejecting (1) the five GCs with extended: structure. (2) all eleven. GCs quoted. above plus Pal 2 (sce below).," We have thus refitted equation \ref{eq:rho_cold}) ) to the observed OH distribution, after rejecting (1) the five GCs with extended structure, (2) all eleven GCs quoted above plus Pal 2 (see below)."1094 These two observed spatial distributions are illustrated in Fie. 4.., These two observed spatial distributions are illustrated in Fig. \ref{fig:rho_OH_acc}.1095 In the first case (filled svmbols). one hardly detects a dillerence with respect to the distribution of the whole OL GCS (as in Figs. 1--3)).," In the first case (filled symbols), one hardly detects a difference with respect to the distribution of the whole OH GCS (as in Figs. \ref{fig:rho_OH_fits}- \ref{fig:rho_OH_modbis}) )."1096 Accordingly. the parameters of the fit are much similar to previously. (compare case (a) ancl case (b) in Table 3)).," Accordingly, the parameters of the fit are much similar to previously (compare case (a) and case (b) in Table \ref{tab:fit_rho_cold}) )."1097 In a second step (open svmbols in Fig., In a second step (open symbols in Fig.1098 4. and case (c) in Table 3)). we have removed. the complete sample of LL OLL GC's suspected of having been accreted as well as Pal 2.," \ref{fig:rho_OH_acc} and case (c) in Table \ref{tab:fit_rho_cold}) ), we have removed the complete sample of 11 OH GCs suspected of having been accreted as well as Pal 2."1099 In the Fe/11] vs UBR diagram. this one is located at the frontier between the OIL and the YII eroups and its nature is thus ill-determined.," In the [Fe/H] vs HBR diagram, this one is located at the frontier between the OH and the YH groups and its nature is thus ill-determined."1100 While we have previously considered it as an OLL GC. we now sort it in the YII group.," While we have previously considered it as an OH GC, we now sort it in the YH group."1101 Removing the so-defined sample of 12 GC's. the changes are more significant but still. the modified observed distribution is satisfactorily fitted by equation (8). the main point being that the model distribution overestimates the mass densitv of the OII subsystem. in the region 20 kkpe.," Removing the so-defined sample of 12 GCs, the changes are more significant but still, the modified observed distribution is satisfactorily fitted by equation (8), the main point being that the model distribution overestimates the mass density of the OH subsystem in the region $D \simeq 20$ kpc."1102 We thus conclude that our fit is robust. even though the actual radial distribution. of the ΟΙ mass. remains slightly uncertain as a few old. clusters may be accreted objects and. thus. interlopers with respect to the genuine initial Galactic GCS.," We thus conclude that our fit is robust, even though the actual radial distribution of the OH mass remains slightly uncertain as a few old clusters may be accreted objects and, thus, interlopers with respect to the genuine initial Galactic GCS."1103 In what follows. we consider the OLL mass distribution clisplaved in Figs.," In what follows, we consider the OH mass distribution displayed in Figs."1104 1 to 3 ancl equation (8) as our fiducial observed and theoretical radial mass density profiles. respectively.," \ref{fig:rho_OH_fits} to \ref{fig:rho_OH_modbis} and equation (8) as our fiducial observed and theoretical radial mass density profiles, respectively."1105 As announced in Section 2. we now put on a firmer foot the hypothesis following which the initial GCS mass. density profile is reasonably approximated by what it presently is.," As announced in Section 2, we now put on a firmer foot the hypothesis following which the initial GCS mass density profile is reasonably approximated by what it presently is."1106 Evolutionary processes act more elliciently upon mass GC's as well as on GCs located in the inner Galactic regions., Evolutionary processes act more efficiently upon low-mass GCs as well as on GCs located in the inner Galactic regions.1107 In other words. the evolution. with time of the GCS has been mostlv. determined. by the initial spatial cistribution of the clusters in the Galactic halo as well as by their initial mass spectrum.," In other words, the evolution with time of the GCS has been mostly determined by the initial spatial distribution of the clusters in the Galactic halo as well as by their initial mass spectrum."1108" In this respect. it is interesting to note that the sharp contrast between the fraction of surviving clusters and the ratio of the final to the initial total mass in clusters as obtained by AleLaughlin (1999) (see Section 2) arises from the choice of a power-law mass spectrum with a steep νο, 2) slope and. probing down to"," In this respect, it is interesting to note that the sharp contrast between the fraction of surviving clusters and the ratio of the final to the initial total mass in clusters as obtained by McLaughlin (1999) (see Section 2) arises from the choice of a power-law mass spectrum with a steep (i.e., –2) slope and probing down to"1109"metal content of galaxies and their environment, is the goal of many recent theoretical efforts.","metal content of galaxies and their environment, is the goal of many recent theoretical efforts."1110" Hydrodynamical simulations of galaxy formation over large cosmological volumes, while advancing rapidly over the past few years (e.g.Cen,Nagamine,&Ostriker2005;Oppen-etal.2010;Cen&Chisari 2011),, typically suffer from poor spatial and mass resolution — which limits their ability to follow self-consistently the venting of metals by small galaxies and the transport of heavy elements from their production sites into the environment."," Hydrodynamical simulations of galaxy formation over large cosmological volumes, while advancing rapidly over the past few years \citep[e.g.][]{Cen05,Oppenheimer08,Oppenheimer09, Wiersma09,Shen10,Cen11}, typically suffer from poor spatial and mass resolution – which limits their ability to follow self-consistently the venting of metals by small galaxies and the transport of heavy elements from their production sites into the environment."1111" And while with suitable wind prescriptions these simulations have shown that galactic outflows can potentially enrich the IGM to approximately the observed levels, it is not clear whether the same prescriptions are also able to produce realistically looking galaxies."," And while with suitable wind prescriptions these simulations have shown that galactic outflows can potentially enrich the IGM to approximately the observed levels, it is not clear whether the same prescriptions are also able to produce realistically looking galaxies."1112" To shed some light on the objects and processes responsible for seeding the circumgalactic and intergalactic medium with nuclear waste, we follow here a different approach."," To shed some light on the objects and processes responsible for seeding the circumgalactic and intergalactic medium with nuclear waste, we follow here a different approach."1113" We present results from a new cosmological N-body/smooth particle hydrodynamic (SPH) simulation of extreme dynamic range, a twin of the “Eris” simulation (Guedesetal. 2011)."," We present results from a new cosmological $N$ -body/smooth particle hydrodynamic (SPH) simulation of extreme dynamic range, a twin of the “Eris"" simulation \citep{Guedes11}."1114". Termed “ErisMC”, the simulations follows the assembly of a massive galaxy halo with a spline softening length of 120 pc and 26 million dark matter and SPH particles in the high-resolution region."," Termed “ErisMC"", the simulations follows the assembly of a massive galaxy halo with a spline softening length of 120 pc and 26 million dark matter and SPH particles in the high-resolution region."1115" The feedback from an active galactic nucleus is neglected, and a star formation recipe is adopted based on a high gas density threshold."," The feedback from an active galactic nucleus is neglected, and a star formation recipe is adopted based on a high gas density threshold."1116" Heating by supernovae occurs in a clustered fashion, and the resulting pressure-driven outflows at high redshifts remove low angular momentum metal-enriched gas."," Heating by supernovae occurs in a clustered fashion, and the resulting pressure-driven outflows at high redshifts remove low angular momentum metal-enriched gas."1117" As shown below, our detailed study of ErisMC’s CGM at z=3 reveals that late galactic superwinds — the result of recent star formation — account for only a small fraction of all the metals found between ~ 100 and 200 kpc from the center of the main host, and resolves (proper)two other sources for the heavy elements found in ErisMC’s environment: its satellite progenitors — which deposit their metals before and during infall — and the orbiting nearby dwarfs."," As shown below, our detailed study of ErisMC's CGM at $z=3$ reveals that late galactic superwinds – the result of recent star formation – account for only a small fraction of all the metals found between $\sim$ 100 and 200 (proper) kpc from the center of the main host, and resolves two other sources for the heavy elements found in ErisMC's environment: its satellite progenitors – which deposit their metals before and during infall – and the orbiting nearby dwarfs."1118 The paper is organized as follows., The paper is organized as follows.1119 In 2 we describe the ErisMC simulation., In 2 we describe the ErisMC simulation.1120 3 presents a detailed study of the galaxy CGM., 3 presents a detailed study of the galaxy CGM.1121 The origin of circumgalactic metals and the role played by satellite progenitors and nearby dwarfs in polluting ErisMC’s environment are discussed in 4., The origin of circumgalactic metals and the role played by satellite progenitors and nearby dwarfs in polluting ErisMC's environment are discussed in 4.1122 The properties of the supernova-driven outflows from the main host are analyzed in 5., The properties of the supernova-driven outflows from the main host are analyzed in 5.1123" Finally, in 6 we summarize our main results."," Finally, in 6 we summarize our main results."1124" Throughout this work, we use the metal mass fraction Zo=0.0142 for the solar abundance (Asplundetal.2009)."," Throughout this work, we use the metal mass fraction $Z_\odot=0.0142$ for the solar abundance \citep{Asplund09}."1125. The Eris suite of simulations was performed in aProbe 3-year cosmology running the parallel TreeSPH code (Wadsleyetal. 2004)., The Eris suite of simulations was performed in a 3-year cosmology running the parallel TreeSPH code \citep{Wadsley04}.1126". Details of the main simulation are given in Guedesetal.(2011),, and are briefly summarized here."," Details of the main simulation are given in \citet{Guedes11}, and are briefly summarized here."1127" The high-resolution region, 4 comoving Mpc on a side, is embedded in a low-resolution, dark matter-only, periodic box of 90 comoving Mpc on a side, and contains 13 million dark matter particles and an equal number of gas particles, for a final dark and gas particle mass of mpm=9.8x104Meo and mgpy=2x104 Mo, respectively."," The high-resolution region, 4 comoving Mpc on a side, is embedded in a low-resolution, dark matter-only, periodic box of 90 comoving Mpc on a side, and contains 13 million dark matter particles and an equal number of gas particles, for a final dark and gas particle mass of $m_{\rm DM}=9.8\times 10^4\,\msun$ and $m_{\rm SPH}=2\times 10^4\,\msun$ , respectively."1128" The gravitational softening length, eg, was fixed to 120 physical pc for all particle species from z=9 to the present, and evolved as 1/(1+2) from z=9 to the starting redshift of z=90."," The gravitational softening length, $\epsilon_G$ , was fixed to 120 physical pc for all particle species from $z=9$ to the present, and evolved as $1/(1+z)$ from $z=9$ to the starting redshift of $z=90$."1129" Compton cooling, atomic cooling, and metallicity-dependent radiative cooling at low temperatures (Mashchenkoetal. are included."," Compton cooling, atomic cooling, and metallicity-dependent radiative cooling at low temperatures \citep{Mashchenko06}1130 are included."1131 Auniform UVbackground modifies the 2006)ionization and excitation state of the gas and is implemented using, Auniform UVbackground modifies the ionization and excitation state of the gas and is implemented using1132"Thus, strong K-H instabilities do develop, despite the factor of 10 in the density contrast.","Thus, strong K-H instabilities do develop, despite the factor of 10 in the density contrast."1133" Although the contact discontinuity is distorted by the K-H instability, the overall shape of this surface is in fairly good agreement with the one predicted by the ? solution (Eq. 1))"," Although the contact discontinuity is distorted by the K-H instability, the overall shape of this surface is in fairly good agreement with the one predicted by the \cite{Wil96} solution (Eq. \ref{eq: ram}) )"1134" for small polar angles (0), and the stand-off distance matches the position of the contact discontinuity well (see Fig."," for small polar angles $\theta$ ), and the stand-off distance matches the position of the contact discontinuity well (see Fig."1135 6 [top])., \ref{fig: adanal} [top]).1136" In the thin-shell limit, the ISM and wind material are assumed to cool instantaneously and mix fully."," In the thin-shell limit, the ISM and wind material are assumed to cool instantaneously and mix fully."1137" In contrast, the forward and reverse shock of the adiabatic model have finite width, ~0.2 ppc, and as expected, there is little mixing between these two layers of gas."," In contrast, the forward and reverse shock of the adiabatic model have finite width, $\sim$ pc, and as expected, there is little mixing between these two layers of gas."1138" Overall, the adiabatic model demonstrated that the set-up and code can achieve results that are consistent with both theoretical expectations and previous studies."," Overall, the adiabatic model demonstrated that the set-up and code can achieve results that are consistent with both theoretical expectations and previous studies."1139 The models that include radiative cooling exhibit a global structure that is similar to the structure of the adiabatic model., The models that include radiative cooling exhibit a global structure that is similar to the structure of the adiabatic model.1140" In Fig. 7,,"," In Fig. \ref{fig: series},"1141" we show the density, temperature, and emissivity for models A-D. The four regions described above can easily be distinguished: the unshocked ISM and similarly unshocked, free-flowing stellar wind, separated by a layer of hot, shocked ISM (the forward shock) and shocked stellar wind (the reverse shock)."," we show the density, temperature, and emissivity for models A-D. The four regions described above can easily be distinguished: the unshocked ISM and similarly unshocked, free-flowing stellar wind, separated by a layer of hot, shocked ISM (the forward shock) and shocked stellar wind (the reverse shock)."1142" Based on their flow characteristics and morphology, models A-D can be grouped into two classes: the ‘slow’ models, A-C, with ISM densities z1 ccm""? and peculiar velocities x40 ss~', and the ‘fast’ model, D, with an ISM density of 0.3 cm? and a stellar velocity 73 ss~!."," Based on their flow characteristics and morphology, models A-D can be grouped into two classes: the `slow' models, A-C, with ISM densities $\gtrsim$ $^{-3}$ and peculiar velocities $\lesssim$ $^{-1}$, and the `fast' model, D, with an ISM density of 0.3 $^{-3}$ and a stellar velocity $\sim$ $^{-1}$."1143 The fluid properties of the slow models show significant departures from the adiabatic case (see Fig. 8));, The fluid properties of the slow models show significant departures from the adiabatic case (see Fig. \ref{fig: nprof}) );1144" the most obvious being that the radiative cooling results in lower post- temperatures, e.g. 7, for model B is an order of"," the most obvious being that the radiative cooling results in lower post-shock temperatures, e.g. $T_{\rm s}$ for model B is an order of"1145properties of ITZRC: are a consequeuce (at least in the UV. rest frame) of the nuclear activity.,properties of HzRG are a consequence (at least in the UV rest frame) of the nuclear activity.1146 Iu spite of this complexity. ITZRC. stil have a crucial role iu the uuderstaudius of galaxy formalon. since they are the oulv wav we have to study tliccarly stages ofcllipticals.," 	In spite of this complexity, HzRG still have a crucial role in the understanding of galaxy formation, since they are the only way we have to study the early stages of."1147 Since all powerful radio σαaxies at low redshift are ejut ellipticals and there is evidence that this is also the case at Dol (Best. Longair RGtteering 1998)). it is believed hat the host galaxies of rac10 galaxies at higher redshifts are also eiaut ellipticals (Peiterieci et al. 1999) ).," Since all powerful radio galaxies at low redshift are giant ellipticals and there is evidence that this is also the case at $z\sim$ 1 (Best, Longair Rötttgering \cite{best98}) ), it is believed that the host galaxies of radio galaxies at higher redshifts are also giant ellipticals (Pentericci et al. \cite{pente99}) )."1148 It is however necessary to understand how the nuclear activity influences what we see 1n orer to make a correct interpretation of he observed properlos., It is however necessary to understand how the nuclear activity influences what we see in order to make a correct interpretation of the observed properties.1149 Another interesting aspect is the relationship between the rapk phase(s) of star formation aud the formation and fuebiug of a massive black hole «ming the formation of these galaxies which are destined tc» become the massive ellipticals we see today., Another interesting aspect is the relationship between the rapid phase(s) of star formation and the formation and fuel`ing of a massive black hole during the formation of these galaxies which are destined to become the massive ellipticals we see today.1150 Exactly what his relationship is is no clear although interactions/merecrs are likely ο play a role., Exactly what this relationship is is not clear although interactions/mergers are likely to play a role.1151 It has been proposed that ultralumunous infrared ealaxies (ULIRGSs) are progenitors of the eiu cllipticals of today (INormendy Sanders. 1992).," 	It has been proposed that ultraluminous infrared galaxies (ULIRGs) are progenitors of the giant ellipticals of today (Kormendy Sanders, 1992)."1152 On the other liaud. Saucers et al. (1988))," On the other hand, Sanders et al. \cite{sand88}) )"1153 suggested that ULIRGs will evolve into quasars., suggested that ULIRGs will evolve into quasars.1154 While there is a continuius devate about what powers these galaxies (starbur‘st or active galactic nuclei (ACUON)) it is clear that 1) ULIRGS show clear delence for interactious/merects Borne e al. 1999..," While there is a continuing debate about what powers these galaxies (starburst or active galactic nuclei (AGN)) it is clear that 1) ULIRGS show clear dence for interactions/mergers Borne et al. \cite{borne99},"1155 Sauders et al. 1988)), Sanders et al. \cite{sand88}) )1156. 2) some. at least. contain powerful ACN (e.g. Sanders et al. 1988)).," 2) some, at least, contain powerful AGN (e.g. Sanders et al. \cite{sand88}) )."1157 A population of luminous galaxies iu the subinillimetre wavelengths has been discovered iji recent vears., A population of luminous galaxies in the submillimetre wavelengths has been discovered in recent years.1158 ddies of the spectral cuerey distrinitions suggest that these are the analogues at high recIshift of ULIRG at low redshift Saers Mirabe 1996))., dies of the spectral energy distributions suggest that these are the analogues at high redshift of ULIRGs at low redshift Sanders Mirabel \cite{sand96}) ).1159 As for may local ULIRGS. it is not clear whether hese galaxies are powerred by sarburst or active iealactic nuclei.," As for many local ULIRGS, it is not clear whether these galaxies are red by starburst or active galactic nuclei."1160 The study of distaut ultraluniuous subi sources cau provide nuportaut information about the nature of ULIRCs aux how a iassive black hole forms at high redshift iix coexists with a powerful starburst., The study of distant ultraluminous submm sources can provide important information about the nature of ULIRGs and how a massive black hole forms at high redshift and coexists with a powerful starburst.1161 The interpretation of the ISAL properties of these objects ju the context of distant radio galaxies then becomes crucial., The interpretation of the ISM properties of these objects in the context of distant radio galaxies then becomes crucial.1162 We study here the UV. (rest frame) spectra of three καν aud the hyperluuinous subimilliiietre source SALA J02399-0136 (2 —2.8)., 	We study here the UV (rest frame) spectra of three HzRG and the hyperluminous submillimetre source SMM J02399-0136 $z=$ 2.8).1163 Our smuple consists ofthe following objects:, 	Our sample consists ofthe following objects:1164explain the observed offset.,explain the observed offset.1165 They fouud that inodels which include a iinorityv population of metal poor stars of the total) aud that are based on empirical. rather than theoretical. libraries better represcut the observe colors at all redshifts considered.," They found that models which include a minority population of metal poor stars $\%$ of the total) and that are based on empirical, rather than theoretical, libraries better represent the observed colors at all redshifts considered."1166 However. ? Were unable to explain the LRC colors using these differeu stellar libraries.," However, \citet{Conroy10} were unable to explain the LRG colors using these different stellar libraries."1167 LIustead. it has also heen sugecstec hat a small population of vounecr or relatively meta oor stars in these ealaxies could explain the SDSS observations.," Instead, it has also been suggested that a small population of younger or relatively metal poor stars in these galaxies could explain the SDSS observations."1168 The addition of other stellar populatious such as bluc/extreme horizontal brauch (IIB) stars aux due strageler stars (typically excluded iun such models) could also help to explain the offsets., The addition of other stellar populations such as blue/extreme horizontal branch (HB) stars and blue straggler stars (typically excluded in such models) could also help to explain the offsets.1169 Some discrepancies wave also beeu suggested between SSP models auc the (resolved) SDSS color magnitude diagrams of 17 of he Alilky Wavs globular clusters (?).., Some discrepancies have also been suggested between SSP models and the (resolved) SDSS color magnitude diagrams of 17 of the Milky Way's globular clusters \citep{An08}.1170 By comparing heir fiducial sequences with the theoretical isochrones of 7.. ?.— demonstrated that the models could not be simultaneously fit to both the main sequence and red giant branch accurately.," By comparing their fiducial sequences with the theoretical isochrones of \citet{Girardi04}, \citet{An08} demonstrated that the models could not be simultaneously fit to both the main sequence and red giant branch accurately."1171 Iu this paper. we compare SSP iuodels with the inteerated colors of globulay clusters in M31.," In this paper, we compare SSP models with the integrated colors of globular clusters in M31."1172 This cluster system is the only large sample of clusters for which accurate inteerated photometiy is available from the SDSS survey., This cluster system is the only large sample of clusters for which accurate integrated photometry is available from the SDSS survey.1173 It therefore provides a unique location iu which to compare SDSS photometry of star clusters witli SSP models., It therefore provides a unique location in which to compare SDSS photometry of star clusters with SSP models.1174 In section 2 we describe the (previously published) data available for λος clusters., In section 2 we describe the (previously published) data available for M31's clusters.1175 The SPS models considered im this paper are introduced iu section 3., The SPS models considered in this paper are introduced in section 3.1176 Section 1 compares these models with the cluster data and section 5 cousiders possible explanations for observed offsets., Section 4 compares these models with the cluster data and section 5 considers possible explanations for observed offsets.1177 ADAE hosts the largest cluster population iu our local eroup., M31 hosts the largest cluster population in our local group.1178 The galaxy proximity fo us (780 kpc) means its clusters’ stellar populationss are unresolved in typical eround based. observations., The galaxy's proximity to us $\sim$ 780 kpc) means its clusters' stellar populations are unresolved in typical ground based observations.1179 UWowever. the integrated cluission. of these clusters cau be accurately measured with relatively short exposures (compared with cluster systems around more distant galaxics).," However, the integrated emission of these clusters can be accurately measured with relatively short exposures (compared with cluster systems around more distant galaxies)."1180 In this paper. we consider all globular clusters listed as class 1° in the catalog of ?archive...," In this paper, we consider all globular clusters listed as `class 1' in the catalog of \citet{Peacock10}."1181 These clusters have non-stellar surface brightuess profiles aud are confirmed to be at the distance of AIL based ou spectroscopy or high resolution UST observatious., These clusters have non-stellar surface brightness profiles and are confirmed to be at the distance of M31 based on spectroscopy or high resolution HST observations.1182 This dataset is therefore unlikely to contain siguificaut contanunation from non-cluster sources., This dataset is therefore unlikely to contain significant contamination from non-cluster sources.1183 This sample also excludes the recently identified population of vouug clusters iu MOI., This sample also excludes the recently identified population of young clusters in M31.1184 Young clusters are classified as those clusters having colors bluer than the old Milkv Wax elobular clusters (2%) or based on spectroscopy of tlie clusters €2)..," Young clusters are classified as those clusters having colors bluer than the old Milky Way globular clusters \citep{Fusi_Pecci05,Peacock10} or based on spectroscopy of the clusters \citep{Caldwell09}."1185 The catalog includes. 116 class 1 clusters. located across the galaxy.," The catalog includes 416 class 1 clusters, located across the galaxy."1186 The SDSS covered a large region of sky in the direction of M31 as part of a supplemental run to the main survey., The SDSS covered a large region of sky in the direction of M31 as part of a supplemental run to the main survey.1187 These observations have previously been used to provide mand griz-baud photometry for 73% and 924 of NDS known globular clusters. respectively (fewer clusters are detected iu the 4—band due to the increased extinction and lower seusitivitv of the survey at these waveleugtlis).," These observations have previously been used to provide and -band photometry for $\%$ and $\%$ of M31's known globular clusters, respectively (fewer clusters are detected in the -band due to the increased extinction and lower sensitivity of the survey at these wavelengths)."1188 The photometry was performed usingSENTRACTOR., The photometry was performed using.1189". The colors were obtained through I apertures aud total magnitudes measured using au aperture in the rause 2.8-10.6"". depeudiug on the size of the cluster."," The colors were obtained through $\arcsec$ apertures and total magnitudes measured using an aperture in the range $\arcsec$, depending on the size of the cluster."1190 The photometry was calibrated using the standard SDSS pipeline calibration., The photometry was calibrated using the standard SDSS pipeline calibration.1191 As such. the colors are on the AB 1uagnitude system.," As such, the colors are on the AB magnitude system."1192 The SDSS calibration is thought to be in good agrecinent with this svstem for the griz--bauds. although a sheht correction has been proposed for the u--boaud (suchthatνο Uspss-0.02).," The SDSS calibration is thought to be in good agreement with this system for the -bands, although a slight correction has been proposed for the -band \citep[such that $u_{\rm{AB}}$ $u_{\rm{SDSS}}$."1193 For full details of these data we refer the reader to ?.., For full details of these data we refer the reader to \citet{Peacock10}.1194 Spectroscopic inetallicities for 200 of these clusters are available from the collated catalog of 7.., Spectroscopic metallicities for 200 of these clusters are available from the collated catalog of \citet{Fan08}.1195 These lucasurements are taken from the studies of ?7.. 7. aud νι," These measurements are taken from the studies of \citet{Huchra91}, \citet{Barmby00} and \citet{Perrett02}."1196 Where clusters are present iu more than one of these studies; preference was given to the larecr aud more recent catalog of ?7..," Where clusters are present in more than one of these studies, preference was given to the larger and more recent catalog of \citet{Perrett02}."1197 ? also provide estimates for the reddening of these clusters., \citet{Fan08} also provide estimates for the reddening of these clusters.1198 The extinction in cach baud was calculated for cach cluster using the ? reddening values aud the extinction curves of 7.. as preseuted im table 6 of ?..," The extinction in each band was calculated for each cluster using the \citet{Fan08} reddening values and the extinction curves of \citet{Cardelli89}, as presented in table 6 of \citet{Schlegel98}."1199 The reddening of individual clusters in AD) ds variable across the salaxw., The reddening of individual clusters in M31 is variable across the galaxy.1200 Therefore. we oulv consider those clusters for which reddening estinates are available.," Therefore, we only consider those clusters for which reddening estimates are available."

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