ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2predominantly southwest of the core. similar to previous results.,"predominantly southwest of the core, similar to previous results."3 The cloud rotation shown in Figure 17. is perpendicular to the outflow. which is a prediction of current theories of star formation.," The cloud rotation shown in Figure \ref{sfoxx_centroid} is perpendicular to the outflow, which is a prediction of current theories of star formation."4 SFO 25. also known as ΕΠΗΠΕ (Walshetal.1992)... is a tvpe D cloud located at a distance of 780 pe.," SFO 25, also known as HH124 \citep{wor}, is a type B cloud located at a distance of 780 pc."5 Embedded within is the IRAS source 06382+1017. whose spectral energy distribution is consistent with an embedded YSO (SEO).," Embedded within is the IRAS source 06382+1017, whose spectral energy distribution is consistent with an embedded YSO \citetext{SFO}."6. SFO 25 is associated with the ILLI region NGC 2264., SFO 25 is associated with the HII region NGC 2264.7 Figure 2. shows a moderately curved rim arching northwards away [rom (he ionizing source., Figure \ref{sfo_dss} shows a moderately curved rim arching northwards away from the ionizing source.8 Llowever. although there is no tail visible in the optical. there is also no point al which the structure beeins to broaden away [rom (he ionizing source as there is for the other tvpe D sources. SFO 4 and SFO 13.," However, although there is no tail visible in the optical, there is also no point at which the structure begins to broaden away from the ionizing source as there is for the other type B sources, SFO 4 and SFO 13."9 Às we shall see below. this source bears many similarities with twpe C cometary clouds as well.," As we shall see below, this source bears many similarities with type C cometary clouds as well."10 The CO 130) and CO 1) profiles show a wide (5 !)). slightly asvmmnietric line (Figure I5bb).," The CO ) and CO ) profiles show a wide (5 ), slightly asymmetric line (Figure \ref{sfoxx_xxv_profiles}b b)."11 The CO 1) profile shows line wings. which may indicate the presence of an outflow.," The CO ) profile shows line wings, which may indicate the presence of an outflow."12 The line profiles show some asymmetry as well. mostly in (he form of a redshifted (ail.," The line profiles show some asymmetry as well, mostly in the form of a redshifted tail."13 The CO 0) inteerated intensitv map (Figure 19)) shows a low level of CO emission throughout the region. which may be the result of contamination [rom nearby diffuse clouds which lie αἱ nearly the same velocity of SFO 25.," The CO ) integrated intensity map (Figure \ref{sfoxxv_itty}) ) shows a low level of CO emission throughout the region, which may be the result of contamination from nearby diffuse clouds which lie at nearly the same velocity of SFO 25."14 The 0) traces the higher column density gas which forms the SFO 25 core and shows a molecular gas tail to the north., The ) traces the higher column density gas which forms the SFO 25 core and shows a molecular gas tail to the north.15 This morphology is more typical of a type C cometary cloud than a tvpe D source., This morphology is more typical of a type C cometary cloud than a type B source.16 The CO 1) emission peaks stronglv at the IRAS source position., The CO ) emission peaks strongly at the IRAS source position.17 The ο) integrated intensity map (Figure 19)) shows strongly peaked emission around the core with much weaker emission along the tail to the north of the source. which looks more (wpical of a type C source than a tvpe D source.," The ) integrated intensity map (Figure \ref{sfoxxv_itty}) ) shows strongly peaked emission around the core with much weaker emission along the tail to the north of the source, which looks more typical of a type C source than a type B source."18 The 2) emission is also strongly peaked at the IRAS source position., The ) emission is also strongly peaked at the IRAS source position.19 We detected no coherent velocity structure in the 0) velocity centroid observations., We detected no coherent velocity structure in the ) velocity centroid observations.20 However the CO I) line wings show evidence of an outflow oriented in the north-south direction (Figure 20)). though unrelated emission near the same velocity makes it hard to separate (he outflow emission from emission due to other sources.," However the CO ) line wings show evidence of an outflow oriented in the north-south direction (Figure \ref{sfoxxv_outflow}) ), though unrelated emission near the same velocity makes it hard to separate the outflow emission from emission due to other sources."21confirmed these results based on direct N-body simulations.,confirmed these results based on direct N-body simulations.22 They also showed that multi-mass clusters with IMBHs are mass-segregated in their centers and that main-sequence stars have cusps that are significantly flatter than —1.75., They also showed that multi-mass clusters with IMBHs are mass-segregated in their centers and that main-sequence stars have cusps that are significantly flatter than $-$ 1.75.23 They found that the IMBHs appear to produce shallow central cusps on the projected density profiles of bright main sequence stars for these clusters. with slopes of ~—0.2. as opposed to steep power-laws (Baumgardtetal.2005).," They found that the IMBHs appear to produce shallow central cusps on the projected density profiles of bright main sequence stars for these clusters, with slopes of $\sim24-$ 0.2, as opposed to steep power-laws \citep{bau05}."25. Noyola&Gebhardt(2006) and Noyola&Gebhardt(2007) (hereafter called respectively NGO6 and NGO7) obtained surface brightness profiles from images for Galactic. LMC. SMC and Fornax dwarf galaxy globular clusters.," \citet{noy06} and \citet{noy07} (hereafter called respectively NG06 and NG07) obtained surface brightness profiles from images for Galactic, LMC, SMC and Fornax dwarf galaxy globular clusters."26 They found that about of the globular clusters in their sample show central slopes in this intermediate range., They found that about of the globular clusters in their sample show central slopes in this intermediate range.27 The surface density profile shape can also be affected in the size of its core when a central black hole is present.Trential.(2007) estimated the value of ἐπ (ratio of core radius to half-light radius) for N-body simulated star clusters containing central black holes., The surface density profile shape can also be affected in the size of its core when a central black hole is \citet{tre07} estimated the value of $r_c/r_h$ (ratio of core radius to half-light radius) for N-body simulated star clusters containing central black holes.28 They used a density averaged radius as a measure for the core radius., They used a density averaged radius as a measure for the core radius.29 They found that the ratio tends to reach values around 0.3 for these cases. while the value is considerably smaller (< 0.1) for clusters without black holes.," They found that the ratio tends to reach values around 0.3 for these cases, while the value is considerably smaller $< 0.1$ ) for clusters without black holes."30" On the other hand. Hurley(2007) finds similarly large r./7, values for N-body simulations evolved including 0.5—10% primordial binaries. but without a central black hole."," On the other hand, \citet{hur07} finds similarly large $r_c/r_h$ values for N-body simulations evolved including $0.5-10$ primordial binaries, but without a central black hole."31 In this case. the Casertano Hut method (Casertano&Hut1985) was used to obtain the three-dimensional core radius.," In this case, the Casertano Hut method \citep{cas85} was used to obtain the three-dimensional core radius."32 The different way in which the core radius was measured from the N-body simulations differs between the two results., The different way in which the core radius was measured from the N-body simulations differs between the two results.33 Recently. Vesperint&Trenti(2010) analyzed. direct body models with and without IMBHs.," Recently, \citet{ves10} analyzed direct N-body models with and without IMBHs."34 They found that shallow cusps with logarithmic slopes as steep as —0.3 are present in various models. not only the ones containing black holes.," They found that shallow cusps with logarithmic slopes as steep as $-$ 0.3 are present in various models, not only the ones containing black holes."35 The apparent discrepancies between different models using different analysis techniques stresses the importance of performing meaningful measurements on models so they can be properly compared with observational data., The apparent discrepancies between different models using different analysis techniques stresses the importance of performing meaningful measurements on models so they can be properly compared with observational data.36 Direct. dynamical evidence for the existence of central black holes using velocity dispersion measurements has been put forward for three nearby globular clusters., Direct dynamical evidence for the existence of central black holes using velocity dispersion measurements has been put forward for three nearby globular clusters.37 MIS was the first case (Gerssenetal.2002.2003).. but alternative models without black holes were also shown to be good fits to the data (Baumgardtetal.2003b)..," M15 was the first case \citep{ger02,ger03}, but alternative models without black holes were also shown to be good fits to the data \citep{bau03a}."38 The latest detailed dynamical measurement and model finds non-conclusive evidence for the presence of a central black hole in this cluster (vandenBoschetal.2006)., The latest detailed dynamical measurement and model finds non-conclusive evidence for the presence of a central black hole in this cluster \citep{bos06}.39. Gl. a large globular cluster in Andromeda. has stronger observational evidence to support the presence of a central black hole. from integrated kinematical measurements (Gebhardtetal.2003.2005).. as well as from X-ray (Pooley&Rappaport2006) and radio (Ulvestadetal.2007) observations. but alternative scenarios have also been presented for this case (Baumgardt 2003c).," G1, a large globular cluster in Andromeda, has stronger observational evidence to support the presence of a central black hole, from integrated kinematical measurements \citep{geb03,geb05}, as well as from X-ray \citep{poo06} and radio \citep{ulv07} observations, but alternative scenarios have also been presented for this case \citep{bau03c}."40. Omega Centauri is the most recent case for which line-of-sight velocity dispersion measurements appear to support the existence of a central black hole of 40.000 M. (Noyolaetal.2008).. but proper motion measurements from HST images find different results (Anderson&vanderMarel2010;vanderMarel&Anderson 2010).," Omega Centauri is the most recent case for which line-of-sight velocity dispersion measurements appear to support the existence of a central black hole of 40,000 $M_\odot$ \citep{noy08}, but proper motion measurements from HST images find different results \citep{and10,vdm10}."41. Evidence has also surfaced for intermediate mass black holes in extra-galactie disk galaxies based on X-ray observations., Evidence has also surfaced for intermediate mass black holes in extra-galactic disk galaxies based on X-ray observations.42 Ultra luminous X-ray. sources. (ULXs) have X-ray luminosities higher than the Eddington limit for a stellar mass black hole., Ultra luminous X-ray sources (ULXs) have X-ray luminosities higher than the Eddington limit for a stellar mass black hole.43 One of the possible explanations for this emission is that it comes from accretion onto an IMBH., One of the possible explanations for this emission is that it comes from accretion onto an IMBH.44 For example. the galaxy M82 contains an ULX source which ts believed to host an IMBH based on the absolute brightness of the source (Matsumoto&Tsuru1999;Matsumotoetal.2001).. and its radio variability (Strohmayer&Mushotzky2003).," For example, the galaxy M82 contains an ULX source which is believed to host an IMBH based on the absolute brightness of the source \citep{mat99,mat01}, and its radio variability \citep{str03}."45. The position of the X-ray source appears to coincide with the young dense star cluster MGG-11 (MeCradyetal.2003)., The position of the X-ray source appears to coincide with the young dense star cluster MGG-11 \citep{mcc03}.46. There ts also the controversial case of the globular cluster RZ 2109 in NGC 4472. which shows the first clear evidence for a star cluster hosting a black hole (Maccaroneetal.2007).. but the size of the black hole is still under debate (Zepfal. 2008).," There is also the controversial case of the globular cluster RZ 2109 in NGC 4472, which shows the first clear evidence for a star cluster hosting a black hole \citep{mac07}, but the size of the black hole is still under debate \citep{zep08}."47. One more interesting object is the X-ray source CXOJ033831.8-352604. associated with a globular cluster in the Fornax elliptical galaxy NGC 1399.," One more interesting object is the X-ray source CXOJ033831.8-352604, associated with a globular cluster in the Fornax elliptical galaxy NGC 1399."48 Irwinetal.(2010) suggest the emission might come from a tidally disrupted white dwarf around an IMBH., \citet{irw10} suggest the emission might come from a tidally disrupted white dwarf around an IMBH.49 In this paper. we create synthetic HST-like images from N-body simulations with and without IMBHs.," In this paper, we create synthetic HST-like images from N-body simulations with and without IMBHs."50 We measure their surface brightness profiles as we would with observations., We measure their surface brightness profiles as we would with observations.51 We provide an analysis of the detailed shape of central density profiles for these models that helps understand the central state of Galactic globular clusters., We provide an analysis of the detailed shape of central density profiles for these models that helps understand the central state of Galactic globular clusters.52 We describe the N-body models in section 2. the synthetic images in section 3. data analysis in section 4. and discussion in section 5.," We describe the N-body models in section 2, the synthetic images in section 3, data analysis in section 4, and discussion in section 5."53 Itis often challenging to make a direct comparison between the results of N-body models and observations because it is hard to take into account the sources of uncertainty of, It is often challenging to make a direct comparison between the results of N-body models and observations because it is hard to take into account the sources of uncertainty of54Further. in the 1.1 jam bauds of water. the absorption is clearly weaker in the Ls dwarf 2411523 than in L6.5 dwarf 2MI1T11 as cau been seen in Fie.Ll.,"Further, in the 1.1 $\mu$ m bands of water, the absorption is clearly weaker in the L8 dwarf 2M1523 than in L6.5 dwarf 2M1711 as can been seen in Fig.4."55 The spectral types are based ou Ixirkpatrick aud those by Geballeetal.(2002) also show the similar types., The spectral types are based on \citet{kir00} and those by \citet{geb02} also show the similar types.56 This result. is somewhat surprising. since the 1.1 aud 1.4 gan water bands were used as Classification criteria by Geballeetal.(2002).," This result is somewhat surprising, since the 1.1 and 1.4 $\mu$ m water bands were used as classification criteria by \citet{geb02}."57. The water baud strengths. however. do uot necessarily increase linearly with the L types. aud some water bands show little chauge iu their streugtlis between L3 aud Ls (Reidetal.20010:Testial...2001).," The water band strengths, however, do not necessarily increase linearly with the L types, and some water bands show little change in their strengths between L3 and L8 \citep{rei01b,tes01}."58. Although there appears a plateau iu the Zar scale between L6 and LS (81L2). we think that the spectral types basically correspond to a temperature sequence. since tlie present spectral types are determiued by the several temperature seusitive spectral features in the optical region on purely enipirical basis (κρανίοetal.1999). and are consistent mostly with the infrared. spectral indices iucludiug methane bands (Burgasseretal.—2002a:Geballe2002).," Although there appears a plateau in the $T_{\rm eff}$ scale between L6 and L8 4.2), we think that the spectral types basically correspond to a temperature sequence, since the present spectral types are determined by the several temperature sensitive spectral features in the optical region on purely empirical basis \citep{kir99} and are consistent mostly with the infrared spectral indices including methane bands \citep{bur02a,geb02}."59.. Thus. it is »ossible that water bands are weaker in the cooler L dwarls than in the warmer objects. while he methane feature at 2.2 jan is stronger in the L& dwarf than in the L6.5 dwarfs as shown in Fie.," Thus, it is possible that water bands are weaker in the cooler L dwarfs than in the warmer objects, while the methane feature at 2.2 $\mu$ m is stronger in the L8 dwarf than in the L6.5 dwarfs as shown in Fig."60 Il., 1.61 We interpret this result as follows., We interpret this result as follows.62 The water abuudance should be larger iu cooler objects. jut extinction by dust clouds. still located iu the optically thin region in L cdwarls. should be arger iui cooler objects as well.," The water abundance should be larger in cooler objects, but extinction by dust clouds, still located in the optically thin region in L dwarfs, should be larger in cooler objects as well."63 Του the iucreasiug water abundauce aud the iucreasiug dust extinction nay cancel out in L cdwarls., Then the increasing water abundance and the increasing dust extinction may cancel out in L dwarfs.64 For this reasou. water band strengths do not show a large increase in L cdwarls (Tsuji2002).. aud it is no wouder that they show an inversiou somewhere in L types.," For this reason, water band strengths do not show a large increase in L dwarfs \citep{tsu02}, and it is no wonder that they show an inversion somewhere in L types."65 The methane feature at 2.2 san does not necessarily suffer the same effect. either jecause the dust extinction is already not so important in the 2.2 pan reeion than the 1.1 juu 'eeion. as also can be seen in the cloud models by Ackerman&Marley(2001).. or else because ol the more rapid increase of the methane baud strength with the decreasing ellective temperature han in the case of the water bands.," The methane feature at 2.2 $\mu$ m does not necessarily suffer the same effect, either because the dust extinction is already not so important in the 2.2 $\mu$ m region than the 1.4 $\mu$ m region, as also can be seen in the cloud models by \citet{ack01}, or else because of the more rapid increase of the methane band strength with the decreasing effective temperature than in the case of the water bands."66 We have coufirmecd these results by a detailed computation ol the spectral line intensities based on the UCMs by conusicering the effect not only of Tey aud ebulalsoofthechangeofthebauckgroundopacitiesduetoatoms. tons. dustgrains. andguasi ΟΠ μπα CLA (Tsuji.Nakajima&Yanagisawa2001).," We have confirmed these results by a detailed computation of the spectral line intensities based on the UCMs by considering the effect not only of $T_{\rm eff}$ and $g$ but also of the change of the background opacities due to atoms, ions, dust grains, and quasi-continuous sources such as $_2$ CIA \citep{tsu03}."67 The spectra between 0.95 and 1.3 4m are shown in Fig.L., The spectra between 0.95 and 1.3 $\mu$ m are shown in Fig.4.68 First. we confirm that FeH #LA—NLN 0.99 pam bands can be clearly seen not only in L dwarfs but also in T cdwarfs as noted by (2002b).," First, we confirm that FeH $F^4\Delta - X^4\Delta$ 0.99 $\mu$ m bands can be clearly seen not only in L dwarfs but also in T dwarfs as noted by \citet{bur02b}."69. The 0.99 jan FeH bands gradually weaken from L6.5 to L8 aud further to T2. but strengthen again at T3.5.," The 0.99 $\mu$ m FeH bands gradually weaken from L6.5 to L8 and further to T2, but strengthen again at T3.5."70 The weakening of the 0.99 jn. ΕΘΗ bands from L6.5 to LS is consistent with the weakening of the 1.61 aud 1.63 fu features in the same objects. aud this lends a support for the contribution of FeH to these features.," The weakening of the 0.99 $\mu$ m FeH bands from L6.5 to L8 is consistent with the weakening of the 1.61 and 1.63 $\mu$ m features in the same objects, and this lends a support for the contribution of FeH to these features."71 However. the observed strengths of FeH 0.99 jan bands. not ouly of T cdwarls but also of L dwarls. are difficult to interpret directly by our GCMs as well as by other models now available.," However, the observed strengths of FeH 0.99 $\mu$ m bands, not only of T dwarfs but also of L dwarfs, are difficult to interpret directly by our UCMs as well as by other models now available."72 Lu this connection. Burgasseretal.(2002b) suggested tliat the clust cloud breaks in the early T dwarls aud FeH formed deep in the photosphere cau be peered," In this connection, \citet{bur02b} suggested that the dust cloud breaks in the early T dwarfs and FeH formed deep in the photosphere can be peered"73"The starting point for the transit detection is the observe photometric elt curves hat have been detrended iux filtered by one of the two pro)osed iuethods. (see section 2.1). either Fourier or Way""elet transformations.","The starting point for the transit detection is the observed photometric light curves that have been detrended and filtered by one of the two proposed methods (see section 2.1), either Fourier or Wavelet transformations."74 The idea of the method is the searc1 for periodicities ou a selectec wavelet-trausforimation., The idea of the method is the search for periodicities on a selected wavelet-transformation.75 Tus ransformation results im a time-series where trausit-lise Signatures are amplified. aux the first transit-like sigua shiows up in the first bin of the FFT of this time-series. 1idepeudeut of the transits” epochs.," This transformation results in a time-series where transit-like signatures are amplified, and the first transit-like signal shows up in the first bin of the FFT of this time-series, independent of the transits' epochs."76 The xoposed. algoritun. TRUFAS. is therefore conrposed of two principal steDRM An automatic way to reject the alse detections from arbitrary noise has also been developed.," The proposed algorithm, TRUFAS, is therefore composed of two principal steps: An automatic way to reject the false detections from arbitrary noise has also been developed."77 However. as it will be shown. the algorithm's rate of false detectious is so low. at least when the algoritlun is applied to the svuthetic helt curves for the CoRoT mission. that their individual study do not represent any real problem.," However, as it will be shown, the algorithm's rate of false detections is so low, at least when the algorithm is applied to the synthetic light curves for the CoRoT mission, that their individual study do not represent any real problem."78 The svuthetic Πο curves are fully described in DTI in the coutext of a “blind test. giving au unbiased conrparison of transi detection algorithius frou several yarticipating teams.," The synthetic light curves are fully described in BT1 in the context of a “blind” test, giving an unbiased comparison of transit detection algorithms from several participating teams."79 These light[m] curves were [m]seenerated with ano end-to-endc» instrument simulator of CoRoT (Auverene et al. 2003)).," These light curves were generated with an end-to-end instrument simulator of CoRoT (Auvergne et al. \cite{auvergne03}) ),"80 nd melde stellar nuülcro-variability (CÁAierain et al. 2001)), and include stellar micro-variability (Aigrain et al. \cite{aigrain04}) )81 and an assorted set of planetary transits aud transit-like events that were inserted ii some of he 999 curves: twenty one contain planetary transits. eleven have low-depth stellar eclipse signals aud one results from an eclipsing triple stellar svsteni.," and an assorted set of planetary transits and transit-like events that were inserted in some of the 999 curves: twenty one contain planetary transits, eleven have low-depth stellar eclipse signals and one results from an eclipsing triple stellar system."82 The CoRoT svuthetic light curves are saupled every 512 s daring 150 davs with quasi-periodic gaps of 30 mun each L7 h. These eaps simulate the fact that the iustrunenut will cross the South-Atlantic Anomaly (SAA) aud these exposures will not be usable., The CoRoT synthetic light curves are sampled every 512 s during 150 days with quasi-periodic gaps of 30 min each 1.7 h. These gaps simulate the fact that the instrument will cross the South-Atlantic Anomaly (SAA) and these exposures will not be usable.83 The first step in the analysis of the light curves is to remove the earth-scattered light introduced in the siuulated data that varies with the orbital period of he satellitc (1.7 hours) aud which is not uniform over he CCD., The first step in the analysis of the light curves is to remove the earth-scattered light introduced in the simulated data that varies with the orbital period of the satellite (1.7 hours) and which is not uniform over the CCD.84 In fact. the contribution that appears in the svuthetic light curves is a residual of this scattered light. hat may cad to a positive or weeative signal.," In fact, the contribution that appears in the synthetic light curves is a residual of this scattered light, that may lead to a positive or negative signal."85 The effect hat introduces this residual earthi-scattered lelt in cach svuthetie curve. with a period alios following the orbital )oriod. is uot exactly the same for all the stars. tit can )o Interred directly from cach light curve.," The effect that introduces this residual earth-scattered light in each synthetic curve, with a period almost following the orbital period, is not exactly the same for all the stars, but it can be inferred directly from each light curve."86 What we lave done in order to remove this effect is o use each light curve ο obtain the shape of this spurious signal. identifying a uaxinumn around the orbital period aud folding the lig curve with the obtained period.," What we have done in order to remove this effect is to use each light curve to obtain the shape of this spurious signal, identifying a maximum around the orbital period and folding the light curve with the obtained period."87 Before the folding. eac orbit is appropriately set oa common level to avo effect of the low frequency noise present in the data tm has not vet οσα removed.," Before the folding, each orbit is appropriately set to a common level to avoid the effect of the low frequency noise present in the data that has not yet been removed."88 The signal obtained after foldiug is smoothed aud μιbtracted from the original ou, The signal obtained after the folding is smoothed and subtracted from the original one.89 This is done individ or each leh curve., This is done individually for each light curve.90 After this detreudine. he μια]. gaps caused by the crossing of SAA are lnearly 1ed.," After this detrending, the small gaps caused by the crossing of the SAA are linearly interpolated."91 The result of this first s is plotted iu Fie. |.," The result of this first step is plotted in Fig. \ref{figorig168},"92 for the case of the original svuthet helt curve 168 ecnerated by DTI. which is a ligo curve contaiuluga transit.," for the case of the original synthetic light curve 168 generated by BT1, which is a light curve containing a transit."93 The spiky aspect of the original lieht curve is due he orbital 1.7 hotrs period which is corrected after deiding., The spiky aspect of the original light curve is due to the orbital 1.7 hours period which is corrected after detrending.94 However. some low frequency nolse ολαΊαν tha needs to be filtered too. as it is clear from Fie. 1..," However, some low frequency noise remains that needs to be filtered too, as it is clear from Fig. \ref{figorig168}."95" For the filtering ο this ""red noise”. two different types of filters have heei tested: a filter in the Fourier domain and a Wavelet filter."," For the filtering of this “red noise”, two different types of filters have been tested: a filter in the Fourier domain and a Wavelet filter."96" This filtering is uot an iuteeral part of TRUFAS: but for any light curve based on ""real data. it is a required previous step."," This filtering is not an integral part of TRUFAS; but for any light curve based on “real data”, it is a required previous step."97 We cuiploy two cliffercut filter methods in order to test aud verify the reliability of TRUFAS., We employ two different filter methods in order to test and verify the reliability of TRUFAS.98into account the conclusion of Daszvisska-Daszkiewicz οἱ al. (,into account the conclusion of Daszyńsska-Daszkiewicz et al. (992005) that convection is rather ineLlicient in FC Vir.,2005) that convection is rather inefficient in FG Vir.100 We assumed. uniform (solid-bodyv). stellar rotation and conservation of elobal angular momentum. during evolution from the ZAAIS., We assumed uniform (solid-body) stellar rotation and conservation of global angular momentum during evolution from the ZAMS.101 These assumptions were chosen or simplicity reasons., These assumptions were chosen for simplicity reasons.102 The influence of rotation on the evolutionary tracks of 8 Scuti models was demonstrated by Breeer Pamvatuvkh (1998)., The influence of rotation on the evolutionary tracks of $\delta$ Scuti models was demonstrated by Breger Pamyatnykh (1998).103 At relatively low rotational velocities of about 60 to 70 kms. the evolutionary tracks are ocated close to those for non-rotating stellar mocels.," At relatively low rotational velocities of about 60 to 70 km/s, the evolutionary tracks are located close to those for non-rotating stellar models."104 The main elfect of rotation is the splitting of the nonradial mode requencies., The main effect of rotation is the splitting of the nonradial mode frequencies.105 Even for slowly rotating stars. this splitting is asvmmetric due to the second-order. ellect. of rotation.," Even for slowly rotating stars, this splitting is asymmetric due to the second-order effect of rotation."106 The whole frequeney spectrum of unstable nonracdial modes will be considered in the next paper where possibilities to explain the appearance of very close frequencies. will. be cliscussecl., The whole frequency spectrum of unstable nonradial modes will be considered in the next paper where possibilities to explain the appearance of very close frequencies will be discussed.107 Here we note the small. but not negligible effects of rotation on the frequeney spectrum of. racial modes.," Here we note the small, but not negligible effects of rotation on the frequency spectrum of radial modes."108 ‘This is caused. mainly by the small change in the moclel structure due to the variations of the ellective gravity inside the star., This is caused mainly by the small change in the model structure due to the variations of the effective gravity inside the star.109 For example. owe fit the frequeney of the radial funclamental mode to the observed value of 12.154 evele +. we obtain t1ο [frequency of the 6th. radial. overtone equal to 34.536 cvele and 34.345 evele [or the equatorial rotational velocities of 49 Km Land SI kins respectively.," For example, if we fit the frequency of the radial fundamental mode to the observed value of 12.154 cycle $^{-1}$, we obtain the frequency of the 6th radial overtone equal to 34.536 cycle $^{-1}$ and 34.345 cycle $^{-1}$ for the equatorial rotational velocities of 49 km $^{-1}$ and 81 km $^{-1}$, respectively."110 Aclelitional tests show that a small decrease of the heavy element abundance (sav. to Z= 0.015) and a small change of opacities (by using new OP data according to Badnell ct 22005. and. Seaton 2005) may result in a very good fit of the frequency. of the 6th overtone to the observed. frequeney of the close pair at. 34.12 evele 1," Additional tests show that a small decrease of the heavy element abundance (say, to $Z=0.015$ ) and a small change of opacities (by using new OP data according to Badnell et 2005 and Seaton 2005) may result in a very good fit of the frequency of the 6th overtone to the observed frequency of the close pair at 34.12 cycle $^{-1}$."111 Another possibility to achieve such a fit is to suggest still faster rotation (Vj of about 90-100 kms. 1) which is outside the allowed limits., Another possibility to achieve such a fit is to suggest still faster rotation $V_{\rm rot}$ of about 90-100 km $^{-1}$ ) which is outside the allowed limits.112 We also tested many other models with the global parameters inside the error box of the ellective temperature and Luminosity., We also tested many other models with the global parameters inside the error box of the effective temperature and luminosity.113" Again. we have considered only those models which [it the radial fundamental mode frequency. to the observed. frequency of 12.154. evcle i ""'"," Again, we have considered only those models which fit the radial fundamental mode frequency to the observed frequency of 12.154 cycle $^{-1}$."114rrwerelore. only moderate changes in the model parameters and in the input ohvsies were allowed.," Therefore, only moderate changes in the model parameters and in the input physics were allowed."115 We varied the following parameters: (i) mass. (ii) chemical composition (bx choosing Z=) 015. for exemple. ancl by changing the proportions in the abundancies of the elements heavier than helium as were determined very recently by Aspluncl et 22004. 2005 for he Sun). (ii) the opacity data (hy using new OL data according to Badnell et 22005 anc Seaton 2005). and (iv) the cllicieney of overshooting from the convective core.," We varied the following parameters: (i) mass, (ii) chemical composition (by choosing $Z=0.015$ , for example, and by changing the proportions in the abundancies of the elements heavier than helium as were determined very recently by Asplund et 2004, 2005 for the Sun), (iii) the opacity data (by using new OP data according to Badnell et 2005 and Seaton 2005), and (iv) the efficiency of overshooting from the convective core."116 All these models. when fitted to the observed frequency. of 12.154 evele n as the racial. fundamental⋅ mode frequency.⋅ eive very similar freeuencies of radial overtones. which do not differ significantly from the frequencies shown in 55.," All these models, when fitted to the observed frequency of 12.154 cycle $^{-1}$ as the radial fundamental mode frequency, give very similar frequencies of radial overtones, which do not differ significantly from the frequencies shown in 5."117 The figure shows that the majority of the close, The figure shows that the majority of the close118"observations of classical Cepheids iu the LMC. NGC 6822, SMC, and IC 1613 (Fig. 2)).","observations of classical Cepheids in the LMC, NGC 6822, SMC, and IC 1613 (Fig. \ref{fig2}) )."119 The Galactic Wesenheit fictions of (2007)..(2008).. ancl establish a distance scale which is = newer than at Px25d (Fig. 3)).," The Galactic Wesenheit functions of , and establish a distance scale which is $\ga10$ nearer than at $P\simeq25^d$ (Fig. \ref{fig3}) )."120 The Galactic Weseuheit calibration established by).. partly on the basis of infrared surface brightuess and interferometric Daade-Woesselink: parallaxes. matches hybrid UST cluster Cepheid based relation within X (Fig. 3)).," The Galactic Wesenheit calibration established by, partly on the basis of infrared surface brightness and interferometric Baade-Wesselink parallaxes, matches hybrid HST / cluster Cepheid based relation within $\la5$ (Fig. \ref{fig3}) )."121 Lastly. regarding the construction of Fig. 2..," Lastly, regarding the construction of Fig. \ref{fig2},"122 it is noted that the slope characterizing longer period Cepheids in IC 1613 is steeper than that describing the short period regime., it is noted that the slope characterizing longer period Cepheids in IC 1613 is steeper than that describing the short period regime.123 Moreover. the SAIC exhibits a significant break in the Wesenheit functiontherein).," Moreover, the SMC exhibits a significant break in the Wesenheit function."124 The LMC displavs a separate trend. and efforts contiuue to characterize the discrepancy and its source.," The LMC displays a separate trend, and efforts continue to characterize the discrepancy and its source."125 The reader is likewise referred to the research of(1998)..(2006). and sugeest that the classical Cepheid Weseuheit relation exhibits a zZoro-poiut dependence ou metallicity2008).. again introducing a potential source of uncertaiuty into the present analysis since the chemical composition of the Cepheids in NGC 5128 is unknown.," The reader is likewise referred to the research of, and suggest that the classical Cepheid Wesenheit relation exhibits a zero-point dependence on metallicity, again introducing a potential source of uncertainty into the present analysis since the chemical composition of the Cepheids in NGC 5128 is unknown."126 The aforementioned researchers eudeavoured to ascertain the influence of cliemical composition by cranuning the distance offset between classical Cepheids located in the central Guetalzich) aud outer (miretal-poor) regions of a particular galaxy (ΑΠΟ. ATLOG. MIO).," The aforementioned researchers endeavoured to ascertain the influence of chemical composition by examining the distance offset between classical Cepheids located in the central (metal-rich) and outer (metal-poor) regions of a particular galaxy (M101, M106, M33)."127 Tlowever. ai degeneracy emerges (photometric contamination) since the stellar density and surface brightuess often increase toward the ceutral region.," However, a degeneracy emerges (photometric contamination) since the stellar density and surface brightness often increase toward the central region."128 noted that a substantial fraction of the difference in distance moduli between classical Cepheids occupying the inner aud outer regions of MIUI could arie from blending., noted that a substantial fraction of the difference in distance moduli between classical Cepheids occupying the inner and outer regions of M101 could arise from blending.129" aud e1miploved. criteria to iutieate the impact of photometric contamination so to enable an unbiased determination of the classical Cepheid metallicity effect from observations of classical Cepheids in MIOG aud ΑΙ, and the reader is encouraged to consider thei evidence."," and employed criteria to mitigate the impact of photometric contamination so to enable an unbiased determination of the classical Cepheid metallicity effect from observations of classical Cepheids in M106 and M33, and the reader is encouraged to consider their evidence."130 Yot the result inferred from variables in MIOG was provided an alternative rationale bv and(20090)... who roted that the observed offset was too large to be attributed to variations iu chemical composition.," Yet the result inferred from variables in M106 was provided an alternative rationale by and, who noted that the observed offset was too large to be attributed to variations in chemical composition."131 Iudeed. abundiuce gracdieut or MÍIOG nauplies that initial estimates of the classical Cepleid inetallieitv effect (5;=0.3 nag dex1 ) nearly double{ιν or see Table 1220094).," Indeed, abundance gradient for M106 implies that initial estimates of the classical Cepheid metallicity effect $\gamma_i \simeq -0.3$ mag $^{-1}$ ) nearly double, or see Table 12."132 A comparably sizeable result is obtained when cxamining the offset otween distance to classical Cepheids occupying the inner region of MIOL and distauce to classical Cepheids in the outer region of that same galaxy. which sample metal-rich and metal-poor variables," A comparably sizeable result is obtained when examining the offset between distance to classical Cepheids occupying the inner region of M101 and distance to classical Cepheids in the outer region of that same galaxy, which sample metal-rich and metal-poor variables"133produce either too red. galaxies or abundance patterns in disagreement with observations.,produce either too red galaxies or abundance patterns in disagreement with observations.134 In Fig., In Fig.135 G6 we compare the star formation histories of some moclels. chosen as the most representative cases shown in the previous sections.," \ref{fig-sfr} we compare the star formation histories of some models, chosen as the most representative cases shown in the previous sections."136" In the upper panel we compare the ""aper E best model (£)) SEL with those from the casesj.", In the upper panel we compare the Paper I best model ) SFH with those from the cases.137. In the lower panel. the solid line is taken from mocelg:: he main burst clearly dominates over the moerger-induced SE episode.," In the lower panel, the solid line is taken from model: the main burst clearly dominates over the merger-induced SF episode."138 Case 2 is represented by model (dashed line): he two SE episodes have a very close intensity ancl shape., Case 2 is represented by model (dashed line): the two SF episodes have a very close intensity and shape.139 Thev are separated by a short episode of galactic wind., They are separated by a short episode of galactic wind.140 Finally. à model in which the SE during the merger-induced rurst exceeds the initial episode (Case 3. model e)) is shown ον à dotted line.," Finally, a model in which the SF during the merger-induced burst exceeds the initial episode (Case 3, model ) is shown by a dotted line."141 At this point it is useful to discuss other possible sources of degeneraey in our model results., At this point it is useful to discuss other possible sources of degeneracy in our model results.142 First. of all. it should be noticed. that all the cases but the two Paper Es best models (labelled as in Table 1). fail in reproducing the present-day SNla rate observed in ellipticals (0.18x:0.06 SNu. Cappellaro et al 1999).," First of all, it should be noticed that all the cases but the two Paper I's best models (labelled as in Table 1), fail in reproducing the present-day SNIa rate observed in ellipticals $0.18\pm 0.06$ SNu, Cappellaro et al 1999)."143 For example. model predicts an explosion rate of ~ 1 SNu. whereas model gives 0.48 SNu.," For example, model predicts an explosion rate of $\sim$ 1 SNu, whereas model gives 0.48 SNu."144 The latter value. however. is only marginally consistent with the rate of 0.30.1 SNu given by Cappellaro οἱ al. (," The latter value, however, is only marginally consistent with the rate of $0.3\pm 0.1$ SNu given by Cappellaro et al. ("1451999) for cllipticals bluer than the average.,1999) for ellipticals bluer than the average.146 The reason for this disagreement is that the parameter zl in eq. (, The reason for this disagreement is that the parameter $A$ in eq. (1471). is set equal to 0.18 in order to have a good [it in the case of models andf£.,"1), is set equal to 0.18 in order to have a good fit in the case of models and."148 For comparative purposes. we run a model in which we use a lower value for A (CIable 1. moclel A=0.00)). namely we chose A=0.06 so that the present- SNla rate for model becomes 0.19 SNu.," For comparative purposes, we run a model in which we use a lower value for A (Table 1, model ), namely we chose A=0.06 so that the present-day SNIa rate for model becomes 0.19 SNu."149 This leads o luminosities and colours essentially unchanged. whereas he «Alg/Fev] ratio rises toward values more similar o Paper Es best model predictions. as a consequence of the ower Fe abundance in the ISM when the merger-induced starburst occurs.," This leads to luminosities and colours essentially unchanged, whereas the $[<Mg/Fe>_V]$ ratio rises toward values more similar to Paper I's best model predictions, as a consequence of the lower Fe abundance in the ISM when the merger-induced starburst occurs."150 However. a svstematie variation of A with he properties of the merger. would be in contrast with the definition of A. Another possible solution can be a variation either with ime or metallicity of the fraction of binary svstems (e.g. De Donder Vanbeveren. 2002).," However, a systematic variation of A with the properties of the merger, would be in contrast with the definition of A. Another possible solution can be a variation either with time or metallicity of the fraction of binary systems (e.g. De Donder Vanbeveren, 2002)."151 It should be remarked. rowever. that. in elliptical galaxies. the metallicity increases very fast in the first hundreds of Myr. reaching the solar value at the time of the galactic wind. and being higher han solar in the wind regime.," It should be remarked, however, that, in elliptical galaxies, the metallicity increases very fast in the first hundreds of Myr, reaching the solar value at the time of the galactic wind, and being higher than solar in the wind regime."152 Therefore. the (i.c. oesent-dav) value for A should. be achieved well before he occurrence of the second burst of star formation.," Therefore, the (i.e. present-day) value for A should be achieved well before the occurrence of the second burst of star formation."153 In any case. we do not consider this as a viable solution for reconciling star formation histories ruled by random gas accretion episodes with the observed. features of elliptical galaxies.," In any case, we do not consider this as a viable solution for reconciling star formation histories ruled by random gas accretion episodes with the observed features of elliptical galaxies."154 We also tried to change the IME. by using a typical flat one CXrimoto Yoshit LOST. models. and AYS7)).," We also tried to change the IMF, by using a typical flat one (Arimoto Yoshii 1987, models and )."155 This IME. was ruled out in Paper I because it predicts a too high metallicity in the stellar populations and too high values for «Mg/Fev]., This IMF was ruled out in Paper I because it predicts a too high metallicity in the stellar populations and too high values for $[<Mg/Fe>_V]$.156 Phe merging-induced starburst. when occurring in these models. pushes the ςAlg/fes] down toa reasonable value for galaxies of that given mass.," The merging-induced starburst, when occurring in these models, pushes the $[<Mg/Fe>_V]$ down to a reasonable value for galaxies of that given mass."157 On the other hand. the Fe and Mg abundances exceed the solar value by a factor of ten. at variance with what is suggested by the observations.," On the other hand, the Fe and Mg abundances exceed the solar value by a factor of ten, at variance with what is suggested by the observations."158 Phe possibility for a Hat IME curing a merging process was explored by Thomas (1999). who concluded. that a significantly [lat exponent (Le. ur«OLS) is needed to satisfv the constraints on. <Male].," The possibility for a flat IMF during a merging process was explored by Thomas (1999), who concluded that a significantly flat exponent (i.e. $x< 0.8$ ) is needed to satisfy the constraints on $[<Mg/Fe>_V]$ ."159 Finally. we tested the effect ofchanging the composition of the aecreted gas.," Finally, we tested the effect of changing the composition of the accreted gas."160 In. particular we focused on a solar composition. as given by Ancers Crevesse (1980).," In particular we focused on a solar composition, as given by Anders Grevesse (1989)."161 When this option is applied to modelανν the cdillerences. with the predictions shown in Table 1 are negligible.," When this option is applied to model, the differences with the predictions shown in Table 1 are negligible."162" ""This is a straightforward consequence. of the fact. that the total mass of the accreted. eas which is turned into stars is very small compared. to the mass of the dominant. stellar population.", This is a straightforward consequence of the fact that the total mass of the accreted gas which is turned into stars is very small compared to the mass of the dominant stellar population.163" A stronger effect is expected for models in which Al=Alpin,", A stronger effect is expected for models in which $M_{acc}=M_{lum}$.164 In this case. in fact. model would exhibit xΑΙοςV]=0.258. owing to the great. importance of the vouneest stellar populations in the weighted: average.," In this case, in fact, model would exhibit $[<Mg/Fe>_V]=0.258$, owing to the great importance of the youngest stellar populations in the weighted average."165 Also the colours become redder (by ~ 0.1 mag). but. not enough to match the CMLIBs.," Also the colours become redder (by $\sim$ 0.1 mag), but not enough to match the CMRs."166 On the other hand. the ellect of assuming Zo.=Z. on model predictions is milder for the chemical aspect. whereas the colours are  0.1 mag recdcdoer.," On the other hand, the effect of assuming $Z_{acc}=Z_{\odot}$ on model predictions is milder for the chemical aspect, whereas the colours are $\sim$ 0.1 mag redder."167 Vhis implies that. in any case. the time of the occurrence of the moerger-induced starburst is more important than the chemical composition of the accreted. gas.," This implies that, in any case, the time of the occurrence of the merger-induced starburst is more important than the chemical composition of the accreted gas."168 In this paper we completed the analysis of the photochemical properties of elliptical. galaxies begun in Paper LE. where we tested. the high. redshift’ formation. of ellipticals.," In this paper we completed the analysis of the photochemical properties of elliptical galaxies begun in Paper I, where we tested the high redshift formation of ellipticals."169 We showed that a photo-chemical diagnostic. namely the CALRs and the «Λοοον Ab. is needed in order to verify whether star formation histories differing from. one of the simple one lead to the same good. agreement with observations.," We showed that a photo-chemical diagnostic, namely the CMRs and the $[<Mg/Fe>_V]$ $M_V$, is needed in order to verify whether star formation histories differing from one of the simple one lead to the same good agreement with observations."170 This diagnostic allowed: us to quantify Paper Es best model response to perturbations. represented by merger-induced starbursts.," This diagnostic allowed us to quantify Paper I's best model response to perturbations, represented by merger-induced starbursts."171 Here we summarise our main conclusions:, Here we summarise our main conclusions:172The Be/X-ray systems represent the largest. sub-class of all High Mass. N-rav. Binaries (HINMND).,The Be/X-ray systems represent the largest sub-class of all High Mass X-ray Binaries (HMXB).173 A survey of the literature reveals that of the ~240 IINMXDs known in our Galaxy and the Magellanic Clouds (Liu et al..," A survey of the literature reveals that of the $\sim$ 240 HMXBs known in our Galaxy and the Magellanic Clouds (Liu et al.,"174 2005. 2006). alia fall within this class of binary.," 2005, 2006), $\ge$ fall within this class of binary."175 ln fact. in recent. vears it has emerged. that there is a substantial population of LAINBs in the SMC comparable in number to the Galactic population.," In fact, in recent years it has emerged that there is a substantial population of HMXBs in the SMC comparable in number to the Galactic population."176 Though unlike the Galactic population. all except one of the SAIC LAIN are Be star svstems.," Though unlike the Galactic population, all except one of the SMC HMXBs are Be star systems."177 In these systems the orbit of the Be Bsstar and the compact object. presumably a neutron star. is generally wide and eccentric.," In these systems the orbit of the Be star and the compact object, presumably a neutron star, is generally wide and eccentric."178 A-ray outbursts are normally associated. with the passage of the neutron star close to the cireumstellar disk (Okazaki Negueruela 2001). and generally are classified as Types I or HE (Stella. White Rosner. 1986).," X-ray outbursts are normally associated with the passage of the neutron star close to the circumstellar disk (Okazaki Negueruela 2001), and generally are classified as Types I or II (Stella, White Rosner, 1986)."179 The Type LE outbursts occur periodically at the time of the periastron passage of the neutron star. whereas Type LE outbursts are much more extensive and occur when the eieeumstellar material expands to fill most. or all of the orbit.," The Type I outbursts occur periodically at the time of the periastron passage of the neutron star, whereas Type II outbursts are much more extensive and occur when the circumstellar material expands to fill most, or all of the orbit."180 Phis paper concerns itself with Type E outbursts., This paper concerns itself with Type I outbursts.181 General reviews of such IMXD systems mav be found in Negueruela (1998). Corbet et al. (," General reviews of such HMXB systems may be found in Negueruela (1998), Corbet et al. ("1822008) and Coe et al. (,2008) and Coe et al. (1832000. 2008).,"2000, 2008)."184 This paper reports data acceuired over LO vears using the Rossi X-ray Timing Explorer (RXTE) on the HMND population of the SAIC., This paper reports data acquired over 10 years using the Rossi X-ray Timing Explorer (RXTE) on the HMXB population of the SMC.185 During the period. of these observations there have been many opportunities to study spin changes arising from accretion torques., During the period of these observations there have been many opportunities to study spin changes arising from accretion torques.186 This extremely homogeneous population permits the first high quality tests to be carried out. of the work of Ghosh Lamb (1979) and Joss Rappaport (1984)., This extremely homogeneous population permits the first high quality tests to be carried out of the work of Ghosh Lamb (1979) and Joss Rappaport (1984).187 The simplified source naming convention used in this work follows that established by Coe et al (2005)., The simplified source naming convention used in this work follows that established by Coe et al (2005).188contribution to the monochromatie flux is reduced.,contribution to the monochromatic flux is reduced.189" Up to o,%0.2ji the change of the curvature is mainly caused by the changing slope of the spatial intensity distribution. as discussed in refaQ-effects.."," Up to $a_{\rm gr}\approx 0.2\,{\rm\mu m}$ the change of the curvature is mainly caused by the changing slope of the spatial intensity distribution, as discussed in \\ref{a0-effects}."190 At larger gram radii the increase of the radius of the inner boundary and the corresponding reduction of the spatial frequencies. where 1544 approaches a constant value (or zero) becomes more important.," At larger grain radii the increase of the radius of the inner boundary and the corresponding reduction of the spatial frequencies, where $V_{2.11}$ approaches a constant value (or zero) becomes more important."191" The best match of the observed visibility τς is obtained for the model with «4,=0.2sau (dotted line in refsed-DP95-a0)). which is somewhat larger than a.= for the best matching model with the sil-Ow optical constants."," The best match of the observed visibility $V_{\rm obs}$ is obtained for the model with $a_{\rm gr}=0.2\,{\rm\mu m}$ (dotted line in \\ref{sed-DP95-a0}) ), which is somewhat larger than $a_{\rm192 gr}=0.16\,{\rm\mu m}$ for the best matching model with the sil–Ow optical constants."193 The resulting dust mass loss rate is about 30 smaller. and the radius of the inner boundary is about 12 smaller.," The resulting dust mass loss rate is about 30 smaller, and the radius of the inner boundary is about 12 smaller."194 Again the curvature of the model visibility is slightly too strong., Again the curvature of the model visibility is slightly too strong.195 But in contrast to the model with the optical constants. there is now an excess of flux at smaller wavelengths AZ55nn.," But in contrast to the model with the optical constants, there is now an excess of flux at smaller wavelengths $\lambda\la 5\,{\rm\mu m}$."196 As discussed at the end of refaÜ-effects.. we would expect a deficiency of flux if the CDS of has a disk-like structure.," As discussed at the end of \\ref{a0-effects}, we would expect a deficiency of flux if the CDS of has a disk–like structure."197 This suggests that the optical constants from David Péggourté (1995)) underestimate the extinction of the grain material at short wavelengths., This suggests that the optical constants from David Péggourié \cite{DP95}) ) underestimate the extinction of the grain material at short wavelengths.198 However. we also do not know. whether the optical constants from Ossenkopf et ((1992))," However, we also do not know, whether the optical constants from Ossenkopf et \cite{OHM92})"199Dwarf nova oscillations (DNOs) were first. cliscoverecl in outbursts of the dwart novae CN. Ori and Z Cam. and in he nova-like variable UN UMa (Warner Robinson 1972).,"Dwarf nova oscillations (DNOs) were first discovered in outbursts of the dwarf novae CN Ori and Z Cam, and in the nova-like variable UX UMa (Warner Robinson 1972)."200 They have since been observed in about 15 cwarl novae and 4 nova-likes (see Table 8.2 of Warner 19953)., They have since been observed in about 15 dwarf novae and 4 nova-likes (see Table 8.2 of Warner 1995a).201 Phe DNOs are ow amplitude. moderately. coherent. luminosity variations with periods in the range of 540 s. In the same kinds of cataclysmic variable (CV) stars there are also occasionally uminosity moculations of longer period (~501000 s) and »oor coherence. known as quasi-period oscillations (QDOs«). irst cliscussecl by Patterson. Robinson Nather (1977).," The DNOs are low amplitude, moderately coherent luminosity variations with periods in the range of 5–40 s. In the same kinds of cataclysmic variable (CV) stars there are also occasionally luminosity modulations of longer period $\sim$ 50–1000 s) and poor coherence, known as quasi-period oscillations (QPOs), first discussed by Patterson, Robinson Nather (1977)."202 A variety of statistical analyses ancl physical models rave been produced for the DNOs ancl QPOs (see review o» Warner 1995a). but no generally accepted: model vet exists.," A variety of statistical analyses and physical models have been produced for the DNOs and QPOs (see review by Warner 1995a), but no generally accepted model yet exists."203 Although there has been significant progress. at least in restricting possible models. by observations made in the zu ultraviolet. it is still the case that further informative observations are to be encouraged.," Although there has been significant progress, at least in restricting possible models, by observations made in the far ultraviolet, it is still the case that further informative observations are to be encouraged."204 Here we present new and archival observations of VW. Livi. and. discuss. their implications.," Here we present new and archival observations of VW Hyi, and discuss their implications."205 Aut first. in Sections 2 and 3. we review the status quo of observations and interpretations of DNOs and QPOs.," But first, in Sections 2 and 3, we review the status quo of observations and interpretations of DNOs and QPOs."206 In Section 4 we present and analyse the observations of VW Livi., In Section 4 we present and analyse the observations of VW Hyi.207 In a subsequent paper (Warner Woudt 2002. hereafter Paper IH). we develop a model to explain. these observations. and apply it to other systems.," In a subsequent paper (Warner Woudt 2002, hereafter Paper II), we develop a model to explain these observations, and apply it to other systems."208" In the optical. DNOs are low amplitude. usually sinusoical moculations in brightness of mocerate stability CO factors = P! of 10! 10"")."," In the optical, DNOs are low amplitude, usually sinusoidal modulations in brightness of moderate stability (`Q' factors = $\dot{P}^{-1}$ of $^4$ $^6$ )."209 Although not present in alb dwarf nova outbursts. when they are observable they usually appear about midway up the rising branch of outburst and disappear at à comparable brightness on the descending branch.," Although not present in all dwarf nova outbursts, when they are observable they usually appear about midway up the rising branch of outburst and disappear at a comparable brightness on the descending branch."210 Their coherence is maximal at the brightest. phase of outburst: on the descending branch they become Less coherent and dillicult to detect above the noise in the Fourier transforms., Their coherence is maximal at the brightest phase of outburst; on the descending branch they become less coherent and difficult to detect above the noise in the Fourier transforms.211 There is a very strong correlation between oscillation period P and svstem brightness (Warner Robinson 1972). such that 2 passes through a minimum about one day after visual maximum. which is when the EUV luminosity reaches maximum (Alauche 1996a.b).," There is a very strong correlation between oscillation period $P$ and system brightness (Warner Robinson 1972), such that $P$ passes through a minimum about one day after visual maximum, which is when the EUV luminosity reaches maximum (Mauche 1996a,b)."212 At this time. the rate of mass transfor AZ in the inner disc (and onto the white dwarl primary) reaches its maximum. (Cannizzo 1993).," At this time, the rate of mass transfer $\dot{M}$ in the inner disc (and onto the white dwarf primary) reaches its maximum (Cannizzo 1993)."213 Both the short periods of the DNOs and their correlation with AZ in the inner cise show that they have their origin near the surface of the primary (Warner 1995b)., Both the short periods of the DNOs and their correlation with $\dot{M}$ in the inner disc show that they have their origin near the surface of the primary (Warner 1995b).214 The amplitude and. phase changes of the optical DNOs observed. during eclipses show that in the optical the entire disc is in some manner involved., The amplitude and phase changes of the optical DNOs observed during eclipses show that in the optical the entire disc is in some manner involved.215 This is interpreted: as, This is interpreted as216The key to evaluatingR is to perform a geometric series expansion of the numerator of Eq. (14)).,"The key to evaluating$R$ is to perform a geometric series expansion of the numerator of Eq. \ref{eq:R-Formal}) ),"217 r.e.. Note that terminating this series after the first. term corresponds to the usual Maxwell-Boltzmann approximation in which case one obtains R=0 as expected.," i.e., Note that terminating this series after the first term corresponds to the usual Maxwell-Boltzmann approximation in which case one obtains $R=0$ as expected."218" Using the full expansion. however. leads to where the coefficients 7, (for = 2) are given by Eq. (27))"," Using the full expansion, however, leads to where the coefficients $r_{n}$ (for $n\geq2$ ) are given by Eq. \ref{eq:R}) )"219 gives R às a geometric series in. powers. of e: which is usually a small quantity., gives $R$ as a geometric series in powers of $e^{-Q/kT}$ which is usually a small quantity.220" The coefficients r, are also less than unity and they rapidly decrease with increasing 7.", The coefficients $r_{n}$ are also less than unity and they rapidly decrease with increasing $n$.221 Hence. the series given in Eq. (27))," Hence, the series given in Eq. \ref{eq:R}) )"222 is a rapidly convergent one., is a rapidly convergent one.223" However. the explicit determination of the 7, requires some attention to the energy dependence of the cross sections."," However, the explicit determination of the $r_n$ requires some attention to the energy dependence of the cross sections."224 Nevertheless. this task is greatly simplified when the forward thermonuclear reaction rates as a function of temperature (ow(KT); have already been compiled.," Nevertheless, this task is greatly simplified when the forward thermonuclear reaction rates as a function of temperature $\langle\sigma v(kT)\rangle$ have already been compiled."225 Combining the expression for the reaction rate in Eq. (6)), Combining the expression for the reaction rate in Eq. \ref{eq:Forward Reaction Rate}) )226 with Eq. (28)).," with Eq. \ref{eq:rn}) ),"227 the expansion coefficients become This immediately gives the correction factors in. terms of the tabulated rates. or converting to units of 75 and noting that reaction rate compilations are in terms of [N4£oY(75)5] we have These expressions for the reverse rate correction. factor R make clear the physics of the correction. factor.," the expansion coefficients become This immediately gives the correction factors in terms of the tabulated rates, or converting to units of $T_9$ and noting that reaction rate compilations are in terms of $[N_A \langle \sigma v(T_9) \rangle^* ]$ we have These expressions for the reverse rate correction factor $R$ make clear the physics of the correction factor."228 The factors become a sum of the correction factors in. terms of thermonuclear averages over decreasing temperatures. kT»KTín.," The factors become a sum of the correction factors in terms of thermonuclear averages over decreasing temperatures, $kT \rightarrow kT/n$."229 These terms achieve the task of increasing the photodissociation rate due to the fact that the Planck distribution includes many more low energy photons as (E+Q)/kT>0. which is what the [exp(CE-Q)/KT1—1] denominator in Eq. CH) ," These terms achieve the task of increasing the photodissociation rate due to the fact that the Planck distribution includes many more low energy photons as $(E + Q )/kT \rightarrow 0$, which is what the $[\exp{\{(E+Q)/kT\}} -1]$ denominator in Eq. \ref{eq:Inverse Reaction Rate 2}) )"230achieves., achieves.231 Eqs. (30)), Eqs. \ref{eq:RkT}) )232 and (31)) are the key equations for this paper., and \ref{eq:RT9}) ) are the key equations for this paper.233 Note that £ov(&7)) is a rapidly decreasing function as the temperature decreases. as well as the 577 pre-factor and the exp[—(—D)O/AT|.," Note that $\langle\sigma v(kT)\rangle $ is a rapidly decreasing function as the temperature decreases, as well as the $n^{-3/2}$ pre-factor and the $\exp{[-\left(n-1\right)Q/kT]}$."234 These three conditions guarantee that this is a well behaved. rapidly decreasing. convergent series even as ο20.," These three conditions guarantee that this is a well behaved, rapidly decreasing, convergent series even as $Q \rightarrow 0$."235 In what follows. we show some illustrations of the magnitude of these corrections and also derive some alternative analytic forms to illustrate the basic structure of these corrections in more detail.," In what follows, we show some illustrations of the magnitude of these corrections and also derive some alternative analytic forms to illustrate the basic structure of these corrections in more detail."236 Indeed. we show that in practice. only one or two terms are needed in the series and the correction factors are largely independent of the underlying nuclear structure.," Indeed, we show that in practice, only one or two terms are needed in the series and the correction factors are largely independent of the underlying nuclear structure."237 As a caveat to the reader. however. we note that the analytic approximations derived below do not include the stellar enhancement factors. and às such should be used with caution in a real astrophysical plasma.," As a caveat to the reader, however, we note that the analytic approximations derived below do not include the stellar enhancement factors, and as such should be used with caution in a real astrophysical plasma."238 When the projectile particle is charged. the capture reaction must tunnel through the Coulomb barrier at low energy.," When the projectile particle is charged, the capture reaction must tunnel through the Coulomb barrier at low energy."239 To account for this. the cross section can be factored into the following form (Fowler.Caughlan.&Zimmerman1967): Here. SCE) is called the astrophysical S-factor and Ec; is the Gamow energy which characterizes the penetrability: where and µ is in atomic mass units.," To account for this, the cross section can be factored into the following form \citep{fow67}: Here, $S\left(E\right)$ is called the astrophysical $S$ -factor and $E_{G}$ is the Gamow energy which characterizes the penetrability: where and $\hat \mu$ is in atomic mass units."240 The S factor contains information about the detailed nuclear interaction., The $S$ factor contains information about the detailed nuclear interaction.241 Away from resonances. the astrophysical SCE) factor is a slowly varying function. and at low temperatures the integrand is dominated by a small region known as the Gamow window (Fowler.Caughlan.&Zimmerman1967) located at an energy Ey as defined below.," Away from resonances, the astrophysical $S(E)$ factor is a slowly varying function, and at low temperatures the integrand is dominated by a small region known as the Gamow window \citep{fow67} located at an energy $E_0$ as defined below."242 As such. to the desired accuracy it can be replaced with an average effective value. S.CEo). Therefore. at low temperature S(E) will cancel in the ratio given in Eq. (28)).," As such, to the desired accuracy it can be replaced with an average effective value, $S_{\mbox{\footnotesize eff}}(E_0).$ Therefore, at low temperature $S(E)$ will cancel in the ratio given in Eq. \ref{eq:rn}) )."243 However. for higher temperatures the variation of the S-factor with energy over the Gamow window can become relevant.," However, for higher temperatures the variation of the $S$ -factor with energy over the Gamow window can become relevant."244 Hence. following Fowler.Caughlan.&Zimmerman(1967) we write where the dot denotes derivative with respect to energy.," Hence, following \cite{fow67} we write where the dot denotes derivative with respect to energy."245 Inserting this cross section into Eq. (28)), Inserting this cross section into Eq. \ref{eq:rn}) )246" the correction coefficient 7, becomes where the function F is defined as follows: This is a familiar integral in nuclear astrophysics (Fowler.Caughlan.", the correction coefficient $r_{n}$ becomes where the function $F$ is defined as follows: This is a familiar integral in nuclear astrophysics \citep{fow67}. .247&Zimmerman 1967).. Eq. (37), Eq. \ref{eq:Integral}) )248) corresponds to a product of the barrier penetrability times a Maxwellian distribution., corresponds to a product of the barrier penetrability times a Maxwellian distribution.249 It is strongly maximized in the Gamow window and well approximated (Fowler.Caughlan.&Zimmerman as a Gaussian integral near the maximum of the integrand.," It is strongly maximized in the Gamow window and well approximated \citep*{fow67,cla68,Iliadis07} as a Gaussian integral near the maximum of the integrand."250 Hence. we write where is the peak of the Gamow window and," Hence, we write where is the peak of the Gamow window and"251It was once thought that forming galaxies could be strong enuüssiou-liue sources because a large umuber of nlassive stars cause the ionization of gas clouds in the ornüue ealaxies (Partridge Pechles 1967).,It was once thought that forming galaxies could be strong emission-line sources because a large number of massive stars cause the ionization of gas clouds in the forming galaxies (Partridge Peebles 1967).252 Motivated w this idea. many attempts have been made to search or such very strong cluission-line sources at high redshift. mt most these searches failed (Pritchet 1991: Pahre Djorgovski 1995: Thompson. Maunucci. Beckwith 1996).," Motivated by this idea, many attempts have been made to search for such very strong emission-line sources at high redshift, but most these searches failed (Pritchet 1994; Pahre Djorgovski 1995; Thompson, Mannucci, Beckwith 1996)."253 ILowever. for these past several vears. Lya emitters lave con. found around kuown hieh-: objects such as quasars (a MeMalion 1996: Thi. AfcMahou. Eeami 1996: Petitjean et al.," However, for these past several years, $\alpha$ emitters have been found around known $z$ objects such as quasars (Hu McMahon 1996; Hu, McMahon, Egami 1996; Petitjean et al."254 1996: Tn. MeMahlon. Cowie 1999).," 1996; Hu, McMahon, Cowie 1999)."255 Futher. a nuuber of high-: Ίσα oenütters have also been found in several blank sky areas (Pascarelle et al.," Further, a number of $z$ $\alpha$ emitters have also been found in several blank sky areas (Pascarelle et al."256 1996: Cowie IIu 1998: Ίνουί et al., 1996; Cowie Hu 1998; Keel et al.257 1999: Steidel et al., 1999; Steidel et al.258 2000)., 2000).259 These recent surveys have reiuforced the poteutial iuportance of search for hieh-: Lya cutters., These recent surveys have reinforced the potential importance of search for $z$ $\alpha$ emitters.260 Tt is remarkable that two Lya cuutters found in Steidel et al. (, It is remarkable that two $\alpha$ emitters found in Steidel et al. (2612000) are observed to be very extended spatially. ee. ~ 100 kpe: we call these Lyra blobs LAB and LAB2.,"2000) are observed to be very extended spatially, e.g., $\sim$ 100 kpc; we call these $\alpha$ blobs LAB1 and LAB2."262 These two LABs ave found at redshift :z3.1 and show uo evidence for the association withACNs*., These two LABs are found at redshift $z \approx 3.1$ and show no evidence for the association with.263". The observational properties are sumnnarized as below (in this Letter. we adopt an Emstein«doe Sitter cosinology with a IIubble coustaut My=1002 kins ! Lye 1) the observed Lye luninosities are ~100)2 eres s.1. 2) they appear clongated morphologically, 3) their sizes amount to~ 100 5»i kpc. D) the observed line widths amount to ~1000 kus i. and 5) they ire uot associated with strong radio-continuuni sources such as powerful radio galaxies."," Their observational properties are summarized as below (in this Letter, we adopt an Einstein-de Sitter cosmology with a Hubble constant $H_0 = 100 h$ km $^{-1}$ $^{-1}$ ); 1) the observed $\alpha$ luminosities are $\sim 10^{43} h^{-2}$ ergs $^{-1}$, 2) they appear elongated morphologically, 3) their sizes amount to $\sim$ 100 $h^{-1}$ kpc, 4) the observed line widths amount to $\sim 1000$ km $^{-1}$, and 5) they are not associated with strong radio-continuum sources such as powerful radio galaxies."264 As for the origin of the LADs. two alternative ideas have been proposed.," As for the origin of the LABs, two alternative ideas have been proposed."265 One is that these LABs are superwiuds driven by the initial starburst iu galaxies because all the above properies as well as the observed frequency of LABs cau be explained in terms of the superwind model (Taniguchi Shiova 2000)., One is that these LABs are superwinds driven by the initial starburst in galaxies because all the above properties as well as the observed frequency of LABs can be explained in terms of the superwind model (Taniguchi Shioya 2000).266 Taniguchi Shiova (2000) also ciscussed the evolutionary lak from dust-euslirouded (or dusty) submillimeter sources (hereafter DSSs) to LABs because the ceutral starburst region iu a forming elliptical ealaxv could be eushrouded bv a lot of eas with dust erains., Taniguchi Shioya (2000) also discussed the evolutionary link from dust-enshrouded (or dusty) submillimeter sources (hereafter DSSs) to LABs because the central starburst region in a forming elliptical galaxy could be enshrouded by a lot of gas with dust grains.267 Their scenario ix sunuuarized as follows: Step I: The initial starburst occurs in the center of preealactic gas cloud., Their scenario is summarized as follows; Step I: The initial starburst occurs in the center of pregalactic gas cloud.268" Step II: This galaxy may be hidden by surrounding eas clouds for the first several tines 107 vears νο, the DSS phase)."," Step II: This galaxy may be hidden by surrounding gas clouds for the first several times $10^8$ years (i.e., the DSS phase)."269 Step I: The superwiud blows and thus the DSS phase ceases., Step III: The superwind blows and thus the DSS phase ceases.270 The superwind leads ο the formation of extended enissiou-line regions around t16 galaxy. (i.e... the LAB phase).," The superwind leads to the formation of extended emission-line regions around the galaxy (i.e., the LAB phase)."271 This lasts for a duration of ~1«10° vears., This lasts for a duration of $\sim 1 \times 10^8$ years.272 Aud. Step IV: The galaxy evolves to an ordinary elliptical ealaxy ~10? vears after the formation.," And, Step IV: The galaxy evolves to an ordinary elliptical galaxy $\sim 10^9$ years after the formation."273 This superwiud nodel predicts that the LADs are bright. at rest-frame ar-intrared if they are high-redshift. luminous aualogue of nearby superwind galaxies like Arp 220.," This superwind model predicts that the LABs are bright at rest-frame far-infrared if they are high-redshift, luminous analogue of nearby superwind galaxies like Arp 220."274" The other idea is that LABs are cooling radiation from xoto-ealaxies or dark matter halos (Παππά, Spaaus. Quataert 2000: Fardal et al."," The other idea is that LABs are cooling radiation from proto-galaxies or dark matter halos (Haiman, Spaans, Quataert 2000; Fardal et al."275 2001: Fabian et al., 2001; Fabian et al.276 19586: Tn 1992)., 1986; Hu 1992).277 Standard cold dark matter models predict that a aree ΠΟΥ of dark matter halos collapse at high redsΠΠ and they cau eiit significant Lya fluxes through collisional excitation of hydrogeji, Standard cold dark matter models predict that a large number of dark matter halos collapse at high redshift and they can emit significant $\alpha$ fluxes through collisional excitation of hydrogen.278 These Lvo cutting halos are :uso cousisteut with the observed linear sizes. velocity widlis. and να fluxes of the LABs.," These $\alpha$ emitting halos are also consistent with the observed linear sizes, velocity widths, and $\alpha$ fluxes of the LABs."279 However. it is uncertain |Low muuch far infrared and subiuilliueter contiuuiuu enission can be enütted because little is known about the «ust content and its spatial distribution iu such dark mater halos.," However, it is uncertain how much far infrared and submillimeter continuum emission can be emitted because little is known about the dust content and its spatial distribution in such dark matter halos."280 Very recently LABL has been detected at submillimeter, Very recently LAB1 has been detected at submillimeter281Note that both estimators give an information on memory and not on distribution of je process increments.,Note that both estimators give an information on memory and not on distribution of the process increments.282 The analvsis of the data shows that the tails of the underlving distribution conform to je power law., The analysis of the data shows that the tails of the underlying distribution conform to the power law.283 Hence. we model the data by a EARIMA process with Pareto innovations.," Hence, we model the data by a FARIMA process with Pareto innovations."284 As 1e power-law distributions belong to the domain of attraction of stable law (see. eg. 1994)," As the power-law distributions belong to the domain of attraction of stable law (see, eg. )"285) the resulting distribution of the FARIALA process V.i0uld be close to the stable one., the resulting distribution of the FARIMA process should be close to the stable one.286 We applied the MeCulloch. quantile fit to obtain the parameters of the distribution ikzimarkmainBodvCitationstarl29931986))., We applied the McCulloch quantile fit to obtain the parameters of the distribution ).287 The value of a was estimated to be 1.213. see Figure 3..," The value of $\alpha$ was estimated to be $1.213$, see Figure \ref{fig_alpha}."288 One may cheek that the estimated value of a of simulated FARIALA times series with Pareto innovations is usually underestimated. see e. e. Figure 3..," One may check that the estimated value of $\alpha$ of simulated FARIMA times series with Pareto innovations is usually underestimated, see e. g. Figure \ref{fig_alpha}."289 Therefore. we assume that the innovations in our model follow the Pareto law with a=1.25.," Therefore, we assume that the innovations in our model follow the Pareto law with $\alpha = 1.25$."290 According to (2005)). in order to recover both the self-similarity exponent Hi aud the memory parameter d (hence. the distribution parameter à) we can use the following5 BMW? computer test.," According to ), in order to recover both the self-similarity exponent $H$ and the memory parameter $d$ (hence, the distribution parameter $\alpha$ ) we can use the following $^2$ computer test."291 The surrogate5 data are obtained here by a random shullling5 of the original data positions., The surrogate data are obtained here by a random shuffling of the original data positions.292 The above formalism can be easily applied to determine basic features of an empirical data series., The above formalism can be easily applied to determine basic features of an empirical data series.293 Now we employ (lis to study an empirical (me series recorded from the system describing the energy of solar flares (Fig. 2))., Now we employ this to study an empirical time series recorded from the system describing the energy of solar flares (Fig. \ref{data}) ).294 The obtained values of the parameters are listed in Table 1.., The obtained values of the parameters are listed in Table \ref{tab1}.295 Therefore. from the results for (he surrogate data. the corresponding estimates for the parameter 1/6 are: αν=FHp;gp0.8452 and 1/04:=Hyg 0.7722.," Therefore, from the results for the surrogate data, the corresponding estimates for the parameter $1/\alpha$ are: $1/\alpha_{FIRT}=H_{FIRT}=0.8452$ and $1/\alpha_{VR}=H_{VR}=0.7722$ ."296" We observe that the estimators are close to the one assumed in our model: 1/04;í!"":=0.8.", We observe that the estimators are close to the one assumed in our model: $1/\alpha_{MC}=0.8$.297 Moreover. we choose d=0.19 as the highest acinissible value of d. which is close to the one obtained via the R/S method for the original data. see Table 1..," Moreover, we choose $d = 0.19$ as the highest admissible value of $d$ , which is close to the one obtained via the R/S method for the original data, see Table \ref{tab1}."298 One may notice that the estimators of MH obtained via FIRT ancl VI& methods. see Table L.. ave greater than theoretically acimissible in the FARIAITA moclel. ee. Chev exceed one.," One may notice that the estimators of $H$ obtained via FIRT and VR methods, see Table \ref{tab1}, are greater than theoretically admissible in the FARIMA model, e. they exceed one."299 As stated in (2008)). this can be justified by performing simulations ol the FARIMA processesancl estimating the parameter // on the simulated timeseries," As stated in ), this can be justified by performing simulations of the FARIMA processesand estimating the parameter $H$ on the simulated timeseries"300Helioseismic data from the Global Oscillation Network Group (GONG) and the Michelson Doppler Imager (MDI) have been used to infer the rotation rate in the solar interior (Thompson et al.,Helioseismic data from the Global Oscillation Network Group (GONG) and the Michelson Doppler Imager (MDI) have been used to infer the rotation rate in the solar interior (Thompson et al.301 1996; Schou et al., 1996; Schou et al.302 1998)., 1998).303 With the accumulation of these data for about 11 years it has now become possible to study its temporal variations., With the accumulation of these data for about 11 years it has now become possible to study its temporal variations.304 Previous studies have revealed a distinct pattern of bands of faster and slower than average rotation (Howe et al., Previous studies have revealed a distinct pattern of bands of faster and slower than average rotation (Howe et al.305 2000a; Antia Basu 2000). which at low latitudes nove towards the equator. and at high latitudes move towards the poles (Antia Basu 2001).," 2000a; Antia Basu 2000), which at low latitudes move towards the equator, and at high latitudes move towards the poles (Antia Basu 2001)."306 This pattern. which is similar to the torsional oscillations observed at the solar surface (Howard LaBonte 1980). has been well studied. and it has been found to penetrate through much of the body of the convection zone (Vorontsov et al.," This pattern, which is similar to the torsional oscillations observed at the solar surface (Howard LaBonte 1980), has been well studied, and it has been found to penetrate through much of the body of the convection zone (Vorontsov et al."307 2002; Basu Antia 2003: Howe et al., 2002; Basu Antia 2003; Howe et al.308 2005. 2006).," 2005, 2006)."309 Below the convection zone. however. there is no generally accepted unambiguous detection of significant temporal variation in the rotation. although Howe et al. (," Below the convection zone, however, there is no generally accepted unambiguous detection of significant temporal variation in the rotation, although Howe et al. ("3102000b. 2007) have reported a periodic. variation near the equator between r=0.65R. and r=0.75R.. in the early years of the last sunspot cycle with a period of 1.3 years. and Gough (2007) has reported that there is evidence that the oscillation penetrates even more deeply.,"2000b, 2007) have reported a periodic variation near the equator between $r=0.65R_{\odot}$ and $r=0.75R_{\odot}$ in the early years of the last sunspot cycle with a period of 1.3 years, and Gough (2007) has reported that there is evidence that the oscillation penetrates even more deeply."311 Other studies (Antia Basu 2000: Basu Antia 2003) have not confirmed such a variation., Other studies (Antia Basu 2000; Basu Antia 2003) have not confirmed such a variation.312 A direct consequence of temporal variation in the angular velocity is that global quantities like angular momentum and kinetic energy of rotation must also exhibit variation with time (cf..," A direct consequence of temporal variation in the angular velocity is that global quantities like angular momentum and kinetic energy of rotation must also exhibit variation with time (cf.,"313 Komm et al., Komm et al.314 2003)., 2003).315 Because of the steep increase in density with depth. the major contribution to these global quantities 1s naturally likely to arise from the deeper layers. where a temporal variation in rotation has not been determined reliably: and as a result. it is difficult to draw any firm conclusion about the global variation of angular momentum and kinetic energy.," Because of the steep increase in density with depth, the major contribution to these global quantities is naturally likely to arise from the deeper layers, where a temporal variation in rotation has not been determined reliably; and as a result, it is difficult to draw any firm conclusion about the global variation of angular momentum and kinetic energy."316 However. following Komm et al. (," However, following Komm et al. ("3172003). we may consider the contributions from different layers within the convection zone separately.,"2003), we may consider the contributions from different layers within the convection zone separately."318 We hope that such a study will shed some light on the rotational dynamics and its role in driving the activity cycle., We hope that such a study will shed some light on the rotational dynamics and its role in driving the activity cycle.319 Indeed. the study we present here is very similar to that of Komm et al. (," Indeed, the study we present here is very similar to that of Komm et al. ("3202003). although we are now able to extend it over an entire solar cycle.,"2003), although we are now able to extend it over an entire solar cycle."321 Apart from angular momentum and kinetic energy. we study also the temporal variation of the gravitational quadrupole and higher-order multipole moments due to the distortion in the solar figure induced by the known rotation.," Apart from angular momentum and kinetic energy, we study also the temporal variation of the gravitational quadrupole and higher-order multipole moments due to the distortion in the solar figure induced by the known rotation."322 In this work we concentrate principally on temporal variations in angular momentum and rotational kinetic energy in different regions in the solar convection zone. and compare their variation with the solar activity.," In this work we concentrate principally on temporal variations in angular momentum and rotational kinetic energy in different regions in the solar convection zone, and compare their variation with the solar activity."323 We describe in Section 2 the basic technique and the data used in the study: the main results are described in Section 3., We describe in Section 2 the basic technique and the data used in the study; the main results are described in Section 3.324 In Section + we present the variations 1n the multipole moments of the external gravitational potential., In Section 4 we present the variations in the multipole moments of the external gravitational potential.325 We draw our conclusions in Section 5., We draw our conclusions in Section 5.326 We use the data from GONG (Hill et al., We use the data from GONG (Hill et al.327 1996) and SOI/MDI (Schou 1999)., 1996) and SOI/MDI (Schou 1999).328 Each data set consists of meal frequencies of different (1./) multiplets. and the corresponding splitting coefficients.," Each data set consists of mean frequencies of different $(n,l)$ multiplets, and the corresponding splitting coefficients."329 We use 110 temporally overlapping data sets from GONG. each covering a period of 108 days. starting from 1995 May 7 and ending on 2006 May 20. each set being displaced by 36 days from its predecessor.," We use 110 temporally overlapping data sets from GONG, each covering a period of 108 days, starting from 1995 May 7 and ending on 2006 May 20, each set being displaced by 36 days from its predecessor."330 The MDI data consist of 49 contiguous data sets. each covering a period of 72 days. starting from 1996 May | and ending on 2006 May 20.," The MDI data consist of 49 contiguous data sets, each covering a period of 72 days, starting from 1996 May 1 and ending on 2006 May 20."331" These ""Sata cover the rising phase of the cycle 23 and most of the descending phase as well.", These data cover the rising phase of the cycle 23 and most of the descending phase as well.332 We use a 2D Regularized Least-Squares (RLS) inversion technique in the manner adopted by Antia et al. (, We use a 2D Regularized Least-Squares (RLS) inversion technique in the manner adopted by Antia et al. (3331998) to infer the angular velocity in the solar interior from each of the available data sets.,1998) to infer the angular velocity in the solar interior from each of the available data sets.334 To study the temporal variation in the angular velocity we look at the residuals obtained by subtracting from the angular velocity at any given time its temporal mean O6.6) with respect to the spherical polar coordinates (7.8). where tis time.," To study the temporal variation in the angular velocity we look at the residuals obtained by subtracting from the angular velocity at any given time its temporal mean $\Omega_0(r, \theta)$ with respect to the spherical polar coordinates $(r, \theta)$, where $t$ is time."335 The averaging to obtain Qu Is over the entire time interval in which the data are available., The averaging to obtain $\Omega_0$ is over the entire time interval in which the data are available.336 Both the temporal mean and the residuals were calculated separately for the GONG and MDI data., Both the temporal mean and the residuals were calculated separately for the GONG and MDI data.337 The residuals. which essentially. define the," The residuals, which essentially define the"338"selected photometric sample performing the same selection, Eqs. (1,,2,,3)),","selected photometric sample performing the same selection, Eqs. \ref{eq:selection1}, \ref{eq:selection2}, \ref{eq:selection3}) ),"339 in the SDSS spectroscopic sample., in the SDSS spectroscopic sample.340 In addition we only take the spectroscopic objects that overlap with our angular mask., In addition we only take the spectroscopic objects that overlap with our angular mask.341 For the bins where our analysis is performed we find fstars=441%.," For the bins where our analysis is performed we find $f_{stars}=4\pm 1\,\%$."342" That is, a negligible dependence with redshift and a broad agreement with the residual contamination found in comparable clustering studies at these redshifts (??)))."," That is, a negligible dependence with redshift and a broad agreement with the residual contamination found in comparable clustering studies at these redshifts \cite{2009arXiv0912.0511S,2011MNRAS.412.1669T}) )."343 The next step is to estimate what is the impact of this contaminants in the large scale angular clustering signal since they introduce a density gradient through the galactic plane., The next step is to estimate what is the impact of this contaminants in the large scale angular clustering signal since they introduce a density gradient through the galactic plane.344" Hence, we measure the correlation function of stars from the SDSS spectroscopic sample, relaxing the cut in mag50 to have enough statistics."," Hence, we measure the correlation function of stars from the SDSS spectroscopic sample, relaxing the cut in $mag50$ to have enough statistics."345 We also use the publicly available Tychos-2 star catalog (27) cut to the same selection and mask as our LRG sample to obtain a second estimate of the correlation of stars.," We also use the publicly available Tychos-2 star catalog \citep{2000A&A...363..385H,2000A&A...355L..27H} cut to the same selection and mask as our LRG sample to obtain a second estimate of the correlation of stars."346" Both determinations are in perfect agreement, as presented in Figure 12,, where the lines represent the correlation function for the SDSS sample, and the circles correspond to the Tychos-2 catalog."," Both determinations are in perfect agreement, as presented in Figure \ref{fig:stars}, where the lines represent the correlation function for the SDSS sample, and the circles correspond to the Tychos-2 catalog."347" This correlation is then included in the theoretical model for ω(θ) taking into account that LRGs and stars are uncorrelated populations, as (see also ?),,?))): where Wodsjmodel is the model for the “observed” correlation function, Wgai,model is the model for the true correlation function of galaxies, Wstars,fit is a simple fit to the measured correlation function for stars (see Figure 12))."," This correlation is then included in the theoretical model for $w(\theta)$ taking into account that LRGs and stars are uncorrelated populations, as (see also \cite{2006ApJ...638..622M}, \cite{2007ApJ...658...85M}) ): where $w_{obs,model}$ is the model for the “observed” correlation function, $w_{gal,model}$ is the model for the true correlation function of galaxies, $w_{stars,fit}$ is a simple fit to the measured correlation function for stars (see Figure \ref{fig:stars}) )."348 Notice that in Eq. (8)), Notice that in Eq. \ref{eq:wstars}) )349 we have removed the explicit dependence of the star fraction and correlation with redshift for simplicity., we have removed the explicit dependence of the star fraction and correlation with redshift for simplicity.350 In this section we will employ the measured angular correlation function to place joint constraints in the growth rate of structure and bias of the LRG sample., In this section we will employ the measured angular correlation function to place joint constraints in the growth rate of structure and bias of the LRG sample.351 We will use the theoretical model for the angular correlation function presented in ?).., We will use the theoretical model for the angular correlation function presented in \cite{2011MNRAS.tmp..385C}.352 It was extensively tested against mock catalogs of photometric surveys with specifications similar to our present case., It was extensively tested against mock catalogs of photometric surveys with specifications similar to our present case.353 It was shown that the inclusion of redshift-space distortions and bias to linear order together with a model for non-linear matter clustering accurately reproduced the measured angular correlation in scales 0=0.5? for redshift bins centered at z~0.5., It was shown that the inclusion of redshift-space distortions and bias to linear order together with a model for non-linear matter clustering accurately reproduced the measured angular correlation in scales $\theta \gtrsim 0.5^\circ$ for redshift bins centered at $z\sim 0.5$.354 We now recall some basic expressions of the model and refer the reader to 7) for further details., We now recall some basic expressions of the model and refer the reader to \cite{2011MNRAS.tmp..385C} for further details.355" 'The model angular correlation function is given by,"," The model angular correlation function is given by,"356here show that the mass loss rate of AB Dor. aud prestunably also of other rapidl-rotatiug voung late-type stars. is substantially hieher than the solar mass loss rate. and that it could be as high as LOHHAF.Vr tas suggested by Woodetal.(2001.2005).,"here show that the mass loss rate of AB Dor, and presumably also of other rapidly-rotating young late-type stars, is substantially higher than the solar mass loss rate, and that it could be as high as $10^{-12}-10^{-11}\;M_\odot\;Yr^{-1}$ , as suggested by \cite{Wood04,Wood05}."357. These values however. are strongly dependent ou the assunued average coronal base deusitv.," These values however, are strongly dependent on the assumed average coronal base density."358 It secs ikelv. based on measurements based on X-ray spectra xesunablv originating from) plasma im closed loops (Section 2)) that this is ecnerally hieher than the solar case. probably by au order of maeuitucde.," It seems likely, based on measurements based on X-ray spectra presumably originating from plasma in closed loops (Section \ref{sec:Simulation}) ) that this is generally higher than the solar case, probably by an order of magnitude."359 For such a case. he predicted mass loss rate for AB Dor is about 100 nues the solar rate.," For such a case, the predicted mass loss rate for AB Dor is about 100 times the solar rate."360 The amass aud angular momentum loss rates found roni our MIID models here are intriguing for the wider woblem of stellar rotational evolution., The mass and angular momentum loss rates found from our MHD models here are intriguing for the wider problem of stellar rotational evolution.361 Models sich as we present here could. iu principle. be euiployed to map out theoretical AMIL as a fiction of stellar activity and rotation rate.," Models such as we present here could, in principle, be employed to map out theoretical AML as a function of stellar activity and rotation rate."362 While the general picture of stellar spin-down with age as a result of wind-driven AAIL. emerged decades ago (6.8.Schatzmanun1962:Weber&Davis1967:Moestel1968:Slaunanich 1972).. the details of situation has proved somewhat complicated aud rotation rate data amassed in the intervening vears for late-tvpe stars exhibits a complex dispersion over stellar age aud luas.," While the general picture of stellar spin-down with age as a result of wind-driven AML, emerged decades ago \cite[e.g.][]{schatzmann62,weberdavis67,mestel68,skumanich72}, the details of situation has proved somewhat complicated and rotation rate data amassed in the intervening years for late-type stars exhibits a complex dispersion over stellar age and mass."363 Faster rotation curing stellar vouth eugeuders ereater magnetic activity through rotationally-powered dynamo action aud the correlation of magnetic activity indicators such as chromospheric cussion lines and coronal N-vav luminosity with rotation is well-established., Faster rotation during stellar youth engenders greater magnetic activity through rotationally-powered dynamo action and the correlation of magnetic activity indicators such as chromospheric emission lines and coronal X-ray luminosity with rotation is well-established.364 Observations of Lye absorption bv the interactions of stellar winds with the surrounding ISAL - the stellar equivalent of the heliopause - also indicate that stellar wind mass-loss rates are larecr for vounser and more active stars., Observations of $\alpha$ absorption by the interactions of stellar winds with the surrounding ISM - the stellar equivalent of the heliopause - also indicate that stellar wind mass-loss rates are larger for younger and more active stars.365 Woodetal.(2002). estimate a relation MExt2004082 ased on combining inferred mass-loss rates with N-rav activity and observed N-ray activity vs stellar age.," \citet{wood02} estimate a relation $\dot{M}\propto t^{-2.00\pm3660.52}$, based on combining inferred mass-loss rates with X-ray activity and observed X-ray activity vs stellar age."367 Their relation suggests that at very fast rotation rates. mass-loss should approach 1000 times the solar value. though they caution against the reliability of this extrapolation.," Their relation suggests that at very fast rotation rates, mass-loss should approach 1000 times the solar value, though they caution against the reliability of this extrapolation."368 Tere. we fiud that plausible coronal base deusities lead to mass loss rates of LOO times that of the present day Sun.," Here, we find that plausible coronal base densities lead to mass loss rates of 100 times that of the present day Sun."369 Observations of rotation rates for stars in open clusters indicate that very voune stars with ages of up to 100 My or so are not as rapidly spuu-down as would be expected based on extrapolation of the Skuuanich(1972) spin-down relation. aud theoretical spin-down modelling efforts have invoked a magnetic “saturation” that limits he AML for very short rotation periods (egChaboveretal.1995:I&xishnaiaurthi1997:Barnes 2003).," Observations of rotation rates for stars in open clusters indicate that very young stars with ages of up to 100 Myr or so are not as rapidly spun-down as would be expected based on extrapolation of the \citet{skumanich72} spin-down relation, and theoretical spin-down modelling efforts have invoked a magnetic “saturation” that limits the AML for very short rotation periods \citep[eg][]{Chaboyer95,krishnamurthi97,barnes03}."370. While our wind model for a single star. such as AB Dor xeseuted here. cannot in itself be used to validate sucli a saturation approach. the ecueral moethodologv does in principle allow for a more thorough exploration of he relevant parameter space to provide an MIID wind xedietion of angular momentum loss as a fiction of stellar age.," While our wind model for a single star, such as AB Dor presented here, cannot in itself be used to validate such a saturation approach, the general methodology does in principle allow for a more thorough exploration of the relevant parameter space to provide an MHD wind prediction of angular momentum loss as a function of stellar age."371" The requisite mput paraueters here would o the global coronal base density, and the magnetic field streneth."," The requisite input parameters here would be the global coronal base density, and the magnetic field strength."372 The spindown time of AB Dor cau be estimated as (Ud-Doulaetal.2009):: where we used &=0.1 and write 6a=XAR..., The spindown time of AB Dor can be estimated as \citep{UdDoula09}: where we used $k=0.1$ and write $r_A=XR_\star$ .373 Based ou Equation ? aud the caleulated mass loss rates. the spindown time of AB Dor can range from 10?1072yy depending ou the particular case and on the value of X (which ranges from 5-10).," Based on Equation \ref{TauSpindown} and the calculated mass loss rates, the spindown time of AB Dor can range from $10^9-10^{12}\;yr$ depending on the particular case and on the value of $X$ (which ranges from 5-10)."374 A value of 10?yr leads o a rotation of P=05d.c?τιd after 5 illion vears., A value of $10^9\;yr$ leads to a rotation of $P=0.5\;d\cdot e^5=71\;d$ after 5 billion years.375 However. Cohenetal.(2009) have shown hat the augular momentum loss rate can be 3-14 times Heher when the stellar maguetic field is dominated by strong polar spots. as appear to characterize voung. ‘ast-rotating stars.," However, \cite{Cohen09b} have shown that the angular momentum loss rate can be 3-4 times higher when the stellar magnetic field is dominated by strong polar spots, as appear to characterize young, fast-rotating stars."376 We expect the aneular moment oss rate to decrease with time as AB Dor becomes an established main-sequence star with spots at lower atitudes., We expect the angular momentum loss rate to decrease with time as AB Dor becomes an established main-sequence star with spots at lower latitudes.377 Therefore. after five billion vears a rotation siular to that of the Sun nüsht be expected.," Therefore, after five billion years a rotation similar to that of the Sun might be expected."378" Iu Equation 7.. 7 is indepeudent of O,."," In Equation \ref{TauSpindown}, $\tau$ is independent of $\Omega_\star$."379 Therefore. for the sale paranueters but with differeut rotation rates. we have Pyο=25d/0.5d50 aud Jo/.J4=50 (the inverse of τι/72).," Therefore, for the same parameters but with different rotation rates, we have $P_1/P_2=25\;d/0.5\;d=50$ and $\dot{J}_2/\dot{J}_1=50$ (the inverse of $\tau_1/\tau_2$ )."380" The ratios of the aueular moment loss rates for Cases A-C with different rotation rates are 60.68. and SO. respectively,"," The ratios of the angular momentum loss rates for Cases A-C with different rotation rates are 60,68, and 80, respectively."381 This is consistent with the expected idealized value., This is consistent with the expected idealized value.382 Tn principle. we can propose a relation between the wind in our model and the Alfvénu radius as follows.," In principle, we can propose a relation between the wind in our model and the Alfvénn radius as follows."383" Since at the Alfvénu radius. a,=c4 we have: where we asstune that Bir)μα/7c)? aud ∣≻⋃⋖↿⋆∣⋟↖↖↕↑↓↗⋃∙∣≻⋃⋅∡⋯≼∠⇀⋯⋝↸∐↕∶↴∙↑∐⋯≺↧∶↴∙∐↸⊓↸Ra2 ath B and oc beine the ⋠⋉∖⋅⊳ ∐↸∖↕≼↧∙≼∐∖∐↴∖↴↕↑⋅↖⇁∙⋜⋯≼↧⋎∐⋟↖⇁↸↗∖∐∐↴∖↴↻↸∖↸∖≼↧⋜↕↑↑↕∐∖∏∏↘⊓∏⋈∖↴⋝⋜↧↴∖↴↸∖∙ ↥⋅↸∖"," Since at the Alfvénn radius, $u_{sw}=v_A$ we have: where we assume that $B(r)=B_0(R_\star/r)^3$ and $\rho(r)=\rho_0(R_\star/r)^2$ , with $B_0$, $\rho_0$, and $v_{A0}$ being the magnetic field, density, and Alfvénn speed at the flux tube base, respectively."384↴∖↴↻↸∖↸⊳↑↕↖↽↸∖↕⋅↖⇁∙⊟⋅∪⋯⊏≺∣∏⋜↧⊓∪∐≺∖∖↖↖⇁↸∖∶↴∙⊾↸∖↑∶ which could provide. i principle the location of the Alfvénn radius based on a known magnetic field distribution aud surface deusitv.," From Equation \ref{AlfvenRadius1} we get: which could provide, in principle the location of the Alfvénn radius based on a known magnetic field distribution and surface density."385 The relation however. is not trivial due to the non-linear relation between the paralucters that define it.," The relation however, is not trivial due to the non-linear relation between the parameters that define it."386 In reality. the radial fuuctious of the magnetic field aud censity used in Equation 8 are uot necessarily valid within the Alfvéun surface.," In reality, the radial functions of the magnetic field and density used in Equation \ref{AlfvenRadius1} are not necessarily valid within the Alfvénn surface."387 We have carried out a 3-dimeusional elobal umucerical ΑΠΟ suuulatiou of the corona of AB Dor. driven bv a Zecman-Doppler huage imaguetie surface map.," We have carried out a 3-dimensional global numerical MHD simulation of the corona of AB Dor, driven by a Zeeman-Doppler Image magnetic surface map."388 We studied three test cases with different base deusitv. aud we also compared the solutions with fast aud slow stellar rotations., We studied three test cases with different base density and we also compared the solutions with fast and slow stellar rotations.389 We find that the coronal structure of AB Dor is donunated by the azimuthal taugliug of the coronal maguetic feld as a result of rapid τotation., We find that the coronal structure of AB Dor is dominated by the azimuthal tangling of the coronal magnetic field as a result of rapid rotation.390 Based ou the AIIID solution. we calculate a reaistic Alfvéónn surface. which enables us to estiuate the mass and anenlar imonientum loss rates;," Based on the MHD solution, we calculate a realistic Alfvénn surface, which enables us to estimate the mass and angular momentum loss rates."391 Our main fiwingis that the mass loss rate is dependent on the value of the average coronal base density. as well as the corona field and stellar wind topology whichare affected by rapida stellar rotation.," Our main findingis that the mass loss rate is dependent on the value of the average coronal base density, as well as the coronal field and stellar wind topology whichare affected by rapid stellar rotation."392 The, The393e.g. Dickinson 1998).,"e.g., Dickinson 1998)."394 Prior to stacking. we verified that none of the objects being stacked. was individually detected byChandra.," Prior to stacking, we verified that none of the objects being stacked was individually detected by."395 This was accomplished by searching for nearby (within 2”) X-ray sources found by with a false-positive probability threshold of 10.? (Dobrzycki et al., This was accomplished by searching for nearby (within $2\arcsec$ ) X-ray sources found by with a false-positive probability threshold of $10^{-5}$ (Dobrzycki et al.396 1999: Freeman et al., 1999; Freeman et al.397 2001): it led to the removal of one source. the 2=3.179 ΔΝ 123639.62-621230 (associated with CXOHDEN J123639.54-621230).," 2001); it led to the removal of one source, the $z=3.479$ AGN 123639.6+621230 (associated with CXOHDFN J123639.5+621230)."398 We lave also manually inspected. iinages in all tliree of the standard X-ray bauds and removed three Lyinan break galaxies lying near uurelated X-ray sources., We have also manually inspected images in all three of the standard X-ray bands and removed three Lyman break galaxies lying near unrelated X-ray sources.399 123615.323-6211253 was removed because it lies withiu the region covered by the faint extended source CNOHDEN 1236[5.04-621112., 123645.3+621153 was removed because it lies within the region covered by the faint extended source CXOHDFN J123645.0+621142.400" 1236[5.54-621111 and 12365L7+62131| were removed because they lie relatively near (6—8"") the bright X-ray sources CNOHDEN J123616.34-621101 aud CXOHDEN J123655.14621311. respectively. aud. could be affected by the low-level wines of the PPSF."," 123645.8+621411 and 123654.7+621314 were removed because they lie relatively near $\arcsec$ ) the bright X-ray sources CXOHDFN J123646.3+621404 and CXOHDFN J123655.4+621311, respectively, and could be affected by the low-level wings of the PSF."401 While we cousider it uulikely that these three Lyiuau break galaxies suffer (rom significant contamination by the unrelated X-ray sources. we wanted to be cautious :ud avoid any. problems with the backeround estimation.," While we consider it unlikely that these three Lyman break galaxies suffer from significant contamination by the unrelated X-ray sources, we wanted to be cautious and avoid any problems with the background estimation."402 Figure 1 shows the redshilts aud J? inagniides of the 2[ selected objects., Figure 1 shows the redshifts and $R$ magnitudes of the 24 selected objects.403 The inediau recshift of the selected objects is z=2.92. corresponding to a inediau lookback time of 11.9 Cyr ol the age of the Universe) aud a meclau luminosity distance of 25.800 Mpe (Hoge 1999).," The median redshift of the selected objects is $z=2.92$, corresponding to a median lookback time of 11.9 Gyr of the age of the Universe) and a median luminosity distance of 25,800 Mpc (Hogg 1999)."404 These galaxies have a range of morphologies but generally appear nonrelaxed., These galaxies have a range of morphologies but generally appear nonrelaxed.405 They do not show sigus of containing ACN iu their optical spectra. although the spectral quality is often limited.," They do not show signs of containing AGN in their optical spectra, although the spectral quality is often limited."406 They were not identiliel as ΑΝ candidates by Jarvis MacAlpine (1998). Conti et al. (," They were not identified as AGN candidates by Jarvis MacAlpine (1998), Conti et al. ("4071999). or Sarajedini οἱ al. (,"1999), or Sarajedini et al. ("4082000).,2000).409 Two of the 21 galaxies have weak (6-8 Jy) radio emission at 8.5 GHz (Richards et al., Two of the 24 galaxies have weak (6–8 $\mu$ Jy) radio emission at 8.5 GHz (Richards et al.410 1998): this emission could be due to either starburst activity or AC, 1998); this emission could be due to either starburst activity or AGN.411 For each of the standard X-ray. bands. we stacked “cutout” images centered ou the positious ol the 21 selected galaxies.," For each of the standard X-ray bands, we stacked “cutout” images centered on the positions of the 24 selected galaxies."412 For the photometry below we use an approximately circular aperture comprised of the 30 pixels with centers within 175 of the stacking position (see 833.L of Paper IV)., For the photometry below we use an approximately circular aperture comprised of the 30 pixels with centers within $1\farcs 5$ of the stacking position (see 3.4 of Paper IV).413 In Figure 2 we show the nuuber of full-baud aud soft-baud counts obtained for each of the incividual galaxies being stacked. aud in Figure 3 we show images created from the stacked data.," In Figure 2 we show the number of full-band and soft-band counts obtained for each of the individual galaxies being stacked, and in Figure 3 we show images created from the stacked data."414 The stacked ull-band image lias au effective exposure time of 22.11 Ms (260 days)., The stacked full-band image has an effective exposure time of 22.44 Ms (260 days).415 In this image. we find 161 wuts within the 30-pixel aperture. while 131-3 are expected [rom the backgrouud.," In this image, we find 164 counts within the 30-pixel aperture, while 134.3 are expected from the background."416 The Poissou srobability of obtaining 161 counts or more when 131.3 are expected is 5x10oE iudicating a detection at the coulideuce level.," The Poisson probability of obtaining 164 counts or more when 134.3 are expected is $5\times 10^{-3}$, indicating a detection at the confidence level."417 In the stacked soft-band image (22.16 Ms exposure). we iud 13 counts when 26.2 are expected [rom background. corresponding to a detection at the coulicence level.," In the stacked soft-band image (22.46 Ms exposure), we find 43 counts when 26.2 are expected from background, corresponding to a detection at the confidence level."418 The higher sienificauce of the soft-band detection is most likely due to the lower ockgrouud in this baud., The higher significance of the soft-band detection is most likely due to the lower background in this band.419 No significant cletection is obtained in the hard baud: unlortunately. the spectral constraint implied by the bard-baud noucletection is not physically interesting.," No significant detection is obtained in the hard band; unfortunately, the spectral constraint implied by the hard-band nondetection is not physically interesting."420 Following Appendix A of Paper IV. we performed Monte Carlo stacking simulations to assess [alse-detectiou probabilities empirically.," Following Appendix A of Paper IV, we performed Monte Carlo stacking simulations to assess false-detection probabilities empirically."421 For each band. we performed: 100.000. trials where we stacked 21 raudom positions using the same photometry aperture as was used for the Lyman break galaxies above.," For each band, we performed 100,000 trials where we stacked 24 random positions using the same photometry aperture as was used for the Lyman break galaxies above."422" The random positions were chosen to lie within 15""x ""local background regious"" centered on each of the Lyinau break galaxies (avoidiug known", The random positions were chosen to lie within $15\arcsec\times 15\arcsec$ “local background regions” centered on each of the Lyman break galaxies (avoiding known423" The random positions were chosen to lie within 15""x ""local background regious"" centered on each of the Lyinau break galaxies (avoidiug known.", The random positions were chosen to lie within $15\arcsec\times 15\arcsec$ “local background regions” centered on each of the Lyman break galaxies (avoiding known424Taking only groups and clusters with sizes less than 0.11 Mpc h! for Nmem = 3 and less than 0.49 Mpc h! for Nmem = 4 produces a new catalogue which is 94% complete and 88% pure.,Taking only groups and clusters with sizes less than 0.11 Mpc $h^{-1}$ for $_{mem}$ = 3 and less than 0.49 Mpc $h^{-1}$ for $_{mem}$ = 4 produces a new catalogue which is $\%$ complete and $\%$ pure.425 These values were chosen because they provide the best improvement to the purity with minimal change to the completeness., These values were chosen because they provide the best improvement to the purity with minimal change to the completeness.426" Thus, by clipping out clusters with a low richness few members), but relatively large sizes, we can improve the purity by around 36%, while reducing the completeness by only 4%."," Thus, by clipping out clusters with a low richness few members), but relatively large sizes, we can improve the purity by around $\%$ , while reducing the completeness by only $\%$."427" Applying this same clipping procedure to the ‘real’ 2SLAQ groups and clusters, we can separate the catalogue into ‘gold’ and ‘silver’ samples."," Applying this same clipping procedure to the `real' 2SLAQ groups and clusters, we can separate the catalogue into `gold' and `silver' samples."428 Where the gold clusters are those that pass the clipping procedure and therefore are the most likely to be genuine., Where the gold clusters are those that pass the clipping procedure and therefore are the most likely to be genuine.429" The silver clusters fail the clipping procedure and may still be genuine, but the probability is lower."," The silver clusters fail the clipping procedure and may still be genuine, but the probability is lower."430" Out of the 313 total 28LAQ groups and clusters, 70 are gold and the remaining 243 are silver."," Out of the 313 total 2SLAQ groups and clusters, 70 are gold and the remaining 243 are silver."431" The gold sample groups and clusters have an average velocity dispersion of v,=585.28 kms and an average size of Rat=0.90 Mpc h!, while the !silver sample groups and clusters have an average velocity dispersion of συ=430.47 kms and an average size of Rei=0.53 Mpc !."," The gold sample groups and clusters have an average velocity dispersion of $\overline{\sigma}_{v}=585.28$ $^{-1}$ and an average size of $\overline{R}_{clt}=0.90$ Mpc $h^{-1}$, while the silver sample groups and clusters have an average velocity dispersion of $\overline{\sigma}_{v}=430.47$ $^{-1}$ and an average size of $\overline{R}_{clt}=0.53$ Mpc $h^{-1}$."432 The average! properties of the gold sample are larger than those of the full (gold+silver) sample as a large fraction of the small structures have been removed., The average properties of the gold sample are larger than those of the full $+$ silver) sample as a large fraction of the small structures have been removed.433" It should be noted that, although we are confident that our mock is a good representation of the 28LAQ galaxies, it is difficult to interpret how well this clipping procedure will translate to the real 28LAQ catalogue."," It should be noted that, although we are confident that our mock is a good representation of the 2SLAQ galaxies, it is difficult to interpret how well this clipping procedure will translate to the real 2SLAQ catalogue."434" Therefore we do not remove the silver sample clusters, rather we assign to them a lower likelihood of being genuine than the gold sample clusters."," Therefore we do not remove the silver sample clusters, rather we assign to them a lower likelihood of being genuine than the gold sample clusters."435 The final 28LAQ cluster catalogue contains an additional column that identifies each cluster as belonging to the gold or silver samples with ‘G’ or ‘S’ respectively., The final 2SLAQ cluster catalogue contains an additional column that identifies each cluster as belonging to the gold or silver samples with `G' or `S' respectively.436 In fig.8 and 9 the groups and clusters that form part of the gold sample are highlighted with gold asterisks., In \ref{fig:vel_d} and \ref{fig:size_d} the groups and clusters that form part of the gold sample are highlighted with gold asterisks.437 Fig., Fig.438" 13. shows SDSS DR7 optical and i--band colour images of the regions around the DFoF detected clusters CL.0008, CL.2204, CL.1122, CL.0024 and CL.0038."," \ref{fig:images} shows SDSS DR7 optical and -band colour images of the regions around the DFoF detected clusters 008, 204, 122, 024 and 038."439" The yellow circles highlight the positions of the cluster member galaxies, which are labeled with their individual redshifts."," The yellow circles highlight the positions of the cluster member galaxies, which are labeled with their individual redshifts."440 These clusters were chosen to sample serval different redshift bins as listed in Table 2.., These clusters were chosen to sample serval different redshift bins as listed in Table \ref{tab:zrange}.441 These images show that the galaxies are distributed in a small area on the sky and are close in redshift space., These images show that the galaxies are distributed in a small area on the sky and are close in redshift space.442 An important point to notice is that the distribution of cluster LRGs in each image varies from spherical to elongated filamentary structures., An important point to notice is that the distribution of cluster LRGs in each image varies from spherical to elongated filamentary structures.443" This is an advantage of the FoF method, which makes no prior assumptions about the shape of the groups and clusters.CL.2204 (top right panel) in particular shows a rather filamentary distribution of cluster"," This is an advantage of the FoF method, which makes no prior assumptions about the shape of the groups and clusters.204 (top right panel) in particular shows a rather filamentary distribution of cluster"444Strónuugren photometry has been proven to be a reliable metallicity indicator for elobular cluster elauts and subeliants (0.5. Richtler 1989.. Crebel Richtler 1992 and references thereii).,"Strömmgren photometry has been proven to be a reliable metallicity indicator for globular cluster giants and subgiants (e.g. Richtler \cite{rich89}, Grebel Richtler \cite{greb92} and references therein)."445 The location of late type (C aud I) elants in the Srónuugren (bg).nmn diagram is correlated with their metallicitics. especially with their iron aud CN albmucdauces.," The location of late type (G and K) giants in the Strömmgren $(b-y),m_1$ diagram is correlated with their metallicities, especially with their iron and CN abundances."446" Whereas the color (b1) is not sensitive to metallicity. the Strónuueren ο filter includes several on absorption lines as well as the CN band at1215AÀ.. and therefore m,=(eby(bgy) is a metallicitv sensitive index (e.g. BellCorstafsson 1978))."," Whereas the color $(b-y)$ is not sensitive to metallicity, the Strömmgren $v$ filter includes several iron absorption lines as well as the CN band at, and therefore $m_1 =447(v-b) - (b-y)$ is a metallicity sensitive index (e.g. BellGustafsson \cite{bell}) )."448 Within a certain color range. 0.5«(bgy)l.l mag. the loci of constaut irou abundance of eiauts and supergiants cui be approximated by straight lues.," Within a certain color range, $0.5 < (b-y) < 1.1$ mag, the loci of constant iron abundance of giants and supergiants can be approximated by straight lines."449" This is valid for CN-""normal” (= CN-weak) stars.", This is valid for CN-“normal” $=$ CN-weak) stars.450 CN- stars. due to their higher absorption in the ο filter. scatter to higher iy values aud therefore mimic a higher Strónmieren metalliitv than their actual iron abunudauce would correspond to.," CN-strong stars, due to their higher absorption in the $v$ filter, scatter to higher $m_1$ values and therefore mimic a higher Strömmgren metallicity than their actual iron abundance would correspond to."451 If a cluster is κjiowii to have only sunall star-to-star variations du its mon abundance. the scatter to higher iy values cau be used to uncover CN-rich member stars.," If a cluster is known to have only small star-to-star variations in its iron abundance, the scatter to higher $m_1$ values can be used to uncover CN-rich member stars."452 For giants redder han (b.y)=Ll mae the calibration breaks down due to TiO aud Mgll isorption iu the y baid., For giants redder than $(b-y) = 1.1$ mag the calibration breaks down due to TiO and MgH absorption in the $y$ band.453 Until now the calibration for SrOónuueren metallicities is based on the spectroscopically determined iron imudauce of relative yoaetalrich stars. and onlv few metal poorer stars (Caehel Richtler 1992)).," Until now the calibration for Strömmgren metallicities is based on the spectroscopically determined iron abundance of relatively metal-rich stars, and only few metal poorer stars (Grebel Richtler \cite{greb92}) )."454 This calibration worked well for the determunation of cluster abuudances in the Magellanic Clouds (Iilker et al. 1995a.. 1995h::," This calibration worked well for the determination of cluster abundances in the Magellanic Clouds (Hilker et al. \cite{hilk95a}, \cite{hilk95b};"455 Dirsch et al. 19993).," Dirsch et al. \cite{dirs}) ),"456 but filed to reproduce the slope of constant iueallicity lines in the (5g).ng dciaerai for the new data of metal poor globular clusters.," but failed to reproduce the slope of constant metallicity lines in the $(b-y),m_1$ diagram for the new data of metal poor globular clusters."457" Other previous iietallicitv calibrations iu the Stróuuuereu (b.ug).my plane have been published for the regime of F and C dwarfs (Schuster Nissen 19893) aud F dwiufs aud ejauts in the small color range of 0.22<(by)0,38 (Malvuto 199 1))."," Other previous metallicity calibrations in the Strömmgren $(b-y),m_1$ plane have been published for the regime of F and G dwarfs (Schuster Nissen \cite{schu}) ) and F dwarfs and giants in the small color range of $0.22 < (b-y) <4580.38$ (Malyuto \cite{maly}) )."459 Iu the ΠοΠΠ lavee samples of homogeneous Espectroscopic measurements of cluster as πο as field eiauts have become available.," In the meantime, large samples of homogeneous spectroscopic measurements of cluster as well as field giants have become available."460 Furthermore. CCD 5Stroónmiueren plotometiy. in particular in globular clusters. is become popular (e.g. Cauall ct al. 1999)).," Furthermore, CCD Strömmgren photometry, in particular in globular clusters, is becoming popular (e.g. Grundahl et al. \cite{grun99}) )."461 Therefore. a revised. more accurate calibration of the metallicity sensitive Strónuu- eren indices is needed.," Therefore, a revised, more accurate calibration of the metallicity sensitive Strömm- gren indices is needed."462 Iu this paper. the Strouunercu metalicity calibration is extended to more metal-poor stars (Fe/H| »2.0 dex).," In this paper, the Strömmgren metallicity calibration is extended to more metal-poor stars ([Fe/H] $>-2.0$ dex)."463 The sample of stars used for the new calibration contains red elants of w Centar. M55. axd M22. and field eiuts from a homogenized sample of Authouv-Twiuwos Twirog (1998)).," The sample of stars used for the new calibration contains red giants of $\omega$ Centauri, M55, and M22, and field giants from a homogenized sample of Anthony-Twarog Twarog \cite{anth98}) )."464 we Cen is known for its star-to-star variations in CN as well as iron abundances (e... Vaiture et al. 1991))., $\omega$ Cen is known for its star-to-star variations in CN as well as iron abundances (e.g. Vanture et al. \cite{vant}) ).465 For 10 eiauts now Cen. Norris Da Costa (1995)) measured accurate abundances from lüeh resolution spectroscopy.," For 40 giants in $\omega$ Cen, Norris Da Costa \cite{norr95}) ) measured accurate abundances from high resolution spectroscopy."466 M22 also shows CN abuance variations (e.g. Norris Freeman 1982.. 1983)).," M22 also shows CN abundance variations (e.g. Norris Freeman \cite{norr82}, , \cite{norr83}) )."467 Doth, Both468WDE 245770 is the 09.7-BOlLle (Steele et al.,HDE 245770 is the 09.7-B0IIIe (Steele et al.469 1998) optical counterpart ofthe classical Be/X-ray binary svsten 0535|26. discovered in 1975 (Coe et al.," 1998) optical counterpart of the classical Be/X-ray binary system A0535+26, discovered in 1975 (Coe et al."470 1975: Rosenhere et al., 1975; Rosenberg et al.471 1975)., 1975).472 Ehe system contains a 1045s spin-period neutron star in a 111 day orbit (LIavakzwa 1981)., The system contains a 104s spin-period neutron star in a 111 day orbit (Hayakawa 1981).473 Optical/LR spectra have revealed Ho. Bree. Paz and Pad to be in emission. throughout the observed. history of the svstem. suggesting the continuous presence of a circumstellar disc.," Optical/IR spectra have revealed $\alpha$ , $\gamma$ , $\gamma$ and $\delta$ to be in emission throughout the observed history of the system, suggesting the continuous presence of a circumstellar disc."474 However. observations reported. in this paper reveal a period. of absorption in the above LI lines. though Lla has rapidly resumed: very weak emission.," However, observations reported in this paper reveal a period of absorption in the above H lines, though $\alpha$ has rapidly resumed very weak emission."475 Combined with a major decrease in ΕΙ luminosity. the implication is that a large reduction in the circumstellar disc characteristic of this De/X-rav binary must have occurred.," Combined with a major decrease in IR luminosity, the implication is that a large reduction in the circumstellar disc characteristic of this Be/X-ray binary must have occurred."476 Further Hit/optical spectra presented. here reveal the rapid re-growth and ensuing stability of a small. dense inner disc.," Further IR/optical spectra presented here reveal the rapid re-growth and ensuing stability of a small, dense inner disc."477 A summary of all the observations in this workmay be found in Table 1., A summary of all the observations in this workmay be found in Table \ref{tab:obs}.478 Alb data have been reduced: using Starlink software. with the exception of the August and September 1998 spectra (which were reduced within HUE). SAAO It photometry and PCS photometry.," All data have been reduced using Starlink software, with the exception of the August and September 1998 spectra (which were reduced within IRAF), SAAO IR photometry and TCS photometry."479" The latter were obtained from the 1.5m ""Telescopio Carlos Sanchez at Teide observatory. Tenerife using the CVE photometer."," The latter were obtained from the 1.5m Telescopio Carlos Sanchez at Teide observatory, Tenerife using the CVF photometer."480 Data were reduced by means of the procedure described by Manfroid (1993)., Data were reduced by means of the procedure described by Manfroid (1993).481 Instrumental values were transformed to the TCS standard svstem (Alonso et al., Instrumental values were transformed to the TCS standard system (Alonso et al.482 1998)., 1998).483 SAAO LR photometry was reduced. using standard SAAQO software which is also based upon the methocs described in Manfroid (1993)., SAAO IR photometry was reduced using standard SAAO software which is also based upon the methods described in Manfroid (1993).484 Optical spectra have mostly. been obtained: via. the service observing programme of the ING. ancl all include La and Le 667824.," Optical spectra have mostly been obtained via the service observing programme of the ING, and all include $\alpha$ and He $\AA$."485 Raw spectra were de-biased. IEat-fielded. extracted and wavelength calibrated (using Cur and ArNe arc lamp spectra) with the package.," Raw spectra were de-biased, flat-fielded, extracted and wavelength calibrated (using CuAr and ArNe arc lamp spectra) with the package."486 Heliocentric velocity. corrections were applied usingVACILEL., Heliocentric velocity corrections were applied using.487 Further processing made usc ofDIPSO., Further processing made use of.488 Image reduction was undertaken using the package. whilst. subsequent aperture photometry took place withinCALA.," Image reduction was undertaken using the package, whilst subsequent aperture photometry took place within."489 Infrared spectra of the system. were obtained. on 1998 November LO using the COGS4 spectrometer of the United Ixingdom Infrared Telescope (UINIICE). Lasvaii.," Infrared spectra of the system were obtained on 1998 November 10 using the CGS4 spectrometer of the United Kingdom Infrared Telescope (UKIRT), Hawaii."490 Four grating positions were employed (0.05 .— 1.26. 1.19 1.54. 1.40 1.80. and 1.70 — 2.40 p/m). giving almost complete spectral coverage from 0.95 to 2.4 jum. A and a Ci-tvpe standards were also observed in order to allow correction of the spectra for telluric absorption.," Four grating positions were employed (0.95 – 1.26, 1.19 – 1.34, 1.40 – 1.80, and 1.70 – 2.40 $\mu$ m), giving almost complete spectral coverage from 0.95 to 2.4 $\mu$ m. A and a G-type standards were also observed in order to allow correction of the spectra for telluric absorption."491 Initial data reduction was carried out at the telescope using the software (Puxlev. Beard Ramsey 1992).," Initial data reduction was carried out at the telescope using the software (Puxley, Beard Ramsey 1992)."492 Subsequent sky subtraction. optimal extraction. division by the standard star. and wavelength calibration usingarc lampobservations were carried out. using ΑΟ.," Subsequent sky subtraction, optimal extraction, division by the standard star, and wavelength calibration usingarc lampobservations were carried out using ."493time.,time.494" Thus Let 5,=10°s5 represent the typical Lorentz [actors of leptons in the wind zone of the pulsar.", Thus Let $\gamma_w = 10^5 \gamma_5$ represent the typical Lorentz factors of leptons in the wind zone of the pulsar.495 This value should remain roughly constant when Εστω , This value should remain roughly constant when $t \lesssim t_{sd}$.496A quasi-monoenergetic proton/ion wind may also be formed. though we only consider leptonic processes here. ancl furthermore do not treat nonthermal particle acceleration al the PW/SNR shell boundary shock.," A quasi-monoenergetic proton/ion wind may also be formed, though we only consider leptonic processes here, and furthermore do not treat nonthermal particle acceleration at the PW/SNR shell boundary shock."497" When svnchrotron losses dominate. the mean Lorentz [actor of a distribution of leptons with random pitch angles evolves in response to a randomly oriented magnetic field . ∪↓∐∐↲≀↧↴∐⋟∖⊽⊔⋅≼↲∐≸≟⊔↥∐≀↧↴≺∢≺∢∪↕⋅≺∐∐≸≟↥∪⊔∐↲≼↲⇀↸↕↽≻↕⋅≼↲⋟∖⋱∖⊽↕∪↥−↶↴↴↴⊽∖∙∩∣∣∶≼⋡⊔⊳↓⋅∐−∕∕∕↻⊼∣∣∣↙∣↽⇀∣⋝↶↴↴↴−⋅↖≺↽↔↴↕∖↽↕∐≸↽↔↴ . . .»Q2 ∙− where/; is the injection time. 5;=5, is Che injection Lorentz factor. aud Writing the nonthermal electron injection fhunetion as the electron Lorentz [actor distribution «N,(5:/)/d5 canbe solved to give In this expression. ΕΣ and we introduce the dimensionless quantities Γιάμοι and 5—5/5,."," When synchrotron losses dominate, the mean Lorentz factor of a distribution of leptons with random pitch angles evolves in response to a randomly oriented magnetic field ofmean strength $B$ according to the expression $-\dot498\gamma_{syn} = (\sigma_{\rm T} B^2/ 6\pi m_e c)\gamma^2$, giving where$t_i$ is the injection time, $\gamma_i = \gamma_w$ is the injection Lorentz factor, and Writing the nonthermal electron injection function as the electron Lorentz factor distribution $dN_e(\gamma;t)/d\gamma$ canbe solved to give In this expression, $\dot{\cal N}_e \equiv \eta \xi_e L_0/(m_ec^2499\gamma_w)$ and we introduce the dimensionless quantities $\hat t = t/500T_B\gamma_w$ and $\hat \gamma = \gamma/ \gamma_w$."501 Adiabatic losses. which are significant on time scales ~//3. are negligible in comparison with svuchrotvon losses of wind electrons when /« 1l. and we restrict ourselves to this regime.," Adiabatic losses, which are significant on time scales $\sim t/3$, are negligible in comparison with synchrotron losses of wind electrons when $\hat t \ll5021$ , and we restrict ourselves to this regime."503" The pL,/ svnchrotron radiation flux pL""""ctupeapa32131isN (5,). where 5.=∨⇤∶∣⊔∣∕∕∕∕∣∣∣↙∣↽⇀⇉⋅⇤∠↗≽∶∐∕∕∕∐↿⋅∕≖⋅≀↧↴↕∐⇂⊔∐↲≺∢∏∐≺∢≀↕↴↥∐↓≀↧↴↖⊂↽↔↴∐≼↲∐≺∢∐≼↲↥≼⇂∐↿⋅∕↘∶∔⋅∔↓⋗⋖↓∪↓⋮⋡⋝⊂↽⊐⋅↴∏⋯⋟∖⊽ \/e/ep."," The $\nu L_\nu$ synchrotron radiation flux $\nu L_\nu^{syn} \simeq504{1\over 2} u_B c\sigma_{\rm T} \gamma_s^3 N_e(\gamma_s)$ , where $\gamma_s = \sqrt{\e/\e_B }$ , $\e = h\nu/m_ec^2$ , $\e_B = B/B_{cr}$ ,and the critical magneticfield $B_{cr} = 4.41\times 10^{13}$ G. Thus"505Limited. Diffusion. approximation. (speed) and a first order frequency dependent ray-tracing routine (accuracy).,Limited Diffusion approximation (speed) and a first order frequency dependent ray-tracing routine (accuracy).506 The thermal radiation by dust grains ts solved in the static diffusion limit via a FLD solver., The thermal radiation by dust grains is solved in the static diffusion limit via a FLD solver.507 The FLD approximation yields the correct solution in. the optically thin. (free-streaming) and optically thick (diffusion) limit., The FLD approximation yields the correct solution in the optically thin (free-streaming) and optically thick (diffusion) limit.508 But in the transition regions. where also the most important feedback due to stellar radiative forces occurs. the FLD approximation is not valid anymore.," But in the transition regions, where also the most important feedback due to stellar radiative forces occurs, the FLD approximation is not valid anymore."509 This first impact of stellar irradiation is therefore accurately solved in the ray-tracing routine., This first impact of stellar irradiation is therefore accurately solved in the ray-tracing routine.510 Our implementation of the FLD method uses a MPI parallelized implicit GMRES restarted solver yielding high speedup factors comparable to modern hydrodynamies solvers., Our implementation of the FLD method uses a MPI parallelized implicit GMRES restarted solver yielding high speedup factors comparable to modern hydrodynamics solvers.511 The ray-tracing routine. which calculates the distribution of stellar irradiation. benefits from the usage of spherical coordinates. in which the outgoing stellar rays are aligned with the radial coordinate axis.," The ray-tracing routine, which calculates the distribution of stellar irradiation, benefits from the usage of spherical coordinates, in which the outgoing stellar rays are aligned with the radial coordinate axis."512 Expanding this first order ray-tracing to frequency dependent irradiation is therefore convenient and results in higher accuracy of the temperature slope deeply inside the disk due to lower optical depth of the infrared part of the stellar spectrum., Expanding this first order ray-tracing to frequency dependent irradiation is therefore convenient and results in higher accuracy of the temperature slope deeply inside the disk due to lower optical depth of the infrared part of the stellar spectrum.513 A minor flaw of the proposed radiative transfer module is the neglecting of scattering., A minor flaw of the proposed radiative transfer module is the neglecting of scattering.514 But in the radiation benchmark test scattering has only a marginal influence on the resulting temperature distribution: Scattering would increase the temperature in the irradiated. parts (up to an optical depth of about unity) due to higher extinction from about in the optically thin envelope up to a maximum of in the optically thick inner rim of the disks midplane., But in the radiation benchmark test scattering has only a marginal influence on the resulting temperature distribution: Scattering would increase the temperature in the irradiated parts (up to an optical depth of about unity) due to higher extinction from about in the optically thin envelope up to a maximum of in the optically thick inner rim of the disks midplane.515 The more effectively shielded outer regions of the disk would be about cooler., The more effectively shielded outer regions of the disk would be about cooler.516 Secondly. for more massive and luminous stars the effect of scattering. will decrease due to stronger forward scattering of the most energetic UV photons. which is automatically included in our frequency dependent first order ray-tracing routine.," Secondly, for more massive and luminous stars the effect of scattering will decrease due to stronger forward scattering of the most energetic UV photons, which is automatically included in our frequency dependent first order ray-tracing routine."517 ? stated that neglecting the effect of scattering except for the first scattering of the starlight yield a good approximation., \citet{RowanRobinson:1980p3678} stated that neglecting the effect of scattering except for the first scattering of the starlight yield a good approximation.518 Another flaw is the frequency averaged instead of frequency dependent re-emission of once absorbed stellar photons. ie. the expansion of the ray-tracing routine," Another flaw is the frequency averaged instead of frequency dependent re-emission of once absorbed stellar photons, i.e. the expansion of the ray-tracing routine"519The VLT/UVES observations of the pilot study sample were carried out between 7 October and 4 December 2001.,The VLT/UVES observations of the pilot study sample were carried out between 7 October and 4 December 2001.520 We used the BLUE 437 setup. yielding a wavelength coverage of A= 3760—4980À.," We used the BLUE 437 setup, yielding a wavelength coverage of $\lambda=3760$ $4980$."521". A 2"" slit was chosen. typically yielding (seeing-limited) resolving powers of R= 20.000-25.000."," A $2''$ slit was chosen, typically yielding (seeing-limited) resolving powers of $R=20,000$ $25,000$."522 While acquiring the data. the CCD binning was set to ς| pixels. which results in oversampling in combination with the wide slit.," While acquiring the data, the CCD binning was set to $1\times 1$ pixels, which results in oversampling in combination with the wide slit."523 Therefore. we rebinned the pipeline-reduced spectra by a factor of 2.," Therefore, we rebinned the pipeline-reduced spectra by a factor of 2."524" The exposure times were set such that a minimum nominal S/N per rebinned pixel of 20 was reached. assuming a seeing of 2"", full moon. thin cirrus. and pointing.secz=2."," The exposure times were set such that a minimum nominal $S/N$ per rebinned pixel of $20$ was reached, assuming a seeing of $2''$, full moon, thin cirrus, and $\sec z =5252$."526 However. given that the overhead for telescope iimage analysis. target acquisition. CCD readout. ete..," However, given that the overhead for telescope pointing, image analysis, target acquisition, CCD readout, etc.,"527 is IO mmin. the minimum exposure time was set to mmin (except for the very bright stars HD 20 and HD 221170. which would saturate the detector in this exposure time) to avoid ineffecient use of the telescope time.," is $\sim 10$ min, the minimum exposure time was set to min (except for the very bright stars HD 20 and HD 221170, which would saturate the detector in this exposure time) to avoid ineffecient use of the telescope time."528 This resulted in S/N>50 for some of the brighter stars., This resulted in $S/N>50$ for some of the brighter stars.529 Furthermore. if the observing conditions were better than assumed in our exposure time calculations. higher S/N spectra were acquired.," Furthermore, if the observing conditions were better than assumed in our exposure time calculations, higher $S/N$ spectra were acquired."530 The observations are summarized in Table 2.., The observations are summarized in Table \ref{Tab:Observations}.531 CCD BVRI photometry was obtained at the ESO-Danish mm-telescope with DFOSC., CCD $BVRI$ photometry was obtained at the ESO-Danish m-telescope with DFOSC.532 The observers were Beers Rossi (runs in December 1998 and 2000). Holmberg (October 2002 and April 2003 runs) and Zickgraf (October 2003 run).," The observers were Beers Rossi (runs in December 1998 and 2000), Holmberg (October 2002 and April 2003 runs) and Zickgraf (October 2003 run)."533 The data was reduced using the usual IRAF photometry reduction packages., The data was reduced using the usual IRAF photometry reduction packages.534" The stars were calibrated to the BVRec system. (where the subscript ""C"" indicates Kron-Cousins) using a selection of stars from nightly observations of Landolt standard fields."," The stars were calibrated to the $BVR_CI_C$ system, (where the subscript 'C' indicates Kron-Cousins) using a selection of stars from nightly observations of Landolt standard fields."535 Typical errors in the resulting magnitudes are of the order of 0.01—0.02 mmag. 1.e.. better than the uncertainties expected to arise from reddening corrections.," Typical errors in the resulting magnitudes are of the order of $0.01$ $0.02$ mag, i.e., better than the uncertainties expected to arise from reddening corrections."536 For. further details. see Beers et al. (," For further details, see Beers et al. ("5372004. in preparation).,"2004, in preparation)."538 V magnitudes and BV colours are listed in Table 1:: additional colours can be found in Table 3 of Paper II., $V$ magnitudes and $B-V$ colours are listed in Table \ref{Tab:HERESsample}; additional colours can be found in Table 3 of Paper II.539 A guick-look analysis of the snapshot spectrum of CS 29497-004. based on the equivalent ratio of the 4129 and 4132 lines immediately revealed that this star is strongly enhanced in neutron-capture elements.," A quick-look analysis of the snapshot spectrum of CS 29497–004, based on the equivalent ratio of the 4129 and 4132 lines immediately revealed that this star is strongly enhanced in neutron-capture elements."540 We therefore carried out a first abundance analysis., We therefore carried out a first abundance analysis.541 A more comprehensive analysis. based on higher resolution and S/N. data already obtained with VLT/UVES. will be presented in a forthcoming paper (Hill et al.," A more comprehensive analysis, based on higher resolution and $S/N$ data already obtained with VLT/UVES, will be presented in a forthcoming paper (Hill et al."542 2004. in preparation).For the abundance analysis we used spherically symmetric model atmospheres from anew grid of MARCS models (?)..," 2004, in preparation).For the abundance analysis we used spherically symmetric model atmospheres from a new grid of MARCS models \citep{Gustafssonetal:2004}. ."543 Enhancement of &-elements by |0.4 ddex has been taken into account., Enhancement of $\alpha$ -elements by $+0.4$ dex has been taken into account.544 The effective temperature was derived from broad-band visual (V) and infrared (/K) photometry with the methods described in ?.., The effective temperature was derived from broad-band visual $V$ ) and infrared $JK$ ) photometry with the methods described in \citet{KeckpaperI}.545 The /Kmeasurements are from the Two, The $JK$measurements are from the Two54622002).,2002).547 The distinction between Chick ancl thin disk is relatively unimportant here as most of the mass is in the thin disk., The distinction between thick and thin disk is relatively unimportant here as most of the mass is in the thin disk.548 The choice of zy is relevant only to the detailed computation ofVays: plausible deviations are smaller than the variation in interpolation functions., The choice of $z_d$ is relevant only to the detailed computation of; plausible deviations are smaller than the variation in interpolation functions.549 I consider scale lengths in the range 2XRy<4 kpe., I consider scale lengths in the range $2 \leq \rd \leq 4$ kpc.550 At the lower limit of this range. the mmodel has a purely exponential surface density distribution.," At the lower limit of this range, the model has a purely exponential surface density distribution."551 At the upper limit of this range. the Milkv Wav disk starts to become rather low surface brightness.," At the upper limit of this range, the Milky Way disk starts to become rather low surface brightness."552" The scale length suggested by CODE L-band data is 2.5 kpe (Binney. Gerhard. Spergel 1997). while more recently Gerhard (2002. 2006) hassuggested a shorter Ay,=2.1 kpc."," The scale length suggested by COBE $L$ -band data is 2.5 kpc (Binney, Gerhard, Spergel 1997), while more recently Gerhard (2002, 2006) hassuggested a shorter $\rd = 2.1$ kpc."553 In varying the scale length. it turns out that MOND prefers a rather short scale length. consistent. with these estimates. (he precise value depending somewhat on the choice of interpolation Iunction and the choice of comparison data.," In varying the scale length, it turns out that MOND prefers a rather short scale length consistent with these estimates, the precise value depending somewhat on the choice of interpolation function and the choice of comparison data."554 The distribution of bulge mass follow the triaxial distribution determined by Binney ((1997): where 5?=4?ιτ+(z/CY., The distribution of bulge mass follow the triaxial distribution determined by Binney (1997): where $b^2 = x^2+(y/\eta)^2+(z/\zeta)^2$.555" They base this profile on fits to the CODE L-band data. finding 5b,=1.9 kpe. by=0.1 kpe. jj=0.5 and ¢=0.6."," They base this profile on fits to the COBE $L$ -band data, finding $b_m = 1.9$ kpc, $b_0 = 0.1$ kpc, $\eta = 0.5$ and $\zeta = 0.6$."556 In order to compute {for the bulge. I neglect the (triaxialitv. aud use the sphere (hat is geometrically equivalent to equation 13..," In order to compute for the bulge, I neglect the triaxiality and use the sphere that is geometrically equivalent to equation \ref{bulgedist}. ."557" This gives the same run of mass with radius. substituting r/(G)7)! υπ b/b,, ancl similarly for 5/54."," This gives the same run of mass with radius, substituting $r/[(\eta \zeta)^{1/3}b_m]$ for $b/b_m$ and similarly for $b/b_0$."558 Then which I integrate numerically., Then which I integrate numerically.559 The factor pig is scaled (o give the correct bulge mass for each choice of scale length: (Table 1)).," The factor $\rho_{B,0}$ is scaled to give the correct bulge mass for each choice of scale length (Table \ref{MW_mods}) )."560 The ellects of deviations [rom (his particular bulge model are explored in 84.1.., The effects of deviations from this particular bulge model are explored in \ref{bulgscal}.561 These details are fairly unimportant here. as (hey only affect the inner 2 kpe where the motions are non-circular owing to the triaxial distribution of the inner bulge-bar component implied by equation 13..," These details are fairly unimportant here, as they only affect the inner 3 kpc where the motions are non-circular owing to the triaxial distribution of the inner bulge-bar component implied by equation \ref{bulgedist}. ."562 All that matters to the results at/2>3 kpe is the total mass enclosed therein., All that matters to the results at$R > 3$ kpc is the total mass enclosed therein.563The discovery that supermassive black holes are ubiquitous among massive galaxies in the local Universe indicates that the majority of galaxies have passed through an active phase during their evolutionary history.,The discovery that supermassive black holes are ubiquitous among massive galaxies in the local Universe indicates that the majority of galaxies have passed through an active phase during their evolutionary history.564 Moreover. the strong correlation observed between black-hole anc bulge mass (Ixormendy Richstone 1995: Alagorrian ct al.," Moreover, the strong correlation observed between black-hole and bulge mass (Kormendy Richstone 1995; Magorrian et al."565 1998: Gebharelt et al., 1998; Gebhardt et al.566 2000a: Ferrarese Merritt. 2000) indicates that the evolution of the central black hole and its host galaxy are intimately related., 2000a; Ferrarese Merritt 2000) indicates that the evolution of the central black hole and its host galaxy are intimately related.567 Consequently. it is elear that studying the evolution of quasar black-hole masses will provide crucia information concerning the evolution of both quasars anc massive earlv-tvpe galaxies.," Consequently, it is clear that studying the evolution of quasar black-hole masses will provide crucial information concerning the evolution of both quasars and massive early-type galaxies."568 Within this context. the last) few vears have seen renewed interest in the possibilities of estimating the centra black-hole masses of active galactic nuclei (AGN)., Within this context the last few years have seen renewed interest in the possibilities of estimating the central black-hole masses of active galactic nuclei (AGN).569 The major impetus for this has been the results of the recen reverberation mapping programmes carried. out on low-redshift quasars and Sevfert galaxies (Wandel. Peterson Alalkan 1999: Ixaspi et al.," The major impetus for this has been the results of the recent reverberation mapping programmes carried out on low-redshift quasars and Seyfert galaxies (Wandel, Peterson Malkan 1999; Kaspi et al."570 2000)., 2000).571 The measurements of the broad-line. region. (BLIU) radius produced: by these long-term monitoring programmes have allowed: so-called virial black-hole mass estimates to be made for. 34 low-redshift ACGN (Ixaspi et al., The measurements of the broad-line region (BLR) radius produced by these long-term monitoring programmes have allowed so-called virial black-hole mass estimates to be made for 34 low-redshift AGN (Kaspi et al.572 2000)., 2000).573 The principal assumption underlving the virial mass estimate is simply that the dvnanmies of the BLE. are dominated by the gravity of the central supermassive black hole., The principal assumption underlying the virial mass estimate is simply that the dynamics of the BLR are dominated by the gravity of the central supermassive black hole.574" Under this assumption an estimate of the central black-hole mass can be gained fron: Alan~ές where Retry is the radius of the DLlt and V, is the '""Reravy:velocity. of the [ine-emitting gas. as traclitionally estimated from the FWLIAT of the 7:3 emission line."," Under this assumption an estimate of the central black-hole mass can be gained from: $M_{bh}\simeq G^{-1}R_{BLR}V_{g}^{2}$; where $R_{BLR}$ is the radius of the BLR and $V_{g}$ is the velocity of the line-emitting gas, as traditionally estimated from the FWHM of the $H\beta$ emission line."575 Given the large number of physical processes. which could. potentially influence the dynamics of the BLE it is not immediately obvious that the gravitational potential of the central black-hole should be dominant (eg., Given the large number of physical processes which could potentially influence the dynamics of the BLR it is not immediately obvious that the gravitational potential of the central black-hole should be dominant (eg.576 Ixrolik 2001)., Krolik 2001).577 llowever. several lines of evidence have recently shown that the dynamics of the BLR appear to be at least. consistent with the virial assumption.," However, several lines of evidence have recently shown that the dynamics of the BLR appear to be at least consistent with the virial assumption."578 Firstly. for the small number of objects for which it is possible to do so. Peterson Wandel (2000) have shown that the motions of the broad-line gas are consistent with being virialized. with the," Firstly, for the small number of objects for which it is possible to do so, Peterson Wandel (2000) have shown that the motions of the broad-line gas are consistent with being virialized, with the"579taken care to illustrate them extensively in this paper.,taken care to illustrate them extensively in this paper.580 It is probable that there are many more. aspects still to. be uncovered., It is probable that there are many more aspects still to be uncovered.581 The dillerences in behaviour for outbursts of the same dwarf nova have vet to be understood., The differences in behaviour for outbursts of the same dwarf nova have yet to be understood.582 Why. during the December 1975. superouthurst of VW. Livin the  δδ s pDNOs were so prominent for a week but onlv very occasionally and intermittenthy present in other outbursts. is amystery.," Why, during the December 1975 superoutburst of VW Hyi, the $\sim$ 88 s lpDNOs were so prominent for a week but only very occasionally and intermittently present in other outbursts, is a mystery."583 We note. however. that Hartmann et al. (," We note, however, that Hartmann et al. ("5841999) rave found that the X-Ray properties of VW. Evi also diller &reatlv from outburst to outburst.,1999) have found that the X-Ray properties of VW Hyi also differ greatly from outburst to outburst.585 Both the N-Havs and he DNOs are associated with the inner disc and boundary aver. and both will be allected by the structure and strength of magnetic field connections from the disc to the primary and its equatorial belt. so this is where the dillerences are oobablv generated.," Both the X-Rays and the DNOs are associated with the inner disc and boundary layer, and both will be affected by the structure and strength of magnetic field connections from the disc to the primary and its equatorial belt, so this is where the differences are probably generated."586 Inter alia. we have shown that the helium-transferring systems in some cases have DNO and. QPO properties similar to those of the hwdrogen-rich CVs.," Inter alia, we have shown that the helium-transferring systems in some cases have DNO and QPO properties similar to those of the hydrogen-rich CVs."587 This is first evidence. even if indirect. that the helium primaries of the AM CVn stars can be magnetic.," This is first evidence, even if indirect, that the helium primaries of the AM CVn stars can be magnetic."588 An advantage that CV observations have over those of X-Ray binaries is that in CVs we often easily see the trains of DNO and QPO moculations. and can study the properties of individual cvcles; whereas in the X-Ray binaries the ow fluxes require statistical analvses to extract average ooperties.," An advantage that CV observations have over those of X-Ray binaries is that in CVs we often easily see the trains of DNO and QPO modulations, and can study the properties of individual cycles, whereas in the X-Ray binaries the low fluxes require statistical analyses to extract average properties."589 If the phenomena really do have the same origin. hen further observations of CVs even with small aux niodest size telescopes are likely to assist understanding of he higher οποίον analogues.," If the phenomena really do have the same origin, then further observations of CVs even with small and modest size telescopes are likely to assist understanding of the higher energy analogues."590 Ato present the interpretation of the CV. rapi oscillations in terms of magnetic accretion is speculative but we note tha behaviours are so complex tha no simple explanations (c.g. white dwarl or acceretion disc pulsations) are likely to be adequate., At present the interpretation of the CV rapid oscillations in terms of magnetic accretion is speculative – but we note that the behaviours are so complex that no simple explanations (e.g. white dwarf or accretion disc pulsations) are likely to be adequate.591 The. existence. of magnetically controlled accretion is in principle testable by the detection of polarisation modulated at the DNO perice this might be accomplished for VW. Livi during the larec amplitude DNOs a rw end. Of outburst. which occur a roughly monthly intervals. but will probably require the use ofa Large telescope observing in the near infrared.," The existence of magnetically controlled accretion is in principle testable by the detection of polarisation modulated at the DNO period – this might be accomplished for VW Hyi during the large amplitude DNOs at the end of outburst, which occur at roughly monthly intervals, but will probably require the use of a large telescope observing in the near infrared."592 Currently the situation is similar to that of the early developmen of the intermediate polar model (Patterson Price 1981: Warner. O'Donoghue Fairall 1981). which for a lew vears gave satisfactory explanations of the observed multiple photometric moclulations but the magnetic nature was only verified by the discovery of circularly polarisect E band [ux in 1986 (Penning. Schmidt Lichert 1986).," Currently the situation is similar to that of the early development of the intermediate polar model (Patterson Price 1981; Warner, O'Donoghue Fairall 1981), which for a few years gave satisfactory explanations of the observed multiple photometric modulations but the magnetic nature was only verified by the discovery of circularly polarised I band flux in 1986 (Penning, Schmidt Liebert 1986)."593" Van der Elis (2000). in connection with the two-QPO correlation in X-Ray binaries and the concomitant variations of the frequencies. remarked that they ""essentially imply that the phenomena are generated in the accretion disc around. any low-magnetic field compact object”."," Van der Klis (2000), in connection with the two-QPO correlation in X-Ray binaries and the concomitant variations of the frequencies, remarked that they “essentially imply that the phenomena are generated in the accretion disc around any low-magnetic field compact object”."594 Our observations of similar phenomena in (presumably) low-field CVs appear to make this conclusion even more widely applicable., Our observations of similar phenomena in (presumably) low-field CVs appear to make this conclusion even more widely applicable.595 The University of Cape Town funds DW's research: PAW is supported by strategic funds made available to BW hy the University and by research funds [rom the National Research Foundation., The University of Cape Town funds BW's research; PAW is supported by strategic funds made available to BW by the University and by research funds from the National Research Foundation.596 MLP acknowledges financial support from the Department of Labour's National Skills Fund., MLP acknowledges financial support from the Department of Labour's National Skills Fund.597 We thank Dave Iilkenny for alerting us to the potential interest. of EC 2117-54., We thank Dave Kilkenny for alerting us to the potential interest of EC 2117-54.598 T. Belloni kindly sent us the data on X-Ray binaries that appear in Fig.e 30.., T. Belloni kindly sent us the data on X-Ray binaries that appear in Fig. \ref{twoqpo}. .599rate at zo2 which subsequently declines while intermediate mass systems (LOMcAR< 10777). have peak star ormation rates at z~ 1.5 whieh then decline more slowly.,rate at $\sim 2$ which subsequently declines while intermediate mass systems $10^{10.2}<M<10^{10.8}$ ) have peak star formation rates at $\sim$ 1.5 which then decline more slowly.600 Lt should also be noted that. if this ellect is real. then galaxy ormation for the progenitors of SALGs must be sullicientLy rapid to produce massive galaxies at redshifts  2.5.," It should also be noted that, if this effect is real, then galaxy formation for the progenitors of SMGs must be sufficiently rapid to produce massive galaxies at redshifts $\sim$ 2.5."601 Our putative detection of downsizing in the SM »opulation is complicated by the relatively bright Dux limits of the SWIRE TRAC bands. for our SILADISS. sources., Our putative detection of downsizing in the SMG population is complicated by the relatively bright flux limits of the SWIRE IRAC bands for our SHADES sources.602 Llowever. the GOODS-N SMCs discussed by Pope et al. (," However, the GOODS-N SMGs discussed by Pope et al. ("6032006). for which much deeper Ηλ data is available. ollow similar downsizing-like behaviour (see Fig.,"2006), for which much deeper IRAC data is available, follow similar downsizing-like behaviour (see Fig."604 4). while a separate and more sophisticated: analysis of the Lockman SILADIS sources (Dye et al.," 4), while a separate and more sophisticated analysis of the Lockman SHADES sources (Dye et al.,"605 2007) also shows the elfect., 2007) also shows the effect.606 Our examination of ambiguously identified. sources has revealed a number of cases where several objects appear to »e physically associated with the SCUBA source., Our examination of ambiguously identified sources has revealed a number of cases where several objects appear to be physically associated with the SCUBA source.607 Whether lis is as interacting pairs/groups or as close associations in some larger structure such as a cluster is uncertain., Whether this is as interacting pairs/groups or as close associations in some larger structure such as a cluster is uncertain.608 The clearest example of this is SILADIZS-SNDE 47. where iee radio identified sources are candidate identifications. all lving at the same estimated redshift z= 1.3.," The clearest example of this is SHADES-SXDF 47, where three radio identified sources are candidate identifications, all lying at the same estimated redshift z = 1.3."609" SILADES-SXDE 27 has two non-radio LDed objects at z,5,,—2.01 ji appear to be associated with the far-LR source. while SHADES-SADE 28 has two radio identifications both at Zpborml "," SHADES-SXDF 27 has two non-radio IDed objects at $_{phot}$ =2.01 that appear to be associated with the far-IR source, while SHADES-SXDF 28 has two radio identifications both at $_{phot}$ =1.88."610SUADES has two radio counterparts. but one of these 88.has no opticalIA or SWIDBI detection.," SHADES14 has two radio counterparts, but one of these has no optical or SWIRE detection."611 Given the redshift estimated for the other racdiof/SWIRLE source in this system. z—2.24. it is quite possible that the companion to this source lies at the same recshift but is just too faint for us to see in SWIRE or SXDE.," Given the redshift estimated for the other radio/SWIRE source in this system, z=2.24, it is quite possible that the companion to this source lies at the same redshift but is just too faint for us to see in SWIRE or SXDF."612 Pairs of radio sources or ealaxies on a range of scales have been found for à number of SMCs (eg., Pairs of radio sources or galaxies on a range of scales have been found for a number of SMGs (eg.613 Blain et al..," Blain et al.,"614 2004: Swinbank et al..," 2004; Swinbank et al.,"615 2006)., 2006).616 Phe physical sepearations between these sources range in size [rom 8S to 135 kpc. and thus cover the distance scales expected for both interacting pairs of galaxies and galaxy clusters.," The physical sepearations between these sources range in size from 8 to 135 kpc, and thus cover the distance scales expected for both interacting pairs of galaxies and galaxy clusters."617 Lt is thus seers likely that some SLLADES sources are to be found in either interacting groups or in candidate high redshift clusters., It is thus seems likely that some SHADES sources are to be found in either interacting groups or in candidate high redshift clusters.618 Is this true for the rest of the population?, Is this true for the rest of the population?619 For the bulk of the SILADIZS sources. there are no clearly selected companions. on the basis of radio or μαι emission.," For the bulk of the SHADES sources, there are no clearly selected companions, on the basis of radio or $\mu$ m emission."620 We thus investigate the possibility of associations by examining the photometric redshifts of all potential companions within a 107 radius of the sources.," We thus investigate the possibility of associations by examining the photometric redshifts of all potential companions within a 10"" radius of the sources."621 This radius is chosen on the basis of correlation. analysis (Serjeant et. al..," This radius is chosen on the basis of correlation analysis (Serjeant et al.,"622. 2007) which suggests the presence of correlations on these angular scales., 2007) which suggests the presence of correlations on these angular scales.623 We calculate the dillerence in. estimated redshift. Az between cach SLLADES source and. all its companions and compare this to twenty reference fields olfset (rom the SLLADES sources by up to 0.1. degrees.," We calculate the difference in estimated redshift, $\Delta z$ between each SHADES source and all its companions and compare this to twenty reference fields offset from the SHADES sources by up to $\sim$ 0.1 degrees."624 I£ there is a real correlation in redshift between the bulk of the SLLADLES sources and their nearby companions then there will be an enhancement of this statistic at zero redshift dillerence for the SILADES sources that is not apparent for the reference fields., If there is a real correlation in redshift between the bulk of the SHADES sources and their nearby companions then there will be an enhancement of this statistic at zero redshift difference for the SHADES sources that is not apparent for the reference fields.625 Figure 4 shows this comparison., Figure 4 shows this comparison.626 As can be seen. we find no enhancement in the number of objects with As—0 , As can be seen we find no enhancement in the number of objects with $\Delta z = 0$ 627The uature of dark matter is one of the major outstanding problems in curent astrophysics. with the conumuunitv being divided between two woad caups: those who propose particle dark natteor axd those who support massive compact objects.,"The nature of dark matter is one of the major outstanding problems in current astrophysics, with the community being divided between two broad camps: those who propose particle dark matter and those who support massive compact objects."628 Both these ideas have received receut inupetus due to the measurement of a neutrine nass (Fuaudaοal.1998). aud the detection of \IACTIOs in the halo of our Galaxy (Alcocketal. 1993).. although the siuall mass of the neutriuo aud the uicertaintv in the location of the ALACTIO lenses suggests that the nature of majority of the darκ πιάτο still remains hidden frou us.," Both these ideas have received recent impetus due to the measurement of a neutrino mass \citep{fu99}629 and the detection of MACHOs in the halo of our Galaxy \citep{al93}, although the small mass of the neutrino and the uncertainty in the location of the MACHO lenses suggests that the nature of majority of the dark matter still remains hidden from us."630 One recent suggestion proposes hat the variability seen in high redshift quasars. rather than beime intrinsic to the source. ix due to tje microleusimg action of Jupiter mass black holes distributed cosinologically along the line-ofsigit to the quasar (Πας1993.1996).," One recent suggestion proposes that the variability seen in high redshift quasars, rather than being intrinsic to the source, is due to the microlensing action of Jupiter mass black holes distributed cosmologically along the line-of-sight to the quasar \citep{ha93,ha96}."631. To account or the characteristics of the variability. a substantial population of black. holes is required. with a density high erough to account for a significant fraction of the total ass in the universe (Schucider1993).," To account for the characteristics of the variability, a substantial population of black holes is required, with a density high enough to account for a significant fraction of the total mass in the universe \citep{sc93}."632. While there is theoretical support for such a model. wit1 1nodels of the quark-hadron transition produci18o blackholes with a dass of ~107M. lee. Tawlius&Tavlor(1997) |]. observations have züled to reveal their presence imn sudies of the ecpuvaleut widths of distant quasars (Canizares1982:Dalcantonetal. 1991).. although this conchsion has been contested (Wawkins1996).," While there is theoretical support for such a model, with models of the quark-hadron transition producing blackholes with a mass of $\sim10^{-3}{\rm \msun}$ [e.g. \citet{ha97}] ], observations have failed to reveal their presence in studies of the equivalent widths of distant quasars \citep{ca82,da91}, although this conclusion has been contested \citep{ha96}."633. Other oljjections to this microlensing hypothesis have heen raised (Dagaunoff&Alalkau1995:Alexander1995).. with correspouding rebutals (Ilawldus&Tay-lor 1997)..," Other objections to this microlensing hypothesis have been raised \citep{ba95,al95}, with corresponding rebuttals \citep{ha97}. ."634 More recenth. arguments have become quite vitriolic with the sugeestion that this idea," More recently, arguments have become quite vitriolic with the suggestion that this idea"635Comparison of the sspectral maps with the mmaps of OSOS reveals some differences. particularly in the abundance distribution.,"Comparison of the spectral maps with the maps of OS05 reveals some differences, particularly in the abundance distribution."636 The aabundance map contains a ridge of high metallicity extending from the galaxy core northwest along the northern edge of the western jet., The abundance map contains a ridge of high metallicity extending from the galaxy core northwest along the northern edge of the western jet.637 Testing with normal spectral fits to the data confirmed that this region had enhanced abundances. ruling out a problem specitic to the mapping technique.," Testing with normal spectral fits to the data confirmed that this region had enhanced abundances, ruling out a problem specific to the mapping technique."638 OSOS suggested —iat the feature might arise from stripped material or a shock associated with the jet., OS05 suggested that the feature might arise from stripped material or a shock associated with the jet.639 The data show no such ridge. though high abundances are seen in the galaxy core.," The data show no such ridge, though high abundances are seen in the galaxy core."640" To confirm the lack of such a feature in the data we extracted a spectrum in an annular segment defined from —je aabundance map. extending northwest. from —45-120"" ποrom the AGN. equivalent to region | of Figure 9 and Table 5 of OSOS."," To confirm the lack of such a feature in the data we extracted a spectrum in an annular segment defined from the abundance map, extending northwest, from $\sim$ from the AGN, equivalent to region 1 of Figure 9 and Table 5 of OS05."641 This spectrum. and the sspectra extracted from region | by OSOS using SAS 6.0. were then fitted with an absorbed APEC model using a range of energy bands.," This spectrum, and the spectra extracted from region 1 by OS05 using SAS 6.0, were then fitted with an absorbed APEC model using a range of energy bands."642 We find that |) although AACIS and EEPIC-pn and MOS cameras give consistent abundances within the errors. the EPIC-pn spectra produce the highest best-fit abundances and ACIS the lowest. and 2) the abundance measured is affected by the energy band used. with wider energy bands producing higher abundances in this region.," We find that 1) although ACIS and EPIC-pn and MOS cameras give consistent abundances within the errors, the EPIC-pn spectra produce the highest best-fit abundances and ACIS the lowest, and 2) the abundance measured is affected by the energy band used, with wider energy bands producing higher abundances in this region."643 Using our 0.7-7.0 keV energy band. EPIC-MOS produces best-fit abundances very similar to ACTS Wee compared to uit .9) while the EPIC-pn abundance is significantly greater VL," Using our 0.7-7.0 keV energy band, EPIC-MOS produces best-fit abundances very similar to ACIS $^{+0.39}_{-0.28}$ compared to $^{+0.17}_{-0.15}$ ) while the EPIC-pn abundance is significantly greater $^{+0.28}_{-0.24}$."644 This suggests tha mmap feature is either the product of inaccurate calibration of the EPIC-pn. or that it is caused by the presence of some spectrally hard emission component which affects the EPIC-pn more than ACIS or EPIC-MOS because of its greater effective area at high energies.," This suggests that map feature is either the product of inaccurate calibration of the EPIC-pn, or that it is caused by the presence of some spectrally hard emission component which affects the EPIC-pn more than ACIS or EPIC-MOS because of its greater effective area at high energies."645 As discussed in OSOS. the high abundance ridge does no correspond to detector structure on either EPIC-MOS or -pn.," As discussed in OS05, the high abundance ridge does not correspond to detector structure on either EPIC-MOS or -pn."646 The ridge is located on the focal-point CCD of EPIC-pn. at least nnorthwest of the CCD edge at the centre of the array.," The ridge is located on the focal–point CCD of EPIC-pn, at least northwest of the CCD edge at the centre of the array."647 There are no structures in the background data files at the position of the ridge., There are no structures in the background data files at the position of the ridge.648 It therefore seems unlikely that a calibration problem associated with a particular CCD feature could be the cause of the high abundances., It therefore seems unlikely that a calibration problem associated with a particular CCD feature could be the cause of the high abundances.649 Examining hard band images (5-7. keV) some clumps of emission are visible. but these are not consistent between instruments and involve only a handful of counts.," Examining hard band images (5-7 keV) some clumps of emission are visible, but these are not consistent between instruments and involve only a handful of counts."650 Unfortunately neither dataset is sufficiently deep to determine whether an excess is present., Unfortunately neither dataset is sufficiently deep to determine whether an excess is present.651 Fitting the MOS and pn spectra described above using an apect+powerlaw model. we tind neither can constrain the powerlaw index. and that for various fixed index values the MOS abundance is still consistently lower than that measured for the EPIC-pn.," Fitting the MOS and pn spectra described above using an apec+powerlaw model, we find neither can constrain the powerlaw index, and that for various fixed index values the MOS abundance is still consistently lower than that measured for the EPIC-pn."652 This argues against an additional spectral component as the cause of the ridge. and leaves the origin of the feature uncertain.," This argues against an additional spectral component as the cause of the ridge, and leaves the origin of the feature uncertain."653" To check whether any hard emission affeets our mmaps we recreated the maps using 3200 count spectra. and ignoring energies above 3.0 keV. We find that the map structures do not change. suggesting that our results are not affected by extraneous hard emission,"," To check whether any hard emission affects our maps we recreated the maps using 3200 count spectra, and ignoring energies above 3.0 keV. We find that the map structures do not change, suggesting that our results are not affected by extraneous hard emission."654 While we have focussed on the interaction between the AGN and ICM. there are other mechanisms which could affect the abundance and temperature structure of AWM 4.," While we have focussed on the interaction between the AGN and ICM, there are other mechanisms which could affect the abundance and temperature structure of AWM 4."655 Gas motions associated with mergers and interactions are one possibility., Gas motions associated with mergers and interactions are one possibility.656 The overall structure of the temperature map. with higher temperatures to the south of the radio jet than to its north at moderate radii. with the pattern reversed at large radii. could in principle be an indication of gas sloshing.," The overall structure of the temperature map, with higher temperatures to the south of the radio jet than to its north at moderate radii, with the pattern reversed at large radii, could in principle be an indication of gas sloshing."657 The infall of a subcluster can disturb the gravitationa potential of the cluster. causing the cluster core to oscillate around its centre. sloshing the gas of the core back and forth.," The infall of a subcluster can disturb the gravitational potential of the cluster, causing the cluster core to oscillate around its centre, sloshing the gas of the core back and forth."658 Examples of gas sloshing have been documented in a number of clusters (e.g..?22).. and confirm the results of numerical simulations (e.g.2). in that the sloshing is observed to produce a spiral pattern of cold fronts. where cool gas from the cluster core is moved outward and brought into contact with hotter material.," Examples of gas sloshing have been documented in a number of clusters \citep[e.g.,][]{Markevitchetal01,MazzottaGiacintucci08,Simionescuetal10}, and confirm the results of numerical simulations \citep[e.g.,][]{AscasibarMarkevitch06}, in that the sloshing is observed to produce a spiral pattern of cold fronts, where cool gas from the cluster core is moved outward and brought into contact with hotter material."659 The staggered pattern of temperature reversals seen to north and south of the radio source could be produced by such spiral sloshing., The staggered pattern of temperature reversals seen to north and south of the radio source could be produced by such spiral sloshing.660 The apparent southward motion of NGC 6051 could also be explained if the cluster core is in motion relative to its surroundings., The apparent southward motion of NGC 6051 could also be explained if the cluster core is in motion relative to its surroundings.661 Sloshing provides a mechanism or transporting metals out of the cores of clusters (2).. and woulc herefore be relevant to our study of AWM 4.," Sloshing provides a mechanism for transporting metals out of the cores of clusters \citep{Simionescuetal10}, and would therefore be relevant to our study of AWM 4."662 However. there are several strong arguments against sloshing being important in this cluster: We therefore conclude that while NGC 6051 may be in motion. it is unlikely that large-scale gas sloshing is occurring in AWM 4.," However, there are several strong arguments against sloshing being important in this cluster: We therefore conclude that while NGC 6051 may be in motion, it is unlikely that large–scale gas sloshing is occurring in AWM 4."663" Under the assumption that iron is the dominant element determining the measured metallicities. we can estimate the mass of iron required to produce the raised abundance seen along the jets. where p is the density of the gas in the jets. V is the volume of the region. AZ is the excess abundance above the background enrichment level and i, is the solar iron mass fraction (2810.P fortheabundancesof?)."," Under the assumption that iron is the dominant element determining the measured metallicities, we can estimate the mass of iron required to produce the raised abundance seen along the jets, where $\rho$ is the density of the gas in the jets, V is the volume of the region, $\Delta Z$ is the excess abundance above the background enrichment level and $\gamma_{Fe}$ is the solar iron mass fraction \citep[1.28$\times10^{-3}$ for the abundances."664 For this purpose. we define the volume of high enrichment along the jets using the smallest rectangular regions in Fig.," For this purpose, we define the volume of high enrichment along the jets using the smallest rectangular regions in Fig."665 3. and assuming a rotational symmetry, \ref{fig:mapzoom} and assuming a rotational symmetry666"One-halo pairs of galaxies are assigned by satislving the eriterià Xr<ring, and Acf/(lci)X€5442","One-halo pairs of galaxies are assigned by satisfying the criteria $\Delta r_{\perp} \leq r_{\perp,max}$ and $\Delta z/(1+z) \leq \beta_{max}$."667 Galaxies are then grouped by a FoF algorithm., Galaxies are then grouped by a FoF algorithm.668 The challenges in this section are {ο sell-consistentlv incorporate the photometric sample into the calculation of the FoF group multipliity funcüon. aud to marginalize over (inv Dright Star Mask holes in the survey The group finding algorithm first assigns canclicate pairs for each galaxy.," The challenges in this section are to self-consistently incorporate the photometric sample into the calculation of the FoF group multiplicity function, and to marginalize over tiny Bright Star Mask holes in the survey The group finding algorithm first assigns candidate pairs for each galaxy."669 As described in 2.2.1.. the redshift space criterion lor aspectvoscopic-photometric pair becomes €rin(zac)SCnasope) ad =23.2<Mopho—21.2.," As described in \ref{photosamplemodeling}, the redshift space criterion for a spectroscopic-photometric pair becomes $c_{\parallel,min}(z_{spec}) \leq c_{\parallel,photo} \leq c_{\parallel,max}(z_{spec})$ and $-23.2 < M_{g,photo} < -21.2$."670 Using these pair assignments. candidate eroups are formed with the FoF algorithm.," Using these pair assignments, candidate groups are formed with the FoF algorithm."671 Spectroscopic neighbors of photometric objects in the group are kept if they satisfy the Ar) requirement with al least one spectroscopic object in the group. aud photometric neighbors of photometric objects in the eroup are kept if (hev satisly the redshift space criterion with at least one spectroscopic object in the group.," Spectroscopic neighbors of photometric objects in the group are kept if they satisfy the $\Delta r_{\parallel}$ requirement with at least one spectroscopic object in the group, and photometric neighbors of photometric objects in the group are kept if they satisfy the redshift space criterion with at least one spectroscopic object in the group."672 We correct the group assignments for the 23 photometric objects in two distinct eroups., We correct the group assignments for the 23 photometric objects in two distinct groups.673 We first give prelerence to groups in which the photometric object is a liber-collision pair: otherwise. we assien the photometric object to the larger group.," We first give preference to groups in which the photometric object is a fiber-collision pair; otherwise, we assign the photometric object to the larger group."674 We must then reconipute {he membership of the group losing the photometric The last step of the eroup-fincing algorithm accounts lor Dright Star Mask and counts the number of isolated photometric objects. Nppore.iso used 1n Eqn. 5..," We must then recompute the membership of the group losing the photometric The last step of the group-finding algorithm accounts for Bright Star Mask and counts the number of isolated photometric objects, $N_{photo,iso}$ used in Eqn. \ref{isolatedcorrection}."675 The Bright Star Mask catalog. consisting of thera.dec. and radius for each circular masked region. was generatec from the Tveho2 catalog as described by the NYU VAGC (?)..," The Bright Star Mask catalog, consisting of the, and radius for each circular masked region, was generated from the Tycho2 catalog as described by the NYU VAGC \citep{blanton/etal:2005}."676 The Bright Star Mask covers of the total survey region., The Bright Star Mask covers of the total survey region.677" Using the observed pair counts as a function of radius. we estimate the probability per unit area of finding a 1-halo pair satisbving ro<Ποο=0.8 Mpe/h and Az/(1+2)<S44, as a function of r_."," Using the observed pair counts as a function of radius, we estimate the probability per unit area of finding a 1-halo pair satisfying $r_{\perp} \leq r_{\perp,max} = 0.8$ $h$ and $\Delta z/(1+z) \leq \beta_{max}$ as a function of $r_{\perp}$."678 The result is well-fit with a power law: dP/dA=0.027rL (Mpc/h)7., The result is well-fit with a power law: $dP/dA = 0.027 r_{\perp}^{-1.08}$ $h)^{-2}$.679 dP/dA is similar to the projected correlation. function. except that it does not subtract the uncorrelated contribution. and it only includes pairs meeting our Ar) criterion.," $dP/dA$ is similar to the projected correlation function, except that it does not subtract the uncorrelated contribution, and it only includes pairs meeting our $\Delta r_{\parallel}$ criterion."680 For each bright star mask within rye. of an LRG. we evaluate the probability that an unobserved LRG resides uider the Dright Star Mask by integrating our power law [it to dP?/dA over the area of overlap between the Dright Star Mask and the LAG annulus satisfving r.πομε.," For each bright star mask within $r_{\perp,max}$ of an LRG, we evaluate the probability that an unobserved LRG resides under the Bright Star Mask by integrating our power law fit to $dP/dA$ over the area of overlap between the Bright Star Mask and the LRG annulus satisfying $r_{\perp} \leq r_{\perp,max}$."681 We assien (he probability of a given Bright Star Mask to cover aai LRG as the maximum probability: assigned [rom all the observed LRG neighbors of the Brieht Star Mask elements., We assign the probability of a given Bright Star Mask to cover an LRG as the maximum probability assigned from all the observed LRG neighbors of the Bright Star Mask elements.682 We mocdifv the final group multiplicity function to account for each possible unobserved LRG. including its ability to. bridge (wo previously existing eroups.," We modify the final group multiplicity function to account for each possible unobserved LRG, including its ability to bridge two previously existing groups."683 Therefore a single group of LRGs can contribute to many group mulliplicity bins, Therefore a single group of LRGs can contribute to many group multiplicity bins684Studies of Globular Cluster (GC) systems have provided uscful:. constraints. on. the formation ane:considered GCS of their host galaxies (for recent reviews see Llarris B Forbes 2002).,Studies of Globular Cluster (GC) systems have provided increasingly useful constraints on the formation and evolution of their host galaxies (for recent reviews see Harris 2001; Forbes 2002).685 In the case of cD and central cluster In a variety of processes may have contributed to central rich. GC systems., In the case of cD and central cluster galaxies a variety of processes may have contributed to their rich GC systems.686 This variety. has meant that the opposed of the rich CGC systems in these galaxies is still advocated understood., This variety has meant that the origin of the rich GC systems in these galaxies is still poorly understood.687 Ideas to explain the rich GC systems of such galaxies have included cooling Lows (Fabian. Nulsen ancl 1984). biased GC formation (West et al.," Ideas to explain the rich GC systems in such galaxies have included cooling flows (Fabian, Nulsen Canizares 1984), biased GC formation (West et al."688 1995). relative (c.g. Ashman & 1992: Bekki et al.," 1995), mergers (e.g. Ashman Zepf 1992; Bekki et al."689 2002). tidal Sx) and accretion (e.g. Muzzio et al.," 2002), tidal stripping and accretion (e.g. Muzzio et al."690 1984: Forbes et mor, 1984; Forbes et al.691e 1997: Cotte et al., 1997; Côtte et al.692 1998). Zn (e.g. Harris clusters: Forbes. Brodie & LOOT: LLarris. Harris & —," 1998), formation (e.g. Harris 1991; Forbes, Brodie Grillmair 1997; Harris, Harris McLaughlin 1998)."693" Alternatively the — a may. be ""underluminous! for their GC richness stripped (Blakeslee. ‘Tonry Metzger. L997: McLaughlin. 1999: Joaslev et al."," Alternatively the galaxies themselves may be `underluminous' for their GC richness (Blakeslee, Tonry Metzger 1997; McLaughlin 1999; Beasley et al."694 2002)., 2002).695 Lt is perhaps the latter three ideas (Le. accretion. Formation and missing light) that are the⊀∙ current best.∙∙ ∙bets! toqu explain the[ excess of⋅ increasi," It is perhaps the latter three ideas (i.e. accretion, formation and missing light) that are considered the current `best bets' to explain the excess of GCs."696nely evolution2001: this paper we further explore the scenario of a galaxies cluster galaxy accreting GCs via tidal stripping (as ther to the accretion of dwarf galaxies and their GC's as origin. by Cotte et al., In this paper we further explore the scenario of a central cluster galaxy accreting GCs via tidal stripping (as opposed to the accretion of dwarf galaxies and their GCs as advocated by Côtte et al.697 1998)., 1998).698 In recent vears. a number poorly new observational facts have lent support to the stripping in acerction scenario.," In recent years, a number of new observational facts have lent support to the stripping and accretion scenario."699 These include: the correlation of & richness of the GC system (called specific frequency mergers with velocity dispersion (c.g. Blakeslee et al., These include: the correlation of relative richness of the GC system (called specific frequency $_N$ ) with velocity dispersion (e.g. Blakeslee et al.700 LOOT) - stripping stripping should occur in the higher velocity dispersion al. that high Sx galaxies have relatively more metal- 1991:GCs (Forbes et al., 1997) - more stripping should occur in the higher velocity dispersion clusters; that high $_N$ galaxies have relatively more metal-poor GCs (Forbes et al.701 1997) - as metal-poor GC's have AlcLaushlin shallower density profile they are more. likely. to be themselves from a donor galaxy: that Sy values may correlate with projected clustereentric distance (Forbes et al., 1997) - as metal-poor GCs have a shallower density profile they are more likely to be stripped from a donor galaxy; that $_N$ values may correlate with projected clustercentric distance (Forbes et al.702 1997), 1997)703and an outer racially biased envelope (BertinandStiavelli1984: see Bertin1993. and references therein): these families turned out to exhibit the characteristic 2/4 projected density profile aud indeed were shown to match nicely the observed photometric and kinematical characteristics of bright ellipticals.,"and an outer radially biased envelope \citealt{ber84}; see \citealt{ber93}704 and references therein): these families turned out to exhibit the characteristic $R^{1/4}$ projected density profile and indeed were shown to match nicely the observed photometric and kinematical characteristics of bright ellipticals."705" In an attempt at providing a justification of these models (in. particular. of the so-called f, models. constructed initially only from ονποσα. arguments) Iron statistical mechanics. (wo routes were pursued Berlin 1931).."," In an attempt at providing a justification of these models (in particular, of the so-called $f_{\infty}$ models, constructed initially only from dynamical arguments) from statistical mechanics, two routes were pursued \citep{sti87}."706 The first combines an explicit statement of partial relaxation. i.e. of a relaxation process that is expected to be inefficient in the outer regions. ancl (he existence of a suitable weight. related to the orbital period. for the cells (hat make the relevant partition of phase space: il ollows qualitative arguments proposed by Lynelen-Bell(1967). ancl is physically. appealing (see also Tremaine 1936)).," The first combines an explicit statement of partial relaxation, i.e. of a relaxation process that is expected to be inefficient in the outer regions, and the existence of a suitable weight, related to the orbital period, for the cells that make the relevant partition of phase space; it follows qualitative arguments proposed by \citet{lyn67} and is physically appealing (see also \citealt{tre86a}) )."707 It was indeecl shown to lead naturally to the [ς models., It was indeed shown to lead naturally to the $f_{\infty}$ models.708 However. (his route is not fully satisfactory fom the mathematical point of view. especially since il involves an explicit approximation for the orbital period that is applicable only to the low bindinge energve. limit of quasi-INeplerian orbits.," However, this route is not fully satisfactory from the mathematical point of view, especially since it involves an explicit approximation for the orbital period that is applicable only to the low binding energy limit of quasi-Keplerian orbits."709" The second route is straightforwardex from the nalhbematical point of view. being based on the classical Boltzmann entropy and on the assumed explicit conservation of a (hird quantity Q. in addition to the total mass M. and to the total οποιον 7,4."," The second route is straightforward from the mathematical point of view, being based on the classical Boltzmann entropy and on the assumed explicit conservation of a third quantity $Q$, in addition to the total mass $M$ and to the total energy $E_{tot}$."710 It was shown to lead to an analvticallv different. Family. of models (that we may call the f! models: see definition in Sect., It was shown to lead to an analytically different family of models (that we may call the $f^{(\nu)}$ models; see definition in Sect.711 3 below). with qualitative properties very similar to those of the f/4. models.," 3 below), with qualitative properties very similar to those of the $f_{\infty}$ models."712 Those models were not studied much. further and did not receive great attention. nol only because the relevant distribution Iunction is not as simple as that of the f4. models. but especially because the conservation of Q could only be justified approximately by inspection of a number of N-body simulations. without a clear- physical justification (see StiavelliandBerlinLO87:: in contrast. the conservation of the additional A-D invariant sometimes invoked in plasma physics is rather (ransparent: see ChandrasekharanclWoltjer 1953)).," Those models were not studied much further and did not receive great attention, not only because the relevant distribution function is not as simple as that of the $f_{\infty}$ models, but especially because the conservation of $Q$ could only be justified approximately by inspection of a number of N-body simulations, without a clear-cut physical justification (see \citealt{sti87}; in contrast, the conservation of the additional ${\bf A} \cdot {\bf B}$ invariant sometimes invoked in plasma physics is rather transparent; see \citealt{cha58}) )."713 In this paper we take advantage of the simple statistical mechanics foundation of the [7 family of unbound partially relaxed stellar svstems to explore the possibilitv of a (thermodynamical description of stellar-dvnamical models that are endowed with realistic properties., In this paper we take advantage of the simple statistical mechanics foundation of the $f^{(\nu)}$ family of unbound partially relaxed stellar systems to explore the possibility of a thermodynamical description of stellar-dynamical models that are endowed with realistic properties.714 Belore proceeding to illustrate the results of this paper. we make a short digression in order to bring out the connections between (he present analysis and the evolution of elliptical ealaxies.," Before proceeding to illustrate the results of this paper, we make a short digression in order to bring out the connections between the present analysis and the evolution of elliptical galaxies."715 We start by recalling that. formally. the sequence of Ixing(1966) models is one," We start by recalling that, formally, the sequence of \citet{kin66} models is one"716We thank the anonvmots referee for constructive eriticism.,We thank the anonymous referee for constructive criticism.717 We also thank BBrowne and IxIxoopimans for interesting discussions about radio propagation effects., We also thank Browne and Koopmans for interesting discussions about radio propagation effects.718 iis supported by an Astronomy Astrophysics Postdoctoral Fellowship. under NSF grant 0104347.," is supported by an Astronomy Astrophysics Postdoctoral Fellowship, under NSF grant AST-0104347."719 iis supported by NASA ATP erant. NAG5-9265., is supported by NASA ATP grant NAG5-9265.720Erlangen-Nürrnberg. When the first extrasolar planets (“exoplanets”) were discovered. it became necessary to find names for individual objects.,"rnberg,\tikzmark{mainBodyEnd0} \tikzmark{mainBodyStart1}Sternwartstr.\tikzmark{mainBodyEnd1} \tikzmark{mainBodyStart2}7,\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}96049\tikzmark{mainBodyEnd3} \tikzmark{mainBodyStart4}Bamberg,\tikzmark{mainBodyEnd4} \tikzmark{mainBodyStart5}Germany\tikzmark{mainBodyEnd5}721\tikzmark{mainBodyStart6}}\tikzmark{mainBodyEnd6}722 723% \date{Received September, \ldots 2010 / accepted \ldots 2010}724 725\authorrunning{Hessman et al.} 726\titlerunning{Exoplanet Naming Convention}727 728 729% Abstract: Context, Aims, Methods, Results, [Conclusions]730\abstract { % Context731The present naming convention for extrasolar planets used by the vast majority of researchers in the field is based upon an interpretation of the provisional I.A.U. standard for multiple star systems.732%733With the existence of hundreds of exoplanets around single stars named by this convention and a handful of exoplanets around binary stars -- circumbinary planets --734it has become necessary to find a uniform and useful naming convention for the latter which is maximally compatible with the single host-star convention and which captures as much of the dynamical information about the planet as possible.735%736We propose a simple and generic naming convention for all exoplanets which follows the provisional I.A.U. standard but more clearly indicates their dynamical status.737%738The proposed convention is compatible with present usage and easily extendible to exoplanets around stars in systems of arbitrary multiplicity.739%740We invite comments and discussion on the proposed convention, in the hope of a timely adoption by the I.A.U. Commissions 5, 8+24, 26,\tikzmark{mainBodyStart7},\tikzmark{mainBodyEnd7} \tikzmark{mainBodyStart8}42,\tikzmark{mainBodyEnd8} \tikzmark{mainBodyStart9}45\tikzmark{mainBodyEnd9} \tikzmark{mainBodyStart10}and\tikzmark{mainBodyEnd10} \tikzmark{mainBodyStart11}53.\tikzmark{mainBodyEnd11}741\tikzmark{mainBodyStart12}}\tikzmark{mainBodyEnd12}742{ %% Aims 743% With the existence of hundreds of exoplanets around single stars named by this convention and a handful of exoplanets around binary stars -- circumbinary planets --it has become necessary to find a uniform and useful naming convention for the latter which is maximally compatible with the single host-star convention and which captures as much of the dynamical information about the planet as possible.744\tikzmark{mainBodyStart13}}\tikzmark{mainBodyEnd13}745{ %% Methods746% We propose a simple and generic naming convention for all exoplanets which follows the provisional I.A.U. standard but more clearly indicates their dynamical status.747\tikzmark{mainBodyStart14}}\tikzmark{mainBodyEnd14}748{ %% Results749% The proposed convention is compatible with present usage and easily extendible to exoplanets around stars in systems of arbitrary multiplicity.750\tikzmark{mainBodyStart15}}\tikzmark{mainBodyEnd15}751{ %% Conclusions752% We invite comments and discussion on the proposed convention, in the hope of a timely adoption by the I.A.U. Commissions 5, 8+24, 26, 42, 45 and 53.753\tikzmark{mainBodyStart16}} When the first extrasolar planets (“exoplanets”) were discovered, it became necessary to find names for individual objects."754 Given the large numbers of exoplanets. the Solar System convention of using mythological or other real names was utterly impractical and astronomically uninformative.," Given the large numbers of exoplanets, the Solar System convention of using mythological or other real names was utterly impractical and astronomically uninformative."755 There already exist several naming conventions which have grown out of the historical needs of the visual ad spectroscopic binary. communities., There already exist several naming conventions which have grown out of the historical needs of the visual and spectroscopic binary communities.756" For instance. the components of visual binaries tend to be labeled with capital letters (e.g. AA B). whereas spectroscopic binaries tend to be labeled with small-case letters or numbers (e.g. & AA consists of two stars. ""Aa"" and ""Ab"". ""Aa"" being the primary)."," For instance, the components of visual binaries tend to be labeled with capital letters (e.g. $\xi$ A B), whereas spectroscopic binaries tend to be labeled with small-case letters or numbers (e.g. $\xi$ A consists of two stars, “Aa” and “Ab”, “Aa” being the primary)."757 However. this system is neither officially defined sanctioned by the LA.U. nor is its use in the literature uniform.," However, this system is neither officially defined and sanctioned by the I.A.U. nor is its use in the literature uniform."758 Indeed. the notation used in earlier literature was often the opposite convention. Le. capital letters for primary stars =~ the matching lower-case letters for the secondaries.," Indeed, the notation used in earlier literature was often the opposite convention, i.e. capital letters for primary stars and the matching lower-case letters for the secondaries."759" For example. the names of the components of €UUMa used by Aitken (1935: p. 249). “A” ᾱ for the two components of the ""A"" system and ""B"" ""b for the two components of the ""B system. are carefully renamed by Griffin (1998:: p. 276) to ""Αα. “Ab”. “Ba”. and ""Bb""."," For example, the names of the components of $\xi$ UMa used by Aitken \cite{aitken}; p. 249), `A” ”a” for the two components of the “A” system and “B” “b” for the two components of the “B” system, are carefully renamed by Griffin \cite{griffin}; p. 276) to “Aa”, “Ab”, “Ba”, and “Bb”."760" Most authors relieve themselves from the obviously onerous task of giving the components of spectroscopic or astrometric binaries names by using the terms ""primary"" and ""secondary"" or designating them as “17 or 727. particularly as indices of dynamical parameters."," Most authors relieve themselves from the obviously onerous task of giving the components of spectroscopic or astrometric binaries names by using the terms “primary” and “secondary” or designating them as ""1"" or ""2"", particularly as indices of dynamical parameters."761 This problem with the naming of stars in multiple systems is well-known and the source of many discussions in various commissions within the 1.A.U. (see Hartkopf Mason 2010)., This problem with the naming of stars in multiple systems is well-known and the source of many discussions in various commissions within the I.A.U. (see Hartkopf Mason \cite{HandM}) ).762 The provisional working standard adopted during the XXIV LA.U. convention ts that of the Washington Mulitplicity Catalog (WMC). which uses the following system: This system makes no distinction between stellar. sub-stellar. and planetary objects but does express à clear hierarchical. structure.," The provisional working standard adopted during the XXIV I.A.U. convention is that of the Washington Mulitplicity Catalog (WMC), which uses the following system: This system makes no distinction between stellar, sub-stellar, and planetary objects but does express a clear hierarchical structure."763 One problem with this. system. is. that. the discovery hierarchy. is not necessarily identical with the dynamical hierarchy: does CC orbit around 997950À or perhaps around BB?, One problem with this system is that the discovery hierarchy is not necessarily identical with the dynamical hierarchy: does C orbit around 97950Ȧ or perhaps around B?764" Systems of the latter type are often expressed as AAB. i.e. by concatenating the component suthxes: the WMC contains references to things like “A-BC™. re. a triple system consisting of the brightest component. ""A"". orbiting around a fainter binary. ""B-7C."," Systems of the latter type are often expressed as AB, i.e. by concatenating the component suffixes: the WMC contains references to things like “A-BC”, i.e. a triple system consisting of the brightest component, “A”, orbiting around a fainter binary, “B”+“C”."765 However. this nomenclature is also not adequate enough to express the dynamical state of just one of the components.," However, this nomenclature is also not adequate enough to express the dynamical state of just one of the components."766" Another problem with the WMC nomenclature is that the names are purely accidental and/or historical: we are used to referring to ""Sirius B. not ""Sirius ΑΡΤ: and ""Sirius AB"" can be a reference to both stars together or a capitalized misprint of just one."," Another problem with the WMC nomenclature is that the names are purely accidental and/or historical: we are used to referring to “Sirius B”, not “Sirius Ab”; and “Sirius AB” can be a reference to both stars together or a capitalized misprint of just one."767" Finally. the decision. what to call the ""A"" component is arbitrary: historically. visual binaries yielded the widest binaries and so defined the upper-case usage and spectroscopic binaries came later. inducing the lower-case additions. but it is equally possible to first discover a binary using some other method. which then implies ""A"" and ""B"" (instead of ""Aa"" and ""Ab"") components. and to discover a companion system later astrometrically. which then would be the ""C"" component"," Finally, the decision, what to call the “A” component is arbitrary: historically, visual binaries yielded the widest binaries and so defined the upper-case usage and spectroscopic binaries came later, inducing the lower-case additions, but it is equally possible to first discover a binary using some other method, which then implies “A” and “B” (instead of “Aa” and “Ab”) components, and to discover a companion system later astrometrically, which then would be the “C” component"768the relations. for which the standard notation of -=G=1 is- used.,"the relations, for which the standard notation of c=G=1 is used."769 Physical. leugth scales axe inB unitsB of+ coo and angular frequencies Q are iu units of ο /CAL, Physical length scales are in units of $\sp 2$ and angular frequencies $\Omega$ are in units of $\sp 3$ /GM.770 kr = 1 is defined as the gravitational radius ry., $r$ = 1 is defined as the gravitational radius $\sb{\rm g}$ .771 LEquatious 1 to show that O5 decreases monotouicallv with increasing kc aud. ou the contrary. for ¢ fixed. both 0 and Og initially rise with ». eo through a maxinaun and then decrease with increasing r.," Equations \ref{eq:1} to \ref{eq:4} show that $\Omega\sb{\rm{K}}$ decreases monotonically with increasing $r$ and, on the contrary, for $a$ fixed, both $\Omega\sb{\rm{V}}$ and $\Omega\sb{\rm{R}}$ initially rise with $r$, go through a maximum and then decrease with increasing $r$."772 Qy reaches its maxi at a radial distance denoted by r = Ίνμας aud Og takes its mnaxinnun value at a distance denoted by r = TRinax: , $\Omega\sb{\rm{V}}$ reaches its maximum at a radial distance denoted by $r$ = $\sb{\rm{Vmax}}$ and $\Omega\sb{\rm{R}}$ takes its maximum value at a distance denoted by $r$ = $\sb{\rm{Rmax}}$ .773For anv gr ac a fixed Og > Oy > On., For any $r$ and $a$ fixed $\Omega\sb{\rm{K}}$ $>$ $\Omega\sb{\rm{V}}$ $>$ $\Omega\sb{\rm{R}}$ .774" This nemus tha the shortes period is to be assigned to Og, a the sanallest ¢ possible which is the imnermost mareinally stable orbit at r = ry. defined by Qp = 0.", This means that the shortest period is to be assigned to $\Omega\sb{\rm{K}}$ at the smallest $r$ possible which is the innermost marginally stable orbit at $r$ = $\sb{\rm{ms}}$ defined by $\Omega\sb{\rm R}$ = 0.775 For cach of the other frequencies there is a choice between Oytr = Ts} Or = Tvinax}- Or = TRmax}: Ogtr = Iyuaaax and Or(r = της).," For each of the other frequencies there is a choice between $\Omega\sb{\rm{V}}( r$ = $\sb{\rm{ms}})$, $\Omega\sb{\rm{V}}( r$ = $\sb{\rm{Vmax}})$, $\Omega\sb{\rm{V}}(r$ = $\sb{\rm{Rmax}})$, $\Omega\sb{\rm{R}}(r$ = $\sb{\rm{Vmax}})$ and $\Omega\sb{\rm{R}}(r$ = $\sb{\rm{Rmax}})$."776 Or’ = Ya) is excluded Jecause ) per definition., $\Omega\sb{\rm{R}}(r$ = $\sb{\rm{ms}})$ is excluded because $\Omega\sb{\rm{R}}($ r = $\sb{\rm{ms}})$ = 0 per definition.777 Whether Οι could be generated at kr> Tu d hard to tell. but Οντίν = rg) could approach QF = r4) or could even be lower than he frequeney of any other mode. depending on a.," Whether $\Omega\sb{\rm{LT}}$ could be generated at $r >$ $\sb{\rm{ms}}$ is hard to tell, but $\Omega\sb{\rm{LT}}( r$ = $\sb{\rm{ms}})$ could approach $\Omega\sb{\rm{K}}( r$ = $\sb{\rm{ms}})$ or could even be lower than the frequency of any other mode, depending on $a$."778 We rave chosen three radial fix poiuts. ie. r = This) ή Tyuuax aud ατα ," We have chosen three radial fix points, i.e. $r$ = $\sb{\rm{ms}}$, $r$ = $\sb{\rm{Vmax}}$ and r = $\sb{\rm{Rmax}}$."779In principle there is no preference or anv r except for r = ry. but we note that Qy aud On do not vary a great deal with r over a reasonable range of r centered on Tyas OF Tpuaaxe respectively.," In principle there is no preference for any $r$ except for $r$ = $\sb{\rm{ms}}$ but we note that $\Omega\sb{\rm{V}}$ and $\Omega\sb{\rm{R}}$ do not vary a great deal with $r$ over a reasonable range of $r$ centered on $\sb{\rm{Vmax}}$ or $\sb{\rm{Rmax}}$, respectively."780 Iu he following we use periods. ic. Per. Py. Py and Pg instead of angular frequencies OQ.," In the following we use periods, i.e. $\sb{\rm{LT}}$, $\sb{\rm K}$, $\sb{\rm V}$ and $\sb{\rm R}$ instead of angular frequencies $\Omega$."781 Each observed perio Is now assigned to one mode at either ry. Pymax OL Tmuax cach resulting in a relation between Alpi aud e.," Each observed period is now assigned to one mode at either $\sb{\rm{ms}}$, $\sb{\rm{Vmax}}$ or $\sb{\rm{Rmax}}$ each resulting in a relation between $\sb{\rm{BH}}$ and $a$."782 Tf hese assignineuts cover all observed periods within their neasunrenient uncertaiufles. a conumon value of Mp ane « should emerge.," If these assignments cover all observed periods within their measurement uncertainties, a common value of $\sb{\rm{BH}}$ and $a$ should emerge."783" Croup 0 contaius the shortest perio of 110ss: if assigned to P(t luna.) al upper dnt or Apu & 109 exists for a» d]. which is in conflict with the lowest value of S109NE,. publishes o» Schóddel et al. (2002))."," Group 0 contains the shortest period of s; if assigned to $\sb{\rm{K}}($ r = $\sb{\rm{rms}})$ an upper limit for $\sb{\rm{BH}} < $ $\times$ $\sp 6$ $_\odot$ exists for $a \rightarrow$ 1, which is in conflict with the lowest value of $\times$ $\sp 6$ $_\odot$ published by Schöddel et al. \cite{S2002}) )."784 Therefore lis ascriptionu is excluded., Therefore this ascription is excluded.785 The uext longer period of 219ss implies an upper inüt of Mp< \ LOCAL... if assigned to Pj(r = Yuus) ., The next longer period of s implies an upper limit of $\sb{\rm{BH}} < $ $\times$ $\sp 6$ $_\odot$ if assigned to $\sb{\rm{K}}($ r = $\sb{\rm{rms}})$ .786 This is within the range of values Or Mp reported iu the iterative auc suggests the assiguinent to Ixepleriau motion aloug the mareially stable orbit., This is within the range of values for $\sb{\rm{BH}}$ reported in the literature and suggests the assignment to Keplerian motion along the marginally stable orbit.787 The only solutiou of equations 1 to L with consistent values for AlIpy aud α is obtained for the following assignation of period aud group 1. 21955 = Pu(r214) [Is]: group 3. 692ss = Pyten) [VI]: group L. ss = PR(r=rriuas) [B].," The only solution of equations \ref{eq:1} to \ref{eq:4} with consistent values for $\sb{\rm{BH}}$ and $a$ is obtained for the following assignation of period and group 1, s = $\sb{\rm{K}}(r$ $\sb{\rm{ms}})$ [K]; group 3, s = $\sb{\rm{V}}(r$ $\sb{\rm{ms}})$ [V]; group 4, s = $\sb{\rm{R}}(r$ $\sb{\rm{Rmax}})$ [R]."788 With this set equation { predicts Pppz 3208s. Such a period shows up with increased power in the June 16 iufraxrec flare ss. Cenzel et al. 2003..," With this set equation \ref{eq:4} predicts $\sb{\rm{LT}} \approx$ s. Such a period shows up with increased power in the June 16 infrared flare s, Genzel et al. \cite{Ge2003},"789 cf, c.f.790 section 2.3)., section 2.3).791 The four relations for Alpi versus (1a) following from equations d. to Lo are drawn in flere Ts which shows an almost perfect. single intersection of the four relations.," The four relations for $\sb{\rm{BH}}$ versus $(1 - a)$ following from equations \ref{eq:1} to \ref{eq:4} are drawn in figure \ref{fig7}, which shows an almost perfect, single intersection of the four relations."792 The wide error polveon of figure 7. iuclicates the naxinal range for Mp and (1 a) without the 321ss period aud he inner polveon covers the coustraiuts set by the ss iod., The wide error polygon of figure \ref{fig7} indicates the maximal range for $\sb{\rm{BH}}$ and (1 $-~a$ ) without the s period and the inner polygon covers the constraints set by the s period.793 This solution docs not suggest a specific radius Ike Tuas Ένας OF λα for the remaimine period of ~500 s. bu it is consistent with a I&epler oriod expected close ο F—IYpmuax.," This solution does not suggest a specific radius like $\sb{\rm{ms}}$, $\sb{\rm{Vmax}}$ or $\sb{\rm{Rmax}}$ for the remaining period of $\sim$ 500 s, but it is consistent with a Kepler period expected close to $r$ $\sb{\rm{Rmax}}$."794" Che periods of eroup 5 around ss, cannot arise from a region rmua "," The periods of group 5 around s, cannot arise from a region $r <$ $\sb{\rm{Rmax}}$ ."795Uowever. we point out that this period is very close to jitst twice the period of he radial epicyclie mode aud might iudicate a domimauce of the amplitude of every second cvcle;," However, we point out that this period is very close to just twice the period of the radial epicyclic mode and might indicate a dominance of the amplitude of every second cycle."796 We note that the eroup 0 period of —100 ss is close to he first haxinonic (10 = 2) expected for the ‘cdiskoscismoloey” C-node. (Nowak Lehr 1998))., We note that the group 0 period of $\sim$ s is close to the first harmonic (m = 2) expected for the 'diskoseismology' C-mode (Nowak Lehr \cite{NL1998}) ).797" The best-fit solution to the intersections is Mp = 415. ον 109ME. and « = 0.993926,", The best-fit solution to the intersections is $\sb{\rm{BH}}$ = $\sp{+0.12}\sb{-0.19}$ $\times$ $\sp 6$ $_\odot$ and $a$ = $\sp{+0.0026}\sb{-0.0074}$.798 The errors quoted are not lo errors but reflect the maximal possible uncertaintv elven o» the maximal uncertainties of the period determination., The errors quoted are not $\sigma$ errors but reflect the maximal possible uncertainty given by the maximal uncertainties of the period determination.799 The radi involved are lus = 1.971 and μιας —2.507., The radii involved are $\sb{\rm{ms}}$ = 1.371 and $\sb{\rm{Rmax}}$ =2.507.800 The gravitational radius (r=1) aimouuts to L1410 1 Cou., The gravitational radius $r$ =1) amounts to $\times$ $\sp{11}$ cm.801 Tudepenucent measurements of the orbital motion of the star SO-2 (82) around the GC black hole have resulted in ΛΙΠΗ = 3.7THL & LOOM. (Schéedel et al. 2002)), Independent measurements of the orbital motion of the star S0-2 (S2) around the GC black hole have resulted in $\sb{\rm{BH}}$ = $\pm$ 1.5 $\times$ $\sp 6$ $_\odot$ (Schöddel et al. \cite{S2002}) )802 and Mp = LOTEO.62 ον 109MI. (Ro/skpey? (Chez et al. 2003)), and $\sb{\rm{BH}}$ = $\pm$ 0.62 $\times$ $\sp 6$ $_\odot$ $\sb 0/8 {\rm{kpc}})\sp 3$ (Ghez et al. \cite{Gh2003}) )803 with Ry the distauce to the black hole., with $\sb 0$ the distance to the black hole.804 Our result isconsistent with the Schóddel et al., Our result isconsistent with the Schöddel et al.805 nueasuremenut at their 1 σ levelaud it also agrees with Chez et s nieasuremaent at their 26 level and Ry = 8 kpc., measurement at their 1 $\sigma$ leveland it also agrees with Ghez et 's measurement at their $\sigma$ level and $\sb 0$ = 8 kpc.806 Euler Cenzel eal. (19963)., Earlier Genzel etal. \cite{GTK1996}) )807determined Alpy to 2.5 - 340 & LOPAL.. from meastrenents of dark matter couccutration iu the central parsee of the \filsy Way which appears to match our result best.,determined $\sb{\rm{BH}}$ to 2.5 - 3.0 $\times$ $\sp 6$ $_\odot$ from measurements of dark matter concentration in the central parsec of the Milky Way which appears to match our result best.808 The high value of à z 1 iueaus that the emission from the Inner parts of the accretion disk is quite lose to the black hole., The high value of $a$ $\approx$ 1 means that the emission from the inner parts of the accretion disk is quite close to the black hole.809 The event horizon is located at τι = ].112., The event horizon is located at $\sb{\rm H}$ = 1.112.810 With Ty. = 1.5Tl the distance to the marginally stable orbit is just 3.6 light seconds., With $\sb{\rm{ms}}$ = 1.371 the distance to the marginally stable orbit is just 3.6 light seconds.811 Light from matter crossine the marginally stable orbit is therefore very likely to disappear in a fairly short time., Light from matter crossing the marginally stable orbit is therefore very likely to disappear in a fairly short time.812 Is the PDS peak at —11 s the characteristic time scale at which the light is facing away?, Is the PDS peak at$\sim$ 11 s the characteristic time scale at which the light is fading away?813 The Πο curve of the flare (figure laa) shows a deep.very short (< 100 8) cu close to asin," The light curve of the flare (figure \ref{fig1}a a) shows a deep,very short $<$ 100 s) cut close to maximum"814If the shock is not strongly collimated. AG51/17. we can neglect the factor OyInR. different parts of the shock will propagate radially with a different Lorentz factor giveu by the driver or the exterual cleusity inhomogeneity: to? ?.. (2)),"If the shock is not strongly collimated, $\Delta \theta \gg 1/\Gamma^2$, we can neglect the factor $ \partial_\theta \ln R $, different parts of the shock will propagate radially with a different Lorentz factor given by the driver or the external density inhomogeneity: \cite{LaumbachProbstein} \cite{1979ApJ...233..831S}. \ref{Komp})"815 ? (logarithmic.e.g...) ," \cite{1979ApJ...233..831S} \citep[logarithmic, \eg][]{} "816The ionization fraction of molecular clouds is au nportaut xumneter for their ανα]ος hrough its coutrol over the influence of anv maenetic fiek.,The ionization fraction of molecular clouds is an important parameter for their dynamics through its control over the influence of any magnetic field.817 The ionization also has a uajor effect ou the chemistry of molecular clouds because ion-neutral reactions are generally much faster than jeutral-neutral reactions., The ionization also has a major effect on the chemistry of molecular clouds because ion-neutral reactions are generally much faster than neutral-neutral reactions.818 In deuse regious shielded frou direct ultraviolet radiation. the ionization is doniünated ov cosunc ravs.," In dense regions shielded from direct ultraviolet irradiation, the ionization is dominated by cosmic rays."819 However. the rate of this process thas not vet been coustramed directly.," However, the rate of this process has not yet been constrained directly."820 The current best estimate comes from chemical models to reproduce the observed abundances of ΟΠ aud ΠΟ in diffuse interstellar clouds (ITartquist et 11978: van Dishoeck Black 1986: Federman et 11996).. notablv those oward Perseus OB2.," The current best estimate comes from chemical models to reproduce the observed abundances of OH and HD in diffuse interstellar clouds (Hartquist et 1978; van Dishoeck Black 1986; Federman et \nocite{hart78,evd86,feder96}, notably those toward Perseus OB2."821 These models indicate that cey=l0πο1T. 1 per II atom. but with a factor of 10 uncertainty because of uncertainties in temperature. radiation field) and the effects of shocks.," These models indicate that $=10^{-16} - 10^{-17}$ $^{-1}$ per H atom, but with a factor of 10 uncertainty because of uncertainties in temperature, radiation field and the effects of shocks."822 In addition. it is uukuown if yyaries with location in the Calasxy. since the diffuse cloud results are limited to the solar neighbourhood.," In addition, it is unknown if varies with location in the Galaxy, since the diffuse cloud results are limited to the solar neighbourhood."823 Each cosuic-ray ionization of IT» vields one nuiuolecule. so hat hhas a consta concentration which depends oulv on aand the abundances of its main destrovers: CO. O aud clectrous.," Each cosmic-ray ionization of $_2$ yields one molecule, so that has a constant concentration which depends only on and the abundances of its main destroyers: CO, O and electrons."824 Hence. the recent detections of infrared absorption lines by Geballe&Ola(1996) aud MeCalletal.(1999). toward massive protostars provide a novel wav to measure Gcgi.," Hence, the recent detections of infrared absorption lines by \cite{gebal96}825 and \cite{bmcc99} toward massive protostars provide a novel way to measure ."826. We coustrall musing models by vanderTaketal.(2000) of the envelopes of these stars. and use the derived values to inodel observations of.. the abundance of which is also proportional toQc.," We constrain using models by \cite{fvdt00} of the envelopes of these stars, and use the derived values to model observations of, the abundance of which is also proportional to."827.. MeCall et preseut observatious of rovibrational lines of iin absorption against seven huniuous (10!10°L.3), McCall et \nocite{bmcc99} present observations of rovibrational lines of in absorption against seven luminous $10^4-10^5$)828of the absorption lues studied im this paper.,of the absorption lines studied in this paper.829 Both the spectral resolution aud signal-to-noise ratio are better iu the new observations., Both the spectral resolution and signal-to-noise ratio are better in the new observations.830 Analysis of the new data will be published elsewhere., Analysis of the new data will be published elsewhere.831 It is important to stress that the strongest absorptions (Fel. Wel} are disfavored for the optical star radial velocity analysis. as they mos probably formed in the powerful outficdue disk wind (Fabrika et al.," It is important to stress that the strongest absorptions (FeII, HeI) are disfavored for the optical star radial velocity analysis, as they most probably formed in the powerful outflowing disk wind (Fabrika et al."832 1997)., 1997).833 We studied radial velocities of weak absorption lines of metals., We studied radial velocities of weak absorption lines of metals.834ρα It is easy to coifuse the Ine idetification when one traces a line from date to date because of the spectral variability and orbital motion of the companious., It is easy to confuse the line identification when one traces a line from date to date because of the spectral variability and orbital motion of the companions.835 In Fig., In Fig.836 lr we preseut indiviual absorption lijo radial velocities for four best-suited lines averaged over one night. as a function of the orbital pli:we for 6 nights.," \ref{fig:rvindid} we present individual absorption line radial velocities for four best-suited lines averaged over one night, as a function of the orbital phase for 6 nights."837 The radia velocity curve measured bv the inost reliable collection of 22 absorption lines in the spectral range 1200-524300 is shown in Fie. 18.., The radial velocity curve measured by the most reliable collection of 22 absorption lines in the spectral range 4200-5300 is shown in Fig. \ref{fig:radvel}.838 Note that with our spectral resolution the absorption lines iu the $8133 spectimm are seen as blends coutaimius 2-3 lines. and we measured their relative radial velocities because of uukuown “laboratory” waveleugths of the blends.," Note that with our spectral resolution the absorption lines in the SS433 spectrum are seen as blends containing 2-3 lines, and we measured their relative radial velocities because of unknown ""laboratory"" wavelengths of the blends."839 Relative iuteusities of the lines change from date to date through the eclipse., Relative intensities of the lines change from date to date through the eclipse.840 For this reason we neasured only those lines aud on only those dates where the lines were the most convincingly detected., For this reason we measured only those lines and on only those dates where the lines were the most convincingly detected.841"iterations of the method, p(us;x), n—1,...,5, when and 6=2 and M—100.","iterations of the method, $\rho(\mu_n;x)$ , $n=1,\ldots,5$, when and $\delta = \frac{2}{5}$ and $M=100$."842" 'The last two curves are practically undistinguishable, indicating convergence."," The last two curves are practically undistinguishable, indicating convergence."843" Indeed, a distance function is required to assess this fact with rigour."," Indeed, a distance function is required to assess this fact with rigour."844 Various choices are possible., Various choices are possible.845 In Fig., In Fig.846" 2 we plot all of the following distances, in the same case of Fig. 1,,"," \ref{distafig2} we plot all of the following distances, in the same case of Fig. \ref{rofig1},"847 but on a larger number of iterations., but on a larger number of iterations.848 x))=(, =.849"18) In all cases we observe exponential decrease of dj(p(n;x),p(ps1;x)) with n, which implies convergence to a density ρµ(µ;x) that is an approximation of p(u;x)."," In all cases we observe exponential decrease of $d_l(\rho(\mu_n;x),\rho(\mu_{n-1};x))$ with $n$, which implies convergence to a density $\rho_M(\mu;x)$ that is an approximation of $\rho(\mu;x)$."850" Validity of the approximation can be checked by increasing the value of M, see the next Section and Sect. 6,,"," Validity of the approximation can be checked by increasing the value of $M$, see the next Section and Sect. \ref{sec-sobo},"851 where a drastically different case is encountered., where a drastically different case is encountered.852" Opposite to the previous choice of the measure c, lies the Bernoulli measure, that is, an atomic measure composed of two atoms."," Opposite to the previous choice of the measure $\sigma$, lies the Bernoulli measure, that is, an atomic measure composed of two atoms."853" This is the first non-trivial case, because one atom alone leads to the equality yp=c."," This is the first non–trivial case, because one atom alone leads to the equality $\mu = \sigma$."854" Recall the interpretation: if the prey at 6 does not move, the predator x converges to it, and the invariant distribution p is an atomic measure locatedat the same point."," Recall the predator--prey interpretation: if the prey at $\beta$ does not move, the predator $x$ converges to it, and the invariant distribution $\mu$ is an atomic measure locatedat the same point."855 Let therefore σ be defined as the sum of two atomic measures located at, Let therefore $\sigma$ be defined as the sum of two atomic measures located at856WUT spectra of NGC L151 (Bromage et al 1985. Kriss et al 1992).,"HUT spectra of NGC 4151 (Bromage et al 1985, Kriss et al 1992)."857" Bromage et al (1985) argue that absorption by je excited imoetastable level of 1175 and its streneth relative to CTV 1519 require this level to be collisionally »opulated ina Ligh deusity medium with N, = 1019 7.", Bromage et al (1985) argue that absorption by the excited metastable level of $^*$ 1175 and its strength relative to CIV 1549 require this level to be collisionally populated in a high density medium with $_{e}$ = $10^{10}$ $^{-3}$.858 Iu NGC 1151. the EW of 1175 is between 0.7 aud 1.0 iues the EW of CIV1518.1550 while iu PCULOL|226. it is between 0.35 aud 0.5 the EW of CIVI518.1550.," In NGC 4151, the EW of $^*$ 1175 is between 0.7 and 1.0 times the EW of CIV1548,1550 while in PG1404+226, it is between 0.35 and 0.5 the EW of CIV1548,1550."859 Note lat an appealing alternative for the identification of this ine is theblieshifted bv 0.3 with respect to POI011226., Note that an appealing alternative for the identification of this line is the by 0.3 with respect to PG1404+226.860 The line width. ouly z 2A. however. argues against this interpretation.," The line width, only $\approx$ 2, however, argues against this interpretation."861 There is no other candidate for Helly blucshitted absorption lines in the spectra., There is no other candidate for highly blueshifted absorption lines in the spectrum.862 The 2 = 0.098 system is consistent with being produced iu the halo of the host ealaxy or of a nearby companion but could also be intrinsic to the mucleus., The $z$ = 0.098 system is consistent with being produced in the halo of the host galaxy or of a nearby companion but could also be intrinsic to the nucleus.863 As for the 2 = 0.091 system the probable presence (vhich needs verification) of the 1176 line suggests that the svstei is iutriusic to the quasar and forms in an outfowing wind with a velocity of 1900 kin st. (a) The EW of the UV absorption lines expected from the Wari Absorber have been calculated for the statistically acceptable models of the ROSAT data., As for the $z$ = 0.091 system the probable presence (which needs verification) of the $^*$ 1176 line suggests that the system is intrinsic to the quasar and forms in an outflowing wind with a velocity of 1900 km $^{-1}$ (a) The EW of the UV absorption lines expected from the Warm Absorber have been calculated for the statistically acceptable models of the ROSAT data.864 They. are eiveu in Table 5. separately for low- aud ligh-state as a nieasure of the uncertainty arising from the continui variability (non simultaneous UV aud X-observations. non- equilibrium of the gas. ete;," They are given in Table 5, separately for low- and high-state as a measure of the uncertainty arising from the continuum variability (non simultaneous UV and X-observations, non- equilibrium of the gas, etc."865 see Section 1 aud Nicastro et al., see Section 1 and Nicastro et al.866 19994)., 1999a).867 We find that the CIV1550 aud NV. 1210 absorption lines are weak at low state. and uceligible at high state (because € aud N are highlv ionized).," We find that the CIV1550 and NV 1240 absorption lines are weak at low state, and negligible at high state (because C and N are highly ionized)."868 There is always sole absorption by hydrogen due to the laree column density in II. (b) Following Spitzer (1978). a standard curve of erowth was calculated for velocity paramcters b = 20.G0.100.110. kan + (Figs," There is always some absorption by hydrogen due to the large column density in H. (b) Following Spitzer (1978), a standard curve of growth was calculated for velocity parameters b = 20,60,100,140 km $^{-1}$ (Figs."869 ad), 8a-d).870 The predicted equivalent widths of Lvo. CIV. NV. and (OVI) were then compared with those derived from the analysis of the UST spectrin.," The predicted equivalent widths of $\alpha$, CIV, NV, and (OVI) were then compared with those derived from the analysis of the HST spectrum."871 We find that at the ROSAT low-state (which is close to the state of the source at the time of the ASCA observations). and for b about 60 km 5. there is a rough match between the Ίνα aud. CTV absorption lines produced by the wari absorber and the observed lines at 2 = 0.098 and 2 = 0.090 (nodoels c-0).," We find that at the ROSAT low-state (which is close to the state of the source at the time of the ASCA observations), and for b about 60 km $^{-1}$, there is a rough match between the $\alpha$ and CIV absorption lines produced by the warm absorber and the observed lines at $z$ = 0.098 and $z$ = 0.090 (models c-e)."872 The NV absorption line youn the wari absorber is. however. always weaker than observed (N is too highly iomized iu the models).," The NV absorption line from the warm absorber is, however, always weaker than observed (N is too highly ionized in the models)."873 Tucluding au additional EUV bip will further imerease the level of ionization. thus not changing the above conclusious.," Including an additional EUV bump will further increase the level of ionization, thus not changing the above conclusions."874 The differeut models differ most stronely iu OVI. so this may be the most restrictive line. but it falls just outside the IIST Iu conclusion. we find no sinele-plhase medimm which can produce both the UV and the X-ray absorption lines. (," The different models differ most strongly in OVI, so this may be the most restrictive line, but it falls just outside the HST In conclusion, we find no single-phase medium which can produce both the UV and the X-ray absorption lines. ("875c) In the case of the best fit to the ASCA data ie. the model with Fe overabundaut bv a factor of ~ 22 over solar. the degree of ionization is not known but it is likely to be too high for the production of the absorption line.,"c) In the case of the best fit to the ASCA data i.e. the model with Fe overabundant by a factor of $\sim$ 22 over solar, the degree of ionization is not known but it is likely to be too high for the production of the $^{*}$ absorption line."876 Thus. in this case also. the UV ancl X-ray. absorbers are very probably in different gaseous phases.," Thus, in this case also, the UV and X-ray absorbers are very probably in different gaseous phases."877 Tab., Tab.878 5 gives the intensity of the strongest lines euütted bv the absorber within the waveleneth range of the UST spectra and in the optical range., 5 gives the intensity of the strongest lines emitted by the absorber within the wavelength range of the HST spectra and in the optical range.879 The calculations were performed with a density of log nip = 9.5., The calculations were performed with a density of log $_{WA}$ = 9.5.880 The lue NeVIITATT l was added to the list because of receut reports of its detection iu high-: quasars (IEunaun ct al., The line $\lambda$ 774 was added to the list because of recent reports of its detection in $z$ quasars (Hamann et al.881 1997)., 1997).882 Table 5 is neant to provide an order of magnitude estimate of which lines may be important/detectable iu the future., Table 5 is meant to provide an order of magnitude estimate of which lines may be important/detectable in the future.883 The acual streneth of the lines depends ou the covering factor of the wart absorber: total coverage was asstuued for the values in Table 5., The actual strength of the lines depends on the covering factor of the warm absorber; total coverage was assumed for the values in Table 5.884 Iu the recent vears it has been realized that eas outflow is ubiquitous in ACN., In the recent years it has been realized that gas outflow is ubiquitous in AGN.885 It occurs under different gas phases: xoad cussion line eas (tle most lielly ionized lines are the most blueshifted) aud UV/optical absorption lines (always observed at rest or blueshifted)., It occurs under different gas phases: broad emission line gas (the most highly ionized lines are the most blueshifted) and UV/optical absorption lines (always observed at rest or blueshifted).886 It is likely that he X-ray absorption features also. originate in outfiowiug eas at velocities comparable with or ligher than those of the UV emissiou/absorptionu gas - this cannot presently )o ascertained because of the insufficient energy resolution of the N-ray instruments. [," It is likely that the X-ray absorption features also, originate in outflowing gas at velocities comparable with or higher than those of the UV emission/absorption gas - this cannot presently be ascertained because of the insufficient energy resolution of the X-ray instruments. ["887This is not counting here the extraordinary blueshift of the ταν absorbing eas if the L keV features are blucshitted OVIT or OVII edges].,This is not counting here the extraordinary blueshift of the X-ray absorbing gas if the 1 keV features are blueshifted OVII or OVII edges].888 The question as to whether UV absorption lines aud X-ray absorption lnes/edges tend to be preseut together or separately in AGN has been adressed by Ulrich (197.S) id mere recently by a number of authors (e.g. Schartel et ., The question as to whether UV absorption lines and X-ray absorption lines/edges tend to be present together or separately in AGN has been adressed by Ulrich (1988) and more recently by a number of authors (e.g. Schartel et al.889 1997. Crenshaw 1997. Mathliur 1997. Shields Taman 1997. I&xiss et al 1996).," 1997, Crenshaw 1997, Mathur 1997, Shields Hamann 1997, Kriss et al 1996)."890" Iu some ACN, the data appear consistent with a sinele-ase. photoionized plasima producing the UV absorption vd the OVID and OVIIE edges (NGC 3783 Shields Tamann 1997: 3€212. 30351 aud NGC 5518 Mathur 1997)."," In some AGN, the data appear consistent with a single-phase, photoionized plasma producing the UV absorption and the OVII and OVIII edges (NGC 3783 Shields Hamann 1997; 3C212, 3C351 and NGC 5548 Mathur 1997)."891 In contrast. in other ACN. he properties of the UVP and the ταν absorbers imply the presence of nultiphase," In contrast, in other AGN, the properties of the UV and the X-ray absorbers imply the presence of multiphase"892revealed a rather hard spectrum. fitted with a power law of photon index 1.6 and an absorption column density of Ny=6x107 em? (?)..,"revealed a rather hard spectrum, fitted with a power law of photon index 1.6 and an absorption column density of $N_{H} = 6 \times 10^{22}$ $^{-2}$ \citep{swank03}."893 Timing analysis of the RXTE All-Sky Monitor (ASM) data showed an X-ray period of 13.55 days (?).., Timing analysis of the RXTE All-Sky Monitor (ASM) data showed an X-ray period of 13.55 days \citep{corbet04}.894 ? presented high energy spectral analysis of the period of the discovery. concluding that the source manifests two distinct spectral behaviors. the first showing a thermal component in the soft X-ray and hard X-ray tail. the second being harder and possibly originating from thermal Comptonization.," \citet{hann04} presented high energy spectral analysis of the period of the discovery, concluding that the source manifests two distinct spectral behaviors, the first showing a thermal component in the soft X-ray and hard X-ray tail, the second being harder and possibly originating from thermal Comptonization."895 This second. low-luminosity. behavior was confirmed to be the preferred state of the source (?)..," This second, low-luminosity, behavior was confirmed to be the preferred state of the source \citep{rodri05}."896" The optical/infrared counterpart to IGR 191402095] was identified by ?.. from Chandra accurate position determination. as the heavily reddened Figure 6 shows the A, spectrum we obtained. with identified spectral features marked."," The optical/infrared counterpart to IGR J19140+0951 was identified by \citet{int06}, from Chandra accurate position determination, as the heavily reddened Figure \ref{fig:IGRJ19140} shows the $K_{s}$ spectrum we obtained, with identified spectral features marked."897 The spectrum shows emission. absorption at 211126Á. a weak Cor Cup emission feature at 1155A.. and moderately strong Bry absorption. typical features of an early supergiant.," The spectrum shows emission, absorption at 126, a weak (or ) emission feature at 155, and moderately strong $\gamma$ absorption, typical features of an early supergiant."898 By comparison with the atlases from ??.. we can conclude that the star is a BI Lab type.," By comparison with the atlases from \citet{hanson96,hanson05}, we can conclude that the star is a B1 Iab type."899 Together with X-ray properties. this allows us to confirm the nature of the system as an Our preliminary results were published in ?.. and recently confirmed by ?.. who. from K- and H-band spectra. constrained the spectral type to a BO.S supergiant.," Together with X-ray properties, this allows us to confirm the nature of the system as an Our preliminary results were published in \citet{nespoli07}, and recently confirmed by \citet{hann07}, who, from $K$ - and $H$ -band spectra, constrained the spectral type to a B0.5 supergiant."900 They also estimated the distance of the source as 5 kpe., They also estimated the distance of the source as 5 kpc.901 From the identified spectral types. we obtained the intrinsic colors (J—Κο from ?:: we then calculated. from 2MASS photometry. the instrumental colors (J-K)sgass. properly transformed through the formula from ? to the ? homogenized photometric system in order to estimate the infrared color excess E(—K).," From the identified spectral types, we obtained the intrinsic colors $(J-K)_{0}$ from \citet{weg94}; we then calculated, from 2MASS photometry, the instrumental colors $(J-K)_{2MASS}$, properly transformed through the formula from \citet{car01} to the \citet{bebe88} homogenized photometric system in order to estimate the infrared color excess $E(J-K)$."902 Assuming the mean extinction law (Ry= 3.1). from Av/E(/—K)=5.82+0.10 (2).. we obtained the total measured visual extinction Ay and the corresponding hydrogen column density value from Nj/Ay=1.79+0.03x10°! atoms em mag (?)..," Assuming the mean extinction law $R_{V}= 3.1$ ), from $A_{V}/E(J-K) = 5.82 \pm 0.10$ \citep{rieke85}, we obtained the total measured visual extinction $A_{V}$ and the corresponding hydrogen column density value from $N_{H}/A_{V} = 1.79 \pm 0.03 \times 10^{21}$ atoms $^{-2}$ mag \citep{pred95}."903" We were thus able to compare the retrieved interstellar value of Nj, with the one provided by X-ray data.", We were thus able to compare the retrieved interstellar value of $N_{H}$ with the one provided by X-ray data.904 In our calculation. we estimated errors through error propagation.," In our calculation, we estimated errors through error propagation."905 Errors in the final values of Nj; are mainly due to errors in the infrared colors and of the transformation between the two photometric systems., Errors in the final values of $N_{H}$ are mainly due to errors in the infrared colors and of the transformation between the two photometric systems.906 We also estimated the distance of the six sources. applying the relation My=K -5-5log d—Ay.," We also estimated the distance of the six sources, applying the relation $M_{K} = K + 5 - 5$ log $d - A_{K}$."907 For each source. My was obtained from our proposed spectral type (?).. the intrinsic color index (K—V)o from ? was used in order to calculate Mj. and the 2MASS K magnitude was employed.," For each source, $M_{V}$ was obtained from our proposed spectral type \citep{weg06}, , the intrinsic color index $(K-V)_{0}$ from \citet{weg94} was used in order to calculate $M_{K}$, and the 2MASS $K$ magnitude was employed."908 We derived A from the relation A;/E(J—K)= 24607. fort =2.2pm (?)..," We derived $A_{K}$ from the relation $A_{\lambda} / E(J-K) = 2.4 (\lambda)^{-1.75}$ , for $ \lambda = 2.2\ \mu$ m \citep{draine89}."909 The results of our caleulations are given in. Table 3.. together with some crucial quantities used in the calculations or displayed for comparison.," The results of our calculations are given in Table \ref{tab:reddening}, , together with some crucial quantities used in the calculations or displayed for comparison."910 The distance estimation ts mainly affected by the uncertainty in the value of the absolute magnitude My. which is due to two contributions. the errors given in the tabulated values of My and the uncertainty in the spectral classification. from which the absolute magnitude is determined.," The distance estimation is mainly affected by the uncertainty in the value of the absolute magnitude $M_{V}$, which is due to two contributions, the errors given in the tabulated values of $M_{V}$ and the uncertainty in the spectral classification, from which the absolute magnitude is determined."911 The largest role is played by the errors in the mean tabulated values themselves (see Wegner 2006 for more details)., The largest role is played by the errors in the mean tabulated values themselves (see Wegner 2006 for more details).912 The retrieved values for d must thus be assumed with prudence., The retrieved values for $d$ must thus be assumed with prudence.913 Using near-infrared spectroscopy of six high-energy sources. IGR J16207-5129. IGR J16465-4507. IGR J16479-4514. AX J1841.0-0536. 4U 19074097 and IGR J19140+0951. we classified their. counterparts through comparison with published atlases.," Using near-infrared spectroscopy of six high-energy sources, IGR J16207–5129, IGR J16465–4507, IGR J16479–4514, AX J1841.0–0536, 4U 1907+097 and IGR J19140+0951, we classified their counterparts through comparison with published atlases."914 We found thatallthe observed systemshave a supergiant companion., We found thatallthe observed systemshave a supergiant companion.915 Our results. combined with information from X-ray data. 1s able for the first time to firmly," Our results, combined with information from X-ray data, is able for the first time to firmly"916also to the amplitude of individual flickering events.,also to the amplitude of individual flickering events.917 There is thus no compelling evidence for a change iu the state of the system. but this possibility cannot be securely discounted.," There is thus no compelling evidence for a change in the state of the system, but this possibility cannot be securely discounted."918 Assundug it is reasonable to compare the respouses from the three bursts. the major differeuce is that MD2 and ΑΠΟ both have responses that are lageed and sineared by iore.," Assuming it is reasonable to compare the responses from the three bursts, the major difference is that MB2 and MB3 both have responses that are lagged and smeared by more."919 This is as expected from their phases., This is as expected from their phases.920 Fie., Fig.921 1l shows predicted responses as a fiction of orbital pliase caleulated using the methods described by OBiceuetal.(2002) for model 2 (4= 0.2)., \ref{EchotomFig} shows predicted responses as a function of orbital phase calculated using the methods described by \citet{OBrien:2002a} for model 2 $q=0.2$ ).922 Given the Buaitations inherent in approximating the true response with a Gaussian. the range of lags observed is iu reasonable agreement with predictions.," Given the limitations inherent in approximating the true response with a Gaussian, the range of lags observed is in reasonable agreement with predictions."923 Iu particular. we do expect that MD2 (phase 0.38) will have a simular lagdistribution to MD2 (phase ~ 0.5). whereas we expect a quicker and less sineared response from MDI. as observed.," In particular, we do expect that MB2 (phase 0.38) will have a similar lag distribution to MB3 (phase $\sim 0.5$ ), whereas we expect a quicker and less smeared response from MB1, as observed."924 The wuplitude of the difference between responses d8 du approximate agreement with expectations., The amplitude of the difference between responses is in approximate agreement with expectations.925 When plotted in this way. it is apparent that all bursts appear to show the onset of a strong response at approximately the same time. but that MD2 aud MD2 extend for longer bevoud that.," When plotted in this way, it is apparent that all bursts appear to show the onset of a strong response at approximately the same time, but that MB2 and MB3 extend for longer beyond that."926 While the ecnueral picture is im agreement with predictions. it is clear that the data are not precise enough to constrain the system parameters further.," While the general picture is in agreement with predictions, it is clear that the data are not precise enough to constrain the system parameters further."927 Nonetheless. the very. fact that tle response appears to change with phase. aud with such a large amplitude. docs indicate significant response from the companion. and the large sincaring of the response around plase 0.5 indicates that the disk must also contribute.," Nonetheless, the very fact that the response appears to change with phase, and with such a large amplitude, does indicate significant response from the companion, and the large smearing of the response around phase 0.5 indicates that the disk must also contribute."928 We cau also consider these results in the context of other observations of this aud other sources., We can also consider these results in the context of other observations of this and other sources.929 The conrparison of inost interest is with Schoenmbs&Zoeschinger(1990).., The comparison of most interest is with \citet{Schoembs:1990a}.930 These authors found that optical bursts (with uo simultancous X-ray coverage) exhibited longer rise times when the companion star would be ou the far side of the disk., These authors found that optical bursts (with no simultaneous X-ray coverage) exhibited longer rise times when the companion star would be on the far side of the disk.931 This is equivalent to saving that the smearing increases close to pliase 0.5 (as we also fiud). provided that the iutriusiec rise time of the X-ray bursts is the same.," This is equivalent to saying that the smearing increases close to phase 0.5 (as we also find), provided that the intrinsic rise time of the X-ray bursts is the same."932 The large dispersion iu the response that we see close to phase 0.5 is also not unprecedoeuted., The large dispersion in the response that we see close to phase 0.5 is also not unprecedented.933 Similar dispersions were found by Trucniperetal.(1985) for 11636536. aud by Itongetal.(2000). for 11826," Similar dispersions were found by \citet{Truemper:1985a} for 1636–536, and by \citet{Kong:2000a} for 1826--24."934 In all of these cases they likely indicate coutributiou from both disk and companion. as a response from the conrpauion alone will vary iu lag. but should always have a relatively narrow response.," In all of these cases they likely indicate contribution from both disk and companion, as a response from the companion alone will vary in lag, but should always have a relatively narrow response."935 There are two likely sources of simcaring of the reprocessed signal., There are two likely sources of smearing of the reprocessed signal.936 The delays most commonly discussed (and considered so far) are those from Lelt-travel times across the binary: global delavs., The delays most commonly discussed (and considered so far) are those from light-travel times across the binary; global delays.937 A finite reprocessing time could also introduce a local delay., A finite reprocessing time could also introduce a local delay.938 N-ravs above about LkkeV will expericuce a low-enough photo-clectric absorption opacity to deposit cucrey at a siguificaut optical depth in au atmosphere., X-rays above about keV will experience a low-enough photo-electric absorption opacity to deposit energy at a significant optical depth in an atmosphere.939 Reprocessed energy will then have to diffuse outward. aud so there will be au additional diffusion time delay.," Reprocessed energy will then have to diffuse outward, and so there will be an additional diffusion time delay."940 This topic was considered initially bv Pedersenctal.(1982) in analyzing bursts from IU 1636536., This topic was considered initially by \citet{Pedersen:1982a} in analyzing bursts from 4U 1636–536.941 They estimated that the diffusion time for a typical photon would be ~0.688. largely negligible compared to the expected light travel time delavs.," They estimated that the diffusion time for a typical photon would be $\sim0.6$ s, largely negligible compared to the expected light travel time delays."942 Coniusky.Loudon.&klein(1987) examined the problem more rigorously. calculating time-dependent responses of a hot stellar atmosphere to au X-rav burst. includius the effect of the burst ou the atmospheric temperature structure aud opacities.," \citet{Cominsky:1987a} examined the problem more rigorously, calculating time-dependent responses of a hot stellar atmosphere to an X-ray burst, including the effect of the burst on the atmospheric temperature structure and opacities."943 Their results were in agreement with those of Pedersenetal.(1982). and they found that the reprocessed elit is expected within just ss. but that there was also a very extended tail to the response up to ss. Finally. MeCGosvanetal.(2003) also considered this issue. but neglected ploto-absorption. which will be critical for the soft N-ravs produced m a must.," Their results were in agreement with those of \citet{Pedersen:1982a}, and they found that the reprocessed light is expected within just s, but that there was also a very extended tail to the response up to s. Finally, \citet{McGowan:2003a} also considered this issue, but neglected photo-absorption, which will be critical for the soft X-rays produced in a burst."944 Dased ou these calculations. we would not expect our method to be significantly affected. by reprocessing lcs. as a (απορία respouse approximation will be dominated bv when most of the Πο emerges. aud lis occurs within a fraction of a second.," Based on these calculations, we would not expect our method to be significantly affected by reprocessing times, as a Gaussian response approximation will be dominated by when most of the light emerges, and this occurs within a fraction of a second."945 Tadeed. our observations are broadly consistent with expectations roni light travel time delavs and. allowing for the nuitatious of the Caussian approxiuatiou. cdo not obviously require additional local delays.," Indeed, our observations are broadly consistent with expectations from light travel time delays and, allowing for the limitations of the Gaussian approximation, do not obviously require additional local delays."946 This conclusion may be different from that obtained or correlated flickering i the low-mass X-ray binarics Sco N-l aud LMC οι for which AleGowanctal.(2003) arene that lags iav be too long for light ravel times alone.," This conclusion may be different from that obtained for correlated flickering in the low-mass X-ray binaries Sco X-1 and LMC X-2, for which \citet{McGowan:2003a} argue that lags may be too long for light travel times alone."947 More aud better data are needed o confirm this discrepancy. but if confuniued the key difference may be in the nature and spectra of the N-rav irradiation.," More and better data are needed to confirm this discrepancy, but if confirmed the key difference may be in the nature and spectrum of the X-ray irradiation."948 C'onüuskw.Loudon.&Elem(1987) also cousidered harder irradiation than that provided by su N-ray burst. and found that it could increase," \citet{Cominsky:1987a} also considered harder irradiation than that provided by an X-ray burst, and found that it could increase"949Cataclysmic variables (CVs) are interacting binaries composed of a low-mass secondary star which ts transferring material to a white dwarf (WD) primary star.,Cataclysmic variables (CVs) are interacting binaries composed of a low-mass secondary star which is transferring material to a white dwarf (WD) primary star.950 In most cases the mass donor is unevolved and the mass transfer occurs via an aecretion disc surrounding the WD., In most cases the mass donor is unevolved and the mass transfer occurs via an accretion disc surrounding the WD.951 Detailed reviews of the properties of CVs have been given by Warner(1995) and Hellier(2001).., Detailed reviews of the properties of CVs have been given by \citet{Warner95book} and \citet{Hellier01book}.952 We are undertaking a project to characterise the populatior of CVs identified by the Sloan Digital SkySurvey (SDSS) (Gánsickeetal.2009;Southworth2007a.b.2009.2010:Dillonetal. 2008)..," We are undertaking a project to characterise the population of CVs identified by the Sloan Digital SkySurvey (SDSS) \citep{Gansicke+09mn, Me+07mn, Me+07mn2, Me+09aa, Me+10aa, Dillon+08mn}."953 In the course of this work we obtainec time-resolved spectroscopy and photometry of the faint system JJ003941.064-005427.5(hereafter 10039) which shows strong Ha emission with a unique variable triple-peakec structure., In the course of this work we obtained time-resolved spectroscopy and photometry of the faint system $+$ 005427.5(hereafter J0039) which shows strong $\alpha$ emission with a unique variable triple-peaked structure.954 Doppler maps of the spectroscopic emission lines suggest that notonly is the accretion dise. markedly non-circular. but that there is an inner emission feature which defies explanation in the standard framework of the structure of CVs.," Doppler maps of the spectroscopic emission lines suggest that notonly is the accretion disc markedly non-circular, but that there is an inner emission feature which defies explanation in the standard framework of the structure of CVs."955 JJ0039 was discovered to be a CV by Szkodyetal. on the basis of an SDSS spectrum which shows Balmer and Hell emission lines., J0039 was discovered to be a CV by \citet{Szkody+05aj} on the basis of an SDSS spectrum which shows Balmer and I emission lines.956 Very broad absorption is visible around the higher-order Balmer lines. indicating that the WD emits a large fraction of the overall light of the system.," Very broad absorption is visible around the higher-order Balmer lines, indicating that the WD emits a large fraction of the overall light of the system."957 Spectroscopic observations were carried out in 2007 August using the FORS2 spectrograph (Appenzelleretal.1998) at the Very Large Telescope (VLT). Chile reftab:obslog)).," Spectroscopic observations were carried out in 2007 August using the FORS2 spectrograph \citep{Appenzeller+98msngr} at the Very Large Telescope (VLT), Chile \\ref{tab:obslog}) )."958 The .1200R grism was used. giving a wavelength coverage of tto wwith a reciprocal dispersion of aand a resolution of aat Ha with a slit width of aand the CCD binned 2x2.," The 1200R grism was used, giving a wavelength coverage of to with a reciprocal dispersion of and a resolution of at $\alpha$ with a slit width of and the CCD binned $\times$ 2."959 The data were reduced using optimal extraction 1986) as implemented in the code (Marsh1989).. which also makes use of the packages andKAPPA.," The data were reduced using optimal extraction \citep{Horne86pasp} as implemented in the code \citep{Marsh89pasp}, which also makes use of the packages and."960 The wavelength calibration of the spectra was performed using one are lamp exposure for each night., The wavelength calibration of the spectra was performed using one arc lamp exposure for each night.961 Wavelength shifts due to spectrograph flexure were measured and accurately removed using the nnight sky emission line (seeSouthworthetal.2006.2008a.b).," Wavelength shifts due to spectrograph flexure were measured and accurately removed using the night sky emission line \citep[see][]{Me+06mn,Me+08mn,Me++08mn}."962. The observing procedure of FORS2 included obtaining target acquisition images from which photometry can be obtained., The observing procedure of FORS2 included obtaining target acquisition images from which photometry can be obtained.963 The first of these images was taken unfiltered whereas the second Was obtained with a V filter., The first of these images was taken unfiltered whereas the second was obtained with a $V$ filter.964 We extracted differential photometry from these images using the package GAIA., We extracted differential photometry from these images using the package .965 The V-band apparent magnitudes of the comparison stars were calculated from their g and 7 magnitudes using the transformations provided by Jesteretal.(2005)., The $V$ -band apparent magnitudes of the comparison stars were calculated from their $g$ and $r$ magnitudes using the transformations provided by \citet{Jester+05aj}.966 Time-series photometry was obtained during 2007 August using the New Technology Telescope (NTT) at ESO La Silla. Chile. and the SUSI2 CCD mosaic imager (D'Odorico 1998).. CC," Time-series photometry was obtained during 2007 August using the New Technology Telescope (NTT) at ESO La Silla, Chile, and the SUSI2 CCD mosaic imager \citep{Dodorico+98spie}."967D#445 was used. binned by a factor of three in both directions to give a plate scale of -ppx7!.," 45 was used, binned by a factor of three in both directions to give a plate scale of $^{-1}$."968 The observations were performed in white light in order to maximise throughput. and care was taken to place the target and comparison stars on parts of the CCD which were least affected by fringing.," The observations were performed in white light in order to maximise throughput, and care was taken to place the target and comparison stars on parts of the CCD which were least affected by fringing."969 Exposure times of ssgave an observing cadence of ss. Debiasing and flat-fielding of the raw images was performed with the software packages and KAPPA., Exposure times of sgave an observing cadence of s. Debiasing and flat-fielding of the raw images was performed with the software packages and .970 Optimal and aperture photometry (Naylor1998). was obtained from the reduced images with the script (Southworthetal. 2004).. which uses the package (Eatonetal. 1999)..," Optimal and aperture photometry \citep{Naylor98mn} was obtained from the reduced images with the script \citep{Me++04mn}, , which uses the package \citep{Eaton++99}. ."971 The differential photometry was, The differential photometry was972"can determine the curvature parameter and the matter density far better than the SNe technique, whereas the latter needs strong priors onQx and wy, in order to place tight constraints on wo and wa (Linder2005;Knoxetal.2006b).","can determine the curvature parameter and the matter density far better than the SNe technique, whereas the latter needs strong priors on$\Omega_{\rm K}$ and $\omega_{\rm m}$ in order to place tight constraints on $w_0$ and $w_{\rm a}$ \citep{linder05b,knox06c}."973". Another example is that the WL technique is sensitive to the priors on the parameters, which can be calibrated by the BAO technique."," Another example is that the WL technique is sensitive to the priors on the parameters, which can be calibrated by the BAO technique."974" By comparing the results in with the LSST BAO+WL result in4,, one sees that the ideal 10,000 deg? survey can place comparable constraints on the dark energy EOS parameters to the LSST survey."," By comparing the results in with the LSST BAO+WL result in, one sees that the ideal 10,000 $^2$ survey can place comparable constraints on the dark energy EOS parameters to the LSST survey."975" We show the dark energy constraints from combinations of LSST with the KDUST 10,000 deg? survey and LSST with half of the KDUST survey (labeled as KDUST5k) in4."," We show the dark energy constraints from combinations of LSST with the KDUST 10,000 $^2$ survey and LSST with half of the KDUST survey (labeled as KDUST5k) in."976. The KDUST5k survey could be carried out by a 2.5m KDUST pathfinder in 12 years., The KDUST5k survey could be carried out by a 2.5m KDUST pathfinder in 12 years.977" We assume that 5000 deg? high-resolution imaging is enough to reach the systematic calibration floor, so the improvement to the LSST survey outside the KDUST survey area is kept the same for both proposals."," We assume that 5000 $^2$ high-resolution imaging is enough to reach the systematic calibration floor, so the improvement to the LSST survey outside the KDUST survey area is kept the same for both proposals."978" Under this condition, the WL+BAO constraints on wo and w, from LSST+KDUST5kare nearly the same as those from LSST--KDUST; they improve the alone dark energy task force (DETF,Albrechtetal.2006) figure of merit (FOM) by30%."," Under this condition, the WL+BAO constraints on $w_0$ and $w_a$ from LSST+KDUST5kare nearly the same as those from LSST+KDUST; they improve the LSST-alone dark energy task force \citep[DETF,][]{albrecht06b}979 figure of merit (FOM) by."980". Adding SNAP-like SN data from KDUST as well can significantly increase the DETF FOM; theresult is as good as LSST or LSST+KDUST joint BAO and WL constraints on wo and uw, in the absence of systematics.", Adding SNAP-like SN data from KDUST as well can significantly increase the DETF FOM; theresult is as good as LSST or LSST+KDUST joint BAO and WL constraints on $w_0$ and $w_a$ in the absence of systematics.981" The current constraint on (wo,Wa} is (Zhao&Zhang 2010),,—0."," The current constraint on $\{w_0,~w_a\}$ is \citep{zhao09}, ,."982"11—0.22» (4)derived from a joint Markov Chain Monte Carlo 260-98(MCMC) analysis of the recently released SNe “Constitution” sample (Hickenetal2009),, the WMAP five-year data (Komatsuetal2009)°,, and the Sloan Digital Sky Survey Luminous Red Galaxy (LRG) sample (Tegmark(SDSS)etal2006).."," derived from a joint Markov Chain Monte Carlo (MCMC) analysis of the recently released SNe “Constitution"" sample \citep{Hicken:2009dk}, the WMAP five-year data \citep{WMAP5}~, and the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxy (LRG) sample \citep{Tegmark:2006az}."983" In this analysis, the dark energy perturbation (DEP), which is important in the parameter estimation (Zhaoetal.2005;Fang,Hu&Lewis was consistently included based on the treatment 2008),,developed in (Zhaoetal.2005).."," In this analysis, the dark energy perturbation (DEP), which is important in the parameter estimation \citep{Zhao:2005vj,Fang:2008sn}, was consistently included based on the treatment developed in \citep{Zhao:2005vj}."984" The central values indicate that the ‘quintom’ scenario (Feng,Wang&Zhang is mildly favored, namely, the EOS today w(z)|--02005)=wo>—1, while EOS in the far past w(z)|z200=Wo+Wa«—1."," The central values indicate that the `quintom' scenario \citep{quintom} is mildly favored, namely, the EOS today $w(z)|_{z=0}=w_0>-1$, while EOS in the far past $w(z)|_{z=\infty}=w_0+w_a<-1$."985" This is consistent with the recent published result using the ‘Constitution’ SNe sample (Shafieloo,Sahni&Starobinsky2009;BiswasHuangetal. 2009)."," This is consistent with the recent published result using the `Constitution' SNe sample \citep{cfa1,cfa2,cfa3,cfa4,cfa5}."986". The error bars in can be used to estimate the FOM of current data, and we find that KDUST+LSST can improve the current FOM by two orders of magnitude."," The error bars in can be used to estimate the FOM of current data, and we find that KDUST+LSST can improve the current FOM by two orders of magnitude."987" To investigate the constraints on dark energy equation-of-state w(z) from the surveys in a model-independent way, we follow (Huterer&Starkman2002;Critten-den,Pogosian&Zhao2009) and employ a Principle Component Analysis (PCA) approach."," To investigate the constraints on dark energy equation-of-state $w(z)$ from the surveys in a model-independent way, we follow \citep{Huterer:2002hy,Crittenden:2005wj} and employ a Principle Component Analysis (PCA) approach."988" We choose 40 uniform redshift bins, stretching to a maximum redshift of z=3, and allow the high redshift (z>3) equation of state to vary."," We choose 40 uniform redshift bins, stretching to a maximum redshift of $z=3$, and allow the high redshift $z>3$ ) equation of state to vary."989" To avoid w(z) with infinite derivatives, each bin rises and falls following a hyperbolic tangent function with a typical transition width dz of order of the width of a bin."," To avoid $w(z)$ with infinite derivatives, each bin rises and falls following a hyperbolic tangent function with a typical transition width $dz$ of order of the width of a bin."990" We choose a constant w=—1.0 as the fiducial model, which is consistent with all the present data."," We choose a constant $w = -1.0$ as the fiducial model, which is consistent with all the present data."991" Since dark energy perturbations play a crucial role in the parameter estimation, we use a modified version of CAMB which allows us to calculate DEP for an arbitrary w(z) consistently (Zhaoetal.2005).."," Since dark energy perturbations play a crucial role in the parameter estimation, we use a modified version of CAMB which allows us to calculate DEP for an arbitrary $w(z)$ consistently \citep{Zhao:2005vj}."992" For the PCA, we calculate the Fisher matrices based on four kinds of observables: supernovae, CMB anisotropies, galaxy number counts (GC) correlationfunctions and WL observations."," For the PCA, we calculate the Fisher matrices based on four kinds of observables: supernovae, CMB anisotropies, galaxy number counts (GC) correlationfunctions and WL observations."993" We also include all the possible cross-correlations among these, including CMBxgalaxy and CMBxWL, which are sensitive to the integrated Sachs-Wolfe effect, as well as galaxy-weak lensing correlations."," We also include all the possible cross-correlations among these, including $\times$ galaxy and $\times$ WL, which are sensitive to the integrated Sachs-Wolfe effect, as well as galaxy-weak lensing correlations."994" We first calculate the Fisher matrices for each of the observables, F7, where the indices i,j run over the parameters of the theory, in our case the binned w;(z)."," We first calculate the Fisher matrices for each of the observables, $F^{a}_{ij}$, where the indices $i,j$ run over the parameters of the theory, in our case the binned $w_i(z)$."995" We then find the normalized eigenvectors and eigenvalues of this matrix {e;(z),Ai}, and write where the rows of W are the eigenvectors and A is a diagonal matrix with elements A;."," We then find the normalized eigenvectors and eigenvalues of this matrix $\{e_i(z),\lambda_i\}$, and write where the rows of $W$ are the eigenvectors and $\Lambda$ is a diagonal matrix with elements $\lambda_i$."996 The Fisher matrix is an estimate of the inverse covariance matrix we expect the data to give us and the eigenvalues reflect how well the amplitude of each eigenvector can be measured., The Fisher matrix is an estimate of the inverse covariance matrix we expect the data to give us and the eigenvalues reflect how well the amplitude of each eigenvector can be measured.997 The true behavior of the equation of state can be expanded in the eigenvectors as and the expected error in the recovered amplitudes is given by σ(αι)=Aj., The true behavior of the equation of state can be expanded in the eigenvectors as and the expected error in the recovered amplitudes is given by $\sigma(\alpha_i) = \lambda_i^{-1/2}$.998" For our forecasts, we assume three surveys: the ideal 10k survey, LSST and KDUST--LSST as listed in1."," For our forecasts, we assume three surveys: the ideal 10k survey, LSST and KDUST+LSST as listed in."999. We also combine the Planck survey for CMB and SNAP SNe., We also combine the Planck survey for CMB and SNAP SNe.1000 We assume flat universe and marginalize over the intrinsic SN magnitudea M and the galaxy bias parameters., We assume a flat universe and marginalize over the intrinsic SN magnitude $M$ and the galaxy bias parameters.1001" In the upper panel of5,, we show the spectra of eigenvalues σ΄ of the Fisher matrices for these three tomographic (αι)surveys, and in the lower panel, we show the ratio of the eigenvalues for LSST+KDUST to that for LSST."," In the upper panel of, we show the spectra of eigenvalues $\sigma^{-2}(\alpha_i)$ of the Fisher matrices for these three tomographic surveys, and in the lower panel, we show the ratio of the eigenvalues for LSST+KDUST to that for LSST."1002" As we can see, the ideal 10k survey is competitive to LSST on the dark energy constraints, and adding KDUST to LSST can improve the constraints on the first few eigenmodes by as much as 85%."," As we can see, the ideal 10k survey is competitive to LSST on the dark energy constraints, and adding KDUST to LSST can improve the constraints on the first few eigenmodes by as much as ."1003. The first five best determined eigenvectors are shown in 6., The first five best determined eigenvectors are shown in .1004". As we can see, the Nth eigenmode has N—1 nodes in z, and the ‘sweet spot! — the redshift where the uncertainty of w(z) gets minimized— is at z~ 0.2."," As we can see, the $N$ th eigenmode has $N-1$ nodes in $z$ , and the `sweet spot' – the redshift where the uncertainty of $w(z)$ gets minimized– is at $z\sim0.2$ ."1005 This is consistent with the analysis done in (Huterer&Starkman2002; 2009)..," This is consistent with the analysis done in \citep{Huterer:2002hy,Crittenden:2005wj}. ."1006 We also find that adding KDUST to LSST makes the eigenmodes stretch, We also find that adding KDUST to LSST makes the eigenmodes stretch1007on our data to confirm that there were no irregularities in our data or methods.,on our data to confirm that there were no irregularities in our data or methods.1008" We find that nearly all our measured parameters agree with those published elsewhere (o withinerrors"".", We find that nearly all our measured parameters agree with those published elsewhere to within.1009 The one major exception to this is PSR J1701—3006D. (hereafter M62D). where we see a highlv significant change in orbital period and dispersion measure (DM) compared with the results of 7..," The one major exception to this is PSR $-$ 3006B (hereafter M62B), where we see a highly significant change in orbital period and dispersion measure (DM) compared with the results of \citet{pdm+03}."1010" We have also measured the rate of change of the orbital period P,—5.51(62)x10.2,", We have also measured the rate of change of the orbital period $\dot{P}\rmsub{b} = -5.51(62) \times 10^{-12}$.1011" We performed an F-test to determine if the addition of P4, was indeed required by the data.", We performed an F-test to determine if the addition of $\dot{P}\rmsub{b}$ was indeed required by the data.1012" Without it. V?=142.82 with 71 degrees of [reedom. while alter fitting for 2, the V improved to 73.94 with 70 degrees of freedom."," Without it, $\chi^2 = 142.82$ with 71 degrees of freedom, while after fitting for $\dot{P}\rmsub{b}$ the $\chi^2$ improved to $73.94$ with 70 degrees of freedom."1013 The probability (hat (his improvenient is due to chance is <1.32x10.11στι so the measured EN does indeed seem (o be required by the data.," The probability that this improvement is due to chance is $< 1.3 \times101410^{-11}$, so the measured $\dot{P}\rmsub{b}$ does indeed seem to be required by the data."1015 M62D is known to eclipse and has an optical aud. counterpart (?).., M62B is known to eclipse and has an optical and X-ray counterpart \citep{cfp+08}.1016 As 7— discuss. the pulsar and companion star are almost certainlv interacting.," As \citeauthor{cfp+08} discuss, the pulsar and companion star are almost certainly interacting."1017" The contribution to 2, expected from general relativity (GR) is two orders of magnitude smaller than observed. assuming Mj=1.4M. and M,=Mj. and there is no realistic combination of pulsar ancl companion masses and inclination angles that would lead to such a large relativistic 3."," The contribution to $\dot{P}\rmsub{b}$ expected from general relativity (GR) is two orders of magnitude smaller than observed, assuming $M\rmsub{p} = 1.4\; \Msun$ and $M\rmsub{c} = M\rmsub{c,min}$, and there is no realistic combination of pulsar and companion masses and inclination angles that would lead to such a large relativistic $\dot{P}\rmsub{b}$."1018 As such. classic tidal elfects of the extended companion are probably the cause of the change in orbital period.," As such, classic tidal effects of the extended companion are probably the cause of the change in orbital period."1019" Our RAIS timing residuals are over a [actor of (wo smaller than obtained by 2.. which probably explains why we were able to detect J, even though we had a shorter timing baseline."," Our RMS timing residuals are over a factor of two smaller than obtained by \citet{pdm+03}, , which probably explains why we were able to detect $\dot{P}\rmsub{b}$ even though we had a shorter timing baseline."1020 PSR J1T01—3006D (hereafter M62D) is a 3.42ms binary MSP., PSR $-$ 3006D (hereafter M62D) is a $3.42\; \ms$ binary MSP.1021 M62D (alone with E and E) were discovered. and initial orbital solutions were given. by ?..," M62D (along with E and F) were discovered, and initial orbital solutions were given, by \citet{cha03}."1022 The orbital period is L1days. with a smalleccentriciivt!.. e~£12x10.!.," The orbital period is $1.1\;1023\mathrm{days}$, with a small, $e \sim 4.12 \times102410^{-4}$."1025 The minimum companion mass is ~0.12M. (assuming a 1.4M. AISP). ancl is likely a white dwarf.," The minimum companion mass is $\sim 0.12\; \Msun$ (assuming a $1.4\; \Msun$ MSP), and is likely a white dwarf."1026 We see no evidence lor eclipses in M62D. but none of our observations cover conjunction. when an eclipse would be most likely.," We see no evidence for eclipses in M62D, but none of our observations cover conjunction, when an eclipse would be most likely."1027 Three observations do. however. start or end within four hours (e154 of the orbital period) of conjunction. and one observation ends ouly 15 minutes before conjunction.," Three observations do, however, start or end within four hours $\sim 15\%$ of the orbital period) of conjunction, and one observation ends only 15 minutes before conjunction."1028 The lack of eclipses increase our confidence that the companion is a white dwarl. and not a main sequence (MS) star (lor which eclipses should be common).," The lack of eclipses increase our confidence that the companion is a white dwarf, and not a main sequence (MS) star (for which eclipses should be common)."1029 We are unable to measure anv precession in the longitude of periastron. so no further constraints can be placed on the nass or geometryof the svstem at this time.," We are unable to measure any precession in the longitude of periastron, so no further constraints can be placed on the mass or geometryof the system at this time."1030 The acceleration of M62D (PP. 1) is somewhat, The acceleration of M62D $\dot{P} P^{-1}$ ) is somewhat1031the LG frame beyond MMpc.,the LG frame beyond Mpc.1032 The black dotted line in Fig. 9((, The black dotted line in Fig. \ref{fig:compare-lf-magtype}( (1033"b) represents the {-ραπά GLF calculated using the method and sample of BGD08 (DR4) with the NYU-VAGC distances, while for the black dashed line the distances were changed to those derived from the Tonryetal.(2000) model.","b) represents the $i$ -band GLF calculated using the method and sample of BGD08 (DR4) with the NYU-VAGC distances, while for the black dashed line the distances were changed to those derived from the \citet{tonry00} model."1034" The latter model gives on average10%,, and up to30%,, larger distances at z<0.01."," The latter model gives on average, and up to, larger distances at $z<0.01$."1035 The DR4 result with the adjusted distances is in better agreement with the GAMA result., The DR4 result with the adjusted distances is in better agreement with the GAMA result.1036 Note that GAMA galaxies with luminosities ~10°Lo have a median redshift of 0.02 compared to 0.006 for the NYU-VAGC sample., Note that GAMA galaxies with luminosities $\simeq10^{8}\Lsun$ have a median redshift of 0.02 compared to 0.006 for the NYU-VAGC sample.1037 Thus the GAMA result is less sensitive to the flow model at these luminosities., Thus the GAMA result is less sensitive to the flow model at these luminosities.1038 See Lovedayetal.(2012) for more details on the GAMA GLFs., See \citet{loveday12} for more details on the GAMA GLFs.1039 Figure 10 shows the GAMA ;-band GLF with error bars in comparison with the GLF from semi-analytical model., Figure \ref{fig:compare-model} shows the GAMA $i$ -band GLF with error bars in comparison with the GLF from a semi-analytical model.1040 The latter was derived by H. Kim C. Baugha (private communication) using an implementation of similar to Boweretal.(2006)., The latter was derived by H. Kim C. Baugh (private communication) using an implementation of similar to \citet{bower06}.1041 The mass resolution of the halo merger trees was improved and the photo-ionisation prescription was changed so that cooling in haloes with a circular velocity below 30kms! (previously 50km s!) is prevented after reionisation (z— 6)., The mass resolution of the halo merger trees was improved and the photo-ionisation prescription was changed so that cooling in haloes with a circular velocity below $30\kms$ (previously $50\kms$ ) is prevented after reionisation $z=6$ ).1042" There is reasonable agreement between the model and data, however, the model LF is higher particular below 109?Lo."," There is reasonable agreement between the model and data, however, the model LF is higher particular below $10^{8.2}\Lsun$."1043" At low luminosities, it is expected that the GAMA data are incomplete because of surface brightness issues and the LF data points are shown as lower limits In addition to the explicit magnitude limit, there is an implicit and imprecisely-defined surface brightness (SB) limit that plagues measurements of the faint end of a GLF (Phillipps&Disney1986;Cross&Dr"," At low luminosities, it is expected that the GAMA data are incomplete because of surface brightness issues and the LF data points are shown as lower limits In addition to the explicit magnitude limit, there is an implicit and imprecisely-defined surface brightness (SB) limit that plagues measurements of the faint end of a GLF \citep{PD86,CD02}. ."1044"iver 2002).. Blantonetal.(2005a) estimated the impact on the SDSS GLF, determining a completeness of about 0.7 at µγ,δο=23.0magarcsec? where this is the SB within the Petrosian half-light radius."," \citet{blanton05} estimated the impact on the SDSS GLF, determining a completeness of about 0.7 at $\effsb=23.0{\rm\,mag\,arcsec}^{-2}$ where this is the SB within the Petrosian half-light radius."1045" Three sources of incompleteness were considered: photometric incompleteness determined from simulations that put fake galaxies in frames run throughPHOTO,, tiling incompleteness because some of the SDSS area was targeted on versions of where the deblender was not performing optimally, and spectroscopic incompleteness."," Three sources of incompleteness were considered: photometric incompleteness determined from simulations that put fake galaxies in frames run through, tiling incompleteness because some of the SDSS area was targeted on versions of where the deblender was not performing optimally, and spectroscopic incompleteness."1046" The tiling incompleteness is not an issue here, the issues associated with the photometric incompleteness may be less severe at the GAMA faint limit because for a given SB the galaxies are smaller, meaning fewer problems with deblender shredding and sky subtraction, and the spectroscopic incompleteness can be mitigated by repeated observations of the same target where necessary."," The tiling incompleteness is not an issue here, the issues associated with the photometric incompleteness may be less severe at the GAMA faint limit because for a given SB the galaxies are smaller, meaning fewer problems with deblender shredding and sky subtraction, and the spectroscopic incompleteness can be mitigated by repeated observations of the same target where necessary."1047" Figure 11 shows the SB versus stellar mass distribution (with masses from the colour-based M/L relation of Tayloretal.2011,, see 3.3))."," Figure \ref{fig:mass-sb} shows the SB versus stellar mass distribution (with masses from the colour-based M/L relation of \citealt{taylor11}, see \ref{sec:gsmf}) )."1048 It is difficult to determine when the input catalogue becomes incomplete., It is difficult to determine when the input catalogue becomes incomplete.1049" Judging from the slightly higher meanSB at 105-10?M in the GAMA sample compared to SDSS, we expect"," Judging from the slightly higher meanSB at $10^{8}$ $10^{9}\Msun$ in the GAMA sample compared to SDSS, we expect"1050to a strong background of photons in the Lyman-Werner range which can dissociate molecular hydrogen.,to a strong background of photons in the Lyman-Werner range which can dissociate molecular hydrogen.1051 In each case. the dependence. of characteristic stellar. mass on redshift Hlattens for haloes more massive than about LOAL. (see equ. (," In each case, the dependence of characteristic stellar mass on redshift flattens for haloes more massive than about $10^{8} M_{\odot}$ (see equ. ("10525)).,5)).1053 The reason for this change is that for lower mass haloes gas is compressed nearly aciabatically to the virial state: as the redshift’ of virialisation declines. the halo mass increases steeply (Figure 4) and the strongly increasing pressure manifests itself as a marked. decline in the characteristic stellar mass with decreasing redshift.," The reason for this change is that for lower mass haloes gas is compressed nearly adiabatically to the virial state; as the redshift of virialisation declines, the halo mass increases steeply (Figure 4) and the strongly increasing pressure manifests itself as a marked decline in the characteristic stellar mass with decreasing redshift."1054 For haloes more massive than about 107Αι. however. the role of angular momentum in halting the collapse weakens he mass dependence of the virial pressure (equ. (," For haloes more massive than about $10^8 M_{\odot}$, however, the role of angular momentum in halting the collapse weakens the mass dependence of the virial pressure (equ. ("10556)). and he characteristic stellar mass declines more gently with decreasing redshift.,"6)), and the characteristic stellar mass declines more gently with decreasing redshift."1056 Phe ordering by mass of the curves for various overdensities (1.0. higher charactersitic masses. [or ower σ [uctuations) is simply a result of the lower pressure attained. at a given redshift. in ie smaller mass haloes ound in lower e [luctuations.," The ordering by mass of the curves for various overdensities (i.e. higher charactersitic masses for lower $\sigma$ fluctuations) is simply a result of the lower pressure attained, at a given redshift, in the smaller mass haloes found in lower $\sigma$ fluctuations."1057 The dashed. lines correspond to the case that the gas can cool to the maximum of the ambient CALB temperature and a temperature of LOW. this latter representing a fiducial value for cooling by CO and dust in Population L regions not subject to strong radiative heating by massive stars.," The dashed lines correspond to the case that the gas can cool to the maximum of the ambient CMB temperature and a temperature of $10$ K, this latter representing a fiducial value for cooling by CO and dust in Population I regions not subject to strong radiative heating by massive stars."1058 These lines are all terminated at the high redshift end at the point where the predicted: characteristic stellar mass is comparable to the total barvonic mass of the halo concerned. and are evidently only applicable well to the right of these points.," These lines are all terminated at the high redshift end at the point where the predicted characteristic stellar mass is comparable to the total baryonic mass of the halo concerned, and are evidently only applicable well to the right of these points."1059 Unlike the case of molecular hydrogen cooling described above. the temperature along these curves js set o»w the CAB for all z ," Unlike the case of molecular hydrogen cooling described above, the temperature along these curves is set by the CMB for all $z > 2.7$."1060The declino in temperature towards the present epoch accounts for the steeper decline in characteristic stellar mass with decreasing redshift than in the {ο cooling case. although again there is à mild. [lattening of the curves for haloes more massive un 107A4..," The decline in temperature towards the present epoch accounts for the steeper decline in characteristic stellar mass with decreasing redshift than in the $H_2$ cooling case, although again there is a mild flattening of the curves for haloes more massive than $10^8 M_\odot$."1061 At low redshift (2< 2.7). the temperature utains a plateau value of LOW. At these low redshifts. the 2ürial pressure with decreasing redshift. because —10 explicit redshift’ dependence in equation. (6) is more significant than the increase of halo mass with decreasing," At low redshift $z < 2.7$ ), the temperature attains a plateau value of $10$ K. At these low redshifts, the virial pressure with decreasing redshift, because the explicit redshift dependence in equation (6) is more significant than the increase of halo mass with decreasing"1062v choosing these new free functions ancl kernels such that they minimise the integration error. we will arrive at à new scheme that can. with sullicient resolution. correctly resolve multiphase Γιά flow.,"By choosing these new free functions and kernels such that they minimise the integration error, we will arrive at a new scheme that can, with sufficient resolution, correctly resolve multiphase fluid flow."1063 In the previous section. we presented. a standard derivation of the SPL equations of motion.," In the previous section, we presented a standard derivation of the SPH equations of motion."1064 However. this standard derivation leads to a scheme that cannot correctly model [uid mixing processes (see rpe[fsecuntroduction)).," However, this standard derivation leads to a scheme that cannot correctly model fluid mixing processes (see \\ref{sec:introduction}) )."1065 ln this section. we move to a more general derivation of the SPLL equations of motion.," In this section, we move to a more general derivation of the SPH equations of motion."1066 We show that. in general. we have a [ree function for each of the Euler equations: jj. 6 and &. as well as a cillerent smoothing kernel for cach.," We show that, in general, we have a free function for each of the Euler equations: $\phir$, $\phiv$ and $\phiu$, as well as a different smoothing kernel for each."1067 There is also a freedom in the energy equation in the choice of integration variable (energy or entropy: re[sec:engent ))., There is also a freedom in the energy equation in the choice of integration variable (energy or entropy; \\ref{sec:engent}) ).1068 In re[sec:error and refsee:ninimise.. &we will then use an error ancl stability analysis of thesemore general equations of motion to constrain the new functions jg.ó and ¢ and our new kernels.," In \\ref{sec:error} and \\ref{sec:minimise}, we will then use an error and stability analysis of thesemore general equations of motion to constrain the new functions $\phir, \phiv$ and $\phiu$ and our new kernels."1069 v choosing these new free functions and kernels such that they minimise the integration error. we will arrive at à new scheme that can. with sullicient resolution. correctly resolve multiphase Γιά flow.," By choosing these new free functions and kernels such that they minimise the integration error, we will arrive at a new scheme that can, with sufficient resolution, correctly resolve multiphase fluid flow."1070 In general. we have some freedom in how we cdiscretise the Euler equations (equations 4--6)) to obtain the equations of motion for SPILL (seec.g.?2??)..," In general, we have some freedom in how we discretise the Euler equations (equations \ref{eqn:euler1}- \ref{eqn:euler3}) ) to obtain the equations of motion for SPH \citep[see e.g.][]{1992ARA&A..30..543M,2005astro.ph..7472P,2009NewAR..53...78R}."1071 Phe gradients in the Euler equations can be expanded to include a new free function for cach equation: η. ó and C: In the continuum form. above. jg.ó and & cancel.," The gradients in the Euler equations can be expanded to include a new free function for each equation: $\phir$, $\phiv$ and $\phiu$: In the continuum form, above, $\phir, \phiv$ and $\phiu$ cancel."1072 But in the discrete. SPILL form. they remain giving a useful additional freedom (7): where H;;=EHIrol.hi)|Lp.)al/2. N;; and L; are svmimetrised. smoothing. kernels one for each Euler equation.," But in the discrete SPH form, they remain giving a useful additional freedom \citep{2005astro.ph..7472P}: where $\hij = \left[\hi + \hj \right]/2$, $\kij$ and $\lij$ are symmetrised smoothing kernels – one for each Euler equation."1073 Standard SPIEL (PLI from here on) is a special case of the above with y=ó—6= Land Hi;jri;=ijv;Livi)=NV;Mog.," Standard SPH (SPH from here on) is a special case of the above with $\phir = \phiv = \phiu = 1$ and $\hijtild 1074= \kijtild = \lijtild = \dwijtild$ ."1075 Equation 16 casts the continuity equation in dilferential form., Equation \ref{eqn:cont} casts the continuity equation in differential form.1076 Εις is problematic since. in this case. the particles no longer represent the [uid exactly.," This is problematic since, in this case, the particles no longer represent the fluid exactly."1077 Instead they represent a moving mesh on which the Euler equations are solved., Instead they represent a moving mesh on which the Euler equations are solved.1078 This leads to the danger that high density regions will contain few particles leading to large errors (?).., This leads to the danger that high density regions will contain few particles leading to large errors \citep{2003ApJ...595..564M}.1079 For this reason. we use instead a generalisedintegral form for the continuity equation: which. taking the time derivative. gives: where: and ἡ=zu," For this reason, we use instead a generalisedintegral form for the continuity equation: which, taking the time derivative, gives: where: and $\dphir = \frac{d\eta}{dt}$."1080 This reduces to the continuity equation (equation 16)) ander the kernel ΕΠ ΑΠΡ and for c=0.," This reduces to the continuity equation (equation \ref{eqn:cont}) ) under the kernel constraint: $\hijtild = \unabla_i\wijtild$, and for $\epsilon = 0$."1081 The latter can be satisfied by construction if 5g;=gj (as is the case for SPLD. or if ἡ=0.," The latter can be satisfied by construction if $\phir_i = \phir_j$ (as is the case for SPH), or if $\dphir=0$."1082 However. in the continuum limit CN.oex. hs 0). «0 and so c will vanish with increasing resolution.," However, in the continuum limit $N\rightarrow \infty$, $h\rightarrow 0$ ), $\epsilon \rightarrow 0$ and so $\epsilon$ will vanish with increasing resolution."1083 For this reason. equation 190 gives a valid approximation to the continuity equation for any choice of à. with e simply. contributing an additional error Lern.," For this reason, equation \ref{eqn:sphcontgen} gives a valid approximation to the continuity equation for any choice of $\phir$, with $\epsilon$ simply contributing an additional error term."1084 A final freeclom in the equations motion for SPLI comes from the energy equation., A final freedom in the equations motion for SPH comes from the energy equation.1085 Equation 15 is the standard energy form of SPL. but there is also an entropy form (2?:: 7? 2).," Equation \ref{eqn:energy} is the standard energy form of SPH, but there is also an entropy form \bcite{1991ApJ...378..637G}; \bcite{2002MNRAS.333..649S}) )."1086 Instead of the internal energy. à. we evolve a function Ms) the that is à monotonic function of the entropy. 5 defined by the equation of state: Away [rom shocks and in the absence of thermal sources or sinks. 24; is a constant of motion.," Instead of the internal energy, $u$ , we evolve a function $A(s)$ – the – that is a monotonic function of the entropy $s$ defined by the equation of state: Away from shocks and in the absence of thermal sources or sinks, $A_i$ is a constant of motion."1087 Thus. taking the time derivative of equation 22. and substituting for equation 12.. We POCOVOD: bv construction.," Thus, taking the time derivative of equation \ref{eqn:entropystate} and substituting for equation \ref{eqn:state}, , we recover: by construction."1088 Schemes that obey equation 23. are called consistent., Schemes that obey equation \ref{eqn:thermconst} are called .1089 In practice. we find ου the tests presented. in. this paper that the energy and entropy formsof SPLL give near- results. provided that equation 23. is satisfied (for acdiabatic How).," In practice, we find – for the tests presented in this paper – that the energy and entropy formsof SPH give near-identical results, provided that equation \ref{eqn:thermconst} is satisfied (for adiabatic flow)."1090 We use the thermodynamicallvy: consistent, We use the thermodynamically consistent1091We (hank M. Ixrauss and FL Nicastro for their help in solving the calibration problems. J. MacDowell for interesting discussions and A. kong. D. Schwartz and J. Drake for useful commentis on the manuscript.,"We thank M. Krauss and F. Nicastro for their help in solving the calibration problems, J. MacDowell for interesting discussions and A. Kong, D. Schwartz and J. Drake for useful comments on the manuscript."1092" This research has been supported by CNAA and by NASA contract NAS 8-390713 (CXC),", This research has been supported by CNAA and by NASA contract NAS 8-39073 (CXC).1093We complemented these images with new deep Ks band images taken over the GOODS-S field with the near-IR VLT imager Hawk-I. The latter were taken in the framework of a program designed to search for z~7 galaxies (Castellanoetal.,"We complemented these images with new deep $K_S$ band images taken over the GOODS-S field with the near-IR VLT imager Hawk-I. The latter were taken in the framework of a program designed to search for $z\sim 7$ galaxies \citep{castellano10a,castellano10b}."1094" In the Ks band, the surveyed area covers of the WFC3 ERS area."," In the $K_S$ band, the surveyed area covers of the WFC3 ERS area."1095" Owing to this and after excluding image edges of dubious quality, the available area reduces to ~ 33 arcmin”."," Owing to this and after excluding image edges of dubious quality, the available area reduces to $\sim$ 33 $^2$."1096 The data reduction of the Ks images is analogous to the procedure used for other Hawk-I data (Castellano 2010a)., The data reduction of the $K_S$ images is analogous to the procedure used for other Hawk-I data \citep{castellano10a}.1097". The net exposure time is 25200s, with a lo r.m.s."," The net exposure time is 25200s, with a $1 \sigma$ r.m.s."1098 of 1.26 counts per second in a 1” aperture., of 1.26 counts per second in a 1” aperture.1099" The magnitude limit at So is ~25.5, one magnitude deeper than the previous ISAAC Ks band."," The magnitude limit at $5 \sigma$ is $\sim25.5$, one magnitude deeper than the previous ISAAC $K_S$ band."1100 We finally built a multiwavelength GOODS-ERS catalogue adding the other public images available in the GOODS-S field., We finally built a multiwavelength GOODS-ERS catalogue adding the other public images available in the GOODS-S field.1101" They include the ACS images in the BVIz bands etal. 2004),, the deep UR images from VIMOS (Nonino2009) and the four IRAC bandsat 3.6, 4.5, 5.8, and 8.0µπι."," They include the ACS images in the $BVIz$ bands \citep{giavalisco04}, the deep $UR$ images from VIMOS \citep{nonino09} and the four IRAC bandsat 3.6, 4.5, 5.8, and 8.0."1102". With respect to the data set used to assemble our previous GOODS-MUSIC sample (Grazianetal.2006;Santini 2009), the present GOODS-ERS data set benefits not only from the much deeper IR coverage provided by the new WFC3 and Ks band data, but also from a deeper version of the z band image, which nearly doubles the exposure time of the previous image, a deeper U band image, and a brand new R image."," With respect to the data set used to assemble our previous GOODS-MUSIC sample \citep{grazian06,santini09}, the present GOODS-ERS data set benefits not only from the much deeper IR coverage provided by the new WFC3 and $K_S$ band data, but also from a deeper version of the $z$ band image, which nearly doubles the exposure time of the previous image, a deeper $U$ band image, and a brand new $R$ image."1103" In this data set, we extracted a 14 band multiwavelength catalogue using the H band as a detection image."," In this data set, we extracted a 14 band multiwavelength catalogue using the $H$ band as a detection image."1104" Colours were carefully obtained with the same technique used in the GOODS-MUSIC catalogue, where we adopted the PSF-matching code CONVPHOT (DeSantisetal.2007) to accurately deblend objects in the ground-based and Spitzer images."," Colours were carefully obtained with the same technique used in the GOODS-MUSIC catalogue, where we adopted the PSF-matching code CONVPHOT \citep{desantis07} to accurately deblend objects in the ground-based and Spitzer images."1105" We note that the depth of the H band even exceeds the depth of the bluest bands, resulting in very poor quality photometric information about the faintest H-selected objects."," We note that the depth of the $H$ band even exceeds the depth of the bluest bands, resulting in very poor quality photometric information about the faintest $H$ -selected objects."1106 The catalogue was cross-correlated with existing spectroscopic samples., The catalogue was cross-correlated with existing spectroscopic samples.1107" For sources lacking spectroscopic information, photometric redshifts were computed by fiting the 14 band multiwavelength photometry to the PEGASE 2.0 templates (Fioc&Rocca-Volmerange1997,, see details in Grazianetal. 2006))."," For sources lacking spectroscopic information, photometric redshifts were computed by fitting the 14 band multiwavelength photometry to the PEGASE 2.0 templates \citealt{fioc97}, see details in \citealt{grazian06}) )."1108" The accuracy reached by the photometric redshifts is very high, the absolute scatter |Az|/(1+Zspec) being equal to 0.03, with only of severe outliers (JAz|/(1+Zspec)> 0.5)."," The accuracy reached by the photometric redshifts is very high, the absolute scatter $|\Delta z|/(1+z_{spec})$ being equal to 0.03, with only of severe outliers $|\Delta z|/(1+z_{spec}) > 0.5$ )."1109 The statistical error associated with each photometric redshift was used to evaluate the limiting magnitude at which a reliable GSMF can be computed., The statistical error associated with each photometric redshift was used to evaluate the limiting magnitude at which a reliable GSMF can be computed.1110" We found that a limit H~26 (or, equivalently, Ks~ 25.5) is appropriate to maintain the error in the mass estimate (see next section) to within 0.3 dex and the relative scatter in the photometric redshifts |Az|/(1+z)<0.1 for of objects, and we adopt this in the following."," We found that a limit $H\simeq 26$ (or, equivalently, $K_S\simeq 25.5$ ) is appropriate to maintain the error in the mass estimate (see next section) to within 0.3 dex and the relative scatter in the photometric redshifts $|\Delta z|/(1 + z) < 0.1$ for of objects, and we adopt this in the following."1111" From the analysis of the individual photometric-redshift probability distributions, we can compute the fraction of ""reliable"" candidates."," From the analysis of the individual photometric-redshift probability distributions, we can compute the fraction of ""reliable"" candidates."1112" We considered a candidate to safely lie within a given redshift interval when the integral of its probability distribution curve, normalized to unity, over that interval is larger than90%."," We considered a candidate to safely lie within a given redshift interval when the integral of its probability distribution curve, normalized to unity, over that interval is larger than."1113".Moreover, we accepted a certain level of tolerance in the definition of the redshift range to allow for the uncertainty in photometric redshifts.","Moreover, we accepted a certain level of tolerance in the definition of the redshift range to allow for the uncertainty in photometric redshifts."1114" Following this method, for all Ks«25.5 sources with Zpno;>2, we can exclude a secondary redshift solution at zpjí;«1.5 in of the sources."," Following this method, for all $K_S<25.5$ sources with $z_{phot}>2$, we can exclude a secondary redshift solution at $z_{phot}<1.5$ in of the sources."1115" This fraction increases to when only bright sources (K,« 24) are considered.", This fraction increases to when only bright sources $K_s<24$ ) are considered.1116" We also extracted a Ks band detection catalogue, and verified that all the objects detected in the Ks band are also detected in the H one, which is unsurprising giventhe extraordinary quality of the WFC3 data."," We also extracted a $K_S$ band detection catalogue, and verified that all the objects detected in the $K_S$ band are also detected in the $H$ one, which is unsurprising giventhe extraordinary quality of the WFC3 data."1117" On the basis of these results, we decided to restrict our analysis to the Ks<25.5 sample, albeit obtained from the H-selected one, for two reasons: firstly, this selection allows a more robust comparison with previous K-selected surveys; and secondly, a Ks=25.5 threshold is more efficient in detecting low mass objects than a H=26 one."," On the basis of these results, we decided to restrict our analysis to the $K_S\leq 25.5$ sample, albeit obtained from the $H$ -selected one, for two reasons: firstly, this selection allows a more robust comparison with previous $K$ -selected surveys; and secondly, a $K_S = 25.5$ threshold is more efficient in detecting low mass objects than a $H =26$ one."1118" Adopting this cut, we extend by two magnitudes the previous work of Fontanaetal. (2006), who studied the GSMF of the Ks<23.5 GOODS-MUSIC sample."," Adopting this cut, we extend by two magnitudes the previous work of \cite{fontana06}, who studied the GSMF of the $K_S<23.5$ GOODS-MUSIC sample."1119" Our Ks<25.5 sample here includes 3210 objects, 421 of which have spectroscopic redshifts."," Our $K_S\leq 25.5$ sample here includes 3210 objects, 421 of which have spectroscopic redshifts."1120 We plot in Fig., We plot in Fig.1121 1. the redshift distribution of the GOODS-ERS sample used in this work (black solid histogram) compared to that of the GOODS-MUSIC sample adopted by Fontanaetal.(2006) (red dotted histogram)., \ref{fig:redshift} the redshift distribution of the GOODS-ERS sample used in this work (black solid histogram) compared to that of the GOODS-MUSIC sample adopted by \cite{fontana06} (red dotted histogram).1122" Since the area covered by the ERS survey is relatively small, the sample is more sensitive to overdensities."," Since the area covered by the ERS survey is relatively small, the sample is more sensitive to overdensities."1123" The extended overdensities at z50.7 and z« 1, which cover the entire GOODS-S field (Vanzellaetal. 2005,, Salimbenietal.2009a and references therein), are clearly recognizable."," The extended overdensities at $z\simeq 0.7$ and $z\simeq 1$ , which cover the entire GOODS-S field \citealt{vanzella05}, , \citealt{salimbeni09} and references therein), are clearly recognizable."1124" Unfortunately, the northern part of GOODS-S also includes a cluster at z~1.6 (Castellano and various groups at z«2.2—2.3 (Salimbenietal. 2011),, which both affect the overall redshift distribution."," Unfortunately, the northern part of GOODS-S also includes a cluster at $z\simeq 1.6$ \citep{castellano07} and various groups at $z\simeq 2.2-2.3$ \citep{salimbeni09,yang10,magliocchetti11}, , which both affect the overall redshift distribution."1125pecsalsoshowstheSLlandLL2skyspectra forlhesesources(Le.spectralakenof [ —s,"also shows the SL1 and LL2 sky spectra for these sources (i.e. spectra taken off-source, but with the same PSF)."1126ource.buliciththe: levelroquepicelsthatarenolsuf ficientlystrongoullierslohavebeenremoved," As expected, these sky spectra have a median of essentially zero, but do show some `bumps' and `dips' which in addition to the background noise, likely also include some low-level rogue pixels that are not sufficiently strong outliers to have been removed."1127" Suchspurious features, howet lpe"," Such spurious features, however, could influence our source spectra as well."1128"hetherornolitisstrongerlhantheamplitudeoflhefeaturesinthesbgspeetra. 2)iwhetherornolithasconsis refsec,pld"," Essentially there are three considerations for believing the reality of a given feature: 1) whether or not it is stronger than the amplitude of the features in the sky spectra, 2) whether or not it has consistent redshift, and 3) whether or not it has roughly the expected shape (i.e. width)."1129eplh, We address these issues in \\ref{sec_optdepth}.1130",Beforehandhowever. weevaminebriefIylhehotdusteontinuumshapeoflhesources."," Before hand however, we examine briefly the hot dust continuum shape of the sources."1131which is relatively insensitive to redshift for +>1 (Blain Longair 1993).,which is relatively insensitive to redshift for $z\ga1$ (Blain Longair 1993).1132" Frou, the 850442 measurcuient. the inplied intrinsic FIR hunuinositv for SAIALJJ00266|1708 is L(FIR)~10021."," From the $\mu$ m measurement, the implied intrinsic FIR luminosity for J00266+1708 is $L({\rm FIR}) \simeq113310^{13} L_{\sun}$."1134 This corresponds to a star-forination rate of massive stars (AL>SAL: ) of approximately 500900 37; | (Scoville et 11997: Condon 1992)., This corresponds to a star-formation rate of massive stars $M>5M_{\sun}$ ) of approximately 500--900 $M_{\sun}$ $^{-1}$ (Scoville et 1997; Condon 1992).1135 These results are consistent with estinates for other sub-uuau ealaxies (Ivison et 11998. 2000a).," These results are consistent with estimates for other sub-mm galaxies (Ivison et 1998, 2000a)."1136 It is still not known whether or not an ACN contributes significantly to the bolometric huuinositv of SMIALJJO00266|1708., It is still not known whether or not an AGN contributes significantly to the bolometric luminosity of J00266+1708.1137 Tf an AGN Is present. if is uot a strong radio source and must be heavily obscured at optical waveleugths (Fig.," If an AGN is present, it is not a strong radio source and must be heavily obscured at optical wavelengths (Fig."1138 3)., 3).1139 Dased ou its SED. JJ00266|1708. appears most simular to highh-reddened ULIG/starbursts such as Arp 220 or IIRIO at a redslüft of +—3.5.," Based on its SED, J00266+1708 appears most similar to highly-reddened ULIG/starbursts such as Arp 220 or HR10 at a redshift of $z\sim 3.5$."1140 The relative importance of the sub-nuu and Lyuiur-break galaxies to the elobal star-formation rate at hiehl-redshift is an active area of discussion (Ilughes ot 11998: Ciuiderdoni ot 11998: Blain oet al., The relative importance of the sub-mm and Lyman-break galaxies to the global star-formation rate at high-redshift is an active area of discussion (Hughes et 1998; Guiderdoni et 1998; Blain et al.1141 1999). 1999c: Trentham et al.," 1999b, 1999c; Trentham et al."1142 1999: Peacock et al., 1999; Peacock et al.1143 2000)., 2000).1144 Depending on the exact coutribution of ACN in the sub-nun population. the ultralunüunous οατα. galaxies (7 LoL: οδ μια)> Limmdy) accouut for approximately of the total amount of star-formatiou at lieh- (Blain et 119995).," Depending on the exact contribution of AGN in the sub-mm population, the ultraluminous sub-mm galaxies $>10^{12} L_{\sun}$ ; $S[850\mu$ $]\ga1$ mJy) account for approximately of the total amount of star-formation at high-redshift (Blain et 1999b)."1145 It is truly remarkable that such a high fraction of all starformation at high-redshift occurs iu ULIC/sub-aumni galaxies., It is truly remarkable that such a high fraction of all star-formation at high-redshift occurs in ULIG/sub-mm galaxies.1146 Iu contrast. ULICs O-v contribute about of the total amount of star formation im the local universe. based on the survey of Taam Sauders (1998).," In contrast, ULIGs only contribute about of the total amount of star formation in the local universe, based on the survey of Kim Sanders (1998)."1147 ITence. the evolution in the ameouut Estar formation occurring im ULICs is about LOO times stronecr than the elobal increase of the star-formation rate lieh redshift seeu for all galaxies (Macau ct 11996: Steidel et 11999).," Hence, the evolution in the amount of star formation occurring in ULIGs is about 100 times stronger than the global increase of the star-formation rate at high redshift seen for all galaxies (Madau et 1996; Steidel et 1999)."1148 It has been suggested for both the subi aud Lxauian-weak populations that each represents the formative lases of inassive LO galaxies (Blain ot 11999c: Caavalisco et 11998)., It has been suggested for both the sub-mm and Lyman-break populations that each represents the formative phases of massive $L^{*}$ galaxies (Blain et 1999c; Giavalisco et 1998).1149 Iu this secnario. the sub-ini svstenis may be associated with a very luuinous. short-lived and heavily cust cushrouded starburst. while he Lyian-break galaxies would be associated with a onegecr-lived. less huminous phase of star formation (Blain et 119996).," In this scenario, the sub-mm systems may be associated with a very luminous, short-lived and heavily dust enshrouded starburst, while the Lyman-break galaxies would be associated with a longer-lived, less luminous phase of star formation (Blain et 1999c)."1150 If this scenario is correct. we could expect massive rescrvoirs of molecular gas associated with th. populations.," If this scenario is correct, we could expect massive reservoirs of molecular gas associated with both populations."1151 The detection of massive reservoirs of uolecular gas (Fraver ct 11998. 1999) suggests that he sub-uuu population is associated with gas-rich massive galaxies (> L).," The detection of massive reservoirs of molecular gas (Frayer et 1998, 1999) suggests that the sub-mm population is associated with gas-rich massive galaxies $\ga L^{*}$ )."1152 The molecular gas masses of tle sub-uini galaxies are 1050 times greater than that found or the \Glky Way., The molecular gas masses of the sub-mm galaxies are 10–50 times greater than that found for the Milky Way.1153 In coutrast. the best studied Lyiuau-xeak galaxy cBhs (Yee et 11996: Pettiui et 22000: Teplitz et 22000) has even less molecular gas than he Milky Was. after correcting the observed CO upper-iuit for lensing aud suspected metallicity effects (Fraver et 11997).," In contrast, the best studied Lyman-break galaxy cB58 (Yee et 1996; Pettini et 2000; Teplitz et 2000) has even less molecular gas than the Milky Way, after correcting the observed CO upper-limit for lensing and suspected metallicity effects (Frayer et 1997)."1154 Therefore. the Lymau-break sources may be sub-L systems representing the ιάπιο blocks of more nassive galaxies. as sugeested by Lowenthal et (1997).," Therefore, the Lyman-break sources may be $L^{*}$ systems representing the building blocks of more massive galaxies, as suggested by Lowenthal et (1997)."1155" Most πα galaxies are simular to JJ00266|1708 in being too faint and/or too red to be included in the Lyiau-break survevs (6.8ο, Steidel et 11999)."," Most sub-mm galaxies are similar to J00266+1708 in being too faint and/or too red to be included in the Lyman-break surveys (e.g., Steidel et 1999)."1156 The oue notable exception is the παπα selected. galaxy wr. SMNEJJILI01110252. ΕΣ vison et 220002: Adelberger Steidel 2000). but for this svstem most of the blue light arises from J2. while the bulk of the ometriice hDmuuinositv is thought to be due to the much redder JL component (Ivison ot 22001).," The one notable exception is the sub-mm selected galaxy pair, J14011+0252 J1/J2 (Ivison et 2000a; Adelberger Steidel 2000), but for this system most of the blue light arises from J2, while the bulk of the bolometric luminosity is thought to be due to the much redder J1 component (Ivison et 2001)."1157 Due to the vigh levels of dust obscuration for the «ποπ. galaxies. here is very little overlap between the sub-uua selected ealaxies and optically-sclected Evuiui-break «πο».," Due to the high levels of dust obscuration for the sub-mm galaxies, there is very little overlap between the sub-mm selected galaxies and optically-selected Lyman-break systems."1158 Even for the Lymuan-break galaxies. most of their star-Ormation (~s80%)}) is obscured bv dust at observed optical wavelengths (Peacock et 22000: Adelberecr Steidel 2000).," Even for the Lyman-break galaxies, most of their star-formation $\sim$ ) is obscured by dust at observed optical wavelengths (Peacock et 2000; Adelberger Steidel 2000)."1159 Tence. it is clear that most of the star-ormation activity at hieh-redslüft is hidden from view at opticalultraviolet wavelengths.," Hence, it is clear that most of the star-formation activity at high-redshift is hidden from view at optical/ultraviolet wavelengths."1160 The results presented for JJO00266|1708. have iuportaut duplications on our ecucral unuderstandius of star-formation at hieh redshift., The results presented for J00266+1708 have important implications on our general understanding of star-formation at high redshift.1161 Roughly half of the total amount of star formation at high redshift is thought to occur in the subaun galaxies (Blain ct 11999h). while the other half is inferred from the opticallyselected Lyian-break sources (Peacock et 22000).," Roughly half of the total amount of star formation at high redshift is thought to occur in the sub-mm galaxies (Blain et 1999b), while the other half is inferred from the optically--selected Lyman-break sources (Peacock et 2000)."1162 Since about of the sub-nuu galaxies are undetected at optical wavelengths. a sieuificaut fraction (~25% )) of the total amount of lieh+vedshift star formation may occur in the verv faint/red sub-anni galaxies simular to SMNEJJO0266|1708.," Since about of the sub-mm galaxies are undetected at optical wavelengths, a significant fraction $\sim 25$ ) of the total amount of high-redshift star formation may occur in the very faint/red sub-mm galaxies similar to J00266+1708."1163 Similar conclusions have also been reached for a raclio-selected sample of sub-uuau ealaxies with faint optical counter-parts (Darger et 22000)., Similar conclusions have also been reached for a radio-selected sample of sub-mm galaxies with faint optical counter-parts (Barger et 2000).1164 The fact that mauv of the most hunuinous high-redshift starbursts/ACGN are too faint to be studied at optical wavelengths highlielts the importance that future sensitive nuuNIB wavelength mstrineuts will have ou our understanding of the distant mniverse., The fact that many of the most luminous high-redshift starbursts/AGN are too faint to be studied at optical wavelengths highlights the importance that future sensitive mm–NIR wavelength instruments will have on our understanding of the distant universe.1165 We report the ideutification of the sub-nuni source SNAINEJJOO266|1708. with a. faint τος ealaxy {ν= 22.5nuunag) which is uudetected at optical wavelengths. despite very deep observatious.," We report the identification of the sub-mm source J00266+1708 with a faint red galaxy $K=22.5$ mag) which is undetected at optical wavelengths, despite very deep observations."1166 This source las an extremely high Iuninositv of approximaely 101218: even after correcting for lensing., This source has an extremely high luminosity of approximately $10^{13}L_{\sun}$ even after correcting for lensing.1167 The current data for the sub-nun Cluster Leus Survey suggest that of the sub-nun population as a whole are faint/2ed galaxies which are undetected at optical waveleueths., The current data for the sub-mm Cluster Lens Survey suggest that of the sub-mm population as a whole are faint/red galaxies which are undetected at optical wavelengths.1168 These faintred απαι galaxies are thought to contribute significantly to the total amount of star formation at high redshift aud are heuce portant to our understanding of the carly evolution of galaxies., These faint/red sub-mm galaxies are thought to contribute significantly to the total amount of star formation at high redshift and are hence important to our understanding of the early evolution of galaxies.1169 The redshift of SMMILJJOO266|1708 dis. currently uuknown. but the galaxy is expected to be at a redshift +>> 2.," The redshift of J00266+1708 is currently unknown, but the galaxy is expected to be at a redshift $z>2$ ."1170 Obtaining a redshift will be extremely challenging with current mstrumentation., Obtaining a redshift will be extremely challenging with current instrumentation.1171 At 522.5 mae. the galaxy pushes the capabilities of even the largest ground based telescopes.," At $K=22.5$ mag, the galaxy pushes the capabilities of even the largest ground based telescopes."1172 Since the galaxy is relatively red οIV> 3.6). we expect Ta to be the brightest optical emission line. aud perhaps the only optical line currently detectable based on comparisons with the ERO IIR10 (Dey et 11999).," Since the galaxy is relatively red $I-K > 3.6$ ), we expect $\alpha$ to be the brightest optical emission line, and perhaps the only optical line currently detectable based on comparisons with the ERO HR10 (Dey et 1999)."1173 If the redshift is similar to that estimated from the SED of the ealaxy (2~ 3.5). Πο would be shifted redward of A-baud. making eround based observations extremely challenging.," If the redshift is similar to that estimated from the SED of the galaxy $z\sim3.5$ ), $\alpha$ would be shifted redward of $K$ -band, making ground based observations extremely challenging."1174"for this case, and refer to it as model Fnx3.","for this case, and refer to it as model Fnx3."1175" Then, together with the first model presented in 822.2, i.e., model Fnx1, the Ba/Fe evolutions predicted by the three models in total are shown in the upper panel of Figure 4."," Then, together with the first model presented in 2.2, i.e., model Fnx1, the Ba/Fe evolutions predicted by the three models in total are shown in the upper panel of Figure 4."1176" In the middle panel, the resultant ADF summed up by the individual models with a ratio of 0.5/0.3/0.2 is compared with the observation (Battagliaetal.2006)."," In the middle panel, the resultant ADF summed up by the individual models with a ratio of 0.5/0.3/0.2 is compared with the observation \citep{Battaglia_06}."1177". 'The star formation history constructed by these relative contributions from individual models is compared with the observed history (Coleman&deJong2008),, which is reassessed without counting very low star formation for t« 2 Gyr and an age threshold of 13 Gyr, in the lower panel."," The star formation history constructed by these relative contributions from individual models is compared with the observed history \citep{Coleman_08}, which is reassessed without counting very low star formation for $t<$ 2 Gyr and an age threshold of 13 Gyr, in the lower panel."1178 Here we review the adopted models in the context of the integrated galactic IMF (IGIMF) theory (Kroupa&Weidner2003;&Kroupa2005)..," Here we review the adopted models in the context of the integrated galactic IMF (IGIMF) theory \citep{Kroupa_03, Weidner_05}. ."1179" The IGIMF predicts that m, depends on the SFR (or gas mass) in a sense that m, is smaller according to a lower SFR.", The IGIMF predicts that $m_u$ depends on the SFR (or gas mass) in a sense that $m_u$ is smaller according to a lower SFR.1180" 'Thus, our hypothesis that model Fnx2 with a moderate SFR is equipped with a larger m, than models Fnx1&33 with a higher SFR seems at odds with the IGIMF theory."," Thus, our hypothesis that model Fnx2 with a moderate SFR is equipped with a larger $m_u$ than models 3 with a higher SFR seems at odds with the IGIMF theory."1181" As one possible explanation to integrate our models into this theoretical scheme, the stellar population associated with model Fnx2 may be regarded as that originally belonging to more massive building block within a hierarchical galaxy formation scenario."," As one possible explanation to integrate our models into this theoretical scheme, the stellar population associated with model Fnx2 may be regarded as that originally belonging to more massive building block within a hierarchical galaxy formation scenario."1182 T'hat casts our attention on the observed fact that the Fnx dSph obeys the luminosity-metallicity relation etal.2008;Revazetal.2009).," That casts our attention on the observed fact that the Fnx dSph obeys the luminosity-metallicity relation \citep{Kirby_08, Revaz_09}."1183". This property(Kirby may provide the possible path to the complex populations such that the merger event of Fnx2 into Fnx1l induces the enhanced star formation, and yields a metal-rich component Fnx3, then ultimately ends up with a metallicity increase as a whole in accordance with the mass growth of galaxy."," This property may provide the possible path to the complex populations such that the merger event of Fnx2 into Fnx1 induces the enhanced star formation, and yields a metal-rich component Fnx3, then ultimately ends up with a metallicity increase as a whole in accordance with the mass growth of galaxy."1184" On the other hand, a maximal IGIMF model together with the total mass of 1.9x10’Mo for the Fnx dSph (Wooteal.2008) approximately gives m,=30, 20 to model Fnx1, Fnx3, respectively."," On the other hand, a maximal IGIMF model together with the total mass of $1.9\times 10^7$ for the Fnx dSph \citep{Woo_08} approximately gives $m_u$ =30, 20 to model Fnx1, Fnx3, respectively."1185" In this study, we have a close look at the signature of a small m,, at an early evolutionary stage."," In this study, we have a close look at the signature of a small $m_u$ at an early evolutionary stage."1186" However, given that we ignore it, one standard chemical evolution model that predicts a declining gas mass (SFR) as a function of elapsed time etal. is expected to lead to a gradual reduction(Kirby in m,2011) that will make the Ba/Fe ratio mildly upward during late evolution."," However, given that we ignore it, one standard chemical evolution model that predicts a declining gas mass (SFR) as a function of elapsed time \citep{Kirby_11} is expected to lead to a gradual reduction in $m_u$ that will make the Ba/Fe ratio mildly upward during late evolution."1187" Finally, we examine the a/Fe feature in the framework of the models constructed in the previous section."," Finally, we examine the $\alpha$ /Fe feature in the framework of the models constructed in the previous section."1188 The prediction of the ratio of O (or Mg) to Fe ejected from an individual SN II still involves uncertainties because these two elements are synthesized by a completely different process., The prediction of the ratio of O (or Mg) to Fe ejected from an individual SN II still involves uncertainties because these two elements are synthesized by a completely different process.1189" O is produced through hydrostatic nuclear burning during (Mg)stellar evolution before the SN explosion, whereas the amount of Fe is determined by the explosion mechanism and fallback dynamics, which are not fully understood yet."," O (Mg) is produced through hydrostatic nuclear burning during stellar evolution before the SN explosion, whereas the amount of Fe is determined by the explosion mechanism and fallback dynamics, which are not fully understood yet."1190" However, the observed feature of [a/Fe] in dSphs cannot be fully interpreted by the knee structure positioned at a low-metallicity, since there exist stars exhibiting low [a/Fe] ratios at a very low [Fe/H] (seedistinctcaseoftheSextandSph;Aokietal."," However, the observed feature of $\alpha$ /Fe] in dSphs cannot be fully interpreted by the knee structure positioned at a low-metallicity, since there exist stars exhibiting low $\alpha$ /Fe] ratios at a very low [Fe/H] \citep[see the distinct case of the Sextan dSph;][]{Aoki_09}."1191 Suppose that the average a/Fe ratiointegrated 2009)..over 10-25 becomes small with the assumption that an increasing trend of yield against the progenitor mass is more prominent for a-elements than for Fe., Suppose that the average $\alpha$ /Fe ratiointegrated over 10-25 becomes small with the assumption that an increasing trend of yield against the progenitor mass is more prominent for $\alpha$ -elements than for Fe.1192" Here we assume that a truncated IMF in models Fnx1 and Fnx3 invokes that the average O, Mg, and Ti yields from SNe II are reduced by a factor of 0.7/0.5/0.3."," Here we assume that a truncated IMF in models Fnx1 and Fnx3 invokes that the average O, Mg, and Ti yields from SNe II are reduced by a factor of 0.7/0.5/0.3."1193" Together with model Fnx 2, the results predicted by the three models for the Fnx dSph are shown by the colored curves, while the dashed curves represent the Milky Way case."," Together with model Fnx 2, the results predicted by the three models for the Fnx dSph are shown by the colored curves, while the dashed curves represent the Milky Way case."1194" For the nucleosynthesis yields of O and Mg form SNe II, we take their values tabulated in Tsujimotoetal.(1995)."," For the nucleosynthesis yields of O and Mg form SNe II, we take their values tabulated in \citet{Tsujimoto_95}."1195". On the other hand, the Ti yield as a function of progenitor's mass is deduced from theabundance analysis of very metal-poor stars as performed by Tsujimoto&Shigeyama (1998),, because the existing nucleosynthesis models fail to predict the Ti yield which"," On the other hand, the Ti yield as a function of progenitor's mass is deduced from theabundance analysis of very metal-poor stars as performed by \citet{Tsujimoto_98}, , because the existing nucleosynthesis models fail to predict the Ti yield which"1196Several authors (Fosalbaetal.2003;—Vielva2006:Granettetal.2008) have claimed that there is a significant cross-correlation between cosmic microwave background radiation (CMBR) anisotropies and the density of galaxies. which is interpreted as the integrated Sachs-Wolfe (ISW) effect.,"Several authors \citep{Fosalba03,Vielva06,Cabre06,Raccanelli08,Ho08,Granett08} have claimed that there is a significant cross-correlation between cosmic microwave background radiation (CMBR) anisotropies and the density of galaxies, which is interpreted as the integrated Sachs-Wolfe (ISW) effect."1197 An anticorrelation caused by the Sunyaev-Zeldovich effect would also be expected on scales smaller than ~1° but this is negligible when averaging large regions of the sky (Hernández-Monteagudo&Rubino-Martín2004).," An anticorrelation caused by the Sunyaev-Zel'dovich effect would also be expected on scales smaller than $\sim 1^\circ $, but this is negligible when averaging large regions of the sky \citep{Hernandez04}."1198.. The conclusion of these authors is that the measured cross-correlation should be interpreted as a detection of the ISW effect within à ACDM-cosmology and it serves to constrain the value of the cosmological parameters., The conclusion of these authors is that the measured cross-correlation should be interpreted as a detection of the ISW effect within a $\Lambda $ CDM-cosmology and it serves to constrain the value of the cosmological parameters.1199 We reanalyze whether this correlation exists by considering galaxies observed by the Sloan Digital Sky Survey (SDSS). taking particular care in the calculation of the cross-correlation errors.," We reanalyze whether this correlation exists by considering galaxies observed by the Sloan Digital Sky Survey (SDSS), taking particular care in the calculation of the cross-correlation errors."1200 The root mean square (r.m.s.), The root mean square (r.m.s.)

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